{
    "0711/0711.4856_arXiv.txt": {
        "abstract": "According to the standard liquid-water definition, the Earth is only partially habitable.  We reconsider planetary habitability in the framework of energy-balance models, the simplest seasonal models in physical climatology, to assess the spatial and temporal habitability of Earth-like planets. We quantify the degree of climatic habitability of our models with several metrics of fractional habitability. Previous evaluations of habitable zones may have omitted important climatic conditions by focusing on close Solar System analogies. For example, we find that model pseudo--Earths with different rotation rates or different land--ocean fractions have fractional habitabilities that differ significantly from that of the Earth itself.  Furthermore, the stability of a planet's climate against albedo-feedback snowball events strongly impacts its habitability. Therefore, issues of climate dynamics may be central in assessing the habitability of discovered terrestrial exoplanets, especially if astronomical forcing conditions are different from the moderate Solar System cases. ",
        "introduction": "\\label{sec:intro}  Planetary science is being challenged by extrasolar planetary system discoveries \\citep{mayor+queloz1995,marcy+butler1996}.\\footnote{See \\citet{reipurth_et_al2007} for a review.  See also http://exoplanet.eu and http://exoplanets.org.} The orbital architectures of many of these systems, with massive planets located either very near to their parent stars or on highly eccentric orbits, are strikingly unusual. These discoveries have been surprising, however, mostly as a consequence of our own Solar-System centric point of view. In fact, planet searches at multi-AU orbital distances are now providing tentative evidence that Solar System giant planets, with their nearly circular orbits, are more the exception than the norm, by simple comparison with the ensemble properties of known extrasolar planetary systems \\citep{marcy_et_al2005}.  If these preliminary trends hold, they may have a profound impact on our perception of our place in the Universe.   The pace of exoplanet discoveries has accelerated sharply in the last few years and the future looks bright. Two significant developments in exoplanet research have occurred within the last year or so. A micro-lensing discovery of a likely terrestrial planet, with a mass $5.5$ times that of Earth, was reported by \\citet{beaulieu_et_al2006}, making it the first discovery of a planet that is thought to be terrestrial (i.e. rocky). Several months ago, \\citet{udry_et_al2007} announced the Doppler-velocimetry discovery of a potentially habitable terrestrial planet around a low-mass M-dwarf.  In considering future discoveries, of specific interest to this study are two dedicated space missions: {\\it Corot} and {\\it Kepler} \\citep{baglin2003,borucki_et_al2003,borucki_et_al2007}. A key objective of both missions is to monitor a large number of stars to detect the (repeatable) micro-eclipses generated when terrestrial planets transit in front of their host star.  It is expected that a few Earth analogues (i.e. with Earth-like masses and comparable distances to their host stars), and possibly hundreds to thousands of additional terrestrial planets unlike Earth, will be identified by {\\it Corot} and {\\it Kepler} after a few years of operation \\citep{borucki_et_al2007, borucki_et_al2003, basri_et_al2005}. The discovery of an Earth-like exoplanet, potentially hosting life as we know it, is therefore within the 5-year astronomical horizon. Ambitious missions are also being prepared to map in detail the orbits of exoplanets around nearby stars (with {\\it SIM PlanetQuest} \\citep{unwin_et_al2007}) and later to obtain spectra of nearby Earth-like planets in the hope that they would reveal the first unambiguous signatures of life on a remote world (with the {\\it Terrestrial Planet Finder} and {\\it Darwin} \\citep{leger+herbst2007}).   The idea that liquid surface water is a prerequisite for a terrestrial planet to have the ability to host life is widely used as the key concept behind searches for habitable planets around other stars. This is because of the central and critical role that water plays in the biochemistry on Earth.\\footnote{It should however be noted that alternate molecules (e.g., ammonia, methyl alcohol) could conceivably perform equivalent roles in different environments, such as those of lower temperature and higher pressure \\citep{haldene1954,firsoff1963,goldsmith+owen2002}.}  As the current exoplanet census indicates, notions based on an old Solar System centric view may only be relevant to a minority of planets (or planetary systems). What if most terrestrial planets discovered in the future have, like the vast majority of exoplanets currently known, highly eccentric orbits generating large seasonal variations?  What if the atmospheric mass and composition, planetary spin rate, or land-sea coverage of these exo-Earths are generally different from what they are on Earth?  If the last ten years of extrasolar planet discoveries offer any guidance, our Solar System appears to show but little of the general planetary diversity found around other stars.   The classical calculations of habitability on Earth-like planets by, e.g., \\citet{dole1964} and \\citet{hart1979} predated extrasolar discoveries. Apart from a few exceptions (\\citealt{franck_et_al2000b,gaidos_et_al2005}; see also \\citealt{vonbloh_et_al2007,selsis_et_al2007} in the specific context of the Gleise 581 system) the subject of planetary atmospheric habitability has been revisited little since the seminal work of \\citet{kasting_et_al1993}. Given the major developments expected in the next five years and beyond, it is important that the climate regimes expected on exotic versions of the Earth, and their consequences for habitability, be studied and better understood.  This will help interpret upcoming planet detections and will inform future long-term efforts on the best strategies to find robust signatures of life on exo-Earths.  The central role of astronomical forcings in determining the seasons and the climatology of the Earth is well known (e.g., the long-term Milankovitch cycles).  Clearly, obliquity, precession, and eccentricity can all strongly affect global and regional habitability conditions on the Earth and, by extension, on any other potentially habitable terrestrial exoplanet.  Previously, only preliminary investigations of the role of obliquity (\\citealt{williams+kasting1997}, hereafter WK97; and \\citealt{williams+pollard2003}) and eccentricity \\citep{williams+pollard2002}, for very specific Earth-like conditions, have been considered in some detail. The surface habitability of a terrestrial planet, however, must depend on the combination of obliquity and eccentricity with the planetary rotation rate, the continental coverage (from dry Earths to water worlds) and the overall mass and composition of the atmospheric layer (e.g., compare Mars vs. Venus), among other factors.  A thorough exploration of how these various global planetary attributes combine to affect climatology and make a terrestrial planet seasonally or regionally habitable is therefore an important element in the search for signatures of life elsewhere. In the present study, we describe how energy balance climate models can contribute to this understanding.   In \\S~\\ref{sec:model}, we introduce the concept of climate modeling hierarchy, describe our energy balance model and validate it on the Earth.  In \\S~\\ref{sec:Hab Zones}, we reconsider various features of habitability for Earth--like planets with seasonally-forced climates. In \\S~\\ref{sec:def}, we discuss several subtleties in the definition of habitability that emerge from our work on seasonal climates, and we finally conclude in \\S~\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  We have reconsidered the notion of habitability for Earth-like planets with seasonal energy-balance climate models. These models show that the concept of regional and seasonal habitability is generally important to assess the ability of terrestrial exoplanets to host life. We find that previous evaluations of habitable zones may have omitted important climatic conditions by focusing on close Earth analogies. We illustrate this with two specific examples: pseudo--Earths rotating at different rates or possessing a smaller ocean fraction than Earth itself. These pseudo-Earths have quantitatively different climatic habitability properties than the Earth itself.  Comparisons to global radiative balance calculations show that seasonal habitability generally extends the inner orbital range of Earth-like planet habitable zones. The outer orbital range is reduced relative to what global radiative balance calculations would indicate, however, because the climate generally makes a dynamical transition to a globally-frozen snowball state before the outer radiative limit is reached. The stability of a planet's climate against snowball events therefore has a strong impact on its habitability. Since this stability is partly determined by external forcing factors such as the magnitude of insolation and the length of winters, we expect issues of climate dynamics to be central in determining the habitability of terrestrial exoplanets, particularly if their forcing conditions are generally different from the moderate cases encountered in the Solar System."
    },
    "0711/0711.4584_arXiv.txt": {
        "abstract": "The determination of the Earth gravity field from space geodetic techniques now allows us to obtain the temporal variations of the low degree coefficients of the geopotential, combining the orbitography of several satellites (e.g. Lageos1, Lageos2, Starlette). These temporal variations of the Earth gravity field can be related to the Earth Orientation Parameters (EOP) through the inertia tensor. This paper shows these relations and discusses how such geodetic data can contribute to the understanding of the variations in EOP. ",
        "introduction": "The Earth Orientation is generally considered as (i) Earth rotation axis movements in space (precession-nutation), (ii) Earth rotation axis movements in the Earth (polar motion), or (iii) Earth rotation speed variations (exces in the length of the day). These movements come from Earth inside masses distributions.  The Earth gravity field can give us information about this distribution of masses because nowadays we can determine the variations of the Earth gravity field by space geodetic techniques.  Hence, there is a link between the variations of the Earth gravity field and the variations of the Earth Orientation Parameters. And the high accuracy now reached in the VLBI (Very Long Baseline Interferometry) Earth Orientation Parameters (EOP) determination requires looking further at the various geophysical contributions to variations in EOP. So we investigate here if this variable gravity field can be valuable for the improving the modelisation of the Earth rotation. ",
        "conclusions": "The part of the length of the day obtained with the $\\Delta C_{20}$ data corresponds to gravitational terms. Then we have compared $\\Delta(LOD)$ corrected from the movements terms (as atmospheric ones), the zonal tides and the decadal terms (from magnetic effects in the core-mantle boundary). But the residual term has an amplitude of the order of $50$ $\\mu$s, whereas the better precision on these LOD data is of the order of $10$ $\\mu$s.  The effect of the variable gravity field on the polar motion can be investigated now, using Eq.~(\\ref{eq:p}).  Finally, we find a 18.6-yr periodical effect on the precession angle development in longitude with a sinus term of about $105$ $\\mu$as (Bourda \\& Capitaine 2004).  In the future, the static gravity field and its temporal variations coming from the GRACE satellite will be very usefull for these kind of studies, because they are very precise."
    },
    "0711/0711.1704_arXiv.txt": {
        "abstract": "We report the detection of new ammonia masers in the non-metastable (8,6) and (11,9) transitions towards the massive star forming region NGC\\,6334\\,I. Observations were made with the ATCA interferometer and the emitting region appears unresolved in the 2.7\\arcsec $\\times$ 0.8\\arcsec~beam, with deconvolved sizes less than an arcsecond. We estimate peak brightness temperatures of $7.8 \\times 10^5$ and $1.2 \\times 10^5$\\,K for the (8,6) and (11,9) transitions, respectively. The masers appear coincident both spatially and in velocity with a previously detected ammonia (6,6) maser. We also suggest that emission in the (10,9), (9,9) and (7,6) transitions may also be masers, based on their narrow line widths and overlapping velocity ranges with the above masers, as observed with the single-dish Mopra radiotelescope. ",
        "introduction": "Ammonia (NH$_3$) is an extremely useful molecular tool for studying the interstellar medium. Many inversion transitions occur in the easily observable 12\\,mm band. They include transitions from both metastable (J=K) and non-metastable (J$>$K) levels of ammonia. Hyperfine structure in the metastable transitions is commonly seen, which can be used to derive optical depths, whereas comparison of different (J,K) transitions can yield information on the rotational temperature (eg. \\citealt{walmsley83,danby88}).  The first suggestion that ammonia may display maser action in interstellar space was made by \\citet{wilson82} who noted that (3,3) emission toward W33 was not matched by (1,1), (2,2) and (4,4) transitions, which were all in absorption. They suggested that this may be due to a weak population inversion. Maser emission in ammonia transitions was first unambiguously identified by \\citet{madden86} in the (9,6) and (6,3) transitions, shortly followed by \\citet{mauersberger86} in the (3,3) transition of $^{15}$NH$_3$. In addition to these transitions, ammonia masers have been detected in the following transitions: (5,4) (7,5), (9,8), (10,8) \\citep{mauersberger87}, (6,5) \\citep{mauersberger88}, (5,5) \\citep{cesaroni92}, (6,6) \\citep{beuther07} and (1,1) \\citep{gaume96}. $^{14}$NH$_3$ (3,3) masers were first detected by \\citet{zhang95}. Pumping mechanisms for most ammonia masers remain unclear. Metastable transitions of ortho-ammonia, like (3,3) are thought to be collisionally excited \\citep{walmsley83} but such a mechanism will only work for non-metastable transitions where exceedingly high H$_2$ densities between $10^{10} - 10^{12}$ cm$^{-3}$ are found (eg. \\citealt{madden86}). \\citet{madden86} suggest two alternative pumping mechanisms for non-metastable transitions: either pumping by a strong infrared radiation field, such as found around a deeply embedded high mass star, or by a fortuitous overlap of a far-infrared line, which allows a population transfer from a (J,K) inversion level to (J+1, K).  Ammonia masers are found in regions of high mass star formation, with the best known example being W51 \\citep{madden86,mauersberger87}, with other prominent examples being G9.62+0.19 \\citep{cesaroni92,hofner94} and NGC6334I \\citep{kraemer95,beuther07}. NGC6334I forms the focus of this work. It is a nearby (1.7kpc, \\citealt{neckel78}) region of high mass star formation, traced by bright infrared emission, an ultracompact (UC) \\hii region, mm continuum sources and methanol maser sites. See \\citet{beuther07} for a more comprehensive summary of characteristics of the region. ",
        "conclusions": "We have observed the two non-metastable transitions of ammonia: (8,6) and (11,9), using the ATCA. The (8,6) transition has a peak brightness temperature of 7.8 $\\times 10^5$\\,K, whereas the (11,9) transition has a peak brightness temperature of 1.2 $\\times 10^5$\\,K, indicating that both are masers. The (11,9) transition is E$_{\\rm l}/k=1449$\\,K above ground making it the highest energy ammonia maser currently known. The position of both masers is consistent with being coincident, to within positional uncertainties, with maser emission in the previously detected (6,6) transition \\citep{beuther07}.  Single-dish observations of (7,6), (9,9) and (10,9) suggest they may also be masers based on their narrow line widths and overlapping velocity range with the above-mentioned masers, which is distinctly different from the thermal emission systemic velocity. High spatial resolution observations of these transitions are planned to decide whether or not these are masers as well."
    },
    "0711/0711.0702_arXiv.txt": {
        "abstract": "{In the log Age vs. integrated absolute magnitude ($M_V$) plane, the open clusters of the Milky Way form a well-defined band parallel to theoretical sequences decribing the passive evolution of Simple Stellar Populations and display a pretty sharp upper threshold in mass ($M\\sim 2\\times 10^4 ~M_{\\odot}$) over a 4 dex range of ages. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.2477_arXiv.txt": {
        "abstract": "% Given multiband photometric data from the SDSS DR6, we estimate galaxy redshifts. We employ a Random Forest trained on color features and spectroscopic redshifts from 80,000 randomly chosen primary galaxies yielding a mapping from color to redshift such that the difference between the estimate and the spectroscopic redshift is small.  Our methodology results in tight RMS scatter in the estimates limited by photometric errors. Additionally, this approach yields an error distribution that is nearly Gaussian with parameter estimates giving reliable confidence intervals unique to each galaxy photometric redshift. ",
        "introduction": "We are given five bands of photometric data from the Sloan Digital Sky Survey Data Release 6 \\htmladdnormallinkfoot{(SDSS DR6)}{http://www.sdss.org/dr6/}. Associated with each magnitude measurement is an error measurement and an extinction measurement used to correct the effect of Galactic dust.  For some objects we have spectroscopic redshifts, hereon denoted $z_{spec}$ for a particular object.  We wish to estimate photometric redshifts for non-spectroscopic objects which are represented by the spectroscopic sample.  We are currently interested only in galaxies with $10^{-4} \\le z_{spec} \\le 1$, and we exclude other objects which are likely to \"contaminate\" our sample.  We further choose for the moment not to consider Luminous Red Galaxies (LRGs), leaving us with some $400,000$ objects with which to train and test.  In our current work we use magnitudes available in the new Ubercal table (Padmanabhan 2007) for each object.  We subtract sequential extinction-corrected magnitudes to get color features, named $u$, $g$, $r$, and $i$. For instance, our color $u$ is actually the magnitude difference $u - g$ and so on. ",
        "conclusions": "Our RMS error is consistent with results from previous studies using similar datasets, though slightly lower RMS errors from different methodologies have been reported.  We may yet gain by expanding our training data set as we have more than $300,000$ training objects in reserve and the computational complexity of random forests is quite modest (in preliminary testing we trained multiple forests with multiple training sets on a Core 2 Duo 2.0GHz MacBook in less than one week; regression on all of our held-out data took several days using the same machine).  Future work will include better accounting for dependence between trees, investigating more deeply the behavior of our error distributions as a function of $z_{spec}$, addressing $z_{spec}$-dependent bias, and extending estimation to objects not represented by the training data.  As it is, the quality of the estimates, the per-object error distributions, and the computational efficiency make this approach an attractive option for photometric redshift estimation."
    },
    "0711/0711.2688_arXiv.txt": {
        "abstract": "{\\it Very Large Array } and Parkes 64 m radiotelescope 21-cm observations of the starburst dwarf galaxy NGC~5253 reveal a multi-component non-axisymmetric \\HI\\ distribution.  The component associated with the stellar body shows evidence for a small amount of rotational support aligned with the major axis, in agreement with optically measured kinematics and consistent with the small galaxian mass.  Approximately 20--30\\% of the \\HI\\ emission is associated with a second component, an \\HI\\ ``plume'' extending along the optical minor axis to the southeast. We consider outflow, inflow, and tidal origins for this feature. Outflow appears improbable, inflow is a possibility, and tidal debris is most consistent with the observations. Thus, kinematics of the \\HI\\ that include this feature are not indicative of the dynamical mass or the local velocity dispersion of the cold gas.  These \\HI\\ observations also reveal a filamentary third component that includes an 800 pc diameter \\HI\\ shell or bubble to the west of the nucleus, coinciding with an H$\\alpha$ shell.  The mass of \\HI\\ in the shell may be as large as $\\sim$$4\\times10^6$ \\msun.  This large mass, coupled with the lack of expansion signatures in the neutral and ionized gas ($v<30$ \\kms), suggests that this feature may be an example of a starburst-blown bubble stalled by interaction with a massive neutral envelope.  Many other \\HI\\ kinematic features closely resemble those seen in H$\\alpha$ emission from the ionized gas, supporting the interpretation of neutral and ionized gas outflow at velocities of $\\sim$30 \\kms.  Comparison between extinction estimates from the Balmer emission-line decrement and the \\HI\\ column densities suggest a gas-to-dust ratio 2--3 times the Galactic value in this low-metallicity ($Z=1/4~Z_\\odot$) galaxy. ",
        "introduction": "NGC~5253 is a remarkable starbursting dwarf galaxy with a prominent minor axis dust lane located in the nearby Centarus A/M~83 galaxy complex \\citep{kara}  .  At visual wavelengths, the impression of NGC~5253 is dominated by its network of ionized gas filaments which extend across its minor axis and reach beyond the stellar distribution \\citep{hodge, graham81, caldwell, marlowe, martin}.  X-ray observations reveal several small starburst-heated bubbles but no monolithic superbubble or galactic wind \\citep{strickland, summers}.  It harbors an extremely luminous, compact, obscured site of nuclear star formation containing several star clusters with masses from $10^5$ -- $10^6$ $M_\\odot$ having ages $<10^7$ yrs \\citep{gonzalez, calzetti97, vanzi, martin-h}.  Much of the nucleus is obscured by high levels of patchy dust extinction up to $A_V=25$ mag \\citep{aitken, moorwood, alonso}. Radio-wave continuum and recombination line studies have revealed a dominant luminous ($10^9$ $L_\\odot$), compact (r=1 pc), young ($<2.3$ Myr) star cluster that appears to be gravitationally bound and potentially destined to evolve into a globular cluster \\citep{beck, turner98, gorjian, mohan}.  The centimeter-wave spectral energy distribution is consistent with an optically-thick thermal Bremsstrahlung source\\footnote{\\citet{turner00} term this object a ``supernebula'' while \\citet{kj} use the name ``ultra-dense \\HII\\ regions'' (UDHII) by analogy with the Galactic ultra-compact \\HII\\ regions \\citet{wood}.}  having densities exceeding $10^4$ cm$^{-3}$ and emission measures $> 10^8$ cm$^{-6}$ pc powered by $> 10^{52}$ Lyman continuum photons s$^{-1}$. \\citet{caldwell} studied the star formation history of NGC~5253 and determined that the light was dominated by a young component with an age between 10$^8$ and 10$^9$ yrs.  Recent studies have concluded that all of the UV-selected star clusters in NGC~5253 have estimated ages $<$20 Myr, indicating that either the star formation episode began very recently or else that the timescale for dynamical disruption of star clusters is very short \\citep{tremonti, harris}.  The overall gas-phase metallicity of NGC~5253 is 12+$\\log(O/H)\\simeq8.16$ (about 1/4 solar), but some of the central regions contain an excess of nitrogen, consistent with local ``pollution'' from Wolf-Rayet star winds \\citep{wr, k97, ls07}.  When observed at low resolution, the distribution and kinematics of its neutral interstellar medium suggest that the \\HI\\ in NGC~5253 either rotates about the minor axis or flows radially along the minor axis \\citep{ks}.  \\citet{turner97} and \\citet{meier} report the $^{12}CO$ detection of molecular gas clouds having kinematics consistent with infall along the minor axis dust lane.  These peculiar gas kinematics may hold clues to the mechanism that triggered the present burst of star formation.  In this paper we present 21-cm aperture synthesis observations from the National Radio Astronomy Observatory's {\\it Very Large Array}\\footnote{The National Radio Astronomy Observatory is operated by Associated Universities, Inc. (AURA) under cooperative agreement with the National Science Foundation.} (VLA) and the Parkes 64 m radio telescope to investigate the conditions in the interstellar medium that have produced such extreme star formation. NGC~5253 is an ideal environment to seek the causes of extreme star formation because the burst is so young and, presumably, feedback from the energetic star clusters has not yet had a dramatic impact on the surrounding ISM.  The distance to NGC~5253 is somewhat controversial and has been measured in many studies including \\citet{saha} (4.1 $\\pm$0.5 Mpc), \\citet{gibson} (3.3$\\pm0.3$ Mpc), \\citet{freedman} (3.25$\\pm$0.2 Mpc), \\citet{karachentsev} (3.9$\\pm$0.5 Mpc), \\citet{thim} (4.0$\\pm$0.3 Mpc), and \\citet{sakai} (3.8 $\\pm$0.2 Mpc). The distances based on Cepheid variables strongly depend on the sample selection \\citep{gibson, thim} and so we prefer the TRGB distance of 3.8 Mpc \\citep{sakai} which implies an angular scale of 18.4 pc arcsec$^{-1}$. ",
        "conclusions": "Radio-wave aperture synthesis observations in the 21-cm line of neutral hydrogen at 8\\arcsec\\ resolution have revealed a complex non-axisymmetric gas distribution and velocity field in NGC~5253. The total neutral atomic content of the galaxy is $1.4\\times10^8$ \\mo\\ from single-dish observations for a distance of 3.8 Mpc.  The neutral medium in NGC~5253 may be characterized as consisting of three components: a component containing the majority of the neutral medium distributed along the stellar distribution and showing evidence for low-amplitude rotation at the level of 15--20 \\kms; a redshifted \\HI\\ plume extending to the south-southeast along the minor axis containing 20--30\\% of the \\HI; and kinematically cold \\HI\\ shells or filaments extending along the minor axes and to the southwest, containing $\\sim$10\\% of the neutral gas mass. Taken together, these different components result in very complex velocity structures.  Here we will try to summarize out best understanding of these structures, starting with the relatively enigmatic \\HI\\ plume.  \\subsection{The Nature of the \\HI\\ Plume}  The kinematically distinct \\HI\\ plume consists of a significant fraction of the total \\HI\\ emission from NGC~5253.  In this paper we have discussed three possible origins for the \\HI\\ plume: outflowing gas, inflowing gas, and a separate \\HI\\ cloud which is interacting with NGC~5253.  The possibility of outflow was motivated by similarity with velocity fields observed in other dwarf starburst galaxies showing \\HI\\ outflow (e.g., NGC 1569, NGC 625, NGC 1705). However, the large fraction of \\HI\\ ($\\approx$25\\%) in the anamalous velocity gas in NGC~5253 is exceptional. Furthermore, strong outflows are usually accompanied by H$\\alpha$ and X-ray emission which is not observed at this location. Such a massive outflow confined to a small solid angle without corresponding signatures in the ionized gas seems unlikely.  The possibility of inflow was motivated by the radial velocities of CO clouds within the minor-axis dust lane \\citep{meier}. Our high resolution \\HI\\ observations indicate that the plume is a kinematically distinct feature from the \\HI\\ associated with the molecular gas.  The most compelling argument for inflow is the near coincidence in velocity of the \\HI\\ plume with the galaxy's systemic velocity in the region where the two are spatially coincident.  If the \\HI\\ plume is infalling, it must be from the far side of NGC~5253 given the plume's velocity field (i.e., the most redshifted gas lies farthest from the galaxy), consistent with the lack of extinction observed in this vicinity despite the high \\HI\\ column density.  Finally, we have the possibility that the \\HI\\ plume is a distinct structure, perhaps a remnant from a recent interaction, seen in projection against the main body of NGC~5253.  In this scenario, a recent gravitational encounter, perhaps with another dwarf galaxy in the same group, produced the \\HI\\ plume and provided the trigger for the current extreme burst of star formation.  Other dwarf galaxies, both bursting and non-bursting are accompanied by pure \\HI\\ structures.  Given the relatively large gas mass associated with the \\HI\\ plume, this would appear to be the most likely explanation.  \\subsection{A Comprehensive View of the \\HI\\ in NGC 5253}  Our high resolution view of NGC~5253 has shown us that NGC~5253 has many similarities to the best studied dwarf starburst galaxies. \\HI\\ concentrations near the minor axis dust lane are coincident in location and velocity with at least two of the molecular clouds seen in $^{12}$CO \\citep{meier}.  The \\HI\\ extension to the south-southeast, designated here the \\HI\\ plume, appears to be spatially and kinematically distinct from the gaseous medium surrounding the dust lane molecular clouds along the southeast minor axis. We present evidence that the redshifted \\HI\\ plume is likely to be either a dynamically distinct tidal remnant or an inflow on the far side of the galaxy. \\HI\\ extensions to the east, west, and southwest show remarkably similar morphologies as the H$\\alpha$ shells and filaments.  In particular, there is a close morphological similarity between the ``quiescent'' H$\\alpha$ shell to the west of the nucleus noted by \\citet{martin} and an 800 pc diameter shell-like feature seen in \\HI.  This may be an example of a starburst-powered shell stalled by the mass of surrounding neutral gas.  These complex \\HI\\ features and kinematic signatures, when considered together with distributions of the molecular, ionized, and hot coronal media, are consistent with a starbursting low-mass galaxy in the initial phases of generating a galactic wind.  Unlike more massive galaxies such as M~82 or NGC~3079 which show monolithic, well-collimated outflows \\citep[e.g.,][and references therein] {sh07,cbv02} several non-co-spatial star formation events over the last $\\sim20$ Myr appear to be required to explain the multiple outflows observed in NGC~5253 \\citep{summers, strickland, caldwell}.  The initiating and sustaining cause of the recent star formation activity may be an inflow or interaction of atomic gas along the ``plume'', identified here for the first time.  As such, NGC~5253 joins a long list of starburst galaxies where recent gravitational interactions are implicated as the trigger for extreme star formation bursts."
    },
    "0711/0711.0408_arXiv.txt": {
        "abstract": "{} {We study the structure of the magnetic elements in network-cell interiors.} {A quiet Sun area close to the disc centre was observed with the spectro-polarimeter of the Solar Optical Telescope on board the Hinode space mission, which yielded the best spatial resolution ever achieved in polarimetric data of the \\ion{Fe}{i}\\,630\\,nm line pair. For comparison and interpretation, we synthesize a similar data set from a three-dimensional magneto-hydrodynamic simulation.} {We find several examples of magnetic elements, either roundish (tube) or elongated (sheet), which show a central area of negative Stokes-$V$ area asymmetry framed or surrounded by a peripheral area with larger positive asymmetry. This pattern was predicted some eight years ago on the basis of numerical simulations. Here, we observationally confirm its existence for the first time.} {We gather convincing evidence that this pattern of Stokes-$V$ area asymmetry is caused by the funnel-shaped boundary of magnetic elements that separates the flux concentration from the weak-field environment. On this basis, we conclude that electric current sheets induced by such magnetic boundary layers are common in the photosphere.} ",
        "introduction": "With ever increasing polarimetric sensitivity and spatial resolution, it becomes now clear that even the most quiescent areas on the solar surface harbour ample amounts of magnetic flux. This flux becomes visible in small patches of field concentrations, called magnetic elements for short. While horizontal fields tend to occur near the edge of granules, the present investigation focusses on fields predominantly oriented in the vertical direction and mainly occurring in the intergranular space \\citep{lites_etal_07b,lites_etal_07a}. There, magnetic elements are often visible in G-band filtergrams as point-like objects. In more active regions, magnetic elements of extended size appear in the form of `crinkles', `ribbon bands', `flowers', etc. \\citep{berger+al04}, which show a sub-structure consisting of a dark central area surrounded by a striation of bright material and further outside by a downdraft of plasma \\citep{langangen_etal_07}. While Doppler measurements of this downdraft are at the limit of spatial resolution with the best ground-based solar telescopes, spectro-polarimetric measurements have less spatial resolution due to image degradation by atmospheric seeing over the required long integration times.  Measurements with the spectro-polarimeter on board the Hinode satellite have a superior spatial resolution of  approximately $0.3\\arcsec$ thanks to excellent pointing capabilities \\citep{shimizu_etal_07} and the absence of seeing. Using Hinode data, we show examples of unipolar magnetic elements of the network-cell interiors that possess a distinct sub-structure, which is strikingly manifested in maps of the asymmetry of spectral lines in the circularly polarized light --- the Stokes-$V$ area asymmetry.  The areas of the blue and the red lobes of Stokes-$V$ profiles are identical when formed in a static atmosphere, but become asymmetric in the presence of gradients in the velocity and magnetic field strength \\citep{illing+al75,solanki+pahlke88,sanchez-almeida+al89,landolfi+landi96}. Hence, the variation in the Stokes-$V$ asymmetry across magnetic elements gives information on their magnetic field and plasma flow properties. Since this information is involved, we computed synthetic Stokes-$V$ profiles of magnetic elements that occur in a three-dimensional simulation of magneto-convection and compare their asymmetry with the measured ones in order to understand their origin. ",
        "conclusions": "Spectro-polarimetric data of a quiet Sun area close to the disc centre obtained with SOT on board Hinode were analysed. We find magnetic elements of either a roundish (tube) or an elongated (sheet) shape, which show Stokes-$V$ profiles of negative area asymmetry in the centre, surrounded or framed by pixels of larger, positive area asymmetry. A comparison with results from 3-D MHD-simulations suggests that this peculiar pattern in Stokes-$V$ area asymmetry (first predicted by Steiner 1999) is due to the confined nature of the field in magnetic elements with a funnel-shaped boundary layer that gives rise to a steep gradient in field strength along the line of sight.  We also conclude that this kind of magnetic element of the internetwork is accompanied by electric current sheets. While these conclusions are evident from the comparison, we cannot exclude  the hypothesis that a suitable magnetic structuring on scales not resolved by the present observation and simulation would lead to the observed pattern in $\\delta A$."
    },
    "0711/0711.1025.txt": {
        "abstract": "We consider in detail the spectral energy distribution (SED) and multi-wavelength variability of NGC~5548.  Comparison with the SEDs of other AGNs implies that the internal reddening of NGC~5548 is E(B-V) = 0.17 mag. The extinction curve is consistent with the mean curve of other AGNs found by Gaskell \\& Benker, but inconsistent with an SMC-type reddening curve.  Because most IR emission originates exterior to the broad-line region (BLR), the SED seen by the inner BLR is different from that seen by the outer BLR and from the earth. The most likely BLR covering factor is $\\sim 40$\\% and it is not possible to get an overall BLR covering factor of less than 20\\%. This requires that the BLR is not spherically symmetric and that we are viewing through a hole. Line-continuum variability transfer functions are consistent with this geometry. The covering factor and geometry imply that near the equatorial plane the BLR covering approaches 100\\%. The spectrum seen by the outer regions of the BLR and by the torus is thus modified by the absorption in the inner BLR. This shielding solves the problem of observed BLR ionization stratification being much greater than implied by photoionization models.  The BLR obscuration also removes the problem of the torus covering factor being greater than the BLR covering factor, and gives consistency with the observed fraction of obscured AGNs. The flux reduction at the torus also reduces the problem of AGN dust-reverberation lags giving sizes smaller than the dust-sublimation radii. ",
        "introduction": "Because of the difficulty of observing in the extreme ultra-violet (EUV), the true shape of the continuum of AGNs has long been a mystery.  If emission lines arise from photoionization, then the equivalent width of the lines gives information on the number of ionizing photons (Zanstra 1931, Osterbrock 1989) and their total energy (Stoy 1933).  Early photoionization model calculations of the broad-line region (BLR) spectrum (e.g., Davidson 1972, MacAlpine 1972, Shields 1972) {\\it assumed} covering factors ($\\Omega/4\\pi$) and continuum shapes that could account for the observed strengths of the emission lines. However, the first spectrophotometry of the rest-frame UV continuum in high-redshift AGNs showed an apparent turn down of the UV spectrum (Oke 1970, 1974). {\\it IUE} satellite observations of lower-redshift AGNs confirmed this (Green et al.\\@ 1980).  Green et al.\\@ reported that the $\\alpha \\thickapprox 0.5$ spectral slope ($F_{\\nu} \\propto \\nu^{-\\alpha}$) from the optical to the UV steepened to $\\alpha \\thickapprox 2.3$ at wavelengths shorter than $\\lambda$1200. MacAlpine (1981) showed that there was a severe ``energy budget'' problem because such a continuum needed $\\Omega/4\\pi \\gtrsim 9$ to match the observed He II $\\lambda$4686 equivalent width.  The observed ionizing continuum was thus providing an order-of-magnitude fewer ionizing photons than were needed.  The problem was compounded when spectroscopy of the Lyman-limit region showed that rest-frame Lyman continuum absorption was rare (Osmer 1979). Smith et al.\\@ (1981) concluded from observations of the Lyman-limit region that $\\Omega/4\\pi \\lesssim 0.15$ and further analysis by MacAlpine (1981) showed that $\\Omega/4\\pi \\thickapprox 0.05$ was most likely.  Netzer (1985) pointed out that there is a serious energy-budget discrepancy for other lines.  The HST composite spectrum of Zheng et al.\\@ (1997) and the multi-wavelength spectra of a sample of PG quasars (Laor 1997) confirmed the steepening of the observed continuum at shorter wavelengths. Photoionization models (e.g., Shields \\& Ferland 1993; Korista et al.\\@ 1998; Goad \\& Koratkar 1998; Kaspi \\& Netzer 1999; Korista \\& Goad 2000) consistently required covering factors an order of magnitude higher than allowed by the limits on Lyman limit absorption.  An important question is whether the continuum {\\it we} see is the same as that seen by the emission lines.  The continuum could be intrinsically anisotropic, or the continuum we see could be modified by extinction from dust grains along the line of sight, or by absorption lines. MacAlpine (1981) suggested that a marked upturn in dust extinction below $\\lambda$1200 could be causing the apparent turndown in the continuum we see.  Green et al.\\@ (1980), however, felt that the turndown below $\\lambda$1200 was probably intrinsic, and that extinction by dust grains was probably not causing it because of the lack of effect on the spectrum longwards of $\\lambda$1200. McKee \\& Petrosian (1974) had previously argued against the presence of dust in AGNs because of the lack of both the $\\lambda$2175 dust feature and the curvature in the UV spectra that dust should produce.  The apparent match up between UV spectral indices and the X-ray region (Laor 1997) argued against significant extinction.  Various other solutions have been offered to the energy-budget problem. Collin-Souffrin (1986), Joly (1987), and Dumont, Collin-Souffrin, \\& Nazarova (1998) favored non-radiative heating of the BLR as the solution to the energy-balance problem. Binette et al.\\@ (1993) suggested the clouds were being heated by a diffuse continuum very close to them. Korista, Ferland, \\& Baldwin (1997) suggested the the UV-EUV SED might be double peaked. Maiolino et al.\\@ (2001c), on the other hand, suggested that the solution to the conflict between the lack of Lyman limit absorption and the high covering factors implied by photoionization models was that the BLR was not covering the source uniformly.  In this paper we investigate the energy-budget question in the well-studied AGN NGC~5548 and investigate the implications for the BLR and torus structure. NGC~5548 is an important test case for investigating the energy-budget issue (Rokaki et al.\\@ 1994; Dumont et al.\\@ 1998) since NGC~5548 has been particularly well studied at all accessible wavelengths, its variability behavior in most spectral regions is well known, and the line-continuum transfer functions, which give information on the BLR size and radial distribution, are better known than for any other AGN. We focus in this paper on the period 2449115 -- 2449130 when NGC 5548 was intensely monitored by the {\\it Hubble Space Telescope}. We adopt a distance, $R$, of 76 Mpc to NGC~5548. This is based on an H$_0 = 73$ km/s/Mpc, $\\Omega_{matter} = 0.27$, $\\Omega_{vacuum} = 0.73$ cosmology and the local velocity field model of Mould et al.\\@ (2000). ",
        "conclusions": "We have used simultaneous observations and the results of multi-wavelength variability studies to consider in detail the most probable SED of NGC~5548 during the highest state in the HST monitoring in 1993.  We have argued that when the continuum of NGC~5548 (and other AGNs) is corrected for Galactic and internal reddening, detailed photoionization modelling shows that the strengths of all the major emission lines can be explained by photoionization by the inferred AGN continuum so long as the covering factors are large ($\\sim 40$\\%). The long-standing ``energy-budget problem'' can now be seen to have been a consequence of not appreciating the significant reddening of the continua of most AGNs and erroneously assuming a spherical geometry. Because the emission-line strengths can be explained by photoionization from the observed (or extrapolated) continuum, there is no need for the additional heating source suggested by Collin-Souffrin (1986), Joly (1987), and Dumont et al.\\@ (1998).  To reconcile our high covering factor with the lack of absorption along the observer's line of sight we need a picture of the BLR which is different from the widely assumed quasi-spherical distribution of BLR clouds. Instead most of the BLR clouds (especially those far from the AGN) have a flattened distribution that extends out to the inner edge of the dusty torus. We have shown that such a model is consistent with the observed line-continuum transfer functions for NGC~5548.  In the geometry needed to explain the observed line strengths, the cloud covering fraction is close to 100\\% in the equatorial plane. We have therefore proposed a shielding model of the BLR where the inner clouds absorb much of the radiation before it reaches the outer clouds or the torus.  This solves the problem of LOC models not explaining the strong change in cloud conditions with radius. The shielding BLR picture quantitatively reproduces the ionization-dependent lags seen in NGC~5548.  The absorption by BLR gas also solves several problems with the torus.  Allowance for BLR absorption increases the torus covering factor calculated from energetics, and makes it slightly greater than the mean BLR covering factor.  This reconciles the torus covering factor with the fraction of obscured (type-2) AGNs and explains why low-ionization by the BLR is never seen.  The reduced flux seen by the torus also allows the dust to exist at smaller radii and reduces the discrepancy between the calculated dust sublimation radii and the radii of the inner edge of the torus found from reverberation mapping.  Since the line ratios, equivalent widths, and the continuum shape of NGC~5548 are similar to other AGNs (see, for example, Sergeev et al.\\@ 1999) we expect that these results can be generalized to most other AGNs."
    },
    "0711/0711.2311_arXiv.txt": {
        "abstract": "We consider the spectrum of mesons for the gauge theory dual to a supergravity configuration of intersecting D3/D7 branes \\cite{Hovdebo:2005hm}, and use the expression for the Lagrangian of the scalar mesons to compute explicitly the Lagrangian for the lightest states in the infrared limit. Assuming that the matter content of this gauge theory is part of a hidden sector, which interacts with the standard model only via gravity,   we explore the cosmological consequences of these lightest scalar mesons for a FRW universe. We show that phantom fields may appear naturally in this kind of scenarios. ",
        "introduction": "The standard Big-Bang model of cosmology  has been very successful in explaining the cosmological observations  \\ci{WMAP,SN}, however the model needs to be extended to include at least two periods of positive  acceleration of our universe, one at high temperatures denoted by Inflation and another at present time and given in terms of Dark Energy. Scalar fields are perhaps the best candidates to explain these periods of acceleration \\ci{scalar}. The effect of scalar fields in the cosmological evolution of our universe has been the object of extensive studies \\ci{scalar,Q.ax}. The equation of state of Dark Energy $w=p/\\rho$ is very close to -1 but the cosmological data  \\ci{WMAP,SN} seems to favor a $w$   smaller than -1 \\ci{DE}. This can be achieved through the use of the so called phantom fields, which are scalar fields with a negative kinetic term \\ci{phantom}, or alternatively by the inclusion of an interaction term between the Dark Energy  scalar field and some other particle \\ci{wapp,IDE,IDE.ax}. Despite their success, these models have to provide a very particular potential for the scalar fields, and in most of the cases these potentials are not deduced from considerations other than to fit the experimental data of the evolution of our universe. Further more, for the case of the phantom fields, it is in general not even clear how such a field can arise consistently within quantum field theory.  If, instead of defining the Lagrangian of the field guided  by the principle  of reproducing the cosmological data, we could obtain its main properties from a theory intended to describe phenomena beyond gravity, this would be a step forward to incorporate this evolution driving mechanism in the scope of a wider theory.  Natural places to look for these fields are supergravity, string theory, grand  unification theories or other quantum theories of gravity.  Ideally we would like to find a field which is naturally coupled to gravity,  or one of its generalizations, ruled by a defined potential, and see what are its cosmological implications.  Due to the lack of abundance of this kind of scenarios, the choice for this work will be different, and we will look for this type of field in the context of the gauge/gravity correspondence \\cite{'t Hooft:1993gx, Susskind:1994vu}, and more in particular, using the AdS/CFT correspondence \\cite{Maldacena:1997re} as it was applied to a system of D3/D7 branes in \\cite{Hovdebo:2005hm}. This analysis provides a define potential which can be use for the kind of field we are looking for and an example of a scenario where a phantom field can arise in the context of a quantum theory.  The AdS/CFT correspondence establishes a duality between string theory in its supergravity limit and a conformal field theory. The first is a theory which contains gravity, while the second is a field theory in a flat background without gravity. This correspondence has been extended to more general gauge/gravity situations, and applied successfully to obtain physical quantities proper to gauge theories in their non-perturbative limit, in the hope of reaching a better understanding of non-perturbative QCD \\cite{Herzog:2006gh, Gubser:2006bz,CasalderreySolana:2006rq,Herzog:2006se,Caceres:2006dj, Caceres:2006as,Chernicoff:2006hi,CaronHuot:2006te,PS,Mateos:2007yp}. See \\cite{Mateos:2007ay} for a resent review.  By analyzing a particular supergravity configuration in \\cite{Hovdebo:2005hm},  it was possible to determine the spectrum of mesons, that is, quark-antiquark bound states, in a ${\\cal N}$ =2 field theory with fundamental matter. With this result at hand, we compute explicitly the Lagrangian for the lightest two scalar mesons in this scenario, which happen to have the same mass. This Lagrangian is valid for strong gauge coupling in the infrared limit of the field theory, and we will use this result as a motivation to analyze the consequences of two mesonic fields, governed by the Lagrangian we find, if they were present in our universe.  Is necessary for us to mention two reasons why we cannot use this result as a solid prediction, and then argue why it is still a strong motivation for the present analysis. On the one hand, the matter content of the gauge theory, where the mesons studied here live, is not part of the standard model of particles and has not been observed to exist in our universe. On the other hand, as mention before, the mesons we are describing live in a four dimensional Minkowsky space-time, flat and without gravity, hence not suitable for cosmological evolution.  Concerning the matter content of the ${\\cal N}$=2 theory that we are working with, we have to say that it is commonly speculated that there could be in nature a hidden, or dark, sector of particles beyond those included in the standard model. The particles in this sector could have scape observation for either of many reasons, like not being in the scope of energies archived in experiments so far, or simply for being coupled to standard matter exclusively by gravitational interaction. Assuming the matter content of the theory at hand to be part of this hidden or dark sector will permit us, in this work, to perform the analysis we are interested in.  In what respects the lack of gravity in the field theory, we  are assuming that a similar field to the mesons we find, could be present in the theory that should describe the processes of elementary particles in a universe with matter. So we will work taking a variational approach where the contribution of these mesons to the action dictating the cosmological processes is given by the integral, over the space-time, of the Lagrangian we find for the fields, multiplied by the determinant of the metric for a Friedmann-Robertson-Walker universe.  In the next section we will shortly describe the main features of the supergravity configuration used in \\cite{Hovdebo:2005hm}, and how it is related to the spectrum of the scalar mesons and the Lagrangian obtained there. In section \\ref{Lagr}, we will perform the explicit calculation of the relevant terms of the Lagrangian in the infrared limit for the two lightest mesons once they have been embedded in a FRW metric. Section \\ref{potential} will be used to discus the main features of the Lagrangian found in section (\\ref{Lagr}). The possible cosmological implications will be address in section \\ref{cosmo} and we will finish with some concluding remarks in section \\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}  Starting from the spectrum of the scalar mesons in the ${\\cal N}$=2 theory dual to a supergravity configuration of D3 and D7 branes, we were able to extract the Lagrangian governing the behavior of the two lightest of these states in the infrared limit. We considered the possibility for these mesons to live in our universe and obtained an action to describe the cosmological consequences of the presence of these fields.  The resulting effective scalar fields $\\vp$ and $\\phi $ have different cosmological contributions depending on the sign of the kinetic term. For positive kinetic terms these fields oscillate around the minimum of the potential but they have different masses. They  behave as dark matter and therefore their energy densities redshift as $\\rho\\propto a^{-3}$. On the other hand, for negative kinetic terms, the  fields become phantom fields with an equation of state parameter $w<-1$. The universe expands in an accelerating way and the fields $\\vp, \\phi$ could then parameterize the Dark Energy.   Concerning the possibility of phantom fields appearing for the case analyzed here, we have to say we believe it to be a more general phenomenon for this kind of scenarios. The reason for this believe is that given the Lagrangian in \\cite{Hovdebo:2005hm}, we can see that to each order $n$, higher than two, there will be a contribution given by the product of the kinetic term of one of the mesons in the spectrum, $\\partial_\\mu\\Phi_a\\partial^\\mu\\Phi_a$, times a polynomial of homogeneous degree $n-2$ on the other fields. Considering then the Lagrangian to all orders, it would be found that at least some of the kinetic terms would appear multiplying a polynomial of the fields, and these fields could take a wide range of values. It would seam natural to expect then for the kinetic term to be negative in some region of values for the fields, giving rise to the existence of phantom fields. To make this expectation concrete, it would be necessary to consider a setting where the physical properties would make the Lagrangian of \\cite{Hovdebo:2005hm}, or another one similarly deduce, exactly summable so that the appearance of phantom fields could be precisely stated."
    },
    "0711/0711.2596_arXiv.txt": {
        "abstract": "{Most  of the observed emission lines and  continuum excess from young accreting low  mass stars (Classical T Tauri  stars -- CTTSs) take place in  the star-disk or inner disk  region.  These regions have a complex emission topology still largely unknown.} {In this  paper the magnetospheric  accretion and inner wind contributions to the observed permitted He and H near infrared (NIR) lines of the bright southern CTTS RU Lupi are investigated for the first time.} {Previous optical observations of RU Lupi showed a large H$\\alpha$ profile, due to the emission from a wind in the line wings, and a micro-jet detected in forbidden lines. We extend this analysis to NIR lines through seeing-limited high  spectral resolution spectra taken with VLT/ISAAC, and adaptive optics (AO) aided narrow-band imaging and low spectral resolution spectroscopy with VLT/NACO. Using spectro-astrometric analysis we investigate the presence of extended emission down to very low spatial scales (a few AU).} {The HeI \\lam10830 line presents a P Cygni profile whose absorption feature indicates the presence of an inner stellar wind. Moreover the spectro-astrometric analysis evidences the presence of an extended emission superimposed to the absorption feature and likely coming from the micro-jet detected in the optical. On the contrary, the origin of the Hydrogen Paschen and Brackett lines is difficult to address. We tried tentatively to explain the observed line profiles and flux ratios with both accretion and wind models showing the limits of both approaches. The lack of spectro-astrometric signal indicates that the HI emission is either compact or symmetric. Our analysis confirms the sensitivity of the HeI line to the presence of faint extended emission regions in the close proximity of the star.} {} ",
        "introduction": "Classical T~Tauri stars  (CTTSs) are young solar and  lower mass stars presenting  a   rich  emission line  spectrum  in  excess   of the continuum from the  stellar photosphere. This  emission   includes   Hydrogen lines, low energy forbidden lines and Helium lines  together  with a  continuum  excess  --  the veiling.  Optically and NIR  forbidden emission lines, such as the [\\ion{O}{i}], [\\ion{S}{ii}] and  [\\ion{Fe}{ii}] lines, have emission regions large enough  to be resolved  and are typically associated with outflowing jets \\citep[e.g.][]{bacciotti00,dougados00,pyo06}. On the other hand, the atomic permitted lines,  such as the Hydrogen lines, are more difficult to  interpret. The energy released by the accretion of matter onto the star can power the permitted line emission that is believed to arise from magnetospheric    accretion    columns    and/or   accretion    shocks \\citep[e.g.][]{muzerolle01,bouvier07}.    The fact  that  the  emission  component  of  permitted  atomic  lines  is attributed to infalling gas is supported by the presence of inverse P~Cygni profiles, and by the correlation of their luminosities  with  the accretion luminosity, derived from the veiling \\citep{muzerolle98b,muzerolle98c}. On   the   other  hand,   the blueshifted absorption features, when present, are an indicator of high  velocity  inner  winds \\citep[e.g.][]{edwards87,najita00}. Indeed, several recent spectroscopic studies suggest that  the emission of the permitted HI and He lines  can actually include a contribution from both inflowing and outflowing gas \\citep[e.g.][]{beristain01,folha01,takami01,whelan04,edwards06}. This has motivated hybrid models where the Hydrogen emission is due to both the magnetospheric accretion and a wind \\citep{alencar05,kurosawa06}.   The fact that the collimated, spatially resolved jets, traced at large distance from the source by optical, near-infrared forbidden lines, are generated in circumstellar regions (less than 10 AU from the source) is predicted by theoretical models \\citep{anderson03,ferreira06} and observationally proved by the frequent detection of blueshifted absorption features in a number of strong permitted lines (H$\\alpha$, Na, D, CaII, MgII, \\citealt{najita00}), formed close to the star. Signatures of outflowing gas revealed by absorption features can also be searched in the profile of the HeI \\lam10830 line \\citep{takami02, edwards03,dupree05}. \\citet{edwards06} analysed a sample of 31 T Tauri stars and showed that the HeI line is much more sensitive than the usually used \\Ha\\, line to trace inner winds in absorption (in 71\\% of the sources was detected an absorption feature in the HeI line against only 10\\% of detections in the \\Ha\\, line). \\citet{beristain01} showed that the diagnostic power of permitted Helium lines in tracing inner winds, lies in their high excitation potential (20-50 eV), that restricts the line formation to a region of either high temperature or close proximity to a source of ionising radiation.\\\\  In order to investigate the reliability of such scenarios, we have observed the circumstellar region of the T Tauri star RU Lupi in the NIR wavelength range.  This source  is one of  the most active  and well known T  Tauri stars \\citep{lamzin96,stempels02,herczeg05}  and  it  is  a good  target  to investigate the  inner winds  and accretion flows  down to  very small spatial  scales.  RU  Lupi is  located at  a distance  of only  140 pc \\citep{dezeeuw99},  allowing us  to reach  small spatial scales with high angular resolution instrumentation and/or specialised techniques of data analysis. In particular, NACO at the  VLT was used to take high  angular resolution images and low spectral resolution spectra. With VLT/ISAAC seeing limited and high spectral resolution spectra were obtained. A spectro-astrometric analysis was applied. Spectro-astrometry is a powerful technique to obtain positional information on the region originating the emission in the components of a line profile even in seeing-limited conditions. \\citet{takami01}, for example, using spectro-astrometry, found  that the blueshifted and redshifted  wings of the H$\\alpha$ line  emitted by RU Lupi show a spatially symmetric offsets with respect to the source, in the same direction of the jet traced, on higher spatial scales, by the forbidden \\oi\\,  and \\s\\,  lines.  \\citet{whelan04} detected  a similar extended emission  in the wings  of the \\pab\\,  line in three  T Tauri sources (DG Tau, V536 Aql, and LkH$\\alpha$ 321).   In this  paper we  present the results obtained from such a study and, in particular, a detailed analysis about the origin of the  Paschen and Brackett HI lines and the HeI \\lam10830 line. The paper is structured as follows: in Sect.~\\ref{sect:data_reduction} we present the different set of observations and the data reduction process; in Sect.~\\ref{sect:results} the observed line profiles and fluxes are analysed; moreover we present the concepts underlying the spectro-astrometry technique and the obtained position spectra; finally in   Sect.~\\ref{sect:discussion} we discuss the origin of permitted H and He lines by means of a model for wind and accretion, and of the results obtained with spectro-astrometry. In the Conclusions (Sect.~\\ref{sect:conclusions}) we summarize our findings. ",
        "conclusions": "\\label{sect:conclusions}  In this paper we analysed infrared spectra and images of the T Tauri star RU Lupi, in order to find observational constraints for the physics of the circumstellar region. This source was previously observed in the optical wavelength range by other authors, that reported about a strong accretion activity  and the presence of a micro-jet detected in the optical forbidden lines and in the H$\\alpha$ line \\citep{takami01}. We have observed for the first time RU Lupi in the NIR with high angular and spectral resolution. In this paper we report about these observations and, in particular, we investigate the origin of the permitted H and He NIR lines through the analysis of the line profiles and fluxes, and using the spectro-astrometry technique which allows us to constrain the emission region of each velocity component of the detected lines.  The problem of the origin of HI lines was widely debated in the last decades. Many different models were proposed to reproduce their profiles and fluxes: wind models \\citep[e.g., ][]{hartmann90}, accretion models \\citep[e.g., ][]{hartmann94,muzerolle98a,muzerolle01} and also hybrid models accounting for the emission from both the accretion columns and the wind \\citep[e.g., ][]{kurosawa06}. These analyses, however, focused mainly on Balmer lines observed in the optical range. The Paschen and Brackett lines detected in our NIR spectra of RU Lupi present broad, slightly blueshifted and nearly symmetric profiles, which are difficult to reproduce with either a wind or an accretion model. This was already noted, in general, by \\citet{folha01}, \\citet{whelan04}, \\citet{nisini04}. We tentatively used a toy model of a spherical envelope of partially ionised Hydrogen with a  wind or an accretion velocity profile \\citep{nisini04} to constrain the HI line ratios and fluxes. This analysis showed that neither a spherical wind nor spherical accretion can reproduce the Brackett and Paschen decrement and fluxes.  Our spectro-astrometric analysis did not highlight any extended emission in the Paschen and Brackett lines, suggesting that the region emitting the HI NIR lines is very compact ($<$1.3 AU) and/or symmetric. On the other hand, the HeI \\lam10830 line blueshifted absorption feature clearly indicates the presence of an inner wind. The wide range of velocities covered by the absorption feature favours the geometry of a spherical wind emerging from the stellar surface rather than a disk wind as found in many TTS by \\citet{edwards06}. In addition, the spectro-astrometric analysis highlighted the presence of an emission superimposed to the absorption feature that extends up to at least $\\sim$3 AU from the source and covers the same velocity range of the H$\\alpha$ extended emission found by \\citet{takami01}. We suggest that this emission comes from the blue lobe of the micro-jet detected in the optical lines by \\citet{takami01}. Interestingly, the HeI showed to be more sensitive than the HI NIR lines to faint extended emission because the absorption feature increases the contrast between the emission from the extended region and the continuum from the source. This confirms the potential of the HeI \\lam10830 line to investigate the complex geometry of the inner part of the circumstellar region of CTTS, where both the accretion and the ejection processes take place \\citep[see also ][]{edwards06}. On the base of these results, future high angular resolution observations of the permitted H and He lines in CTTS can give useful hints to the understanding of the accretion/ejection mechanisms and the development of theoretical models accounting for the emission from both the outflowing and the accreting gas."
    },
    "0711/0711.0888_arXiv.txt": {
        "abstract": "We present the first results of a combined VLT VIMOS integral-field unit and \\textit{Hubble Space Telescope} (\\textit{HST})/ACS study of the early-type lens galaxy SDSS\\,J2321$-$097 at $z=0.0819$, extending kinematic studies to a look-back time of 1\\,Gyr.  This system, discovered in the Sloan Lens ACS Survey, has been observed as part of a VLT Large Programme with the goal of obtaining two-dimensional stellar kinematics of 17 early-type galaxies to $z \\approx 0.35$ and Keck spectroscopy of an additional dozen lens systems.  Bayesian modelling of both the surface brightness distribution of the lensed source and the two-dimensional measurements of velocity and velocity dispersion has allowed us, under the only assumptions of axisymmetry and a two-integral stellar distribution function (DF) for the lens galaxy, to dissect this galaxy in three dimensions and break the classical mass--anisotropy, mass-sheet and inclination--oblateness degeneracies. Our main results are that the galaxy (i) has a total density profile well described by a single power-law $\\rho \\propto r^{-\\gamma'}$ with $\\gamma' = 2.06^{+0.03}_{-0.06}$; (ii) is a very slow rotator (specific stellar angular momentum parameter $\\lambda_{\\mathrm{R}} = 0.075$); (iii) shows only mild anisotropy ($\\delta \\approx 0.15$); and (iv) has a dark-matter contribution of $\\sim 30$~per cent inside the effective radius. Our first results from this large combined imaging and spectroscopic effort with the VLT, Keck and \\textit{HST} show that the structure of massive early-type galaxies beyond the local Universe can now be studied in great detail using the combination of stellar kinematics and gravitational lensing. Extending these studies to look-back times where evolutionary effects become measurable holds great promise for the understanding of formation and evolution of early-type galaxies. ",
        "introduction": "\\label{sec:Introduction}  Within the hierarchical galaxy formation scenario, early-type galaxies are assumed to be the end-products of major mergers with additional accretion of smaller galaxies (e.g.~\\citealt{Burkert-Naab2004} for a review).  As such the study of their structure, formation and evolution is an essential step in understanding galaxy formation and the standard $\\Lambda$ cold dark matter ($\\Lambda$CDM) paradigm, in which the hierarchical galaxy formation model has its foundations. Early-type galaxies are observed to follow tight scaling relations between their stellar population, dark matter, kinematic and black hole properties, the origins of which are still not well understood \\citep[e.g.][]{Djorgovski-Davis1987, Dressler1987, Magorrian1998, Bolton2007}.  Considerable observational progress has been made in the last decades in our understanding of the relative contributions of baryonic (mostly stellar), dark matter and black hole constituents of early-type galaxies through stellar dynamical tracers and X-ray studies (e.g.\\ \\citealt{Fabbiano1989}; \\citealt{Mould1990}; \\citealt*{Saglia1992}; \\citealt{Bertin1994}; \\citealt*{Franx1994}; \\citealt{Carollo1995}; \\citealt{Arnaboldi1996}; \\citealt{Rix1997}; \\citealt{Matsushita1998}; \\citealt{Loewenstein-White1999}; \\citealt{Gerhard2001}; \\citealt{Seljak2002}; \\citealt*{Borriello2003}; \\citealt{Romanowsky2003}).  More recently, the SAURON collaboration \\citep{deZeeuw2002} has extensively studied a large and uniform sample of early-type galaxies in the local Universe (i.e.\\ $z \\la 0.01$), combining SAURON integral-field spectroscopy (IFS) on the William Herschel Telescope with high-resolution \\textit{Hubble Space Telescope} (\\textit{HST}) imaging to obtain a three-dimensional picture of these galaxies in terms of their mass and kinematic structure \\citep{Emsellem2004, Emsellem2007, Cappellari2006, Cappellari2007}.  \\begin{figure*} \\centering \\resizebox{0.375\\hsize}{!}{\\includegraphics{FIG/HST-435.ps}} \\hspace{0.025\\hsize} \\resizebox{0.375\\hsize}{!}{\\includegraphics{FIG/HST-814.ps}} \\caption{\\textit{HST}/ACS images of \\galaxy, taken through the \\textit{F435W} (left-hand panel) and \\textit{F814W} (right-hand panel) filters. The images are single exposures with an integration time of $420\\,\\mathrm{s}$; the pixel scale is $0.05\\,\\mathrm{arcsec}$. A full description of these observations and the data-reduction process is given in \\citet{Bolton2006}.} \\label{fig:HST_images} \\end{figure*}  Despite making a great leap forward, the complex modelling of the SAURON systems and the unknown dark matter content and total mass of these galaxies still require some assumptions to be made about their gravitational potential [e.g.\\ constant stellar mass-to-light ratio ($M/L$) from stellar population models] and inclination. These issues can be partly overcome in discy systems where the inclination can be determined from the photometry or possibly from dust lanes.  However, systems with lower ellipticity can show stronger degeneracies \\citep[e.g.][]{Cappellari2007}. Also, near the effective radius dark matter becomes a major contributor to the stellar kinematics \\citep{Gerhard2001, Treu-Koopmans2004} and it is not yet clear what impact its neglect has, even though it is estimated to be not too severe \\citep{Cappellari2006}.  At higher redshifts, the extraction of detailed kinematic information from early-type galaxies becomes increasingly more difficult because of the decrease in apparent size and cosmological surface brightness dimming.  However, galaxies at higher redshifts are much more likely to act as gravitational lenses on background galaxies \\citep*[e.g.][]{Turner1984}.  The additional information obtained from gravitational lensing on, for example, the enclosed mass and the density profile has been shown to break some of the classical degeneracies in stellar kinematics and lensing, namely the mass--anisotropy and the mass-sheet degeneracies, respectively \\citep[e.g.][hereafter BK07]{Treu-Koopmans2002a, Koopmans2003, Koopmans-Treu2003, Treu-Koopmans2004, Koopmans2006, Barnabe-Koopmans2007}.  In order to use gravitational lensing and stellar dynamics in a systematic way, the Lenses Structure \\& Dynamics (LSD) Survey was started in 2001, combining \\textit{HST} imaging data on lens systems with stellar kinematics obtained with Keck \\citep{Koopmans-Treu2002, Koopmans-Treu2003, Treu-Koopmans2002a, Treu-Koopmans2002b, Treu-Koopmans2003, Treu-Koopmans2004}. This project and its results, predominantly at $z\\approx 0.5\\!\\!-\\!\\!1.0$, have shown the validity of this methodology and its effectiveness at higher redshifts. These high-redshift systems, however, are relatively rare and hard to follow up, which has limited studies to the use of luminosity-weighted stellar velocity dispersions \\citep[e.g.][]{Koopmans-Treu2002,Treu-Koopmans2002b} or long-slit spectroscopy in a few large apertures along the major axes \\citep[e.g.][]{Koopmans-Treu2003, Treu-Koopmans2004}.  The Sloan Lens ACS Survey \\citep[SLACS,][]{Bolton2005,Bolton2006} was begun to extend the sample of lens galaxies suitable for joint lensing and dynamical studies. SLACS lens candidates were selected from the SDSS Luminous Red Galaxy sample \\citep{Eisenstein2001} and a quiescent subsample [defined by $\\mathit{EW}(\\mathrm{H}\\alpha)<1.5\\,$\\AA] of the MAIN SDSS galaxy sample \\citep{Strauss2002} by an inspection of their SDSS fibre spectroscopy \\citep{Bolton2004}. Galaxies whose spectra presented emission lines at a higher redshift than the redshift of the galaxy itself were observed in an \\textit{HST} snapshot programme to check whether the source of the emission lines was a gravitationally lensed background galaxy. To date, SLACS has confirmed around 80~lens systems (Bolton et al., in preparation). In contrast to most earlier lens surveys which drew their candidate systems from catalogues of potential sources \\citep[e.g.~quasi-stellar objects or radio sources as in the CLASS survey,][]{Browne2003}, SLACS is a \\emph{lens}-selected survey and as such preferentially finds systems with bright lens galaxies but comparatively faint sources. It is therefore an ideal starting point for detailed investigations into the properties of the lens galaxies.  In this paper we introduce a follow-up project to SLACS which aims at combining the lensing information with detailed two-dimensional kinematic information obtained with the VIMOS integral-field unit (IFU) mounted on the VLT.  The full sample comprises 17~systems. In this paper we demonstrate the available data and the analysis methodology on one system, {\\galaxy}. The lens in this system is an early-type galaxy at $z=0.0819$. The SDSS spectrum shows additional [\\ion{O}{ii}] and H$\\beta$ emission from a gravitationally lensed blue galaxy at $z_{\\mathrm{s}} = 0.5342$ \\citep{Bolton2006}. In \\textit{HST} imaging (Fig.~\\ref{fig:HST_images}) the lensed source is shown to form an almost complete Einstein ring of radius $1.68\\,\\mathrm{arcsec}$, corresponding to an Einstein radius of $\\REin=2.6\\,\\mathrm{kpc}$ at the lens redshift. The projected mass enclosed within the Einstein radius is $M(<\\REin) = 1.3\\times10^{11}\\,M_{\\sun}$.  Details on this system are listed in Table~\\ref{tab:basic_data}.  In Section~\\ref{sec:Observations}, we describe the VIMOS and \\textit{HST} observations that form the basis of this analysis. The IFU data are presented in Section~\\ref{sec:data_analysis}, with a description of the data reduction in Section~\\ref{ssec:data_reduction} and of the extraction of two-dimensional kinematic maps in Section~\\ref{sec:kinematic_analysis}.  In Section~\\ref{sec:inferences} we determine several commonly used parameters from the kinematic data alone, in a way that is directly comparable to earlier studies, for example from the SAURON collaboration. We apply the joint lensing and dynamics analysis developed by BK07 in Section~\\ref{sec:analysis}. In Section~\\ref{sec:degeneracies} we present a physical argument that explains why gravitational lensing and stellar dynamics can break the degeneracy between oblateness and inclination. We discuss our results and conclude in Section~\\ref{sec:discussion}. We use a concordance $\\Lambda$CDM model throughout this paper, described by $\\Omega_{\\mathrm{M}}=0.3$, $\\Omega_{\\Lambda} = 0.7$ and $H_{0} = 100\\,h\\,\\mathrm{km\\,s^{-1}\\,Mpc^{-1}}$ with $h=0.7$ unless stated otherwise. At the redshift of \\galaxy, $1\\,\\mathrm{arcsec}$ corresponds to $1.1\\,h^{-1}\\,\\mathrm{kpc} = 1.5\\,\\mathrm{kpc}$.  \\vfill  \\begin{table} \\centering \\begin{minipage}{0.85\\hsize} \\caption{Basic data for \\galaxy. The data are taken from \\citet{Treu2006err} and \\citet{Koopmans2006}.} \\begin{tabular}{@{}l@{\\hspace{12em}}r@{$\\;$}l} \\hline $\\alpha_{\\mathrm{J2000}}$ & $23^{\\mathrm{h}}21^{\\mathrm{m}}20\\fs93$ & \\\\ $\\delta_{\\mathrm{J2000}}$ & $-09\\degr39\\arcmin10\\farcs2$ & \\\\ $z_{\\mathrm{l}}$ & $0.0819$ & \\\\ $z_{\\mathrm{s}}$ & $0.5342$ & \\\\ $\\sigma_{\\mathrm{ap}}$ & $(236\\pm7)$&$\\mathrm{km\\,s^{-1}}$ \\\\ $\\sigma_{\\mathrm{c}}$ & $(245\\pm7)$&$\\mathrm{km\\,s^{-1}}$ \\\\ $\\REin$ & $1.68\\,\\mathrm{arcsec} = 2.59$ & $\\mathrm{kpc}$  \\\\ $R_{\\mathrm{eff},B}$ & $(8.47\\pm0.11)$&$\\mathrm{kpc}$ \\\\ $M_{B}$ & $-21.72\\pm0.05$ & \\\\ $R_{\\mathrm{eff},V}$ & $(7.93\\pm0.07)$&$\\mathrm{kpc}$ \\\\ $M_{V}$ & $-22.59\\pm0.05$ & \\\\ $\\qiso$   & $0.77$ & \\\\ $\\vartheta_{\\mathrm{PA},\\star}$ & $126\\fdg5$ & \\\\ \\hline \\end{tabular} \\textit{Notes:} $z_{\\mathrm{l}}$ and $z_{\\mathrm{s}}$ are the redshifts of \\galaxy\\ and the gravitationally lensed background galaxy, respectively. $\\sigma_{\\mathrm{ap}}$ is the velocity dispersion measured from the 3-arcsec-diameter SDSS fibre, $\\sigma_{\\mathrm{c}}$ is the derived central velocity dispersion \\citep{Treu2006}. $\\REin$ is the Einstein radius. $R_{\\mathrm{eff}}$ and $M$ are the effective radii and absolute magnitudes determined by fitting de Vaucouleurs profiles to the $B$- and $V$-band ACS images \\citep{Treu2006}. $\\qiso$ and $\\vartheta_{\\mathrm{PA},\\star}$ are isophotal axial ratio and position angle of the major axis, respectively \\citep{Koopmans2006}. \\end{minipage} \\label{tab:basic_data} \\end{table} ",
        "conclusions": "\\label{sec:discussion}  In this paper we have presented the first results from an integral-field spectroscopic survey of early-type lens galaxies from SLACS. The combination of integral-field spectroscopy from VIMOS/IFU mounted on the VLT with high-resolution imaging from \\textit{HST}/ACS has enabled us to conduct the first in-depth study of the structure of a luminous elliptical galaxy beyond the local Universe, \\galaxy\\ at $z=0.0819$.  We have applied a new analysis method that combines the kinematic and lensing information in a fully self-consistent way and have shown how this combination breaks some of the degeneracies that limit the separate application of these two methods.  The galaxy that we have studied here turns out to be a fairly ordinary elliptical with properties similar to those of local galaxies of comparable luminosity, such as those studied by the SAURON collaboration \\citep{Emsellem2007, Cappellari2007}.  \\galaxy\\ is a slow rotator in the classification of \\citet{Emsellem2007} with an angular momentum parameter of $\\lambda_{\\mathrm{R}} = 0.075$. The velocity dispersion map is flat to the limit where we were able to measure the kinematic parameters reliably. Using the updated estimator of \\citet{Binney2005} for the ratio between systematic and random velocities, $v/\\sigma_{v}$, that fully exploits the information in integral-field spectroscopic data, we have shown that the ellipticity of the stellar distribution in \\galaxy\\ is due to anisotropy of the velocity distribution rather than rotation, again in line with local galaxies of comparable luminosity \\citep{Cappellari2007}.  We have modelled the galaxy by making use of the {\\dynlen} algorithm which self-consistently combines gravitational lensing and stellar dynamics under the assumption of axisymmetry and a two-integral DF. We adopted for the system the total gravitational potential generated by an axisymmetric power-law mass density distribution of logarithmic slope $\\gamma'$ and axial ratio $q$. The best-fitting model given the data is obtained in the framework of Bayesian statistics by maximizing the evidence merit function. The results for the best-fitting model are summarized as follows. \\begin{enumerate} \\item The logarithmic slope of the total density is $\\gamma' = 2.06^{+0.03}_{-0.06}$ (the error is given within the 68~per cent confidence interval), which is very close to (and consistent with) an isothermal density distribution. \\item The axial ratio of the total density distribution is $q = 0.74^{+0.02}_{-0.05}$. Since in our approach mass is not required to follow light, this $q$ does not have to coincide with the (average) axial ratio of the luminous distribution, which is calculated from the reconstructed DF, giving $\\qlight = 0.85$ for the best model. \\item The inclination angle of the galaxy is $i = 67\\fdg8^{+1.1}_{-7.8}$. The small error bar shows that inclination can be well determined in combination with lensing data. \\item The `maximum bulge' approach, that is, the rescaling of the circularized stellar density profile which maximizes the contribution of the luminous component to the total density profile, prescribes a stellar mass $M_{\\mathrm{eff}} \\simeq 2.0 \\times 10^{11} M_{\\sun}$ inside the effective radius, which corresponds to an average $M/L$ of $5.2\\,(M/L)_{\\sun,B}$. The total mass enclosed in the same region is approximately $2.9 \\times 10^{11} M_{\\sun}$: the non-visible matter, therefore, accounts for about 30 per cent of the total mass within that three-dimensional radius. \\item The local $\\vphi/\\bar{\\sigma}$ ratio on the meridional plane confirms that {\\galaxy} as a whole is a slow rotator, with the random motions becoming less predominant compared to rotation only in the very central regions. The best model yields a global anisotropy parameter $\\delta = 0.14$, fully consistent with the value obtained directly from the data, showing that the galaxy is close to an isotropic rotator.  The other global anisotropy parameters have values $\\beta = 0$ [as a consequence of having assumed a two-integral DF $f(E, L_{z})$] and $\\gamma = -0.32$, indicating a mild tangential anisotropy. \\end{enumerate}  These results are in good agreement with the analysis of {\\galaxy} by \\citet{Koopmans2006}, which we here extend in a more rigorous way. In particular, the results confirm the essentially isothermal profile of the mass density distribution, which appears to be a defining characteristic of early-type galaxies.  The best-fitting model is consistent with a dark matter fraction of 30~per cent within 10\\,kpc (approximately corresponding to the unprojected effective radius), similar to what \\citet{Cappellari2006} determine for the SAURON sample of early-type galaxies by making use of three-integral Schwarzschild dynamical models under the caveat that light traces mass. However, it seems plausible that a constant-M/L model could still reproduce the observed kinematics of {\\galaxy} if one allowed for a more flexible three-integral dynamical model rather than the two-integral model considered in this work. The relative difficulty in unambiguously discriminating between the constant-$M/L$ model and the power-law model is also a consequence of the fact that for the specific lens system studied here, the mass within the Einstein radius is to a large extent dominated by the stellar mass, so that the difference between the two models does not show up dramatically in the lens model. The combined analysis of more distant objects, however, is expected to provide more unequivocal results in this respect. In general, for galaxies at higher redshift the ratio between the Einstein radius and the effective radius becomes larger (see e.g.~the SLACS galaxy sample in \\citealt{Koopmans2006}) and any deviation from a model for which mass follows light would become significantly more prominent, since at least the surface brightness slope would be much steeper than what is allowed by the lensing data.  The dynamical structure of {\\galaxy} (i.e.\\ anisotropy and $\\vphi/\\bar{\\sigma}$ map) is also in good agreement with what is found by \\citet{Cappellari2007} for the most massive ellipticals of the SAURON sample.  The analysis has shown that the combination of gravitational lensing and stellar dynamics is a powerful method which allows the dissection in three dimensions of an elliptical galaxy (assumed to be well described as a two-integral axisymmetric dynamical system), breaking to a significant extent the classical degeneracies between inclination and flattening as well as between mass and anisotropy. The way the degeneracy between inclination and oblateness is overcome can be understood within a simplified physical picture: $q$ and $i$ are coupled in the projected potential which enters in the description of gravitational lensing; when it comes to the dynamics, however, the surface brightness weighted integral along the minor axis over the line-of-sight stellar velocity dispersion is not a function of the inclination, since, due to the properties of the two-integral DF, the intersection of the velocity dispersion tensor with the $(v_{R}, v_{z})$ plane is always a circle, that is, $\\sigma_{R}^{2} = \\sigma_{z}^{2}$ for each position in the galaxy (see Section~\\ref{sec:degeneracies}).  We have shown that the method, in its implementation as the {\\dynlen} algorithm, is robust enough to make use of observational data in order to recover the non-linear parameters which characterize the total gravitational potential and the geometry of the system (i.e.\\ inclination, positional angle and lens centre) with relatively tight error bars (the confidence intervals shown in Table~\\ref{tab:stat}). This first application therefore shows promise for the future study of the other SLACS systems at higher redshift.  In forthcoming papers in this series we will extend this work to the entire sample of 17~SLACS lenses with VLT VIMOS IFS, covering a range of lens galaxy morphology, mass and redshift ($z = 0.08 - 0.35$). The VLT sample will be complemented by a further 13 lenses for which we have obtained long-slit spectroscopy at the Low Resolution Imager and Spectrograph \\citep[LRIS,][]{Oke1995} on the Keck-I telescope. Several slit positions -- aligned with the major axis and offset along the minor axis -- have been obtained for each system in the Keck sample, thus effectively producing two-dimensional kinematic information across most of the lens galaxy."
    },
    "0711/0711.2305_arXiv.txt": {
        "abstract": "{{}{Nearly half the stellar mass of present-day spiral galaxies has formed since $z\\!=\\!1$, and galaxy kinematics is an ideal tool to identify the underlying mechanisms responsible for the galaxy mass assembly since that epoch.} {Here, we present the first results of the ESO large program, ``IMAGES'', which aims at obtaining robust measurements of the kinematics of distant galaxies using the multi-IFU mode of GIRAFFE on the VLT. 3D spectroscopy is essential to robustly measure the often distorted kinematics of distant galaxies (e.g., Flores et al.\\ 2006).  We derive the velocity fields and $\\sigma$-maps of 36 galaxies at $0.4\\!<\\!z\\!<\\!0.75$ from the kinematics of the \\oii\\ emission line doublet, and generate a robust technique to identify the nature of the velocity fields based on the pixels of the highest signal-to-noise ratios (S/N). } {Combining these observations with those of Flores et al., we have gathered a unique sample of 63 velocity fields of emission line galaxies ($W_{0}(\\mbox{{\\sc [Oii]}})\\ge15$\\,\\AA) at $z\\!=\\!0.4$\\,--\\,0.75, which are a representative subsample of the population of $M_{\\rm stellar}\\!\\ge\\!1.5\\!\\times\\!10^{10}M_{\\sun}$ emission line galaxies in this redshift range, and are largely unaffected by cosmic variance. Taking into account all galaxies -with or without emission lines- in that redshift range, we find that at least 41$\\,\\pm\\,$7\\% of them have anomalous kinematics, i.e., they are not dynamically relaxed. This includes 26$\\,\\pm\\,$7\\% of distant galaxies with complex kinematics, i.e., they are not simply pressure or rotationally supported.}  {Our result implies that galaxy kinematics are among the most rapidly evolving properties, because locally, only a few percent of the galaxies in this mass range have complex kinematics. It is well-established that galaxies undergoing a merger have complex large-scale motions and thus are likely responsible for the strong evolution of the galaxy kinematics that we observe.} ",
        "introduction": "The resolved 3D kinematics of distant galaxies are a powerful tracer of the major processes governing star-formation and galaxy evolution in the early universe such as merging, accretion, and hydrodynamic feedback related to star-formation and active galactic nucleus (e.g., Barnes \\& Hernquist 1996; Barton et al.\\ 2000; Dressler 2004). Thus, robustly measuring the internal kinematics of galaxies in the distant universe plays a crucial role for our growing understanding of how galaxies formed and evolved.  Over the last decade, great efforts have been made to study the properties of galaxies in the distant universe (at $z\\!\\sim\\!1$), revealing a strong evolution with cosmic time. For instance, the cosmic star formation rate (SFR) has declined by a factor $\\sim\\!10$ from $z\\!\\sim\\!1$ to the present (Lilly et al.\\ 1996; Madau et al.\\ 1996; Hammer et al.\\ 1997; Cowie et al.\\ 1999; Flores et al.\\ 1999). Such a strong evolution of cosmic SFR is consolidated by subsequent works, e.g., Haarsma et al.\\ (2000), Wilson et al.\\ (2002). Although the conclusions are made from different data, they are consistent within a factor of 3 (Hopkins 2004).  Heavens et al.\\ (2004) suggest that the cosmic SFR may have reached its peak as late as $z\\!\\sim\\!0.6$. Overall, about half of the stellar mass in intermediate-mass galaxies was formed since $z\\!=\\!1$, mostly in luminous infrared galaxies (Hammer et al.\\ 2005).  Galaxy interactions and merging may be mechanisms that played a significantly larger role for star-formation in the distant universe than today.  Le F{\\`e}vre et al.\\ (2000) found that the merger rate in the distant universe was about a factor of 10 times higher than at low redshift (see also Conselice et al.\\ 2003; Bell et al.\\ 2006; Lotz et al.\\ 2006).  The high merger rate detected in the distant universe raises a challenge to the standard scenario of disk formation (Hammer et al.\\ 2007).  If major mergers generate the ellipticals inevitably, we would find a large fraction of elliptical galaxies rather than about $\\sim$70\\% of spiral galaxies among the intermediate-mass galaxies in the local universe.  Similarly, the fraction of luminous compact blue galaxies (LCBGs) increases with redshift by about an order of magnitude out to $z\\!\\sim\\!1$ (Werk et al.\\ 2004; Rawat et al.\\ 2007).  LCBGs may be the progenitors of local spheroidal or irregular galaxies at low redshift (e.g., Koo et al.\\ 1995; Guzman 1999), or of the bulges of massive spirals (Hammer et al.\\ 2001; Noeske et al.\\ 2006).  3D spectroscopy of the internal kinematics of LCBGs suggests that they are likely merger remnants ({\\\"O}stlin et al. 2001; Puech et al.\\ 2006a).  Strong evolution as a function of cosmic time has also been claimed for the Tully-Fisher relationship (TFR, Tully \\& Fisher 1977; Giovanelli et al.\\ 1997), which relates the luminosity and the rotation velocity of disk galaxies. Out to $z\\!\\sim\\!1$, the B-band TFR has been found to have evolved by $\\sim\\!0.2$\\,--\\,$2$ mag (e.g., Portinari \\& Sommer-Larsen 2007 and references therein). This brightening of the B-band TFR can be explained by the enhanced star-formation rates at higher redshifts (Ferreras \\& Silk 2001; Ferreras et al.\\ 2004), but is still a matter of debate. Conselice et al.\\ (2005) do not find significant evolution in either the stellar mass or K-band TFR's slope or zero point. However, the most striking evolution of the TFR is provided by its large scatter at high redshifts (Conselice et al.\\ 2005), which may be related to the disturbed kinematics of distant galaxies (e.g., Kannappan \\& Barton 2004).   The rapid time decrease of cosmic SFR, the role of merging in the early evolution of galaxies, and the possible evolution of the TFR are only examples of why measuring the kinematics of distant galaxies precisely and robustly is a {\\it sine qua non} for studying galaxy evolution. However, this is often beyond what can be achieved with classical long-slit spectroscopy. The morphologies and kinematics of distant galaxies are often complex, and their small sizes make it very difficult to precisely position and align the slit. Both limitations can be overcome with integral-field spectroscopy, although the method is relatively complex and time-consuming.  Flores et al.\\ (2006) presented the first study of a statistically meaningful sample of 35 intermediate-mass galaxies at z=0.4-0.7, using the integral-field multi-object spectrograph GIRAFFE on the ESO-VLT. They defined a classification scheme to distinguish between rotation and kinematic perturbances, which may stem from interactions and mergers, from the 3D kinematics and high-resolution HST imaging. Intriguingly, they find that the large scatter of distant TFR shown in previous studies is due to non-relaxed systems while the pure rotational disks exhibit a TFR that is similarly tight as that of local spirals. Here, we present another sample of 36 galaxies with very similar selection criteria, to enlarge the total sample size and put the conclusions on statistically more robust grounds. This is the first of a series of publications related to the ESO-VLT large program IMAGES, which aims at studying the evolutionary sequence of galaxies over the last 8 Gyrs (see Ravikumar et al.\\ 2007 for more details).  The paper is organized as follows. In Sect.~2 we describe the observations and the sample selection. The methodology to describe and classify the distant galaxy kinematics is shown in Sect.~3, as well as a detailed description of the 36 observed velocity fields (VFs). Sections 4 and 5 include the discussion and the conclusion. In this paper, we adopt the Concordance cosmological parameters of $H_0\\!=\\!70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M\\!=\\!0.3$ and $\\Omega_\\Lambda\\!=\\!0.7$.   {\\scriptsize \\begin{table}[!t] \\caption{Journal of observations.} \\begin{center} \\begin{tabular}{lccr}\\hline Run ID & Setup & Exposure (hr) \\\\\\hline 174.B-0328(A) &   L04 & 10   \\\\ 073.A-0209(A) &   L05 & 4.5  \\\\ 174.B-0328(A) &   L05 & 10.6 \\\\ 174.B-0328(E) &   L05 & 10.4 \\\\\\hline \\end{tabular} \\end{center} \\label{tbobservation} \\end{table} } ",
        "conclusions": "We have been able to measure the VFs of 36 galaxies at \\zzz\\ using deep exposures of the multiplex integral-field spectrograph GIRAFFE at the VLT in the multi-IFU mode, measuring the kinematics of the spatially and spectrally well resolved \\oii\\ emission line doublet. In combination with a similar study by Flores et al.\\ (2006), we have a relatively large and representative sample of 63 galaxies with $M_{\\rm stellar} \\ge 1.5\\times 10^{10}M_{\\sun}$. Thus, our results are representative for the population of intermediate-mass galaxies in this redshift range, and it cannot be affected by cosmic variance. To date, this is the only existing representative sample of distant galaxies with measured VFs.  We confirm and consolidate the results of Flores et al.\\ (2006), that a significant fraction of intermediate mass galaxies had perturbed or complex kinematics 5 Gyrs ago. Our method to classify the kinematics of the galaxies is particularly robust. It attributes a large weight to the velocity dispersion in the central region of the galaxy, where the S/N are the highest. Even if we assume that all quiescent galaxies at \\zzz\\ had well ordered VFs, we find that 41$\\,\\pm\\,$7\\% of the galaxies are not kinematically relaxed, including 26$\\,\\pm\\,$7\\% of galaxies that show complex kinematics. Undoubtedly, galaxy kinematics are evolving very rapidly, since most present-day galaxies in the same mass range are likely to have ordered VFs.  This result may be combined with the fact that anomalous VFs are responsible for most of the large observed dispersion of both the Tully-Fisher and the $j_{\\rm disk}$--$V_{\\rm max}$ relationships (see Flores et al.\\ 2006 and Puech et al.\\ 2007). It suggests a random walk evolution of galaxies related to a high fraction of merging events, including major mergers (Puech et al.\\ 2007). Mergers may indeed reproduce all the peculiar kinematics at \\zzz, as well as being responsible for the dispersion of fundamental relations. Other mechanisms, such as in-fall of high velocity gas clouds, gas outflows or bars may also contribute to the observed evolution in the kinematics. To understand their influences, and moreover the underlying mechanisms that activate them, requires detail analyzes of individual objects and as a prerequisite, a full model of the significance of the GIRAFFE measurements. If major merging is the main mechanism responsible for the large fraction of complex VFs, then this implies that, since $z\\!=\\!1$, from 50\\% to 70\\% of intermediate mass galaxies have experienced a major merger. This is quantitatively in good agreement with the spiral rebuilding scenario as proposed by Hammer et al.\\ (2005)."
    },
    "0711/0711.4715_arXiv.txt": {
        "abstract": "%  Here we describe the Red MSX Source (RMS) survey which is the largest, systematic, galaxy-wide search for massive young stellar objects (MYSOs) yet undertaken. Mid-IR bright point sources from the MSX satellite survey have been followed-up with ground-based radio, millimetre, and infrared observations to identify the contaminating sources and characterise the MYSOs and UCHII regions. With the initial classification now complete the distribution of sources in the galaxy will be discussed, as well as some programmes being developed to exploit our sample. ",
        "introduction": "Massive stars ($M$$>$8~\\msun, $L$$>$10$^4$~\\lsun) play a fundamental role in many areas of astrophysics. They are the principal source of UV radiation and heavy elements in galaxies, and are responsible for injecting huge amounts of kinetic energy into the ISM through powerful molecular outflows, strong stellar winds and supernova explosions. The momentum imparted through these processes provides an important source of mixing and turbulence within the ISM. They are also thought to play a key role in regulating star formation, either by disrupting molecular clouds before stars have been able to form, or constructively through the expansion of their HII regions (e.g., collect and collapse) or the propagation of strong shocks into their surroundings (e.g., radiatively driven implosion). Massive stars have an enormous influence not only on the physical structure and chemistry of their local environment, but also on the structure and evolution of their host galaxies.  Despite the importance of massive stars the processes involved in their formation and the early stages of their evolution are still poorly understood. Massive stars reach the main sequence while still deeply embedded within their natal molecular clouds, and therefore their entire pre-main sequence formation and evolution is hidden behind high levels of extinction. They are extremely rare and as a consequence are generally located farther away than regions of low-mass star formation. The large distances are compounded by the fact that massive stars are through to form exclusively in clusters and limited spatial resolution makes it difficult to identify and attribute derived quantities to individual sources. Furthermore, the evolution of massive stars is extremely rapid, which means that key stages in their early evolution are short lived. Due to these observational difficulties, until relatively recently, the only catalogue of massive young stellar objects (MYSOs) had been limited to 30 or so serendipitously detected sources (\\citealt{henning1984}) most of which are nearby. ",
        "conclusions": "The RMS survey aims to produce a large, unbiased sample of massive young stellar objects (MYSOs) located throughout the Galaxy. We have colour-selected approximately 2000 MYSO candidates from the MSX and 2MASS point source catalogues that have colours similar to known MYSOs. We have almost completed a multi-wavelength programme of follow-up observations to distinguish between genuine MYSOs and types of embedded or dust enshrouded objects that have similar mid-infrared colours that contaminate our sample. So far using these observations we have unambiguously classified $\\sim$90\\% of our sample, approximately half of which are identified as either YSOs and UCHII regions. The luminosities of a sub-sample of these sources, for which good IRAS data are available, confirms that we are detecting a significant number of genuine MYSOs. When complete, we estimate our database will contain a well selected sample of $\\sim$500 bona fide MYSOs and a further 500--600 UCHII regions, along with complementary multi-wavelength observations.  Now the classification is almost complete we have begun a number of observational programmes to exploit our sample and to address many open questions concerned with massive star formation. These include: high resolution observations of potentially triggered regions to investigate different triggering mechanisms; chemical surveys using millimetre lines and IR absorption features to investigate the chemistry and possible chemical evolution of MYSOs; investigation of molecular outflows associated with MYSOs; and high resolution spectroscopy of accretion disks. All of these programmes are being conducted on sub-samples selected as a function of luminosity, distance and location. The whole sample and all of the observational results of our multi-wavelength campaign are available at www.ast.leeds.ac.uk/RMS."
    },
    "0711/0711.1076_arXiv.txt": {
        "abstract": "{In recent papers convincing evidence has been presented for chemical stratification in Ap star atmospheres, and surface abundance maps have been shown to correlate with the magnetic field direction. Radiatively driven diffusion, which is known to be sensitive to the magnetic field strength and direction, is among the processes responsible for these inhomogeneities.} {Here we explore the hypothesis that equilibrium stratifications -- such that the diffusive particle flux is close to zero throughout the atmosphere -- can in a number of cases explain the observed abundance maps and vertical distributions of the various elements.} {An iterative scheme adjusts the abundances in such a way as to achieve either zero particle flux or zero effective acceleration throughout the atmosphere, taking strength and direction of the magnetic field into account.} {The investigation of equilibrium stratifications in stellar atmospheres with temperatures from 8500 to 12000\\,K and fields up to 10\\,kG reveals considerable variations in the vertical distribution of the 5 elements studied (Mg, Si, Ca, Ti, Fe), often with zones of large over- or under-abundances and with indications of other competing processes (such as mass loss). Horizontal magnetic fields can be very efficient in helping the accumulation of elements in higher layers.} {A comparison between our calculations and the vertical abundance profiles and surface maps derived by magnetic Doppler imaging reveals that equilibrium stratifications are in a number of cases consistent with the main trends inferred from observed spectra. However, it is not clear whether such equilibrium solutions will ever be reached during the evolution of an Ap star.} ",
        "introduction": "Atomic diffusion in stars, when efficient enough to overcome mixing processes, leads to inhomogeneous abundance distributions of chemical elements. This is probably what happens in the atmospheres of upper main-sequence, chemically peculiar (CP) stars which exhibit a wide variety of strong abundance anomalies, since the outer layers of these stars can be considered stable enough to allow element diffusion to take place. A considerable number of papers, starting with Michaud (\\cite{mic70}), have examined to what extent the diffusion model is able to explain the observed anomalies. Often, but not always, these anomalies appear to be quite well correlated with the respective radiative accelerations (frequently the leading contribution to the diffusion velocity) of each element. However, the CP phenomenon involves so many complex processes that a direct comparison of observed apparent abundance anomalies and calculated radiative accelerations will not be sufficient for fully describing the build-up of abundance peculiarities. On the one hand, abundance determinations in the past did not take stratification of the chemical elements into account; on the other, these stratifications are built up by a time-dependent, non-linear diffusion process which is quite sensitive to magnetic fields and to macroscopic motions (residual turbulence or stellar wind for instance) and depends not only on radiative accelerations.  Despite the fact that full modelling of atmospheres of CP stars is still out of reach, notable progress has been made in the field of diffusion. A first study addressing the special behaviour of silicon in magnetic atmospheres was carried out by Vauclair et al. (\\cite{vhp79}), followed by the quantitative modelling of Si stratification by Alecian \\& Vauclair (\\cite{alv81}). A theoretical prediction of manganese accumulation in hot CP stars was proposed by Alecian \\& Michaud (\\cite{alm81}), and a first attempt at detecting such a stratification was carried out by Alecian (\\cite{ale82}) in the HgMn star $\\upsilon$\\,Her using a method based on the curve of growth of Mn resonance lines. Abundances of iron peak elements in several HgMn stars were analysed in detail by Smith \\& Dworetsky (\\cite{smi93}). In their study, these authors used what were state of the art methods at that time, applying schematic corrections for the chemical stratifications. Within the framework of this approximate treatment, they found that their results were in excellent agreement with the predictions of the diffusion model. A detailed study of the stratifications of several metals in the magnetic star 53 Cam, assuming the presence of a stellar wind, is due to Babel \\& Michaud (\\cite{bam91}) and to Babel (\\cite{bam92}). Their main conclusion for that particular star was that diffusion alone cannot account for the observations and that more sophisticated models have to be developed, including mass loss confined by magnetic fields. These early studies were limited by technical constraints, such as insufficient computing power, and by the lack of atomic and observational data. Therefore no firm conclusions could be reached concerning the chemical stratifications produced by diffusion processes. Fortunately, the situation has improved drastically. Thanks to high performance detectors, inhomogeneous element distributions appear to be established beyond reasonable doubt; see e.g. Kochukhov et al. (\\cite{koc04}) for horizontal distributions in the magnetic atmosphere of 53 Cam and Kochukhov et al. (\\cite{koc06}) for vertical distributions in the Ap star HD\\,133792. Significant progress has also been made in modelling, since self-consistent atmospheric models for non-magnetic stars, including abundance stratifications compatible with the amount of elements which can be supported by radiative forces, are in an advanced stage of development (Hui-Bon-Hoa et al. \\cite{hbh02}).  In the present paper, we attempt to advance one further step on the long path towards the complete modelling of the migration process of chemical elements, by looking for an equilibrium solution to the stratification of metals in magnetic atmospheres. We rely on the physics and methods presented in previous papers, where we have computed the Zeeman amplification of radiative accelerations in detail (Alecian \\& Stift \\cite{als04}), and studied diffusion velocities in magnetic atmospheres (Alecian \\& Stift, \\cite{als06}). The present work still assumes LTE and the temperature/pressure structure of the atmosphere is computed with solar abundances (ATLAS9, Kurucz, \\cite{kur93}). The CARAT code is  used in the improved version discussed by Alecian \\& Stift (\\cite{als06}). It carries out full opacity sampling of Zeeman split spectral lines, determines the radiative flux by solving the polarised equation of radiative transfer, includes radiative accelerations due to bound-free transitions, and takes the redistribution of momentum among ions into account. In Section \\ref{sec:buildup} we discuss theoretical aspects of the element stratification process and in Section \\ref{sec:numerics} we present the numerical method used to obtain solutions for equilibrium stratifications. In Section \\ref{sec:resdisc} we discuss our results in view of recent observations. ",
        "conclusions": "\\label{seq:conclusion}  In this work we present the first detailed numerical calculations of element stratifications due to atomic diffusion in magnetic atmospheres. This study addresses the abundance anomalies observed in Ap stars, either magnetic or non-magnetic (HgMn stars). Radiative accelerations are based on full opacity sampling of Zeeman split spectral lines, and diffusion velocities are obtained by the methods described in Alecian \\& Stift (\\cite{als06}). The stratifications derived in the present study correspond to equilibrium solutions: one looks for abundance stratifications such that diffusion velocities are close to zero everywhere in the atmosphere.  Because of the high computational cost, we restricted our study to a few metals and stellar atmospheres. We want to explore those cases where our results can be compared with observed abundance maps and stratifications. We did not try to use atmospheric models corresponding perfectly to those derived from the observations of 53\\,Cam and of HD\\,133792, but instead used standard solar abundance ATLAS9 models that are close enough to allow a meaningful discussion. It emerges from this comparison that, in several cases, equilibrium solutions are consistent with stratifications reconstructed from observed Stokes spectra by means of magnetic Doppler imaging. However, significant differences also exist. Some of these differences may be partly due to possible shortcomings in the inversion of the observational material, but one has to keep in mind that our modelling approach, with its search of equilibrium stratifications, is more likely to be responsible. Indeed, equilibrium solutions can be very different from the stratifications encountered in real stars, which result from non-linear, time-dependent processes, with a competition between several physical processes (atomic diffusion, inhomogeneous mass-loss, complex magnetic geometries, NLTE effects, turbulence, etc.). The present study constitutes one more step toward the detailed modelling of magnetic atmospheres, but for the first time in this quest, the results we obtain can be compared with observations. Time-dependent atomic diffusion is the next step, which we shall address in the near future."
    },
    "0711/0711.4465_arXiv.txt": {
        "abstract": "% I outline a method for estimating astrophysical parameters (APs) from multidimensional data.  It is a supervised method based on matching observed data (e.g.\\ a spectrum) to a grid of pre-labelled templates. However, unlike standard machine learning methods such as ANNs, SVMs or k-nn, this algorithm {\\em explicitly} uses domain information to better weight each data dimension in the estimation.  Specifically, it uses the sensitivity of each measured variable to each AP to perform a local, iterative interpolation of the grid. It avoids both the non-uniqueness problem of global regression as well as the grid resolution limitation of nearest neighbours. ",
        "introduction": "Consider the problem of estimating the astrophysical parameters (APs) of a star from its spectrum using a grid of pre-labelled spectra. Let ${\\vec p} = \\{p_i\\}$, $i=1 \\ldots I$ be the data vector (spectrum or multiband photometry) and ${\\vec \\phi} = \\{\\phi_j\\}$, $j=1 \\ldots J$ be the AP vector (e.g.\\ \\teff, \\logg, etc.).  Standard approaches involve performing a global regression on the grid to infer the mapping $\\vec \\phi = g(\\vec p)$, using, for example, a artificial neural network (ANN) (e.g.\\ Bailer-Jones et al.\\ 1998) or a support vector machine (SVM) (e.g.\\ Tsalmantza et al.\\ 2006).  Although these methods meet with reasonable success, they have problems when it comes to estimating multiple APs, in particular if some APs have a relatively weak signature (as is the case with \\logg\\ and [Fe/H]: compare the vertical scales in Fig.~\\ref{formod}). To overcome this we should weight the variables according to their sensitivity with respect to the APs of interest. In principle, ANNs and SVMs implicitly learn this weighting from the data, but this is difficult with many noisy variables. Furthermore, a global regression approach is strictly flawed, because while the photon counts in a band varies uniquely with the APs, the converse is not true (Fig.~\\ref{formod}).  The global regression is trying to solve an inverse problem and the lack of uniqueness could lead to a poor fit.  This degeneracy problem is exacerbated at low spectral resolution and by noise.   \\begin{figure}[t] \\epsscale{0.80} \\plotone{O9.3_1.eps} \\caption{Variation of photon counts with \\logteff\\ (top) and \\logg\\ (bottom) in three filters (bands). A $J$-dimensional fit to each band (independently) is a forward model. This is a true function (unlike a fit to the inverse).} \\label{formod} \\end{figure} ",
        "conclusions": "The algorithm currently performs slightly worse than one of the best generic regression algorithms available (an SVM), yet this is not bad considering that (a) it is in an early stage of development, and (b) the results were obtained using a (suboptimal) {\\em linear} forward model.  Unlike ANNs, the dimensionality of the algorithm's fitting depends on the number of APs ($J$), not the number of data dimensions ($I$), so it should scale well to typical spectral problems ($J$ is a few, $I$ is a few thousand).  Moreover, the method has the ability to detect and report multiple solutions which arise from degeneracies in the data (see Fig.~\\ref{formod})."
    },
    "0711/0711.1240_arXiv.txt": {
        "abstract": "{ In supersymmetric extensions of the Standard Model, the lightest neutralino, the gravitino, and the axino can appear as the lightest supersymmetric particle and as such provide a compelling explanation of the non-baryonic dark matter in our Universe. For each of these dark matter candidates, I review the present status of primordial production mechanisms, cosmological constraints, and prospects of experimental identification. \\PACS{ {95.35.+d}{Dark matter}   \\and {12.60.Jv}{Supersymmetric models}   \\and {04.65.+e}{Supergravity} } % } % ",
        "introduction": "\\label{intro}  Numerous astrophysical and cosmological considerations point to the existence of non-baryonic dark matter in our Universe~\\cite{Bergstrom:2000pn,Bertone:2004pz}. In fact, based on observations of supernovae, galaxy clusters, and the cosmic microwave background (CMB), we believe today that our Universe is flat with about 76\\%, 20\\%, and 4\\% of the critical energy density $\\rho_c$ provided in the form of dark energy, non-baryonic dark matter, and baryons, respectively~\\cite{Spergel:2006hy,Yao:2006px}. A nominal ``$3\\sigma$'' range% \\footnote{Note that the nominal ``$3\\sigma$'' range is derived assuming a restrictive six-parameter ``vanilla'' model. A larger range is possible---even with additional data from other cosmological probes---if the fit is performed in the context of a more general model that includes other physically motivated parameters such as a nonzero neutrino mass~\\cite{Hamann:2006pf}. Thereby, the range $0.094 < \\Omega_{\\CDM} h^2 < 0.136$ has been obtained in Ref.~\\cite{Hamann:2006pf}.} of the dark matter density $\\Omega_{\\CDM}=\\rho_{\\CDM}/\\rho_c$ can be inferred from measurements of the CMB anisotropies by the Wilkinson Micro\\-wave An\\-iso\\-tropy Probe (WMAP) satellite~\\cite{Spergel:2006hy} \\begin{equation} \\Omega_{\\CDM}^{3\\sigma}h^2=0.105^{+0.021}_{-0.030} \\label{Eq:OmegaDM} \\end{equation} with $h=0.73^{+0.04}_{-0.03}$ denoting the Hubble constant in units of $100\\,\\km\\,\\Mpc^{-1}\\seconds^{-1}$.  Relying on the pieces of evidence, we think that a particle physics candidate for dark matter has to be electrically neutral, color neutral,% \\footnote{A colored dark matter candidate is disfavored by severe limits from searches for anomalous heavy nuclei~\\cite{Yao:2006px}.} and stable or have a lifetime $\\tau_{\\CDM}$ that is not much smaller than the age of the Universe today $t_0 \\simeq 14~\\mathrm{Gyr}$. Moreover, the species providing the dominant contribution to $\\Omega_{\\CDM}$ have to be sufficiently slow to allow for structure formation. For example, since the neutrinos of the Standard Model are too light, $\\sum_i m_{\\nu_i} \\lesssim \\Order(1~\\eV)$~\\cite{Lesgourgues:2006nd}, they were too fast at early times. Accordingly, they are classified as hot dark matter which can constitute only a minor fraction of $\\Omega_{\\CDM}$ since otherwise structure formation cannot be understood~\\cite{Hannestad:2007dd}. Thus, the observationally inferred dark matter density can be considered as evidence for physics beyond the Standard Model.  Supersymmetric (SUSY) extensions of the Standard Model are an appealing concept because of their remarkable properties, for example, with respect to gauge coupling unification, the hierarchy problem, and the embedding of gravity~\\cite{Wess:1992cp,Nilles:1983ge,Haber:1984rc,Martin:1997ns,Drees:2004jm,Baer:2006rs}. As superpartners of the Standard Model particles, new particles appear including fields that are electrically neutral and color neutral. Since they have not been detected at particle accelerators, these sparticles must be heavy or extremely weakly interacting.  Because of the non-observation of reactions that violate lepton number $L$ or baryon number $B$, it is often assumed---as also in this review---that SUSY theories respect the multiplicative quantum number \\begin{align} \\mathrm{R}=(-1)^{3B+L+2S} \\ , \\end{align} known as R-parity, with $S$ denoting the spin. Since Standard Model particles and superpartners carry respectively even (+1) and odd (-1) R-parity, its conservation implies that superpartners can only be produced or annihilated in pairs and that the lightest supersymmetric particle (LSP) cannot decay even if it is heavier than most (or all) of the Standard Model particles.% \\footnote{While R-parity conservation is assumed in this review, its violation is a realistic option; see, e.g., \\cite{Dreiner:1997uz,Allanach:2007vi,Takayama:2000uz,Buchmuller:2007ui,Ibarra:2007jz}.} An electrically neutral and color neutral LSP can thus be a compelling dark matter candidate. For the lightest neutralino, the gravitino, and the axino, which are well-motivated LSP candidates, this is shown below. For each scenario, I will address implications for cosmology and experimental prospects. Note that the discussion of gravitino/axino dark matter in Sects.~\\ref{sec:GravitinoDM} and~\\ref{sec:AxinoDM} will be more extensive than the one of neutralino dark matter in Sect.~\\ref{sec:NeutalinoDM}, for which numerous excellent reviews exist such as~\\cite{Jungman:1995df,Drees:2004jm,Olive:2007hm,Bertone:2007ki}. ",
        "conclusions": "\\label{sec:Conclusion} \\begin{table*} \\caption{Supersymmetric dark matter candidates, their identity, and key properties. With the listed production mechanisms, $\\Omega_{\\LSP}=\\Omega_{\\CDM}$ is possible. The respective production leads typically to a cold, warm, or hot dark matter component as indicated. Quantities marked with `(?)' seem to be unaccessible in light of the current understanding of cosmological constraints within a standard thermal history.} \\label{tab:SUSYDMCandidates} \\begin{tabular}{lllllll} \\hline\\noalign{\\smallskip} LSP & identity & mass & interactions & production & constraints & experiments\\\\ \\noalign{\\smallskip}\\hline\\noalign{\\smallskip} $\\widetilde{\\chi}^0_1$ & lightest neutralino & $\\Order(100\\,\\GeV)$ & g, g', $y_i$ & therm.~relic & $\\leftarrow$ cold & indirect searches \\vspace{0.1cm} \\\\ & (spin 1/2)& & weak & & & direct searches \\vspace{0.1cm} \\\\ & mixture of& & {\\tiny $M_{\\mathrm{W}}\\sim 100~\\GeV$} & & & collider searches \\vspace{0.1cm} \\\\ & $\\Bino$, $\\Wino$, $\\HiggsinoUp$, $\\HiggsinoDown$& & & & & \\\\ \\noalign{\\smallskip}\\hline\\noalign{\\smallskip} $\\widetilde{G}$ & gravitino & eV--TeV & $(p/\\MPl)^{n}$ & therm.~prod. & $\\leftarrow$ cold & $\\stau$ prod. at colliders \\vspace{0.1cm} \\\\ & (spin 3/2)& & extremely weak & NLSP decay & $\\leftarrow$ warm & + $\\stau$ collection \\vspace{0.1cm} \\\\ & superpartner& & {\\tiny $\\MPl = 2.4\\!\\times\\! 10^{18}\\,\\GeV$} & & & + $\\stau$ decay analysis \\vspace{0.1cm} \\\\ & of the graviton& & & & BBN& $\\hookrightarrow$ $\\mgravitino$, $\\MPl$ (?) \\\\ \\noalign{\\smallskip}\\hline\\noalign{\\smallskip} $\\widetilde{a}$ & axino & eV--GeV & $(p/f_a)^{n}$ & therm.~relic. & $\\leftarrow$ hot/warm & $\\stau$ prod. at colliders \\vspace{0.1cm} \\\\ & (spin 1/2)& & extremely weak & therm.~prod. & $\\leftarrow$ cold/warm & + $\\stau$ collection \\vspace{0.1cm} \\\\ & superpartner& & {\\tiny $f_a\\gtrsim 10^{9}\\,\\GeV$} & NLSP decay & $\\leftarrow$ warm/hot& + $\\stau$ decay analysis \\vspace{0.1cm} \\\\ & of the axion& & & & BBN& $\\hookrightarrow$ $m_{\\ax}$ (?), $f_a$ \\\\ \\noalign{\\smallskip}\\hline \\end{tabular} \\end{table*}  Dark matter is strong evidence for physics beyond the Standard Model. Extending the Standard Model with SUSY, an electrically neutral and color neutral LSP becomes a dark matter candidate for conserved R-parity. I have shown that the neutralino $\\neutralino$, the gravitino $\\gravitino$, and the axino $\\axino$ can be the LSP and as such explain the non-baryonic dark matter in our Universe.  The neutralino $\\neutralino$ is already part of the MSSM which provides a solution of the hierarchy problem and allows for gauge coupling unification. Being the superpartner of the graviton and the gauge field associated with supergravity, the gravitino $\\gravitino$ is equally well motivated with a mass $\\mgr$ that reflects the SUSY breaking scale. As the superpartner of the axion, also the axino $\\axino$ appears naturally once the strong CP problem is solved with the Peccei--Quinn mechanism in a SUSY setting.  While mass values and interactions can be very different for the $\\neutralino$, $\\gravitino$, and $\\axino$, I have illustrated for each of these LSP candidates that natural regions in the parameter space exist in which $\\Omega_{\\LSP}=\\Omega_{\\CDM}$. These regions are limited by bounds from electroweak precision observables, B-physics observables, Higgs and sparticle searches at LEP, and by BBN constraints. The constraints from $\\Omega_{\\CDM}$ and BBN also imply serious upper limits on the reheating temperature after inflation $\\TR$ which can be relevant for models of inflation and baryogenesis.  Most promising are the experimental prospects in the case of the $\\neutralino$ LSP. Being a WIMP, the $\\neutralino$ LSP should be accessible in direct and indirect dark matter searches. Indeed, first hints might have already been found in the EGRET data~\\cite{deBoer:2005bd,Elsaesser:2004ap}. With ongoing indirect searches, the increasing sensitivity of direct searches, and the advent of the LHC at which $\\neutralino$ dark matter could be produced, we will be able to test whether these hints are indeed the first evidence for the existence of SUSY dark matter. While an excess in missing transverse energy is expected to be the first evidence for SUSY at the LHC already within the next three years, the identification of the $\\neutralino$ being the LSP will require the reconstruction of the SUSY model realized in nature. If superparticles are within the kinematical reach, precision studies at the ILC will be crucial for this endeavor.  In the $\\gravitino/\\axino$ LSP scenarios with conserved R-parity, no dark matter signal should appear in direct or indirect searches. However, since an electrically charged LOSP such as the $\\stau$ is viable in the $\\gravitino/\\axino$ LSP scenarios, (quasi-) stable $\\stau$'s might occur as muon-like particles instead of an excess in missing transverse energy. Indeed, an excess of (quasi-) stable $\\stau$'s could appear as an alternative first evidence for SUSY at the LHC in the next three years. Because of the severe upper limits on the abundance of stable charged particles~\\cite{Yao:2006px}, one would then expect that the $\\stau$ is the NLSP that decays eventually into the $\\gravitino/\\axino$ LSP or that R-parity is broken. A distinction between these scenarios will require the analysis of the $\\stau$ decays. For this challenge, the ILC with its tunable beam energy seems crucial~\\cite{Hamaguchi:2004df,Feng:2004yi,Martyn:2006as,Hamaguchi:2006vu,Martyn:2007mj}.  Table~\\ref{tab:SUSYDMCandidates} presents an overview of the SUSY dark matter candidates discussed in this review. As the LSP, each of them---the lightest neutralino $\\neutralino$, the gravitino $\\gravitino$, or the axino $\\axino$---could provide $\\Omega_{\\CDM}$ and could be produced and identified at colliders in the near future.  \\bigskip  I would like to thank the organizers of SUSY 2007 for inviting me to an exciting and stimulating conference. I am grateful to A.~Brandenburg, L.~Covi, A.~Freitas, K.~Hamaguchi, G.~Panotopoulos, T.~Plehn, J.~Pradler, L.~Rosz\\-kowski, S.~Schilling, N.~Tajuddin, Y.~Y.~Y.~Wong, D.~Wyler, and M.~Zagermann for valuable discussions and collaborations on the topics covered in this review."
    },
    "0711/0711.1289_arXiv.txt": {
        "abstract": "We present the properties of the discrete X-ray sources detected in our monitoring program of the `typical' elliptical galaxy, NGC 3379, observed with {\\em Chandra} ACIS-S in five separate pointings, resulting in a co-added exposure of 324-ks. From this deep observation, 132 sources have been detected within the region overlapped by all observations, 98 of which lie within the $D_{25}$ ellipse of the galaxy. These 132 sources range in \\LX\\ from 6$\\times 10^{35}$ \\ergps\\ (with 3$\\sigma$ upper limit $\\le$4$\\times 10^{36}$ \\ergps) to $\\sim2~\\times 10^{39}$\\ergps, including one source with \\LX$>1~\\times 10^{39}$\\ergps, which has been classified as a ULX. From optical data, 10 X-ray sources have been determined to be coincident with a globular cluster, these sources tend to have high X-ray luminosity, with three of these sources exhibiting \\LX$> 1\\times10^{38}$\\ergps. From X-ray source photometry, it has been determined that the majority of the 132 sources that have well constrained colors, have values that are consistent with typical LMXB spectra. Additionally to this, a sub-population of 10 sources has been found to exhibit very hard spectra and it is expected that most of these sources are absorbed background AGN. There are 64 sources in this population that exhibit long-term variability, indicating that they are accreting compact objects. 5 of these sources have been identified as transient candidates, with a further 3 possible transients. Spectral variations have also been identified in the majority of the source population, where a diverse range of variability has been identified, indicating that there are many different source classes located within this galaxy. ",
        "introduction": "Low-mass X-ray binaries (LMXBs) are the only direct fossil evidence of the formation and evolution of binary stars in the old  stellar populations of early-type galaxies. First discovered in the Milky Way (see Giacconi 1974), these binaries are composed of a compact accretor, neutron star or black hole, and a late-type stellar donor. The origin and evolution of Galactic LMXBs has been the subject of much discussion, centered on two main evolution paths (see Grindlay 1984; review by Verbunt \\& Van den Heuvel 1995): the evolution of primordial binary systems in the stellar field, or formation and evolution in Globular Cluster (GC).  With the advent of {\\em Chandra} (Weisskopf \\etal\\ 2000), many LMXB populations have been discovered in early-type galaxies (see review Fabbiano 2006), and the same evolutionary themes (field or GC formation and evolution) have again surfaced, supported and stimulated by a considerably larger and growing body of data. These {\\em Chandra} observations  have provided important results on the spatial distributions and X-ray luminosity functions of  LMXB populations (e.g., Kim \\& Fabbiano 2004; Gilfanov 2004), on their average spectra, and on their association with GCs (e.g. Angelini \\etal\\ 2001; Kundu \\etal\\ 2002; White \\etal\\ 2002; Sivakoff \\etal\\ 2006). However,  most of the {\\em Chandra} observations of LMXB systems so far consist of fairly shallow individual snapshots for each observed galaxy, with limiting luminosity ($\\sim0.3-$8.0~keV) of a few $10^{37}$~erg~s$^{-1}$. These data give us information on the high luminosity LMXB sources, but do not cover  the typical luminosity range of the  well studied LMXB populations of the Galaxy and M31, which extends down a decade towards dimmer luminosities. Moreover, apart from rare exceptions, these observations do not have the time sampling that would permit variability studies and the identification of X-ray transients. Although, from this limited sample of multi-epoch observations with higher limiting luminosities, already a variety of different variability behaviours of LMXBs have been observed (e.g. Irwin 2006; Sivakoff \\etal\\ 2007). Both multi-epoch observations and low luminosity thresholds are important aspects of the observational characteristics of Galactic LMXBs and are needed for constraining the evolution of these populations (e.g., Piro \\& Bildsten 2002, Bildsten \\& Deloye 2004). For these reasons we proposed (and were awarded) a very large program of monitoring observations of nearby elliptical galaxies with {\\em Chandra} ACIS-S3.   NGC3379, in the nearby poor group Leo (D=10.6~Mpc (Tonry \\etal\\ 2001)) was chosen for this study because  is a relatively isolated unperturbed `typical' elliptical galaxy, with an old stellar population (age of 9.3~Gyr, Terlevich \\& Forbes 2002) and a poor globular cluster system ($S_{GC}=1.3 \\pm 0.7$, Harris 1991; where $S_{GC}$ = No.GC$\\times 10^{(0.4(M_V+15))}$). These characteristics make NGC3379 ideal for exploring the evolution of LMXB from primordial field binaries.  Observationally, NGC3379 is an ideal target for LMXB population studies, because of its proximity, resulting in a resolution of $\\sim30$~pc with {\\em Chandra}, and the lack of a prominent hot gaseous halo, demonstrated by a previous short {\\em Chandra} observation (David \\etal\\ 2005). These characteristics optimize the detection of fainter LMXBs, and minimize source confusion; because of its angular diameter ($D_{25} = 4.6$ arcmin, RC3), NGC3379 is entirely contained in the ACIS-S3 CCD chip, and is not affected by the degradation of the {\\it Chandra} PSF at large radii.  Here we publish the catalog of LMXBs with their properties resulting from the entire observational campaign of NGC3379 (four observations between January 2006 and January 2007, for a total of $\\sim300$~ks), which has been recently completed, and including the first 30~ks observation taken in 2001, from the {\\it Chandra} archive. In the companion paper (Fabbiano \\etal\\ 2007) we summarize our results relative to GC-LMXB associations and discuss their implications for our understanding of LMXB formation.  In addition to these two papers, further highlights from the X-ray binary population of NGC 3379 will be presented in Brassington \\etal\\ (2008, in prep), where the properties of the transient population of NGC 3379 will be presented. Forthcoming papers will also present: the properties of the ULX, the X-ray luminosity function and the diffuse emission of the galaxy, as well as the properties of the nuclear source and the intensity and spectral variability of the luminous X-ray binary population. Preliminary results from the first {\\it Chandra} observations of our program were reported in Kim \\etal\\ (2006) and Fabbiano \\etal\\ (2006).  This paper is organized as follows: \\S 2. details the observational program and describes the data analysis methods and results, including pipeline processing of the data, source detection, astrometry and matching of sources from the different observations, X-ray photometry and overall population results, variability analysis and optical counterpart matching (GC and background objects); \\S 3. is the source catalog, including the results from the individual observations and the co-added data; \\S 4 presents the discussion of the properties of the sources catalog; \\S 5 summarizes the conclusions of this work. ",
        "conclusions": "We have presented a source catalog and variability atlas resulting from our monitoring deep observations of the nearby elliptical NGC 3379 with \\CHANDRA\\ ACIS-S. Our results can be summarized as follows:  \\begin{itemize}  \\item{132 X-ray point sources have been detected within NGC 3379, ranging in luminosity from 6$\\times 10^{35}$\\ergps\\  (with 3$\\sigma$ upper limit $\\le$4$\\times 10^{36}$ \\ergps) to $\\sim2~\\times 10^{39}$\\ergps, with 98 of these sources residing within the $D_{25}$ ellipse of the galaxy.}  \\item{Only one ULX has been identified within this galaxy, with a galactocentric radius of 6.7\\arcsec, and a peak luminosity of \\LX$\\sim3\\times10^{39}$\\ergps.}  \\item{Ten globular clusters have been identified to be coincident with X-ray sources, all of which lie within the $D_{25}$ ellipse of the galaxy. These GC-LMXB associations tend to have high X-ray luminosities, with three of these sources exhibiting \\LX$>~1\\times10^{38}$\\ergps.}  \\item{From source photometry, it has been determined that the majority of source with well constrained colors have values that are consistent with a typical LMXB spectrum of $\\Gamma=1.5-2.0$, with no intrinsic absorption.}  \\item{A sub-population of 10 sources has been found to exhibit very hard spectra. These objects are distributed uniformly in the sky and it is expected that most of these sources are absorbed background AGN.}  \\item{64 sources, 48\\% of the X-ray source population, have been found to exhibit some type of long-term variability, which clearly identifies them as accreting compact objects. 5 of these variable sources have been identified as transient candidates, with a further 3 identified as possible transients.}  \\item{Spectral variability analysis has revealed that the sources within NGC 3379 exhibit a range of variability patterns, where both high/soft$-$low/hard and low/soft$-$high/hard spectral transitions have been observed, as well as sources that vary in luminosity, but exhibit no spectral variation, indicating that there are many different source classes within this galaxy.}  \\end{itemize}  In addition to this catalog paper, our companion paper Fabbiano \\etal\\ (2007) discusses the dearth of low-luminosity GC-LMXBs within this galaxy. Further highlights from the X-ray binary population of NGC 3379 will also be presented in Brassington \\etal\\ (2008, in prep), where the properties of the transient population of NGC 3379 will be presented.  Forthcoming papers will also present: the properties of the ULX, the X-ray luminosity function and the diffuse emission of the galaxy, as well as the properties of the nuclear source and the intensity and spectral variability of the luminous X-ray binary population. The results from this deep observation will then be compared to the X-ray source catalog of the old, GC rich elliptical galaxies NGC 4278, which has also recently been the subject of a deep \\CHANDRA\\ observation.   \\appendix"
    },
    "0711/0711.2655_arXiv.txt": {
        "abstract": "We use the large cosmological Millennium Simulation (MS) to construct the first all-sky maps of the lensing potential and the deflection angle, aiming at gravitational lensing of the CMB, with the goal of properly including small-scale non-linearities and non-Gaussianity. Exploiting the Born approximation, we implement a map-making procedure based on direct ray-tracing through the gravitational potential of the MS. We stack the simulation box in redshift shells up to $z\\sim 11$, producing continuous all-sky maps with arcminute angular resolution. A randomization scheme avoids repetition of structures along the line of sight and structures larger than the MS box size are added to supply the missing contribution of large-scale (LS) structures to the lensing signal.  The angular power spectra of the projected lensing potential and the deflection-angle modulus agree quite well with semi-analytic estimates on scales down to a few arcminutes, while we find a slight excess of power on small scales, which we interpret as being due to non-linear clustering in the MS.  Our map-making procedure, combined with the LS adding technique, is ideally suited for studying lensing of CMB anisotropies, for analyzing cross-correlations with foreground structures, or other secondary CMB anisotropies such as the Rees-Sciama effect. ",
        "introduction": "\\label{i} \\renewcommand{\\thefootnote}{\\fnsymbol{footnote}} \\footnotetext[1]{E-mail: carbone@ieec.uab.es} \\footnotetext[2]{E-mail: volker@MPA-Garching.MPG.DE} \\footnotetext[3]{E-mail: bacci@sissa.it} \\footnotetext[4]{E-mail: mbartelmann@ita.uni-heidelberg.de} \\footnotetext[5]{E-mail: sabino.matarrese@pd.infn.it} \\renewcommand{\\thefootnote}{\\arabic{footnote}}   The cosmic microwave background (CMB) is characterized both by primary anisotropies, imprinted at the last scattering surface, and by secondary anisotropies caused along the way to us by density inhomogeneities and re-scatterings on electrons that are freed during the epoch of reionization, and heated to high temperature when massive structures virialize.  One of the interesting effects that can generate secondary anisotropies is the weak gravitational lensing of the CMB, which arises from the distortions induced in the geodesics of CMB photons by gradients in the gravitational matter potential \\citep{Matthias_rew,Lewis06}.  Forthcoming CMB probes do have the sensitivity and expected instrumental performance which may allow a detection of the lensing distortions of the primary CMB anisotropies, which would then also provide new insights and constraints on the expansion history of the universe and on the process of cosmological structure formation \\citep{acquaviva_baccigalupi_2006,hu_etal_2006}. However, accurate predictions for the expected anisotropies in total intensity and polarization are clearly needed for analyzing this future data, which demands for detailed simulated maps.  The increasing availability of high-resolution N-body simulations in large periodic volumes makes it possible to directly simulate the CMB distortions caused by weak lensing using realistic cosmological structure formation calculations. This work represents a first step in that direction.  Existing studies already give access to statistical properties of the expected all-sky CMB lensing signal, such as the two-point correlation function and the power spectrum of the lensing potential and deflection angle, see e.g. \\citep{Lewis05} and references therein. This is based on `semi-analytic' calculations that use approximate parameterizations of the non-linear evolution of the matter power spectrum. On the other hand, up to now N-body numerical simulations have been used to lens the CMB only on small patches of the sky in order to exploit the practicality of the flat-sky approximation, see e.g. \\citep{Amblard et al 04} and references therein.  However, our approach of propagating rays through the forming dark matter structures gives access to the full statistics of the signal, including non-linear and non-Gaussian effects. Furthermore, it allows the accurate characterization of correlations of CMB lensing distortions with the cosmic large-scale structure, and with other foregrounds such as the Sunyaev-Zeldovich and Rees-Sciama effects. Hopefully this will allow improvements in the methods for separating the different contributions to CMB anisotropies in the data, which would be of tremendous help to uncover all the cosmological information in the forthcoming observations. \\begin{figure*} \\begin{center} \\resizebox{8cm}{!}{\\includegraphics{stacking.ps}} \\caption{Sketch of the adopted stacking and randomization process. The passage of CMB photons through the dark matter distribution of the Universe is followed by stacking the gravitational potential boxes of the MS, which are $500\\,h^{-1}{\\rm Mpc}$ on a side (comoving). Shells of thickness $500\\,h^{-1}{\\rm Mpc}$ are filled with periodic replicas of the box. All boxes (squares) that fall into the same shell are randomized with the same coordinate transformation (rotation and translation), which, in turn, differs from shell to shell.} \\label{stacking} \\end{center} \\end{figure*}  From an experimental point of view, the improved precision of the CMB observations, in particular that of the next generation experiments\\footnote{See \\textsf{lambda.gsfc.nasa.gov} for a complete list of operating and planned CMB experiments}, may in fact require an accurate delensing methodology and a detailed lensing reconstruction. CMB experiments targeting for instance the CMB polarization, and in particular the curl component of the polarization tensor, the so called $B$-modes from cosmological gravitational waves, may greatly benefit from a precise knowledge of the lensing effects in order to separate them from the primordial cosmological signal \\citep{seljak_hirata_2004}.  In particular, for a correct interpretation of the data from the forthcoming Planck satellite\\footnote{www.rssd.esa.int/PLANCK}, it will be absolutely essential to understand and model the CMB lensing, as the satellite has the sensitivity and overall instrumental performance for measuring the CMB lensing with good accuracy. We note that a first detection of CMB lensing in data from the Wilkinson Microwave Anisotropy Probe (WMAP\\footnote{See \\textsf{map.gsfc.nasa.gov}}) combined with complementary data has already been claimed by \\citep{smith_etal_2007} and \\citep{Seljak0108}.  In this study we introduce a new methodology for the construction of all-sky lensing-potential and deflection-angle maps, based on a very large cosmological simulation, the {\\it Millennium} run \\citep{Springel2005}.  As a first step in the analysis of the maps produced using the MS dark matter distribution, we have determined the interval of angular scales on which these maps match the semi-analytical expectations, since we expected a lack of lensing power on large scales, due to the finite volume of the N-body simulation.  To compensate for this effect, we have implemented a method for adding large-scale power which allows to recover the correct lensing signal on the scales outside this interval, i.e. on scales larger than the MS box size.  At the other extreme, at the smallest resolved scales, we are interested in the question whether our maps show evidence for extra lensing power due to the accurate representation of higher-order non-linear effects in our simulation methodology. On these small scales, the impact of non-Gaussianities from the mapping of non-linear lenses is expected to be largest.  This paper is organized as follows. In Section \\ref{lmvtba}, we briefly describe the basic aspects of lensing relevant to our work. In Section \\ref{mmpwtms}, we describe the N-body simulation and the details of our map-making procedure. In Section~\\ref{tslpada}, we present the lensing-potential and deflection-angle maps, and study the distribution of power in the angular domain. In Section~\\ref{c} we provide a summary and discussion. ",
        "conclusions": "\\label{c}  We constructed the first all sky maps of the cosmic microwave background (CMB) weak-lensing potential and deflection angle based on a high-resolution cosmological N-body simulation, the Millennium Run Simulation (MS). The lensing potential and deflection angle are evaluated in the Born approximation by directly ray-tracing through a three-dimensional, high-resolution mesh of the evolving peculiar gravitational potential and its gradient. The time evolution is approximated by 53 simulation outputs between redshift $z=0$ and $z\\simeq 11$, each used to cover a thin redshift interval corresponding to a shell in the past light-cone around the observer. To prevent artificial repetition of structures along the line-of-sight, while at the same time avoiding discontinuities in the force transverse to a line-of-sight, we tessellate shells of comoving thickness corresponding to the size of the box ($500\\,h^{-1}{\\rm Mpc}$) with periodic replicas which are coherently rotated and translated within each shell by a random amount.  Moreover, in order to include the contribution to the lensing signal from the scales larger than the MS box size, we have implemented a method for adding large-scale structure as described in the text.  Using the Hierarchical Equal Area Latitude Pixelization (\\textsf{HEALPIX}) package for obtaining a uniform sky-coverage, we have constructed simulated CMB lensing maps with $\\sim 5$ million pixels and an angular resolution of $\\sim 1.72^\\prime$, based on potential fields calculated on $2560^3$ meshes from the Millennium simulation.  In the present study, we analyze the power spectrum of the lensing potential and the deflection angle, and compare it with predictions made by semi-analytic approaches. We note that our general approach for map-making can be extended to other CMB foregrounds, including the Integrated Sachs-Wolfe (ISW) and Rees-Sciama effects at low redshifts, as well as estimates of the Sunyaev Zel'dovich (SZ) effects, or of the X-ray background.  This will in particular allow studies of the cross-correlation of the lensing of CMB temperature and polarization with these effects, which will be the subject of a forthcoming study.  In our approach we do not take into account the contributions of the baryonic physics to the lensing effects on the CMB. We expect in fact that these contributions could be non-negligible only on the typical scales of cluster cores and below, thus well above $l\\sim 3000$. Our comparison of the angular power spectrum of the lensing-potential and the deflection-angle with semi-analytic expectations reveals two different regimes in our results. First, for multipoles up to $l\\sim 2500$, our simulated maps produce a lensing signal that matches the semi-analytic expectation.  Second, we find evidence for a slight excess of power in our simulated maps on scales corresponding to few arcminutes and less, which we attribute to the accurate inclusion of non-linear power in the Millennium simulation. It will be especially interesting to study the non-Gaussianities in the signal we found and its implied consequences for CMB observations.  The new method proposed here demonstrates that an all-sky mapping of CMB lensing can be obtained based on modern high-resolution N-body simulations. This opens the way towards a full and accurate characterization of CMB lensing statistics, which is unaccessible beyond the power spectrum with the existing semi-analytical techniques. This is relevant in view of the forthcoming CMB probes, both as a way to detect, extract and study the CMB lensing signal, which carries hints on the early structure formation as well as the onset of cosmic acceleration, and as a tool to distinguish CMB lensing from the Gaussian contribution due to primordial gravitational fluctuations."
    },
    "0711/0711.2180_arXiv.txt": {
        "abstract": "We present internal surface-brightness profiles, based on HST/ACS imaging in the $F606W$ bandpass, for 131 globular cluster (GC) candidates with luminosities $L\\simeq 10^4$--$3\\times10^6\\,L_\\odot$ in the giant elliptical galaxy NGC 5128. Several structural models are fit to the profile of each cluster and combined with mass-to-light ratios from population-synthesis models, to derive a catalogue of fundamental structural and dynamical parameters parallel in form to the catalogues recently produced by \\citeauthor{mcl05} and by \\citeauthor{barmby07} for GCs and massive young star clusters in Local Group galaxies. As part of this, we provide corrected and extended parameter estimates for another 18 clusters in NGC 5128, which we observed previously. We show that, like GCs in the Milky Way and some of its satellites, the majority of globulars in NGC 5128 are well fit by isotropic \\citeauthor{wil75} models, which have intrinsically more distended envelope structures than the standard \\citeauthor{king66} lowered isothermal spheres. We use our models to predict internal velocity dispersions for every cluster in our sample. These predictions agree well in general with the observed dispersions in a small number of clusters for which spectroscopic data are available. In a subsequent paper, we use these results to investigate scaling relations for GCs in NGC 5128. ",
        "introduction": "\\label{sec:intro}  The spatial structures and internal stellar kinematics of old globular clusters (GCs) contain information on both their initial conditions and their dynamical evolution over a Hubble time. An efficient way of extracting this information is to fit detailed models to the surface-brightness profiles and (where available) velocity-dispersion data of individual clusters, and then to look for possible correlations between the physical properties of large numbers of GCs. This has been done for most of the $\\simeq\\! 150$ globulars in the Milky Way, yielding comprehensive catalogues of their structural and dynamical parameters \\citep{djo93,pry93,har96,mcl05}. Much of this work has traditionally started from the assumption that individual globulars are well described by the classic \\citet{king66} models of single-mass, isotropic, modified isothermal spheres, although recently alternative models have also been employed \\citep{mcl05}.  Explorations of numerous scaling relations and interdependences between the properties of Galactic GCs \\citep[e.g.,][]{djo94} have led to the definition of a fundamental plane for globulars that is analogous to but physically distinct from that for early-type galaxies and bulges \\citep{djo95,bur97,mcl00}. Understanding the GC fundamental plane in full detail is still not a completely solved problem, but important advances have been made in recent years as it has become possible to measure the internal properties of GCs in many other galaxies. High-resolution Hubble Space Telescope (HST) imaging has been used to fit \\citet{king66} and other models to the surface-brightness profiles of scores of globulars in the Large and Small Magellanic Clouds and the Fornax dwarf spheroidal \\citep{mg03a,mg03b,mg03c,mcl05}, M31 \\citep[e.g.,][]{barmby02,barmby07}, M33 \\citep{lar02}, and the giant elliptical galaxy NGC 5128 = Centaurus A (\\citealt{hol99,har02}). Internal velocity dispersions and dynamical mass estimates are also available for smaller but growing numbers of GCs in these systems \\citep{djo97,dub97,lar02,marho04,rej07}.  Here we add to this database with structural measurements of 131 GCs in NGC 5128. This galaxy is an attractive target for such studies in part because of its large GC population, estimated by \\citet{har06} at ${\\cal N}_{\\rm GC} \\simeq 1500$. It thus contains many objects at the high end of the star-cluster mass range ($10^6$--$10^7\\ M_{\\odot}$), which is largely unprobed in the ten-times smaller GC system of the Milky Way but where it is increasingly suggested that the cluster population may encompass a variety of objects including classic globulars, the compact nuclei of dwarf elliptical galaxies, and the new class of ultra-compact dwarf galaxies \\citep{hilker99,drink00,hasegan05}. In addition, the proximity of NGC 5128 ($D=3.8$~Mpc; see below) makes it possible to resolve the core radii as well as just the half-light radii of GCs over nearly their full mass range ($M\\ga 10^4\\ M_\\odot$), and thus to fit them rigorously with detailed structural models.  This paper is the third in a series of four dealing with HST observations of GCs in NGC 5128. In \\citet[][$\\equiv$~\\citetalias{har02}]{har02} the Space Telescope Imaging Spectrograph (STIS) and Wide Field Planetary Camera 2 (WFPC2) aboard HST were used to measure surface-brightness profiles for 27 very bright GCs in NGC 5128. \\citet{king66} models were fitted to these profiles to derive a structural fundamental plane that could be compared directly to that of the Milky Way globulars. In \\citet[][$\\equiv$~\\citetalias{har06}]{har06} we published the first results from a new HST-based survey of GCs in NGC 5128, using the Advanced Camera for Surveys (ACS) in its Wide Field Channel (WFC) to image a total of 131 GC candidates at a resolution of 0\\farcs05 (linear resolution $\\simeq\\! 0.9$ parsec). \\citetalias{har06} gives the full description of the cluster sample along with some rough overall characteristics of the ensemble of objects. In the present paper, we derive surface brightness profiles for all of these clusters and fit each of them with a number of different structural models. The final paper in this series \\citep[][$\\equiv$~\\citetalias{mcl07}]{mcl07} uses these results to examine a number of structural correlations for GCs in NGC 5128, which are then compared to the globulars in the Milky Way and to various other types of massive clusters. In related work, \\citet{barmby07} present the results of a similar ACS study of GCs in M31 and compare the fundamental planes of old GCs in that galaxy, the Milky Way, NGC 5128, the Large and Small Magellanic Clouds, and the Fornax dwarf spheroidal.  In the next Section we describe the steps we have taken to derive surface brightness profiles for the GC candidates from \\citetalias{har06}, to characterise the point-spread function (PSF) that blurs the very central regions ($R\\la 2$--3~pc) of these profiles, to transform the surface-brightness data from their native HST filter to the standard $V$ bandpass, and to estimate metallicities for the clusters from separate, ground-based Washington photometry.  In \\S\\ref{sec:models}, we apply publicly available population-synthesis models to estimate individual $V$-band mass-to-light ratios for the clusters, given their metallicities and assuming various (old) ages. Following this, we summarise the main properties of each of three structural models (those of \\citealt{king66}, \\citealt{wil75}, and \\citealt{sersic}) that we have convolved with the ACS/WFC PSF and fit to every observed surface-brightness profile.  Section \\ref{sec:fits} gives the results of these fits and uses them to infer a wide range of structural and dynamical parameters, including total cluster luminosities and masses, effective and core radii and stellar densities, concentration indices, relaxation times, total binding energies, predicted central velocity dispersions, and $\\kappa$-space parameters for the fundamental plane \\citep{bbf92}. We present these in tables that are available in machine-readable format either online\\footnotemark \\footnotetext{See http://www.astro.keele.ac.uk/$\\sim$dem/clusters.html} or upon request from the first author. Note that our measurements of GC luminosities and intrinsic sizes in particular supersede the recent estimates of \\citet{vdb07}, who based his numbers on a less detailed analysis of some very basic cluster characteristics given in \\citetalias{har06}.  In \\S\\ref{sec:fits} we also address the question of whether the standard \\citet{king66} model specifically gives the best possible fit to GC surface-brightness profiles in NGC 5128. We then extend the range of physical parameters calculated for the smaller sample of clusters previously fitted with \\citeauthor{king66} models in \\citetalias{har02}, and we provide important corrections to some of the more basic parameters (in particular, the intrinsic central surface brightnesses) already published in that earlier work. The results of this re-analysis are tabulated in Appendix \\ref{sec:STIStables}.  In \\S\\ref{sec:apvel} we combine our structural modeling and population-synthesis mass-to-light ratios to predict line-of-sight velocity dispersions within a series of circular apertures with physical radii suited to realistic observational set-ups. We compare these predictions with spectroscopic data from \\citet{marho04} and \\citet{rej07} for some of our current cluster sample. Finally, \\S\\ref{sec:summary} summarises the paper.  Our modeling analysis in this paper is in all respects very similar to that undertaken by \\citet{mcl05} for a sample of 103 old GCs and 50 young massive clusters drawn from the Milky Way, the Large and Small Magellanic Clouds, and the Fornax dwarf spheroidal. The catalogues of structural and dynamical properties that we produce here for GCs in NGC 5128 are likewise very close in form and content to those in \\citeauthor{mcl05}. We have recently completed the same type of modeling and produced parallel catalogues for a further 93 GCs in M31 \\citep{barmby07}. As we mentioned above, \\citeauthor{barmby07} combine results to compare the fundamental planes of the old GCs in all six galaxies. In \\citetalias{mcl07} \\citep{mcl07} we directly compare GC structural correlations only between the Milky Way and NGC 5128, but we also examine how they relate to other kinds of massive star clusters.  In all of what follows, we adopt a distance of 3.8 Mpc to NGC 5128. This value is representative of recent measurements based on the tip of the red-giant branch [$(m-M)_0 = 27.98 \\pm 0.15$], the planetary nebulae luminosity function [$(m-M)_0 = 27.97 \\pm 0.14$], surface-brightness fluctuations [$(m-M)_0 = 27.78 \\pm 0.10$], Mira variables [$(m-M)_0 = 27.96 \\pm 0.11$], and Cepheids [$(m-M)_0 = 27.67 \\pm 0.20$]; see \\citet{har99}, \\citet{rej04}, and \\citet{ferr07}. The nominal average of these five, reasonably high-precision distances is $(m-M)_0 = 27.88 \\pm 0.06$, or $3.76 \\pm 0.11$~Mpc. All these methods have undergone recent calibration revisions of various kinds \\citep[cf.][]{ferr07} but the net results have been to shift the mean up or down by amounts at the level of only $0.1$~mag. At a distance of 3.8 Mpc, 1 arcsecond is subtended by 18.4 parsec. One ACS/WFC pixel (0\\farcs05) then corresponds to 0.92~pc. ",
        "conclusions": ""
    },
    "0711/0711.0466_arXiv.txt": {
        "abstract": " ",
        "introduction": "We start with a CDM like power law in the quantity $\\rho/\\sigma^{3}$ where $\\rho$ is the local space density and $\\sigma$ is the local radial velocity dispersion. The above quantity is often loosely referred to as the phase space density (as it is here) but it is actually a `pseudo' phase space density (e.g. see Dekel $\\&$ Arad (2004) for a discussion of the true 6-D phase space density). This power law is maintained in the outer regions of the model but with a continued rise at sufficiently small radius, the phase space density is assumed to reach the Tremaine \\& Gunn (1979) limit in the absence of other lower and less fundamental limiting effects. This quantum statistical upper limit on the phase space density applies to thermal particles as well as fermions. Hence, as long as the particles are not bosons we expect an eventual cap/core in the central phase space density. Here a simple model is proposed in order to mimic the lingering effects of a putative but as yet unknown primordial phase space density bound (Q$_{p}\\equiv\\bar{\\rho}/{\\bar{\\sigma}^{3}}$ where $\\bar{\\rho}$ is the mean density and $\\bar{\\sigma}$ is the one dimensional velocity dispersion). With a constant central phase space density ($Q_{o}$) core of size r$_{c}$, the following profile is defined:  \\begin{equation} \\rho/\\sigma^{3}\\equiv  Q(r) = \\frac{Q_{o}}{(1+(r/r_{c})^{\\alpha})^{1.92/\\alpha }} \\end{equation}  The choice of a model independent power law index of 1.92 (close to that found by Taylor \\& Navarro) comes from the work of Dehnen \\& McLaughlin (2005). All but one of the models presented here were computed with the shape parameter $\\alpha=1.92$. Lower values of $\\alpha$ result in a more gradual transition to the outer power law and as will be shown produce very similar results.  The above expression for Q(r) is substituted into the Jeans equation after eliminating $\\sigma^{2}$, thereby allowing the determination of the density profile $\\rho$. (Note the assumption of spherical symmetry).  \\begin{equation} \\frac{d\\log \\ \\rho}{d\\log \\ r} = -0.6\\frac{GM_{r}}{r}\\left(\\frac{Q}{\\rho} \\right)^{2/3}-1.2\\beta +0.4\\frac{d\\log \\ Q}{d\\log \\ r} \\end{equation}  with  \\begin{equation} \\frac{d M_{r}}{d\\log \\ r} = \\ln(10)  4  \\pi  r^{3}  \\rho \\end{equation} where $\\beta$ is the anisotropy parameter (Binney \\& Tremaine, 1987) and is positive for a predominantly radial anisotropy. The following dimensionless number involving the initial conditions was found to lie between 2 and 3 for all models presented here.  \\begin{equation} \\gamma = 4\\pi G r^{2}_{c} Q_{o}(\\sigma^{2}_{o})^{1/2}=4\\pi G r^{2}_{c} Q^{2/3} _{o}\\rho^{1/3}_{o} \\end{equation}  Integrations can be carried out with $\\beta = 0$, but for the above range of $\\gamma$, the dispersion is found to increase outwards around the point of gradient change in Q(r). As CDM simulations indicate a small radial anisotropy in the $outer$ parts of a halo, at each step in the integration, trial values of $\\beta$ are stepped through (in units of 0.001) in order to determine that value which makes the logarithmic gradient in dispersion have the shallowest $negative$ value. In this way a $\\beta$ profile is obtained starting at zero in the center (actually 0.001 for computational reasons) remaining small throughout the core and usually ending up $\\sim0.2-0.3$. At some point further out the scheme tries to make $\\beta$ decrease but it is constrained to remain at its maximum value. The above value of $\\beta_{max}$ is a rough average of the outermost values determined for CDM halos (Fig. 3 of Dehnen \\& McLaughlin (2005)). Unlike models computed without the above simple, well defined prescription, those here exhibit consistent scaling relations (see discussion below). Furthermore, the initial rise of $\\beta$ is very nearly a linear function of the logarithmic density gradient as advocated by Hansen \\& Moore (2006). Beyond $\\beta\\sim0.3$ however, both the Hansen \\& Moore and the Dehnen \\& McLaughlin results show a large scatter in $\\beta$ and we have elected to keep it constant in this region.  The core can be considered isothermal (i.e. constant velocity dispersion and negligible velocity anisotropy) within the region where the logarithmic density gradient is greater than -0.1. This radius is $\\sim0.2$~r$_{c}$.  Depending on what observations are given (i.e. initial rotation curve slope, the location of the bend in the rotation curve or its amplitude) determines which of the parameters $\\rho_{o}$, r$_{c}$ or Q$_{o}$ one chooses to fix initially. A model is constructed by integrating equations (2) and (3) while systematically varying the other two parameters until the logarithmic density gradient becomes -4.000 at M$_{vir}$\\footnote{ Approximate solutions which obviate the need for the trial and error procedure are given in the appendix.}. M$_{vir}$ and R$_{vir}$ are defined as the values of mass and radius where the mean density  becomes $100\\times\\rho_{c}$ (h$_{100}=0.7$ assumed throughout).  This outer boundary condition was determined empirically (e.g. by fitting the outer profile to the observations of the cluster A1689 as described below). The density profile derived from equation (2) is quite different from an NFW profile. It possesses a constant density core followed by a relatively steep radial fall-off. A steep fall-off in density appears to be demanded by the observations of galaxy cluster A1689 (Broadhurst et al. 2005) and is shown in Fig. 1. The value of $\\rho_{o}$ determined for a converged model corresponds to a particular value of $\\beta_{max}$. Remarkably, any other model with {\\it{the same}} $\\rho_{o}$ can then be obtained from the following scaling relations (i.e. Q$_{o}\\propto M^ {-1}$, Q$_{o}\\propto \\sigma_{o}^{-3}$ and Q$_{o}\\propto r^{-3}_{c}$) and hence is a member of a {\\it{one parameter}} family of models. Interestingly, these scaling relations are identical to those discussed by Dalcanton \\& Hogan (2001) to describe the results of `gentle' merging given that during a merger Q cannot increase. Based on this discussion, the inverse relation between Q$_{o}$ and M found here suggests that equation (1), for all of its simplicity, is consistent with {\\it{a form of}} hierarchical structure formation. As discussed by the above authors, it is decidedly not compatible with `phase packing' where one expects more massive objects (with higher central velocity dispersions) to have smaller core radii.  In what follows, models are specified by the four parameters Q$_{o}$ (M$_{\\odot}pc^{-3}(km \\ sec^{-1})^{-3}$), r$_{c}$ (kpc), $\\rho_{o}$ (M$_{\\odot}pc^{-3}$) and $\\alpha$ and are enclosed in brackets.  It is useful to express the equivalent {\\it{gas}} entropy in terms of Q. We do this by evaluating the adiabatic constant $K=kTn_{e}^ {-2/3}$. With Q in the same units as above we obtain  \\begin{equation} K=\\frac{8.84\\times10^{-7}\\mu\\mu_{e}^{2/3}((3-2\\beta(r))/3)}{Q^{2/3}}=K_{o} (3-2\\beta(r))/3)(1+(r/r_{c})^{\\alpha})^{1.28/\\alpha} \\end{equation} and $K_{o}=8.84\\times10^{-7}\\mu\\mu_{e}^{2/3}Q_{o}^{-2/3}$ kev cm$^{-2}$  Note that the entropy of the gas is initially constant at $K_{o}$ and then increases as a power law with index 1.28 as long as $\\beta$ remains constant in the outer region (see Fig. 1). We emphasize that inherent in equation (5) is the assumption that the entropy of the gas is the same as the entropy of the dark matter and that as merging continues the increase in entropy is the same for both components. Thus this value of K must be a lower limit to the actual gas entropy and as such it provides a floor on which gas physics processes (i.e. cooling, heating, astration etc.) can be played out.  The characteristics of a representative model (in this case for the galaxy cluster A1689) are shown in Fig. 1. ",
        "conclusions": "A new model for dark matter halos has been proposed and is successfully applied to observations of objects with masses ranging from $\\sim10^{9}$ to $\\sim10^{15}$ M$_{\\odot}$. Fig. 3 is a graphical summary of the results which shows the scaling relations discussed earlier. It is important to note key differences in structure occur with different scaling normalizations. The filled symbols are structures with relatively high values of $\\rho_{o}$ ($\\sim 0.1 M_{\\odot}~ pc^{-3}$) while the open symbols are structures with lower values of $\\rho_{o}$ ($\\sim 0.01 M_{\\odot} ~ pc^{-3}$) and lower $\\beta_{max}$'s. Close examination of Fig. 3 reveals a real systematic shift between these two groups of objects. Objects with identical central densities would lie essentially dispersionless along a line of the indicated slope. Further, for two models with the same mass, the one with the higher $Q_{o}$ has the higher $\\rho_{o}$ and smaller $r_{c}$ (i.e. quantitatively $\\partial$~log Q$_{o}$/$\\partial$~log $ \\rho_{o}=0.67$ and $\\partial$~log Q$_{o}$/$\\partial$~log r$_{c}=-1.57$ for a fixed halo mass). While the density is not expected to increase during merging, Hernquist et al. (1993) propose a scheme whereby the density decreases while the dispersion remains constant. Dalcanton \\& Hogan interpret this as a result of more violent merging so that different merging histories at earlier times could account for the variations in central density seen now. More observations are required to determine if the apparent dichotomy in central density is real and if so what its origin is. For example, if we assume that the dichotomy extends to galaxy clusters, then a model with the same mass as A1689 but with one tenth the central density will have its gas entropy floor raised by $\\sim2.8$. Such a change in central gas entropy floor is one characterization differentiating cooling flow clusters from non cooling flow clusters.  The error bars in each panel of Fig. 3 were calculated by assuming an empirical approximation to $\\gamma$ (equation (4)) in terms of the central density (i.e. $\\gamma\\sim1.80\\rho^{-0.0753}_{o}$) and variables $\\rho_{o}$ and r$_{c}$ were then treated as independent with estimated uncertainties of $\\pm0.25$ and $\\pm0.10$ respectively.  Rotation curves derived from the model and the NFW profile can be very similar (e.g. the two A class galaxies in \\S3.1) while according to the model the data appear incompatible with standard CDM cosmological simulations which do not predict `observable-sized' cores. The similarity is most pronounced among objects in the high central density group. One property of our cored models which may be relevant to the `missing satellite' problem is that these structures are more vulnerable to tidal disruption than the NFW models, especially those halos with low central densities.  From this work a relation has been determined between Q$_{o}$ and the mass of dark matter halos over a range from $\\sim10^{9}$ to $\\sim10^{15}$ M$_{\\odot}$ . An observational challenge is to find the lower mass limit to objects with dark matter halos thus providing an estimate of the primordial value of Q (Q$_{p}$). Knowledge of Q$_{p}$ allows the determination of the mass of the dark matter particle (assuming that the particles are thermal) since then Q is proportional to the fourth power of the particle mass (e.g. equating the value of Q$_{o}$ found above for Fornax with Q$_{p}$ provides a lower limit on this mass of 431~ev). An additional constraint comes from an analysis of the power spectrum of the Ly$\\alpha$ forest. From this one can determine the free streaming length ($\\lambda_{fs}$) of the dark matter particle. This quantity in turn is simply related (in the case of thermal particles) to the particle mass. A recent determination of a limit on $\\lambda_{fs}$ by Seljak et al. (2006) implies a thermal dark matter particle mass limit of $>10$ kev (i.e. Q$_{p}>1$ which according to our scaling relation above implies dark matter halos with masses as low as $\\sim10^3$~M$_{\\odot}$).  An attractive alternative to the above `classical' WDM scenario has been proposed by Strigari et al. (2007). If the particles are non-thermally produced by the decay of a supersymmetric particle for example and if they are born sufficiently late then the initial velocities of the resulting daughter particles can be sufficiently high to yield a free streaming length comparable to that found from the Ly$\\alpha$ forest analysis but with Q$_{p}$ orders of magnitude lower than above (i.e. Q$_{p}\\sim10^{-5}-10^{-6}$) and a correspondingly much higher dark halo mass limit. This picture has the additional feature that it is hierarchical in the conventional CDM sense since the parent particles are born cold and being bosons they are not subject to the ultimate phase space density restriction.   Future results from experimental particle physics and even more sophisticated cosmological simulations should lead to a fuller understanding of the dark matter problem and the viability of the model."
    },
    "0711/0711.0520_arXiv.txt": {
        "abstract": "We study the cosmological stability of a class of theories with a dynamical preferred frame. For a range of actions, we find cosmological solutions which are compatible with observations of the recent history of the Universe: a matter dominated era followed by accelerated expansion. We then study the evolution of linear perturbations on these backgrounds and find conditions on the parameters of the theory which allow for the growth of structure sourced by the new degrees of freedom. ",
        "introduction": "Theories of modified gravity can be invoked as viable solutions to the problem of missing mass. Though initially restricted to making predictions regarding quasistatic, weak field configurations \\cite{milgrom,angus,smgaugh}, more recent models are derivable from generally covariant actions \\cite{bek1,navvan,bruneton,dub} and so can, in principle, be compared with traditional tests of general relativity (for instance see \\cite{kahya,seif,zhang}).  The best known proposal, the Tensor-Vector-Scalar (TeVeS) theory of gravity proposed by Bekenstein has been studied in some depth \\cite{bek2,fam,skor,mavro}, and has been found to explain a plethora of observations, from galactic rotation curves to the growth of structure during recombination without invoking dark matter. We have shown \\cite{zfs1} that it can be seen as a particular case of Einstein-Aether theories of gravity \\cite{ae} with non-canonical kinetic terms and that one can consider a different subclass of these models which seem to be observationally viable \\cite{zfs2}. For convenience and to distinguish them from standard Einstein-Aether theories, we have named these theories Generalized Einstein Aether Theories.  In this paper we wish to explore, in as general terms as possible, the cosmological stability of these theories. It has been shown that the ``vector'' part of TeVeS plays an essential role in the growth of structure that leads to the formation of galaxies \\cite{dodel}. Perturbations in the time-like vector field are gravitationally unstable and grow in a manner akin, but not identical, to that of normal, pressureless matter. We wish to find if this phenomenon is present in Generalized Einstein Aether Theories.  In this paper we present the following results: \\begin{itemize} \\item we present the full set of background, and linear perturbation, equations for Generalized Einstein Aether Theories; \\item  for a broad class of models there is accelerated expansion following a period of matter domination; \\item  in a pure baryonic universe, the growth of structure can be sustained and enhanced by the presence of the Aether vector field; \\item although the growth of structure is driven by the vector field, unlike in TeVeS, it remains strictly unit timelike leading to qualitatively different features; \\item there are choices of parameters in the Einstein-Aether theories that give a reasonable fit to current measurements of large scale structure \\item a few notable features can be used to distinguish it from conventional gravity and dark matter, in particular a different growth rate for structure and the presence of non zero, scalar anisotropic stresses (``$\\Phi-\\Psi$''). \\end{itemize}  There are a number of steps we must go through before we can assess if there is the desired instability of the vector field, or aether, in these theories. We are looking for it in the cosmological setting and hence we must identify the restricted class of theories that give us acceptable cosmologies. We do so in Section \\ref{sect:background}. Once we have established what region of parameter space we want to explore, we then check if perturbations in the vector field grow in a matter dominated era. We shall see, in Section \\ref{sect:vector} that this defines a subset of parameters. We then consider the full coupled system of metric, matter, and vector field perturbations for modes outside the horizon well before recombination and trace their evolution up to the present time.  As will become clear, we are able to identify the subset of theories which could possibly lead to realistic growth of structure in a purely baryonic universe. In Section \\ref{sect:conc} we discuss if the constraints we find in this paper are compatible with a host of constraints that have been identified in \\cite{zfs2}. ",
        "conclusions": "\\label{sect:conc}  We have shown that the vector field in the Generalized Einstein Aether theory can play the same role as in TeVeS and the same role as dark matter in sustaining the growth of structure during recombination. A growing mode for the vector field exists in the matter era for a wide variety of parameters and, through its associated energy density and anisotropic stress, can become a dominant source term in the Poisson equation, thus sustaining gravitational potentials which further act as sources for the evolution of the baryonic density contrast. The shape of the power spectrum is dictated predominantly by the $k$ dependence of the growing mode which may either be determined by dominant initial data for the vector field (Case 2) or the manner in which an initially very small vector field reaches the growing mode solution via coupling to evolving gravitational potentials (Case 1). As in the case of dark matter, a turnover in the power spectrum exists in the latter case due to the comparitive lack of damping of metric modes entering the horizon after radiation domination.  Though the vector's stress energy tensor does contain terms proportional to the gravitational potentials (leading to an effective time dependent rescaling of Newton's constant in Poisson's equation) we find that this is a subdominant effect. At late times and on large scales Poisson's equation (\\ref{eqn:Poisson})  takes the approximate form:  \\begin{equation} \\label{eq:cosmol} \\nabla^{2}\\Phi\\sim 4\\pi G\\rho_{B}+\\frac{l_{E}}{2}{\\cal F}_{K}(\\eta){\\cal H}\\nabla^{2}V \\end{equation} where $\\rho_{B}$ is the baryonic energy density.  The dominant component of the right hand side at late times is found to be the term associated with the field $V$. In so far as the evolution of $V$ is sourced by terms in the metric but not entirely fixed by it via constraints, its contribution to the Poisson equation is more akin to dark matter.  It is interesting to contrast this with the weak-field quasistatic limit of the theory discussed in \\cite{zfs2}, wherein $A^{\\mu}$'s stress energy tensor is determined \\textit{entirely} by the metric and thus all observed mass discrepancies must be attributed to a modification of Poisson's equation which takes the form:  \\begin{equation} \\label{eq:weakfi} \\nabla.(\\mu(|\\nabla\\Phi|/M)\\nabla\\Phi)=4\\pi G\\rho_{B} \\end{equation}  Therefore in this regime mass discrepancies can be interpreted as arising from a change in the relationship between the dominant component of the matter energy density (taken to be baryonic matter) and the gravitational field. In producing a realistic power spectrum, we find that the effective energy density of the growing mode $V$ is considerable at the scale where overdensities become nonlinear ($k\\sim 2h Mpc^{-1}$). A similar situation exists in the TeVeS theory. This raises the issue, as yet unaddressed, of the degree to which $V$, instrumental in the growth of large scale structure, contributes to the mass of bound structure and, if so, whether it remains akin to a dark source in the Poisson equation (as in \\ref{eq:cosmol}) rather than a modification of how the gravitational field is sourced by baryonic matter (as in \\ref{eq:weakfi}). In other words, covariant realizations of the MOND program may well have unintentionally reintroduced dark matter,  albeit non-particulate dark matter, via the back door. However, it is by no means obvious that the collective picture given by a universe where a vector field gives rise to mass discrepancies would be degenerate with a particle dark matter model i.e the theory may resemble different types of dark matter in different regimes but not a single type across all regimes. Indeed we have seen that the cosmological case involves significant anisotropic stresses at late times while the quasistatic case considered in  \\cite{zfs2} does not.  We were able to clarify the analogy to the cosmological perturbations of Bekenstein's TeVeS theory. While the vector field is also nonvanishing in the cosmological background in that case, it is seen that the growth of structure requires the vector field be of non-fixed norm. As in the model discussed here, a growing mode in the vector field eventually dominates the evolution of the baryonic matter, its scale dependence being determined by the influence of metric source terms in its equation of motion. However, in TeVeS the growing mode and vector field perturbed stress energy tensor will tend to have different time dependences and so make for a potentially differing evolution of perturbations over cosmic time. In both theories the quantity $\\delta\\tilde{T}^{0}_{\\phantom{0}0}$ will generally retain a time dependence during the matter era, in contrast to the perturbed dark matter density in $\\Lambda$CDM models.  We now consider the parameter space explored in this paper. We have found that a power spectrum with a turnover and giving a reasonable fit to data tend to follow when the vector field and its time derivative are initially small. Following this, the degree to which the power spectrum today can agrees with the data is dictated the values of the parameters $n$ and $\\gamma$ (recall that ${\\cal F}=\\gamma(-{\\cal K})^{n}$) and the $c_{i}$ (see section \\ref{sect:Case1}). In the background we have seen (see figure \\ref{fig1}) through that some values of $n$ and $\\gamma$ in the modified Friedmann equation permit late time acceleration. To first order in perturbations we have seen that the influence of $n$ on the evolution of the matter and metric fields is via its combined influence on the exponent of the vector field growing mode and, along with $\\gamma$, the time dependence of the quantity $F_{K}$ to zeroth and first order which themselves act as sources in the Poisson equation and difference between Newtonian gauge potentials. In particular we were able to find values of $n$ consistent with late time acceleration and a growing mode. We were unable to produce a realistic power spectrum in the absence of dark matter for the value $n=1$ (as hinted at by the lack of growing modes in the homogeneous solution to the vector equation for this value).  The $c_{i}$, which determine the vector field's kinetic term, affect the degree to which the Friedmann equation is modified in the cosmological background. For instance, if $c_{1}+3c_{2}+c_{3}=0$ then there is no effect on the background evolution. To first order in perturbations, the $c_{i}$ determine the nature of growing mode. We have restricted ourselves to values of $b_{3}$ (see equation (\\ref{bees})) which led to an approximate monomial growth of the vector field. Deviations from this behaviour will have an integrated impact on the evolution of the other fields again, for instance, acting as a dark source in the Poisson equation but it is to be expected that different $n,c_{i}$ will match observations of the matter power spectrum.  It is worth noting that in this case as well as TeVeS a power-law growing mode emerges from a non-Maxwellian kinetic term (i.e. ${\\cal K} \\neq -F^{\\mu\\nu}F_{\\mu\\nu}$;$c_{1}+c_{3}\\neq0$); it has been shown in \\cite{zfs1,tartag,bf2} that such kinetic terms can arise from a change of variables at the level of the action. For instance in TeVeS the vector field kinetic term is Maxwellian when written as a bimetric theory but not when written as a single metric theory. Our choices for $F({\\cal K})$ were dictated by a desire for simplicity. It is to be hoped that a deeper grounding of the ideas discussed herein would fix the expected form of ${\\cal F}$. As all contributions of the vector field to the matter and metric evolution equations at 1st order in perturbations are proportional to the background value of ${\\cal F}_{K}$, the time dependence of this quantity can be expected to have a significant impact. We have seen that this will tend to produce enhanced growth for viable forms of a monomial ${\\cal F}(\\cal K)$.  Though we have chosen the sign of the term $\\alpha=c_{1}+3c_{2}+c_{3}$ to be consistent with subluminal propagation of gravitational waves in the limit far from matter, in generic backgrounds the speed of gravitational waves will depend on $F({\\cal K})$ and it is not obvious that the theory would pass tests such as those discussed in \\cite{moore}. Recent work \\cite{li,garfinkle} on the consequences of the theory with $F \\propto {\\cal K}$ have chosen the opposite sign of $\\alpha$ where gravity propagates superluminally with respect to the preferred frame, thus avoiding any constraints from the above tests \\cite{PPN} ; such a choice would change the requirements for a growing vector mode for more complicated $F({\\cal K})$ as well as leading to the background cosmology and weak field limit being described by the same sign ${\\cal K}$, rendering them no longer independent regimes.  Finally, the smallness of the mass scale appearing in the vector action $M\\sim cH_{0}<<M_{Pl}$ is purely phenomenological. Though the value lends itself to modifications of gravity associated with low gravitational fields in the weak field limit and at late times in the cosmological background (i.e. as $H\\rightarrow M$), it is fixed as a number in the action entirely by the former and lacks a deeper theoretical justification. A model wherein $M$ itself is a function of cosmic time (through, say, a background time dependence on $A^{2}$) may have considerable effects on the background cosmology as well as time evolution of the perturbed vector field\u2019s stress energy tensor.  As the size of perturbations is constrained on large scales by the baryon power spectrum at late times and the CMB anisotropy at recombination (allowing for the ISW effect), a full simulataneous modelling of each in the presence of the vector field should allow further discrimination between models and the degree to which they can give an account of mass discrepancies compatible with data in precision cosmology.  In summary we have considered a set of models of a universe with a timelike vector field with a noncanonical kinetic term. We have found that the model on one hand exhibits behaviour typical of `modified gravity' theories in high symmetry (for instance modifying the Friedmann equation in the FRW background and the Poisson equation in the static spherically symmetric case \\cite{zfs2}). Additionally we have found that the extra degrees of freedom introduced through covariant realization of the former behaviour behave in a manner more akin to dark matter in perturbations around the FRW background. In particular, we find that at late times the stress energy tensor components of the field $V$ can become the dominant component in the cosmological Poisson equation. It is not clear, for instance, that this dominant component should be found to align with displacement of the subdominant baryonic matter in large scale structure. Even in so far as the field $V$ may resemble dark matter in the linear regime we have found a number of aspects in which its account differs from that of cold dark matter. We hope that this may help point the way towards a general approach to observationally distinguishing between the concordance model of cosmology and modified gravity.  {\\it Acknowledgments}: We thank Ted Jacobson, Constantinos Skordis, and David Mota for discussions. TGZ is supported by a PPARC studentship. GDS was supported by Guggenheim and Beecroft Fellowships, and by a grant from the US Department of Energy to the particle-astrophysics theory group at CWRU.  GDS and PF thank the Galileo Galilei Institute of Florence for their hospitality.  GDS thanks Oxford Astrophysics for its hospitality, and Mapelsoft for the use of Maple software. \\vspace{-.3in}"
    },
    "0711/0711.1526_arXiv.txt": {
        "abstract": "We present $FUSE$ observations of the hot white dwarfs in the post-common envelope binaries Feige~24, EUVE~J0720$-$317, BPM~6502, and EUVE~J2013$+$400. The spectra show numerous photospheric absorption lines which trace the white dwarf orbital motion. We report the detection of \\ion{C}{3}, \\ion{O}{6}, \\ion{P}{5}, and \\ion{Si}{4} in the spectra of Feige~24, EUVE~J0720$-$317 and EUVE~J2013$+$400, and the detection of \\ion{C}{3}, \\ion{N}{2}, \\ion{Si}{3}, \\ion{Si}{4}, and \\ion{Fe}{3} in the spectra of BPM~6502. Abundance measurements support the possibility that white dwarfs in post-common envelope binaries accrete material from the secondary star wind. The $FUSE$ observations of BPM~6502 and EUVE~J2013$+$400 cover a complete binary orbit. We used the {\\it FUSE} spectra to measure the radial velocities traced by the white dwarf in the four binaries, where the zero-point velocity were fixed using the ISM velocities in the line of sight of the stellar systems. For BPM~6502 we determined a white dwarf velocity semi-amplitude of $K_{\\rm WD} = 18.6\\pm0.5$km s$^{-1}$, and with the velocity semi-amplitude of the red dwarf companion ($K_{\\rm RD} = 75.2\\pm3.1$km s$^{-1}$), we estimate the mass ratio to be $q = 0.25\\pm0.01$. Adopting a spectroscopic mass determination for the white dwarf, we infer a low secondary mass of $M_{\\rm RD} = 0.14\\pm0.01\\ M_\\odot$. For EUVE~J2013$+$400 we determine a white dwarf velocity semi-amplitude of $K_{\\rm WD} = 36.7\\pm0.7$km s$^{-1}$. The $FUSE$ observations of EUVE~J0720$-$317 cover approximately 30\\% of the binary period and combined with the $HST$ GHRS measurements (Vennes et al. 1999, ApJ 523, 386), we update the binary properties. {\\it FUSE} observations of Feige~24 cover approximately 60\\% of the orbit and we combine this data set with $HST$ STIS (Vennes et al. 2000, ApJ, 544, 423) data to update the binary properties. ",
        "introduction": "Post common-envelope binaries consist of an evolved primary (white dwarf or sdB) and a late-type main-sequence secondary in close orbit. These binaries are thought to have evolved from wide binary systems, where the more massive star evolved off the main sequence filling its Roche lobe and beginning mass transfer onto its less massive companion. If the transfer is dynamically unstable, a common-envelope (CE) is formed and friction between the stellar components and the CE decreases the orbital separation and induces the ejection of the CE. Depending on the separation of the components, some of these systems further evolve to become cataclysmic variables. From the sample of well studied post-CE binaries \\citep{sch2003}, approximately half will evolve into cataclysmic variables within a Hubble time.  The atmosphere of a white dwarf a in close binary system usually displays enhanced traces of heavy elements. \\citet{ven1999} found that the large abundance of carbon and helium as well as time variable helium abundance in some white dwarfs provide evidence for on-going accretion from the red dwarf.  We present high-resolution far-ultraviolet (FUV) spectroscopic observations of Feige~24 \\citep[PG~0232$+$035:][]{ven1994b,ven2000}, WD~0718$-$316 \\citep[EUVE~J0720$-$317, 2RE~J0720$-$318:][]{ven1994a,bar1995b}, WD~1042$-$690 \\citep[BPM~6502:][]{kaw2000} and WD~2011$+$398 \\citep[EUVE~J2013$+$400, 2RE~J2013$+$400:][]{tho1994,bar1995a} in \\S 2. We present the analyses of the {\\it FUSE} spectra in \\S 3, determine new orbital and stellar parameters of the four binary systems in \\S 4 and measure the heavy element abundance in the white dwarf atmospheres in \\S 5. We summarize in \\S 6. ",
        "conclusions": "We measured the white dwarf radial velocities in the close binary systems Feige~24, J0720$-$317, BPM~6502 and J2013$+$400 using {\\it FUSE} spectra. Combined with previous optical studies of the binaries, we have updated the binary properties. Of the four systems, only J0720$-$317 is likely to come into contact within a Hubble time and is therefore a representative progenitor of the presently observed cataclysmic variables. We have also reanalyzed the optical data of \\citet{kaw2002} and \\citet{mor2005} and showed that the correct orbital period of BPM~6502 is $0.3367849\\pm0.0000006$ days and that the period reported by \\citet{mor2005} should be discarded.  We measured the abundance of trace elements in the atmospheres of Feige~24, J0720$-$317, BPM~6502 and J2013$+$400 and found that the observed abundance pattern in BPM~6502 can be explained by steady accretion at a very low rate ($\\sim 1 \\times 10^{-17}\\ M_\\odot$ yr$^{-1}$). In the cases of J0720$-$317 and J2013$+$400 accretion from the secondary at much lower rates ($\\sim 10^{-19}\\ M_\\odot$ yr$^{-1}$) can explain the observed abundances, however the higher abundances of P in these stars cannot be explained by accretion from the secondary alone."
    },
    "0711/0711.3523_arXiv.txt": {
        "abstract": "{Recent analyses of low-mass eclipsing binary stars have unveiled a significant disagreement between the observations and the predictions of stellar structure models. Results show that theoretical models underestimate the radii and overestimate the effective temperatures of low-mass stars but yield luminosities that accord with observations. A hypothesis based upon the effects of stellar activity was put forward to explain the discrepancies.} {In this paper we study the existence of the same trend in single active stars and provide a consistent scenario to explain systematic differences between active and inactive stars in the H-R diagram reported earlier.} {The analysis is done using single field stars of spectral types late-K and M and computing their bolometric magnitudes and temperatures through infrared colours and spectral indices. The properties of the stars in samples of active and inactive stars are compared statistically to reveal systematic differences.} {After accounting for a number of possible bias effects, active stars are shown to be cooler than inactive stars of similar luminosity therefore implying a larger radius as well, in proportions that are in excellent agreement with those found from eclipsing binaries.} {The present results generalise the existence of strong radius and temperature dependences on stellar activity to the entire population of low-mass stars, regardless of their membership in close binary systems.} ",
        "introduction": "Structure and evolution of stars in the main sequence is, in principle, well understood. Theoretical models have made very significant progress in the past decades and are able to successfully fulfill most of the observational diagnostics (e.g., Lebreton 2000; Gallart et al. 2005). However, there are still a few open issues that taint the optimistic picture. Most notably, the theoretical modeling of the structure and evolution of stars on both ends of the main sequence is not yet fully resolved (e.g., Ribas 2006a).  Model testing in the low-mass regime has been usually restricted to the evaluation of the mass-luminosity relationship because of the relatively large availability of observational data. Studies such as Delfosse et al. (2000) concluded that models agree well with empirical measurements, especially when the tests are performed in near-infrared bands, which are less affected by metallicity (Chabrier \\& Baraffe 2000; Bonfils et al. 2005). The detailed study of the M-type eclipsing binary YY Gem by Torres \\& Ribas (2002) contributed to the tests by providing simultaneous and accurate measurements of the component masses, radii, effective temperatures and luminosities. The surprising result was that the components of the system were found to be significantly larger and cooler than the predictions of stellar structure models, but yet with luminosities in agreement with theoretical calculations. The same trend was later confirmed by a number of subsequent analyses of other eclipsing binaries (see Ribas 2006b and L\\'opez-Morales 2007 for a complete summary). Thus, a growing body of observational evidence is challenging the adequacy of current stellar models in reproducing the overall properties of low-mass stars.  Using the available data and analysing the possible scenarios, Ribas (2006b) and Torres et al. (2006) concluded that a plausible explanation for the observed discrepancies could be related to stellar activity. In this respect, the active components of eclipsing binaries appear to be larger and cooler than inactive single stars (i.e., those reproduced by theoretical models) while keeping similar luminosity. To first approximation, this means that, regardless of changes in the stellar outer layers, the rate of nuclear burning in the core is not modified by activity and therefore the overall flux is conserved. The mechanism responsible for the observed radius and temperature differences has not been identified yet, although the efficiency of convection in strong magnetic fields (Mullan \\& MacDonald 2001), or the simple effect of flux conservation in a spot-covered stellar surface (L\\'opez-Morales \\& Ribas 2005) could explain the observed changes in the stellar properties. Note that a possible correlation of the radius differences with stellar metallicity could also be present (Berger et al. 2006).  The stellar activity hypothesis has been built in the context of close binary systems, where orbital synchronization forces the components to rotate fast, therefore triggering high levels of magnetic activity. With a much longer history, the question of the possible differences between active and inactive stars has been a recurrent one (Kuiper 1942; Joy \\& Abt 1974). Recent studies have been generally focused on the comparison of radiative properties such as spectral types or photometric colours. The conclusions have been quite diverse, both in favour (Hawley et al. 1996; Amado \\& Byrne 1997) and against (West et al. 2004; Bochanski et al. 2005) the existence of systematic colour differences. One of the most conclusive analyses is that of Stauffer \\& Hartmann (1986), where the authors identified a separation of the sequences of active and inactive stars in a luminosity-colour plot. The reasons for the observed separation in the sequences could not be unambiguously identified since it could both come from luminosity differences at constant effective temperature, effective temperature differences at constant luminosity, or a combination of the two.  The results from eclipsing binaries provide a new perspective to analyse this long-standing issue since temperatures, radii and luminosities can be determined in a fundamental manner. As discussed above, the components of eclipsing binaries seem to be cooler and larger than model predictions but with luminosities in agreement. Very importantly, this observation is confirmed by the theoretical study of Chabrier et al. (2007), who have recently modified stellar evolution codes to include the effects of stellar activity. These authors also report systematic differences between the properties of active and inactive stars in similar amount and trend as those found from eclipsing binary studies.  With improved statistics with respect to Stauffer \\& Hartmann (1986) and the context provided by the new evidence discussed above, we analyse the existence of differences between active and inactive stars of spectral types late-K and M. In the present study we use luminosities directly determined from accurate trigonometric parallaxes and carry out a thorough analysis of possible biases, such as the effect of pre-main sequence (PMS) and binary stars. Also, we interpret the results in physical terms, i.e., effective temperature and radius variations. If active stars were indeed cooler and larger than their inactive counterparts, while keeping similar luminosities, this should be observable in single field stars (in addition to close binaries) thus generalizing the proposed stellar activity scenario to all low-mass stars. ",
        "conclusions": "The results of the analysis fully agree with the findings from low-mass eclipsing binaries. Therefore, the radius and temperature discrepancies found in the latter are not a consequence of their mutual interactions as being in close binaries but can be generalised to the entire population of low-mass stars. By this scenario, {\\em any} active star (either single or binary) would be equally affected by the mechanism responsible for lowering its photospheric temperature and increasing its radius, while keeping its luminosity roughly constant. In this work we thus provide solid ground for the interpretation of the differences between active and inactive stars pointed out already by Kuiper (1942), and later studied in more detail by Stauffer \\& Hartmann (1986). The observed separation between the active and inactive star sequences in the H-R diagram is caused by systematic temperature (or spectral type) rather than luminosity deviations. Our conclusions stem from a comprehensive analysis using a large sample of stars with accurate trigonometric parallaxes, and therefore the results are statistically relevant.  In a recent study, Torres (2007) has determined the mass and radius of the M2.5 transiting planet host GJ 436 from the stellar radiative properties and constraints from the transit. The resulting radius also seems to be larger than model predictions by about 10\\% yet the star does not show any strong activity. Note that the masses and radii determined in this way are not fundamental (as opposed to those coming from eclipsing binaries) and that they heavily rely on the assumed value of the effective temperature, which is not well established for M-type stars as discussed above. Thus, the comparison with stellar models may not be as meaningful as that resulting from classical eclipsing binaries. However, with the expected increase in the number of M stars with transiting planets, this method can provide valuable data to further evaluate the activity effects.  If we generalise the activity-related differences found in single and binary stars we conclude that current stellar structure and evolution models are not adequate to describe the physical (radii) and radiative (effective temperature, colour indices) properties of {\\em all} active low-mass stars. The problem is particularly severe for the analysis of open clusters or star forming regions, where low-mass stars may display high-levels of magnetic activity. This is because there exists a good correlation between activity and age, in the sense that younger stars rotate faster than their older counterparts (e.g., Hartmann et al. 1984) and thus have more efficient mechanisms to generate magnetic energy. In the particular case of M-type stars (i.e., masses below $\\sim$0.6 M$_{\\odot}$), which stay active up to relatively long rotation periods (Pizzolato et al. 2003) and, conversely, old ages (up to 1 Gyr or more), the discrepancy between models and observations may apply to the entire young disk population. We are currently working on a realistic implementation of the effects of magnetic activity in theoretical structure and evolution models to resolve the observed shortcomings. The appropriate framework in this direction is being established by Chabrier et al. (2007)."
    },
    "0711/0711.4590_arXiv.txt": {
        "abstract": "Helioseismology has allowed us to study the  structure of the Sun in unprecedented detail. One of the triumphs of the theory of stellar evolution was that helioseismic studies had shown that the structure of solar models is very similar to that of the Sun. However, this agreement has been spoiled by recent revisions of the solar heavy-element abundances. Heavy element abundances  determine the opacity of the stellar material and hence, are an important input to stellar model calculations. The models with the new, low abundances do not satisfy helioseismic constraints. We review here how heavy-element abundances affect solar models, how these models are tested with helioseismology, and the impact of the new abundances on standard solar models. We also discuss the attempts made to improve the agreement of the low-abundance models with the Sun and discuss how helioseismology is being used to determine the solar heavy-element abundance. A review of current literature shows that attempts to improve agreement between solar models with low heavy-element abundances and seismic inference have been unsuccessful so far.  The low-metallicity models that have the least disagreement with seismic data require changing all input physics to stellar models beyond their acceptable ranges. Seismic determinations of the solar heavy-element abundance yield results that are  consistent with the older, higher values of the solar abundance, and hence,  no major changes to the inputs to solar models are required to make higher-metallicity solar models consistent with helioseismic data. ",
        "introduction": "\\label{sec:intro}  Arthur Eddington began his book ``The Internal Constitution of the Stars'' saying that {\\it ``At first sight it would seem that the deep interior of the sun and stars is less accessible to scientific investigation than any other region of the universe. Our telescopes may probe farther and farther into the depths of space; but how can we ever obtain certain knowledge of that which is hidden behind substantial barriers? What appliance can pierce through the outer layers of a star and test the conditions within?''} (Eddington 1926). Eddington went on to say that perhaps we should not aspire to directly ``probe'' the interiors of the Sun and stars, but instead use our knowledge of basic physics to determine what the structure of a star should be. This is still the predominant approach in studying stars today. However, we now also have the ``appliance'' that can  pierce through the outer layers of the Sun and give us detailed knowledge of what the internal structure of the Sun is. This `appliance' is helioseismology, the study of the interior of the Sun  using solar oscillations. While solar neutrinos can probe the solar core, helioseismology provides us a much more detailed and nuanced picture of the entire Sun.  The first definite observations of solar oscillations were made by Leighton et al.\\ (1962), who detected roughly periodic oscillations in Doppler velocity with periods of about 5 minutes. Evans \\& Michard (1962) confirmed the initial observations. The early observations were of limited duration, and the oscillations were generally interpreted as phenomena in the solar atmosphere. Later observations that resulted in power spectra as a function of wave number (e.g., Frazier 1968) indicated that the oscillations may not be mere surface phenomena. The first major theoretical advance in the field came when Ulrich (1970) and Leibacher \\& Stein (1971) proposed that the oscillations were standing acoustic waves in the Sun, and predicted that power should be concentrated along ridges in a wave-number v/s frequency diagram. Wolff (1972) and Ando \\& Osaki (1975) strengthened the hypothesis of standing waves by showing that oscillations in the observed frequency and wave-number range may be linearly unstable and hence, can be excited. Acceptance of  this interpretation of the observations as normal modes of solar oscillations was the result of the observations of Deubner (1975), which first showed ridges in the wave-number v/s frequency diagram. Rhodes et al.~(1977) reported similar observations. These observations did not, however, resolve the individual modes of solar oscillations, despite that, these data were used to draw initial inferences about solar structure and dynamics. Claverie et al.~(1979) using Doppler-velocity observations, integrated over the solar disk were able to resolve the individual modes of oscillations corresponding to the largest horizontal wavelength. They found a series of almost equidistant peaks in the power spectrum just as was expected from theoretical models. However, helioseismology as we know it today did not begin till Duvall \\& Harvey (1983) determined frequencies of a reasonably large  number of solar oscillation modes covering a wide range of horizontal wavelengths. Since then many sets of solar oscillations frequencies have been published. A lot of early helioseismic analysis was based on frequencies determined by Libbrecht et al.~(1990) from observations made at the Big Bear Solar Observatory in the period 1986--1990. Accurate determination of solar oscillations frequencies requires long, uninterrupted observations of the Sun, that are possible only with a network of ground based instruments or from an instrument in space. The Birmingham Solar Oscillation Network (BiSON;  Elsworth et al.~1991; Chaplin et al.~2007a) was one of the first such networks. BiSON, however, observes the Sun in integrated light and hence is capable of observing only very large horizontal-wavelength modes. The Global Oscillation Network Group (GONG), a ground based network of telescopes,  and the Michelson Doppler Imager (MDI) on board the Solar and Heliospheric Observatory (SOHO) have now collected data for more than a decade and have given us an unprecedented opportunity to determine the structure and dynamics of the Sun in great detail. Data from these instruments have also allowed us to probe whether or not the Sun changes on the time scale of a solar activity cycle.  Helioseismology has proved to be an extremely important tool in studying the Sun. Thanks to helioseismology, we know the most important features of the structure of the Sun extremely well. We know what the sound-speed and density profiles are (see e.g., Christensen-Dalsgaard et al.~1985, 1989; Dziembowski et al.~1990; D\\\"appen et al.~1991; Antia \\& Basu~1994a; Gough et al.~1996; Kosovichev et al.~1997; Basu et al.~1997, 2000; etc.), which in turn means that we can determine the radial distribution of pressure. We can also determine the profile of the adiabatic index (e.g., Antia \\& Basu 1994a; Elliott 1996; Elliott \\& Kosovichev 1998). Inversions of solar oscillations frequencies have allowed us to determine a number of other fundamental facts about the Sun. We know, for instance, that the position of  base of the solar convection zone can be determined precisely (Christensen-Dalsgaard et al.~1991; Basu \\& Antia 1997; Basu 1998). Similarly, we can determine the helium abundance in the solar  convection zone (D\\\"appen \\& Gough 1986; Christensen-Dalsgaard \\& P\\'erez Hern\\'andez 1991; Kosovichev et al.~1992; Antia \\& Basu 1994b). In addition to these structural parameters, helioseismology has also revealed what the rotational profile of the Sun is like. It had been known for a long time that the rotation rate at the solar surface depends strongly on latitude, with rotation being fastest at the equator and slowest at the poles. Only with helioseismic data however, we have been able to probe the rotation of the Sun as a function of depth (Duvall et al.~1986; Thompson et al.~1996; Schou et al.~1998b; etc.).  The ability of helioseismology to probe the solar interior in such detail has allowed us to use the Sun as a laboratory to test different inputs that are used to construct solar models. For instance, helioseismic inversions have allowed us to study the equation of state of stellar material (Lubow et al.~1980; Ulrich 1982; \\jcd\\ \\& D\\\"appen 1992; Basu \\& Christensen-Dalsgaard 1997; Elliott \\& Kosovichev 1998; Basu et al.~1999) and to test opacity calculations (Korzennik \\& Ulrich 1989; Basu \\& Antia 1997; Tripathy \\& \\jcd~1998). Assuming that opacities, equation of state, and nuclear energy generation rates are known, one can also infer the temperature and hydrogen-abundance profiles of the Sun (Gough \\& Kosovichev 1988; Shibahashi 1993; Antia \\& Chitre 1995, 1998; Shibahashi \\& Takata 1996; Kosovichev 1996 etc.). These studies also provide a test for nuclear reaction rates (e.g., Antia \\& Chitre 1998; Brun et al.~2002) and the  heavy-element abundances in the convection zone (e.g., Basu \\& Antia 1997; Basu 1998) and the core (e.g., Antia \\& Chitre 2002).  One of the major inputs into solar models is the abundance of heavy elements. The heavy-element abundance, $Z$, affects solar structure by affecting radiative opacities. The abundance of some specific elements, such as oxygen, carbon, and nitrogen can also affect the energy generation rates through the CNO cycle. The effect of $Z$ on opacities changes the boundary between the radiative and convective zones, as well as the structure of radiative region;  the effect of $Z$ on energy generation rates can change the structure of the core. The heavy element abundance of the Sun is believed to be known to a much better accuracy than that of other stars, however, there is still a lot of uncertainty and that results in uncertainties in solar models. It is not only the total $Z$ that affects structure, the relative abundance of different elements has an effect as well. Elements that affect core opacity  are, in the order of importance, iron, sulfur, silicon and oxygen. The elements that contribute to opacity in the region near the base of the convection zone and thereby affect the position of the base of the solar convection zone are, again in the order of importance, oxygen, iron and neon. Although, the main effect of heavy element abundances is through opacity, these abundances also affect the equation of state. In particular, the adiabatic index $\\Gamma_1$ is affected in regions where these elements  undergo ionization. This effect is generally small, but in the convection zone where the stratification is adiabatic and hence the structure is determined by equation of state rather than opacity, this effect can be significant.  The importance of the solar heavy-element abundance does not merely lie in being able to model the Sun correctly, it is  often used as the standard against which heavy-element abundances of other stars are measured. Thus the predicted structure of those stars too become uncertain if the solar heavy-element abundance is uncertain. Given that for most stars other than the Sun, we usually only know the position on the HR diagram, an error in the solar  abundance could lead to errors in the predicted mass and age of the stars. Stellar evolution calculations are used throughout astronomy to classify, date, and interpret the spectra of individual stars and of galaxies, and hence errors in metallicity affect age determinations, and other derived parameters of stars and star clusters. The exact value of the solar heavy-element abundance determines the amount of heavy elements that had been present in the solar neighborhood when the Sun was formed. This, therefore, determines the chemical evolution history of galaxies.  Solar models in the 1990's were generally constructed with the solar heavy element mixture of Grevesse \\& Noels (1993). The ratio of the mass fraction of heavy elements to hydrogen in the Sun was determined to be $Z/X= 0.0245$. Grevesse \\& Sauval (1998; henceforth GS98) revised the abundances of oxygen, nitrogen, carbon and some other elements, and that resulted in $Z/X=0.023$. In a series of papers Allende-Prieto et al.~(2001, 2002) and Asplund et al.~(2004, 2005a) have revised the spectroscopic determinations of the solar photospheric composition. In particular, their results indicate that  carbon, nitrogen and oxygen abundances are lower by about 35\\% to 45\\% than those listed by GS98.  The revision of the oxygen abundance leads to a comparable change in the abundances of neon and argon since these abundances are generally measured through the abundances ratio for Ne/O and Ar/O. Additionally, Asplund (2000) also determined a somewhat lower value (by about 10\\%) for the photospheric abundance of silicon compared with the GS98 value. As a result, all the elements for which abundances are obtained from meteoritic measurements have seen their abundances reduced by a similar amount. These measurements have been summarized by Asplund, Grevesse \\& Sauval (2005b; henceforth AGS05). The net result of these changes is that $Z/X$ for the Sun is reduced to $0.0165$ (or $Z=0.0122$), about 28\\% lower than the previous value of GS98 and almost 40\\% lower than the old value of Anders \\& Grevesse (1989). The change in solar abundances implies large changes in solar structure as well as changes in quantities derived using solar and stellar models, and therefore, warrants a detailed discussion of the consequence of the changes, and how one can test the new abundances. In this paper we review  the effects of solar abundances on  solar models and how the models with lower abundances stand up against helioseismic tests.  The review is written in a pedagogical style with  detailed explanations of how the analysis is done. However, the review is organized in such  a manner that not all readers need to read all sections unless they want to. We start with a description of how solar models are constructed (\\S~\\ref{sec:models}), this section also describes the sources of uncertainties in solar models (\\S~\\ref{subsec:ssmerr}). In \\S~\\ref{sec:helioseis} we describe how helioseismology is used to test solar models as well as how helioseismology can be used to determine solar parameters like the convection-zone helium abundance, the depth of the convection zone and how input physics, like the equation of state, can be tested. In \\S~\\ref{sec:res} we describe what helioseismology has taught us about the Sun  and inputs to solar model thus far. Thus readers who are more interested in helioseismic results rather than techniques can go directly to this section. In \\S~\\ref{sec:solarz} we give a short review of how solar abundances are determined and the results obtained so far. The consequences of the new abundances are described in \\S~\\ref{sec:abund}. This section also includes a brief summary of the changes in the solar neutrino outputs (\\S~\\ref{subsec:radint}) and some consequences of the new abundances on models of stars other than the Sun (\\S~\\ref{subsec:stellar}), though the latter discussion is by no means complete and comprehensive. Readers who are only interested in how the lower solar abundances affect the models would perhaps wish to go straight to this section. The next section, \\S~\\ref{sec:recon}, is devoted to reviewing the numerous attempts that have been made to reconcile the low-metallicity solar models with the Sun by changing different physical inputs. In \\S~\\ref{sec:seisz} we describe attempts that have been made to determine solar metallicity using helioseismic techniques. In \\S~\\ref{sec:problems} we discuss some possible reasons for the discrepancy between helioseismically determined abundances  and the new spectroscopic abundances  and we present some final thoughts in \\S~\\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  The discussions presented in the  previous sections clearly indicate that there is a large discrepancy between the heavy-element abundances needed to make solar models consistent with helioseismic data and the lowered solar heavy element abundances as compiled by AGS05. The disagreement between the models and the Sun, as determined from helioseismology, covers almost the whole of the solar interior. As discussed in \\S~\\ref{sec:abund}, the most obvious discrepancies between these models and the Sun are the position of the base of the convection zone and the helium abundance in the convection zone. Careful investigations, however, have shown that the models are discrepant in the ionization zones in the upper part of the convection zone, and in the core as well. Various suggestions have been put forward to bring the models using reduced abundances back in concordance with helioseismic data, none of the suggestions however, seem to work very well.  Since the publication of the AGS05 abundances, there have been attempts to determine solar abundances from helioseismic data, as has been discussed in \\S~\\ref{sec:seisz}. All the seismic estimates point towards a high solar $Z$ compared with AGS05. Since the $Z$-dependence of solar  oscillations frequencies is not direct, but occurs through the  influence of $Z$ on input microphysics such as opacities, equation of state, nuclear reaction rates etc., it is  possible, in principle,  that the errors in the input physics results in discrepant values of $Z$ obtained from helioseismology on one hand and the analysis technique used by Allende Prieto et al.~(2001, 2002) and Asplund et al.~(2004, 2005a) on the other. However, the seismic $Z$ determinations have been done using techniques that depend on different inputs, and despite the differences in techniques and dependences on different inputs, all seismic estimates of $Z/X$ are consistent with the higher GS98 abundances, and they agree with each other as well. Thus the question arises as to  what could be causing the discrepancy between the seismic abundances on one hand and the abundances determined from spectral lines listed in AGS05.  If the convection-zone abundances of the Sun are indeed consistent with the low abundances compiled by AGS05, then almost all the input physics that goes into construction of stellar models are much more uncertain than they are normally assumed to be. On the other hand, if the GS98 abundances are correct then the currently known input physics is consistent with seismic data. Thus  it is either a  sheer coincidence that the errors in the input physics cancel each other when GS98 abundances are used, or the recently revised abundances due to AGS05 need to be revised upwards. The only other option is that some fundamental process is missing in the theory of stellar structure and evolution, but it is difficult to speculate what that could be. It is easier to find reasons  that could cause the thermal stratification of the new three-dimensional model atmospheres to be modified and give results different from AGS05.  The serious disagreements between helioseismic estimates and recent spectroscopic estimates of the solar heavy-element abundance requires a careful examination of solar  atmospheric models, and also  models of the solar interior. This is important not only to resolve the discrepancy, but also because solar abundances are used as the scale for abundances of other astronomical objects. A lot of improvements in solar models during the last three decades of the previous century were fueled by the solar neutrino problem. The discrepancy between the observed neutrino fluxes and those predicted by the solar models led to critical examination of all input physics that goes in construction of solar models. These developments, including the seismic studies, finally led to confirmation of new physics beyond the standard model of particle physics. We can expect that the discrepancy caused by revision of solar heavy element abundances will lead to further improvements in  models of the solar atmosphere and perhaps of the solar interior as well."
    },
    "0711/0711.3465_arXiv.txt": {
        "abstract": "The Balloon-borne Large Aperture Submillimeter Telescope (\\blast) is a sub-orbital surveying experiment designed to study the evolutionary history and processes of star formation in local galaxies (including the Milky Way) and galaxies at cosmological distances. The BLAST continuum camera, which consists of 270 detectors distributed between 3 arrays, observes simultaneously in broad-band (30\\%) spectral-windows at 250, 350, and 500\\,\\micron. The optical design is based on a 2\\,m diameter telescope, providing a diffraction-limited resolution of 30\\arcsec\\ at 250\\,\\micron. The gondola pointing system enables raster mapping of arbitrary geometry, with a repeatable positional accuracy of ${\\sim}$\\,30\\arcsec; post-flight pointing reconstruction to $\\lesssim 5$\\arcsec\\ rms is achieved.  The on-board telescope control software permits autonomous execution of a pre-selected set of maps, with the option of manual override.  In this paper we describe the primary characteristics and measured in-flight performance of \\blast\\@. \\blast\\ performed a test-flight in 2003 and has since made two scientifically productive long-duration balloon flights: a 100-hour flight from ESRANGE (Kiruna), Sweden to Victoria Island, northern Canada in June 2005; and a 250-hour, circumpolar-flight from McMurdo Station, Antarctica, in December 2006. ",
        "introduction": "\\label{sec:intro}  We have built and flown a balloon-borne submillimeter observatory, designed to study the star formation in our Galaxy, in nearby resolved galaxies, and in high-redshift starburst galaxies by observing at several hundred microns, a wavelength range which is not available from the ground.  The Balloon-borne Large Aperture Submillimeter Telescope (\\blast) is a stratospheric balloon-borne 2\\,m telescope which observes the sky with bolometric detectors operating in three 30\\% wide bands at 250, 350, and 500\\,\\micron.  The diffraction-limited optics are designed to provide \\blast\\ with a resolution of $30\\arcsec$, $42\\arcsec$, and $60\\arcsec$ at the three wavebands, respectively.  The detectors and cold optics are adapted from those to be used on the SPIRE instrument on {\\sl Herschel} \\citep{grif03}.  \\blast\\ addresses important Galactic and cosmological questions regarding the formation and evolution of stars, galaxies and clusters \\citep{devlin04} by providing the first large-area (${\\sim}$\\,0.8--200\\,deg$^2$) surveys of unique spectral-coverage, angular resolution and sensitivity.  The primary scientific goals of \\blast\\ are: (i) to conduct confusion-limited and shallower wide-area extragalactic surveys to constrain the redshift-distribution, star formation history, and evolution of optically-obscured luminous galaxies by measuring photometric-redshifts \\citep[derived from the \\blast\\ colors and complementary data,][]{hughes02}; (ii) to study the spatial clustering of this population; (iii) to improve our understanding of the earliest stages of star formation by determining the physical properties and mass-function of cold pre-stellar cores and the efficiency of star formation within different Galactic environments; and (iv) to investigate the nature and structure of the interstellar medium by making the highest resolution maps to date of diffuse Galactic emission at these wavelengths.   The \\blast\\ bands bracket the peak of the thermal radiation emitted by dust at temperatures of around 10\\,K\\@.  Assuming that the temperature range of dust in submillimeter galaxies is ${\\sim}$\\,30--60\\,K, \\blast\\ can explore redshifts from 1 to 5 \\citep{hughes02}.  The primary advantage of \\blast\\ over existing submillimeter bolometer arrays such as {SCUBA} and {SHARC} is its enhanced sensitivity at wavelengths $\\le 500$\\,\\micron\\ (where the far-IR background peaks) due to the increased atmospheric transmission at balloon altitudes. With sensitivities of 250\\,mJy\\,s$^{1/2}$, the \\blast\\ mapping speed is 10 times faster than the design goal for {SCUBA}\\,2 at 500\\,\\micron\\ \\citep{holland06} and more than 100 times the mapping speed of the pioneering flights of the FIR balloon-borne telescope {PRONAOS} \\citep{lam98}.  \\blast\\ is complemented at shorter wavelengths by surveys from {\\it IRAS}, {\\it ISO}, {\\it Spitzer} and {\\it Akari}, and at longer wavelengths by SCUBA~2, LMT, and ALMA, to constrain spectral energy distributions\\@. {\\it Spitzer\\/}'s higher resolution will also provide accurate astrometry for many of the \\blast\\ sources.  \\blast\\ complements the {\\it Herschel}\\ satellite by testing detectors and filters which are similar to those to be used in the SPIRE instrument \\citep{grif03}. Furthermore, the results from \\blast\\ will be available early enough to influence the design (depth and area) of future {SPIRE} surveys.  In addition, \\blast\\ will complement the large-area spectroscopic Galactic surveys of {\\it SWAS}, provide submillimeter targets for the Fabry-Perot spectrograph {SAFIRE} on {\\it SOFIA}, and impact the design of the scientific case for the next generation of submillimeter surveys from space (e.g.\\ {\\it SPICA} and {\\it SAFIR}, and later {\\it SPIRIT} and {\\it SPECS}).  \\blast\\ has made three flights to date. A 24-hour test flight (\\blastfs) from Fort Sumner, NM, in September 2003 demonstrated the performance of the instrument sub-systems.  \\blast\\ made its first science flight (\\blastkir) in June of 2005, flying on a $1.1 \\times 10^6$\\,m$^3$ balloon from the Swedish Space Corporation base of ESRANGE, near Kiruna, Sweden to Victoria Island in northern Canada.  During this 4.5-day flight at an average altitude of 38\\,km, the instrument performed well, except for degraded optical performance. This was possibly due to a failure of structural elements in the carbon fiber mirror. The result was decreased resolution, which had a significant impact on our ability to achieve our extragalactic science goals\\@.  \\blast\\ acquired 100 hours of data on Galactic targets, providing some of the first arcminute resolution images at these wavelengths \\citep{Chapin2008}.  In this flight, surveys of five star forming regions were conducted, including the well-studied Cygnus-X field, three regions of intermediate/high-velocity cirrus, the Cassiopeia-A supernova remnant, and several individual bright targets.  \\blast\\ made its second science flight (\\blastmcm) in December, 2006, flying from the Williams Field Long Duration Balloon (LDB) facility near McMurdo Station in Antarctica.  In this flight, the instrument met its performance goals in resolution, sensitivity, and pointing. The flight was terminated after 250 hours of data were acquired. Unfortunately, an anomaly with the parachute separation system prevented the parachute from being released from the gondola after landing.  The parachute dragged \\blast\\ for 200\\,km before it came to rest in a crevasse field on the Antarctic Plateau.  While most of the instrument was destroyed, the hard drives containing the data were found nearby, and the complete dataset recovered.  In this flight, \\blast\\ conducted shallow (10\\,deg$^2$) and deep, confusion-limited (0.8\\,deg$^2$) extragalactic surveys of the Chandra Deep Field South (CDF-S) as well as a 10\\,deg$^2$ region near the South Ecliptic Pole. A Galactic survey in Vela mapped 52\\,deg$^2$ and 200\\,deg$^2$ of sky in deep and shallow observations, respectively.  Surveys of several Galactic and extragalactic selected targets were also conducted.  The telescope and optics are described in Section~\\ref{sec:telescope}, while Sections~\\ref{sec:rec} and~\\ref{sec:cryo} focus on detector arrays and cryogenics.  Sections~\\ref{sec:gondola} through \\ref{sec:psens} provide detailed information regarding the gondola, command and control, and pointing systems.  Sections \\ref{sec:thermal} through \\ref{sec:performance} provide overviews of the unique thermal and power requirements, and a description of the pointing performance. ",
        "conclusions": "\\blast\\ is an instrument designed to probe the local and high-z Universe at short submillimeter wavelengths.  Over the course of two science flights, \\blast\\ has addressed a broad range of Galactic and extragalactic topics. The balloon environment requires a novel combination of ground and satellite technology and techniques.  This paper is intended to serve as a reference for a suite of science papers from the \\blastkir\\ and \\blastmcm\\ flights. \\\\  The \\blast\\ collaboration acknowledges the support of NASA through grant numbers NAG5-12785, NAG5-13301 and NNGO-6GI11G, the Canadian Space Agency (CSA), the Science and Technology Facilities Council (STFC), Canada's Natural Sciences and Engineering Research Council (NSERC), the Canada Foundation for Innovation, the Ontario Innovation Trust, the Puerto Rico Space Grant Consortium, the Fondo Istitucional para la Investigacion of the University of Puerto Rico, and the National Science Foundation Office of Polar Programs; C.~B. Netterfield also acknowledges support from the Canadian Institute for Advanced Research. L.~Olmi would like to acknowledge Pietro Bolli for his help with Physical Optics simulations during the testing phase of the \\blastmcm\\ telescope. We would also like to thank the Columbia Scientific Balloon Facility (CSBF) staff for their outstanding work, the Precision Machining Group at Lawrence Livermore Laboratory, the support received from Empire Dynamic Structures in the design and construction of the gondola, Daniele Mortari for helpful discussions in the development of the Pyramid code, Dan Swetz for buliding the Fourier transform spectrometer, and  Luke Bruneaux, Kyle Lepage,  Danica Marsden, Vjera Miovic, and James Watt for their contribution to the project."
    },
    "0711/0711.3186_arXiv.txt": {
        "abstract": "{We present the results of realistic, 3D, hydrodynamical, simulations of surface convection in red giant stars with varying effective temperatures and metallicities. We use the convection simulations as time-dependent, hydrodynamical, model atmospheres to compute spectral line profiles for a number of ions and molecules under the assumption of local thermodynamic equilibrium (LTE). We compare the results with the predictions of line formation calculations based on 1D, hydrostatic, model stellar atmospheres in order to estimate the impact of 3D models on the derivation of elemental abundances. We find large negative 3D$-$1D LTE abundance corrections (typically $-0.5$ to $-1$ dex) for weak low-excitation lines from molecules and neutral species in the very low metallicity cases. Finally, we discuss the extent of departures from LTE in the case of neutral iron spectral line formation. ",
        "introduction": "Classical spectroscopic abundance analyses of late-type stars normally make use of theoretical 1D model atmospheres constructed under the assumptions of plane-parallel geometry or spherical symmetry, hydrostatic equilibrium, and flux constancy. In late-type stars, the convection zone reaches and affects the stellar surface layers from which the outgoing stellar flux emerges; one-dimensional model stellar atmospheres, however, inherently rely on rather crude implementations of convective energy transport such as the mixing length theory \\citep{boehm-vitense58} or analogous alternative formulations \\citep[e.g.][]{canuto91}. In view of the dynamic and multi-dimensional nature of convection, the application of time-independent hydrostatic 1D model atmospheres to spectroscopic analyses is therefore a potential source of severe systematic errors. Moreover, analyses based on 1D models cannot predict strengths and profiles of spectral lines without invoking ad-hoc fudge parameters (e.g. micro- or macro-turbulence). In recent years, on the other hand, it has become possible to perform realistic 3D hydrodynamical simulations of stellar surface convection \\citep[e.g.][]{nordlund82,nordlund90,stein98,asplund99,freytag02,ludwig02,carlsson04,voegler04}. In the case of solar-like stars, such 3D simulations have been proven successful in reproducing self-consistently % observational constraints such as the topology of the granulation pattern and the detailed shape of spectral lines \\citep[see review by][]{asplund05}. Here, we will present some results of our recent 3D simulations of convection at the surface of red giant stars \\citep{collet06,collet07} focusing in particular on their application to spectral line formation. ",
        "conclusions": "We have presented here some illustrative results of the application of 3D red giant surface convection simulations to spectral line formation in LTE. The differences between the temperature stratifications of 3D simulations and corresponding 1D model atmospheres as well as the 3D temperature and density inhomogeneities and velocity gradients can significantly affect the predicted line strengths and, consequently, the value of elemental abundances inferred from spectral lines. At very low metallicities, where the deviations of the mean 3D thermal structure from the classical 1D stratification are largest, the 3D$-$1D LTE abundance corrections are negative and considerable for lines of neutral species (about $-0.8$~dex for weak low-excitation \\ion{Fe}{i} lines). Corrections to CNO abundances derived from CH, NH, and OH weak low-excitation are also found to be typically in the range $-0.5$~dex to $-1.0$~dex in very metal-poor giants \\citep[for further details see][]{collet07}. Finally, we have also examined possible departures of \\ion{Fe}{i} line formation from LTE; such non-LTE corrections are opposite to and, according to preliminary 1D test calculations of ours, of the same order of magnitude as the ones due to granulation."
    },
    "0711/0711.3653_arXiv.txt": {
        "abstract": "{Scalar--tensor gravity is the simplest and best understood modification of general relativity, consisting of a real scalar field coupled directly to the Ricci scalar curvature. Models of this type have self-accelerating solutions. In an example inspired by string dilaton couplings, scalar--tensor gravity coupled to ordinary matter exhibits a de Sitter type expansion, even in the presence of a {\\it negative} cosmological constant whose magnitude exceeds that of the matter density. This unusual behavior does not require phantoms, ghosts or other exotic sources. More generally, we show that any expansion history can be interpreted as arising partly or entirely from scalar--tensor gravity. To distinguish any quintessence or inflation model from its scalar--tensor variants, we use the fact that scalar--tensor models imply deviations of the post-Newtonian parameters of general relativity, and time variation of the Newton's gravitational coupling $G$. We emphasize that next-generation probes of modified GR and the time variation of $G$ are an essential complement to dark energy probes based on luminosity-distance measurements. } ",
        "introduction": "There are three phenomenological approaches to explaining the accelerating expansion of the universe. The first is a positive cosmological constant whose magnitude (considered as vacuum energy) somewhat exceeds the current matter density \\cite{Weinberg:1988cp}. The second is to introduce dynamical dark energy, usually in the form of an ultralight quintessence scalar field \\cite{Ratra:1987rm}-\\cite{Copeland:2006wr} The third is to modify gravity so as to produce self-accelerating solutions \\cite{Dvali:2000hr}-\\cite{Carena:2006yr}; usually such solutions can be regarded as modifying the left-hand side of the Einstein equations of motion in such a way as to simulate a dark energy source on the right-hand side of these equations.  Scalar--tensor models \\cite{Brans:1961sx}-\\cite{Nordtvedt:1970uv} can be regarded as a combination of the latter two approaches. General relativity is modified by a real scalar field $\\theta (x)$ that couples directly to the Ricci scalar curvature $R$. If the vacuum expectation value of $\\theta $ is dynamically evolving today, then the Einstein equations are modified in a nonlinear way and exhibit new types of solutions. In the absence of sources scalar--tensor models are classically equivalent to higher derivative modified gravity models based on a nonlinear function $f(R)$, but this equivalence almost certainly does not hold in realistic contexts. In fact scalar--tensor models have a big advantage over other approaches to modified gravity, in that it is transparent to identify regimes in these models that are weakly--coupled and free of ghosts, violations of the dominant energy condition, and other pathologies.  The existence of self-accelerating solutions of scalar--tensor gravity models has already received considerable attention \\cite{EspositoFarese:2000ij}-\\cite{Demianski:2007mz}. Strong observational constraints on such scalar--tensor cosmologies have been derived and discussed in the literature \\cite{Riazuelo:2001mg}-\\cite{Coc:2006rt}. Our purpose here is to exhibit some simple analytic solutions that demonstrate the promise, weaknesses and generality of accelerated expansion from scalar--tensor gravity. Although we will focus on the connection to dark energy, most of our analysis is also relevant for building models of primordial inflation.  In all of our examples the real scalar is ultralight. There are two known motivations for such fields. The first is string theory, in which the low energy effective action can exhibit a massless dilaton and other massless moduli fields. These generally have exponential couplings to the Ricci scalar. They may or may not have direct couplings to matter; when they do have such couplings, these may be sufficiently universal to satisfy the very strong constraints from equivalence principle tests \\cite{Carroll:1998zi}. The second motivation for ultralight scalars (or pseudoscalars) are the pseudo-Nambu Goldstone bosons of spontaneously broken global symmetries that have an additional breaking due to nonperturbative effects or to a weak explicit breaking \\cite{Hill:1988vm}-\\cite{Frieman:1995pm}. In simple examples the scalar potential of such PNGB's respects a discrete periodic remnant of their original shift symmetry. We will assume that scalar--tensor gravity implementations of this idea imply periodic functions of $\\theta$ coupling to $R$. ",
        "conclusions": "We have seen that scalar-tensor gravity models can have remarkably simple self-accelerating solutions, without resorting to phantoms or other ghosts. These solutions are made possible by the fact that the conserved energy associated to the scalar is not the conventional energy that one would read off from the lagrangian in the flat space limit.  As we have seen, this self-accelerating feature can even overcome, for a time, the effects of a negative cosmological constant of similar magnitude. In such a scenario, our immediate cosmological future exactly contradicts what one would predict from a naive extrapolation of the current expansion. Although the model we exhibited was not entirely realistic, it has an intriguing connection to previous attempts to construct realistic string models with ultralight moduli.  Scalar-tensor models are very constrained by observational data. It does not appear likely that one could account for dark energy entirely from the effects of scalar-tensor gravity, especially in a framework where the gravity sector is manifestly weakly coupled and ghost--free.  However we gave two examples of realistic models where the novel properties of scalar-tensor gravity play an important role in the current accelerated expansion. These models are somewhat tuned in order to satisfy present day limits on the time variation of $G$. They predict interesting variations of $G$, on the order or 10\\% , at earlier times. Improved CMB \\cite{Zahn:2002rr} or neutron star \\cite{Thorsett:1996fr} constraints would directly test such scenarios, as would galaxy cluster based tests of modified gravity \\cite{Zhang:2007nk,Schmidt:2007vj}.  More generally, our examples show the importance of using a broad-based observational approach to dark energy. The ambitious Stage IV dark energy probes currently planned \\cite{Albrecht:2006um} are certainly not sufficient by themselves \\cite{Boisseau:2000pr,Capozziello:2007iu} to disentangle scalar-tensor effects from quintessence and other scenarios.  The self-accelerating properties of scalar-tensor models look promising for models of primordial inflation. It appears that such models could have virtues similar to models of hybrid inflation \\cite{GarciaBellido:1995fz,Linde:1993cn}. Variations of $G$ at times prior to Big Bang Nucleosynthesis are hardly constrained. In this arena it also seems more promising to forge a concrete link between scalar-tensor gravity and string theory.  \\subsection*"
    },
    "0711/0711.0051_arXiv.txt": {
        "abstract": "We present the {\\it Spitzer}  Space Telescope InfraRed Array Camera (IRAC) and Multiband Imaging Photometer (MIPS) observations of the elliptical galaxy NGC~315.  After removal of the host galaxy's stellar emission, we detected for the first time an infrared-red nucleus in NGC~315. We measured the spectral energy distribution (SED) for this active nucleus with wavelength range covering from radio to X-ray, and obtained the bolometric luminosity of $\\rm L_{bol} \\approx 1.9 \\times 10^{43} ergs \\ s^{-1}$, corresponding to an extremely low Eddington ratio (L/L$_{\\rm Edd}$) of 4.97 $\\times$ 10$^{-4}$. Our results confirm that the physical nature of the nucleus of NGC 315 is a low-luminosity AGN, consistent with the recent optical and {\\it Chandra} X-ray observations. ",
        "introduction": "The recent discovery of a tight correlation between black hole (BH) mass and the bulge's mass of host galaxy suggests that the creation and growth of central black holes are an integral part of the formation of galactic bulges (Kormendy \\& Richstone 1995; Kauffmann \\& Heckman 2005). On the other hand, it also implies that supermassive black holes (SMBHs) are ubiquitous at the centers of elliptical galaxies. Although indeed most host galaxies of radio-loud and the brightest quasars at high-redshift are elliptical galaxies (Hutchings \\& Morris 1995; Bahcall, Kirkakos \\& Schneider 1996; Falomo et al. 2005), the majority of SMBHs in nearby ellipticals are not very active (Ho, Filippenko \\& Sargent 1997), and  therefore hardly detectable.  The low state of SMBHs in nearby ellipticals is most probably due to the lack of enough fuel supply - cold gas inflow into the central engine. It is generally accepted that elliptical galaxies are spheroidal distributions of older stars and lack substantial cold intersteller material (ISM). However, an H$_\\alpha +$ [NII] narrow-band imaging survey has shown that elliptical galaxies certainly have plenty of extended warm ($\\sim 10^4$ K) gas (Shields 1991). By detecting positive signal from co-adding images of the {\\it Infrared Astronomy Satellite} (IRAS) survey scans, Knapp et al. (1989) discovered that $\\sim$ 50 percent of optically-selected ellipticals with visual magnitude brighter than V=14 contained cool interstellar matter. Temi et al. (2004) estimated that 16 elliptical galaxies observed with the {\\it Infrared Space Observatory} (ISO) have dust masses on average 10 times greater than those previously detected by IRAS. Most recently, Kaneda, Onaka, \\& Sakon (2005)  detected even polycyclic aromatic hydrocarbon (PAH) emission features in four elliptical galaxies with the {\\it Spitzer} Infrared Spectrograph (IRS). Recent neutral hydrogen observations also show that plenty of extended gas is clearly present in early-type galaxies (Morganti et al. 2006; Noordermeer 2006 and references therein). All these results suggest that the presence of a significant amount of cold/warm interstellar gas at centers of elliptical galaxies is quite common, although the origin of such amount of dust and cold gas is still uncertain (Caon et al. 2000).  AGN signature of SMBHs has been seen in elliptical galaxies. In the recent Palomar spectroscopic survey, Ho, Filippenko \\& Sargent (1997) found that about 50\\% elliptical galaxies show detectable emission-line nuclei,  most of which ($\\sim$ 87\\%) are classified as LINERs. Although it is still a subject of hot debate whether LINERs are genuine low-luminosity AGNs (see the recent review by Ho 2003), AGN activity is confirmed by radio and X-ray observations (Terashima et al. 2002; Filho et al. 2004).  Typical AGNs are also strong infrared sources. Ho et al. (1997) found that about 55\\% of elliptical galaxies with emission lines were detected to have infrared emission by IRAS. Due to the lower spatial resolution ($\\sim$ 3'$-$4' ) of IRAS, it is impossible to determine whether  the far-infrared emission from elliptical galaxies arises from AGN point sources at the center, or from extended dust distribution seen as dust-lane in optical images (Lauer et al. 2005). One example is NGC 315, a well studied elliptical galaxy classified as a LINERs by Ho et al. (1997). The HST image shows a clear dusty disk with 2.5\" diameter in its center (Verdoes Kleijn et al. 1999) and the galaxy is also detected by IRAS at 12, 25, 60, and 100 \\micron\\ with flux densities of 0.081, 0.150, 0.270, and 0.706 Jy, respectively (IRAS Faint Source Catalogue, Moshir et al. 1990), corresponding to $L_{\\rm FIR}=2.3 \\times 10^9 L_\\odot$ at the distance of 65.8 Mpc (Ho, Filippenko \\& Sargent 1997). Clearly, a high-resolution infrared image of NGC 315 is required to determine whether its infrared emission comes from the central 2\" region.  The {\\it Spitzer Space Telescope} (Werner et al. 2004) has 2\" spatial resolution in mid-infrared (3.6\\micron\\ $-$ 8.0 \\micron) and 5\" at 24\\micron. This makes it possible to study the infrared emission from central regions of  elliptical galaxies. In this work we present a study of the nuclear activity in the elliptical galaxy NGC 315 based on observations by the Infrared Array Camera (IRAC, Fazio et al. 2004) and the Multiband Imaging Photometer for Spitzer (MIPS, Reike et al. 2004) aboard the {\\it Spitzer} Space Telescope. This paper is organized as follows: the IRAC and MIPS data reduction, flux calibration and photometry are in Sect. 2.  We present multi-wavelength physical properties of nucleus of NGC 315  in Sect. 3. ",
        "conclusions": ""
    },
    "0711/0711.0267_arXiv.txt": {
        "abstract": "{} {Study the contradictory magnetic field strength distributions retrieved from independent analyses of spectropolarimetric observations in the near-infrared ($1.56$ $\\mu$m) and in the visible (630 nm) at internetwork regions.} {In order to solve this apparent controversy, we present simultaneous and co-spatial $1.56$ $\\mu$m and 630 nm observations of an internetwork area. The properties of the circular and linear polarization signals, as well as the Stokes V area and amplitude asymmetries, are discussed. As a complement, inversion techniques are also used to infer the physical parameters of the solar atmosphere. As a first step, the infrared and visible observations are analysed separately to check their compatibility. Finally, the simultaneous inversion of the two data sets is performed.} {The magnetic flux densities retrieved from the individual analysis of the infrared and visible data sets are strongly correlated. The polarity of the Stokes V profiles is the same at co-spatial pixels in both wavelength ranges. This indicates that both \\hbox{$1.56$ $\\mu$m} and 630 nm observations trace the same magnetic structures on the solar surface. The simultaneous inversion of the two pairs of lines reveals an internetwork full of sub-kG structures that fill only 2\\% of the resolution element. A correlation is found between the magnetic field strength and the continuum intensity: equipartition fields ($B\\sim 500$ G) tend to be located in dark intergranular lanes, whereas weaker field structures are found inside granules. The most probable unsigned magnetic flux density is \\hbox{10 Mx/cm$^2$}. The net magnetic flux density in the whole field of view is nearly zero. This means that both polarities cancel out almost exactly in our observed internetwork area.} {} ",
        "introduction": "Solar magnetic fields as seen in magnetograms form a particular pattern outside active regions, in the so-called quiet Sun. The largest polarimetric signals are confined to the supergranular boundaries forming the photospheric network. Inside these large-scale convective cells it is possible to detect magnetic fields with a more disperse character and with smaller polarimetric signals. These areas are called the internetwork. The magnetic structures of the network can be modeled either with a flux tube model or with the MISMA hypothesis \\citep[MIcroStructured Magnetic Atmosphere,][]{jorge_96}. Both models reproduce the observed Stokes profiles and agree on the existence of nearly vertical magnetic fields of kG strength occupying only some 10-20 \\% of the resolution element. Observations at $1.56$ $\\mu$m \\citep[][]{muglach_92, lin_95}, \\hbox{525 nm} \\citep[][]{stenflo_73} and 630~nm \\citep[][]{jorge_00} give consistent results. Unfortunately, this is not the case for the internetwork, with recent studies based on high quality spectro-polarimetric data at $1.56$ $\\mu$m and $630$ nm yielding contradictory results.  The first observation of Stokes V spectra of an internetwork region was carried out by \\cite{keller_94}. These authors observed the pair of Fe\\,{\\sc i} lines at \\hbox{$525$ nm} and analysed four average Stokes V spectra using the line ratio technique \\citep{howard_71}. Their conclusion was that the magnetic field strength was below \\hbox{$1$ kG (500 G)} with a probability of $95$\\% (68\\%). Later, \\cite{lin_95} measured directly the Zeeman splitting of the Fe\\,{\\sc i} line pair at $1.56$ $\\mu$m to infer the magnetic field strength. He showed that the typical magnetic field strength in the internetwork was well below kG. The retrieved kG structures were correlated with network patches. A few years later, \\cite{lin_99}, using the same pair of infrared lines, deduced field strengths between $0.2$ and $1$ kG.  The success of the Advanced Stokes Polarimeter \\citep{lites_93} filled the literature with works concerning the internetwork magnetism using the Fe\\,{\\sc i} pair of lines at 630 nm. Most of these studies obtained that the field strength distribution at the internetwork presents a peak at kG. These kG fields, even if they occupy a small fraction of the Sun's surface (0.1-1\\%), contained most of the magnetic flux of the internetwork. Applying a Principal Component Analysis (PCA) inversion, \\cite{hector_02} also recovered magnetic field strengths of the order of kG. The database for this PCA algorithm consisted of synthetic profiles generated under the MISMA hypothesis. The most surprising fact is that, analysing the same data in terms of a Milne-Eddington inversion, \\cite{lites_02} obtained field strengths clearly below 1 kG. Using the line ratio technique, \\cite{ita_jorge_03} confirmed the presence of ubiquitous kG fields in the internetwork. They used speckle reconstructed data in order to achieve a spatial resolution of $0.5''$. %  \\cite{khomenko_03} carried out the study of an internetwork region at $1.56$ $\\mu$m, measuring the Zeeman splitting to deduce the magnetic field strength. These authors obtained that the internetwork magnetic field strength was well below the kG regime, with a probability distribution function that could be reproduced by a decreasing exponential with an e-folding factor of $250$ G.  From all these works, one can summarize that studies using the 630 nm or the 1.5 $\\mu$m lines have led to contradictory results: the infrared lines indicate fields below 1 kG, while the visible lines indicate a dominance of kG fields. Two plausible explanations have been presented to explain such a discrepancy. One is based on the idea that visible and near-infrared observations contain incompatible information about the internetwork magnetism \\citep{hector03, jorge_ita_03}. The other suggests that the line pair at $630$ nm is not reliable for retrieving the magnetic field strength at the internetwork \\citep{luis_03, marian_06, arturo_07}.  \\cite{hector03} suggested that, if kG and sub-kG fields co-exist in the same resolution element, the visible lines would be sensitive to the kG contribution, whereas the infrared lines would trace the sub-kG structures. The explanation comes from the different sensitivity of the spectral lines to the magnetic field. According to \\cite{luis_03}, this may be the situation in the case of having kG fields occupying $\\sim$ 20\\% of the resolution element and sub-kG fields filling the rest. If the percentages change, the premise seems not valid anymore. However, \\cite{hector_04} found observational evidence indicating the coexistence of sub-kG and kG fields in the internetwork. Their analysis was based on a three-component inversion, two of the components being magnetic with fixed field strengths of 1700 and 500 G and the other one being field-free. The observations were compatible with $\\sim$ 20\\% of the resolution element filled by the \\hbox{1700 G} field and the rest by the $500$ G field.  \\cite{luis_03} performed numerical simulations of Stokes profiles at 630 nm and at 1.56 $\\mu$m, assuming a magnetic field filling 5\\% of the resolution element. For consistency with the results by \\cite{khomenko_03}, they parameterized the field strength distribution as a decreasing exponential with a mean value of 250 G. After adding some noise (to a level of 10$^{-3}$ I$_\\mathrm{c}$, with I$_\\mathrm{c}$ the continuum intensity) to the synthetic profiles and inverting them, the resulting magnetic field strength distribution retrieved from the synthetic visible data set was centered at kG. On the contrary, the distribution inferred from the synthetic infrared spectra was close to the original one. They concluded that, for a sufficiently small noise level, both infrared and visible observations should agree and retrieve the initial distribution. However, and advancing some of the results that we will show in this paper, other effects apart from the noise level make the \\hbox{630 nm} magnetometry unreliable. This follows from the results by \\cite{marian_06} that, for the present observing conditions (with magnetic flux densities of $\\sim$ 10 Mx/cm$^2$, noise level of \\hbox{$6.5 \\times 10^{-5}$ I$_\\mathrm{c}$}, spatial resolution from 0.5$''$ to 1$^{\\prime\\prime}$), the 630 nm pair of lines does not carry enough information to retrieve the magnetic field strength.  \\begin{figure} \\centering \\includegraphics[width=7cm]{calcio.eps} \\caption{Ca\\,{\\sc ii} K line core image from the VTT slit jaw system. Bright features are good indicators of magnetic activity. The black square marks the area scanned for this study. The lack of intense bright points in the observed area indicates a very quiet region without intense magnetic activity.} \\label{slitjaw} \\end{figure}  Apart from numerical simulations, simultaneous and co-spatial observations in both spectral ranges should give the final answer to the 1.56 $\\mu$m-630 nm controversy. The first work with this aim was carried out by \\cite{jorge_ita_03} using data from the Vacuum Tower Telescope (VTT) and THEMIS telescopes at the Observatorio del Teide. This yielded two data sets in the $630$ nm and \\hbox{$1.56$ $\\mu$m} lines, but with a significantly different spatial resolution since the THEMIS telescope was not equipped with an image stabilization system at that time. The separate inversion of both data sets with a Milne-Eddington approximation led to field strengths of about $300$ G for the infrared lines and $1100$ G for the visible ones. More surprisingly, $25$\\% of the analysed signals showed opposite polarities in the same pixel at both wavelength ranges. \\cite{khomenko_shelyag_05} showed, using magnetohydrodynamic simulations, that a different spatial resolution in both wavelengths can result in apparent opposite polarities at the same spatial point. In a subsequent publication, \\cite{ita_06} degraded the spatial resolution of the VTT data used by \\cite{jorge_ita_03} to match that of the THEMIS observations.% From a simultaneous MISMA inversion they also obtained an important fraction of kG fields in the visible and weaker fields in the infrared. Consequently, they concluded that the information contained in the infrared data was incompatible to that contained in the visible data set. They assert that kG and weaker magnetic fields were coexisting in the resolution element, the visible being sensitive to the kG fields and the infrared to the weaker ones.  Recently, the polarimeter on-board HINODE has provided spectro-polarimetric observations at 630 nm with the highest spatial resolution ever achieved with this kind of data (0.3$''$). If the occupation fraction of the magnetic elements in the resolution element would have increased at least one order of magnitude from 1$''$ to 0.3$''$ one would have been capable of separating the effects of the thermodynamics and the magnetic field strength \\citep{andres_marian_07_astroph}. \\cite{david_07_astroph} perform a Milne-Eddington inversion of HINODE's data on a quiet Sun region. The noise level is around one order of magnitude higher than the one of the observations presented in this paper. In order to perform the inversions they assume one single magnetic atmosphere in the resolution element and some contamination of stray light. The retrieved magnetic field strength distribution has a peak at very weak fields, in disagreement with the previous results using the pair of spectral lines at 630 nm. However, even if this result is compatible with the magnetic fields obtained using the infrared lines at 1.5 $\\mu$m, the validity of the visible spectral lines at the HINODE observational conditions should be tested.  In this work we present simultaneous and co-spatial $1.56$ $\\mu$m and 630 nm observations taken {\\it at the same telescope, under very similar seeing conditions and using image stabilization systems}. Spurious effects arising from a different spatial resolution in both data sets, due to different observing conditions, are thus avoided. Prior to the analysis, noise is reduced to very low values using a PCA procedure to ensure that it does not bias the study of the compatibility of the internetwork observations at these two wavelengths.     \\begin{figure} \\includegraphics[width=4.3cm]{mapai_tip.eps} \\hspace{0.3cm} \\includegraphics[width=4.3cm]{mapav_tip.eps}\\\\ \\\\ \\includegraphics[width=4.3cm]{mapai_polis.eps} \\hspace{0.3cm} \\includegraphics[width=4.3cm]{mapav_polis.eps} \\caption{Maps of continuum intensity and integrated absolute circular polarization for the aligned infrared (top) and visible (bottom) data. A grid is superposed for an easier visibility of the co-spatiality of structures.} \\label{mapas_alineados} \\end{figure} ",
        "conclusions": "Simultaneous and co-spatial observations of Fe\\,{\\sc i} $1.56$ $\\mu$m and Fe\\,{\\sc i} $630$ nm pair of lines have been analysed. A total of 93\\% of the observed field of view has polarimetric signals above three times the noise level. This means that at 1$''$ spatial resolution the internetwork is full of magnetic fields.  In order to study the internetwork magnetism, the information of the full Stokes vector has to be taken into account. At this point the $1.56$ $\\mu$m lines present a big advantage when compared to the 630 nm ones. Both pairs of lines are equally sensitive to the circular polarization. However, the infrared lines show conspicuous linear polarization signals while the visible ones show these signals only in few weak features.  The Stokes V profiles are quite asymmetric at both spectral ranges, having similar distributions of area and amplitude asymmetries. Most of the observed two-lobed profiles have an amplitude asymmetry of about 15\\% and an area asymmetry mainly zero.  We have performed the separate inversion of the infrared and visible data in order to check their compatibility. Our conclusions are: \\begin{itemize} \\item The magnetic flux density recovered from the independent inversion is the same at both spectral ranges (both spatial distributions have a strong correlation). Small differences can be produced by the different formation heights of the two pairs of lines. \\item There are no systematic opposite polarities for co-spatial Stokes V profiles in both observations. \\item The magnetic field strength distributions are very different. The infrared observations result in equipartition or even weaker fields while the visible lines show fields in the kG regime. \\end{itemize}  These three points are compatible since the validity of the magnetic field strength distributions retrieved from the 630 nm data has been questioned. We conclude that the information of our observations is compatible in both spectral ranges. This means that both the $1.56$ $\\mu$m and 630 nm lines trace the same magnetic structure on the solar surface. Anyway, it does not mean that the infrared distribution of magnetic field strengths is the real one. One should keep in mind that the Zeeman effect  is only capable of detecting magnetic fields in $\\approx \\sim$2\\% of the resolution element. What we really know is that our observations lead to reliable results and that the magnetic field strengths that we infer are really present on the internetwork. Ram\\' irez V\\'elez, L\\'opez Ariste \\& Semel (private comunication), with observations of a Mn\\,{\\sc i} line in the visible that is sensitive to the magnetic field strength, have concluded that the magnetic field strengths on the internetwork are well below the kG regime. This supports the fact that the infrared is not biased towards weak fields.  In order to put more physical constraints, we have performed the simultaneous inversion of both data sets. The inferred magnetic field strength distribution is dominated by sub-kG fields. The occupation fraction of those fields peaks at $2$\\% but with a broad distribution. Two different magnetic field strength distributions are present on granules and intergranules. The granular one has intrinsically weaker field strengths than the intergranular one, which is centered on the equipartition field of the photosphere. The magnetic flux density distribution is the same on granules and intergranular lanes. There is no net magnetic flux density in the studied field of view, implying that there is a perfect cancellation of positive and negative magnetic flux density.  The unsigned magnetic flux density histograms peak at \\hbox{$\\sim 10$ Mx/cm$^2$}. \\cite{martin_87} and \\cite{wang_95} obtained the same value using magnetograms with $2-3''$ spatial resolution. This is agreement with \\cite{lites_04}, who show that there is no apparent increment in the magnetic flux density when improving the spatial resolution from 1 to 0.6$''$.  Since the Zeeman effect allows us to study $\\sim$ 2\\% of the resolution element, what happens in the rest of the surface? Is it field-free? If not, which kind of magnetic fields can we find in it? \\cite{andres_07} show the first experimental evidence of magnetic flux cancellation at the internetwork. This supports the fact that the peak of the circular and linear polarization histograms well above the noise level is produced by cancellations in the resolution element. This means that there are still undetected magnetic fields at a spatial resolution of 1$''$ (mixed polarities or extremely weak fields). The study of the internetwork magnetism has become a very interesting field that deserves more efforts to know if it can play an important role in the solar global magnetism as claimed by \\cite{javier_04}."
    },
    "0711/0711.2923_arXiv.txt": {
        "abstract": "We present a supplement to the Macquarie/AAO/Strasbourg H$\\alpha$ planetary nebulae (PNe) catalogue (MASH), which we denote MASH-II. The supplement consists of over 300 true, likely and possible new Galactic PNe found after re-examination of the entire AAO/UKST H$\\alpha$ survey of the southern Galactic Plane in digital form. We have spectroscopically confirmed over 240 of these new candidates as bona-fide PNe and we include other high quality candidates awaiting spectroscopic confirmation as possible PNe. These latest discoveries largely comprise two distinct groups: small, star-like or moderately resolved PNe at one end and mostly large, extremely low surface brightness PNe at the other. Neither group were easy to discover from simple visual scrutiny of the original survey exposures as for MASH but were relatively straightforward to uncover from the digital images via application of semi-automated discovery techniques. We suspect the few PNe still hidden in the H$\\alpha$ survey will lie outside our search criteria or be difficult to find. ",
        "introduction": "Planetary Nebulae (PNe) are a very brief ($\\sim$10,000--50,000 year) phenomenon exhibited by most stars of low to intermediate mass after their post AGB phase when the previously ejected shroud of enriched and processed stellar material becomes ionised by the hot, remnant stellar core as it evolves towards a white dwarf. PNe are an essential means of understanding the late evolution of low-mass stars like our Sun. PNe hold the key to determining the physics of mass loss and time-scales of low mass stellar evolution. Since low mass stars can each recycle a significant fraction of their mass into the ISM, PNe are vital probes of stellar nucleosynthesis processes and ISM enrichment. Their strong emission line spectra make them visible at large distances, enabling kinematics and abundances of elements generally difficult to measure in stars to be determined.  The properties of their progenitor stars can be investigated and can yield important information on core mass, temperature and stellar evolutionary history which ties into their observed nebula properties. If distances can be determined for example, from measured parallaxes (Harris et al. 2007) or photometric parallaxes (Ciardullo et al. 1999), by assuming a fixed distance in other galaxies (e.g. Reid \\& Parker 2006), or more generally via statistical methods such as the new H$\\alpha$ surface brightness radius relation of Frew \\& Parker (2006), then observed properties can unlock true physical parameters. These include the size of the inner shell and outer faint AGB halo (Corradi et al. 2003), total nebular mass and central star brightness and evolutionary state.  The PNe population of our Galaxy from observation based estimates is 23,000$\\pm$6,000 disc PNe (Zijlstra \\& Pottasch 1991) or 28,000$\\pm$5,000 in total (Frew \\& Parker 2006). A larger value of 46,000$\\pm$13,000 PNe has been obtained from stellar population synthesis models (Moe \\& De Marco 2006), though this is only for PNe with radii $<$ 0.9 pc and so is not strictly comparable with the aforementioned estimates. It is however clear that all estimates far exceed the number of PNe actually catalogued so far ($\\sim$2,500) and that Galactic numbers are dwarfed by large numbers of known extragalactic PNe (Magrini 2006). Such disparity can be attributed to a combination of heavy obscuration by dust in the Galactic Plane, the short lifetime of PNe, and the extraordinary diversity in observed PN properties. The observables which exhibit the most diversity, and thus influence any survey aiming for high completeness the most, are the surface brightness, angular extent and morphology (that may have been modified by an interaction with the ISM). To address the discrepancy in PN numbers sky surveys with significant areal coverage, sensitivity and angular resolution over suitable passbands are thus required to uniformly sample the whole population.  Historically, surveys have had difficulty meeting these requirements simultaneously and were often restricted to one wavelength domain (for an extensive overview see Parker et al. 2006a). This has led to significant biases and incompleteness in PNe catalogues limiting the range of observed evolutionary diversity. Substantial progress has been made recently in addressing these biases with the advent of modern H$\\alpha$ surveys such as the SuperCOSMOS H$\\alpha$ survey (see Section \\ref{sec:mash}) and the INT Photometric H$\\alpha$ Survey of the Northern Galactic Plane (IPHAS; Drew et al. 2005). A search for PNe within IPHAS is well underway (e.g. Mampaso et al. 2006, Viironen et al. 2006). For these reasons we expect large numbers to have eluded detection in previous surveys, and coupled with the aforementioned scientific drivers, we are strongly motivated to search for undiscovered PNe in our Galaxy.  This paper is structured as follows. Section \\ref{sec:mash} introduces the MASH PNe catalogue focusing on previous discovery techniques. Section \\ref{sec:mashii} describes our candidate selection and visualisation techniques. Section \\ref{sec:spec} discusses the spectroscopic confirmation of candidates which leads to the presentation of the MASH-II PNe catalogue in Section \\ref{sec:cat}. Completeness is discussed in Section \\ref{sec:comp} and we conclude in Section \\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc} The MASH-II catalogue of over 300 true, likely and possible PNe has been presented. We provide accurate positions, angular diameters and morphological classifications for all catalogue entries. Over 80 per cent of entries have been observed spectroscopically. The new PNe have been discovered by application of semi-automated discovery techniques applied to the online digital data of all 233 SHS fields for the first time. We used IAM photometry to target high surface brightness, small PNe and blocked-down quotient images of SHS fields to target low surface brightness, large PNe. Candidates were visualised using multi-wavelength data, making particular use of UKST SuperCOSMOS B$_\\mathrm{J}$ images at the earliest stages to increase our detection capability. This enabled many PNe to be found that may have escaped detection with simpler `two-colour' discovery techniques. MASH-II will strengthen the many science programs already underway with MASH (see Paper~I for details)."
    },
    "0711/0711.3322_arXiv.txt": {
        "abstract": "{} {We present the results of a contemporaneous photometric and spectroscopic monitoring of two RS~CVn binaries, namely \\object{$\\lambda$~And} and \\object{II~Peg}. The aim of this work is to investigate the behavior of surface inhomogeneities in the atmospheres of the active components of these systems which have nearly the same temperature but different gravity.} {The light curves and the modulation of the surface temperature, as recovered from line-depth ratios (LDRs), are used to map the photospheric spots, while the H$\\alpha$ emission has been used as an indicator of chromospheric inhomogeneities.} {The spot temperatures and sizes were derived from a spot model applied to the contemporaneous light and temperature curves. We find larger and cooler spots on II~Peg ($T_{\\rm sp} \\simeq 3600$\\,K) compared to $\\lambda$~And ($T_{\\rm sp} \\simeq 3900$\\,K); this could be the result of both the different gravity and the higher activity level of the former. Moreover, we find a clear anti-correlation between the H$\\alpha$ emission and the photospheric diagnostics (temperature and light curves). We have also detected a modulation of the intensity of the \\ion{He}{i} D$_3$ % line with the star rotation, suggesting the presence of surface features also in the upper chromosphere of these stars. A rough reconstruction of the 3D structure of their atmospheres has been also performed by applying a spot/plage model to the light and temperature curves and to the H$\\alpha$ flux modulation. In addition, a strong flare affecting the H$\\alpha$, the \\ion{He}{i} D$_3$, and the cores of \\ion{Na}{i} D$_{1,2}$ lines has been observed on II~Peg.} {The spot/plage configuration has been reconstructed in the visible component of $\\lambda$~And and II~Peg which have nearly the same temperature but very different gravity and rotation periods. A close spatial association of photospheric and chromospheric active regions, at the time of our observations, has been found in both stars. Larger and cooler spots have been found on II~Peg, the system with the active component of higher gravity and higher activity level.   The area ratio of plages to spots seems to decrease when the spots get bigger. Moreover, with the present and literature data, a correlation between the temperature difference $\\Delta T = T_{\\rm ph}-T_{\\rm sp}$ and  the surface gravity has been also suggested.} ",
        "introduction": "\\label{sec:Intro} The rotational modulation of brightness and other photospheric diagnostics in late-type stars is commonly attributed to starspots on their photospheres. The presence of chromospheric inhomogeneities similar to the solar plages is pointed out by several studies based on the H$\\alpha$ and/or other diagnostics \\citep[see, e.g., ][]{Stra93a, Cata96, Cata00, Fra98, Fra00a}.  The simultaneous study of photospheric (spots) and chromospheric (plages) active regions (ARs) is important for a better understanding of the physical processes occurring during the emersion of magnetic flux tubes from below the photosphere. Indeed, the relative position and size of ARs at different atmospheric levels can provide information on the magnetic field topology. In the past ten years, rotational modulation due to surface inhomogeneities at photospheric and chromospheric level has been revealed in several RS~CVn binaries, \\citep{Fra98, Cata00, Bia06}.  Simultaneous photometric and spectroscopic observations have frequently shown a spatial association between spots and plages in RS~CVn systems  \\citep{Rodo87,Cata96,Fra98} as well as in young mildly-active solar-type stars \\citep{Fra00a,Bia07}. A  close plage-spot association has been also detected in very active main-sequence single stars, like the rapidly rotating star LQ~Hya \\citep{Stra93a}. Also in close binary systems \\citep[e.g. TZ~CrB,][]{Fra97} and in extremely active stars, like the components of the contact binary VW~Cep \\citep{Fra96}, there are evidences of spot-plage associations. Moreover, for some RS~CVn binary systems there is some indication of a systematic longitude lag of 30$^\\circ$--50$^\\circ$ between the plages and spots \\citep{Cata96,Cata00}.   In recent works \\citep[][ hereafter Paper I and Paper II, respectively]{Cata02,Fra05}, we presented the results of a contemporaneous spectroscopic and photometric monitoring of three active single-lined (SB1) RS~CVn binaries (\\object{VY~Ari}, \\object{IM~Peg}, and \\object{HK~Lac}), showing that it is possible to recover fairly accurate spot temperature and size values by applying a spot model to contemporaneous line-depth ratios (LDRs) variation and light curves. Subsequently we have investigated the active region topology in these stars at both chromospheric and photospheric layers \\citep[][ hereafter Paper III]{Bia06}.  In the present paper, a similar analysis is applied to two other active SB1 RS~CVn's, namely $\\lambda$~And and II~Peg, which are stars with similar effective temperatures ($T_{\\rm eff}\\simeq$\\,4700 and 4600\\,K, respectively) but with different gravities,  being $\\log g\\simeq$\\,2.5 for $\\lambda$~And \\citep{Dona95} and $\\log g\\simeq$\\,3.2 for II~Peg \\citep{Berdy98b}.  $\\lambda$~And (HD~222107) is a bright ($V=3\\fm82$) and active giant, classified as G8 IV-III. \\citet{Calder38} was the first to discover its photometric variability, whose amplitude sometimes reaches 0$\\fm$30. Six years later, \\citet{Walker44} showed that $\\lambda$~And is a SB1 with an almost circular orbit of period 20$\\fd$5212. It is an atypical member of the RS~CVn class, because it is largely out of synchronism, its rotational period being 53$\\fd$952 \\citep{Stra93b}. Indeed, in most RS~CVn binaries the rotational period of both components is equal to the orbital period of the system within a few percent. Thus $\\lambda$~And, whose orbit is very close to be circular, is a puzzle for the theory of tidal friction \\citep{zahn77}, which predicts that rotational synchronization in close binaries should precede orbit circularization. It is one of the brightest of all known chromospherically active binaries. Photoelectric light curves of $\\lambda$~And \\citep{BoppNoah80,PoeEaton85} show that its photometric period is somewhat variable, probably due to differential rotation and latitude drift of the spots during magnetic cycles. In addition, there is a long-term cycle of about 11 yr in the mean brightness. Photometric long-term studies have been performed also by \\citet{Henry95}. Moreover, the H$\\alpha$ emission has been found to be rotationally modulated and anti-correlated with the light curve \\citep{BalDup82}.  II~Peg (HD 224085) has been classified as a K2-3 IV-V SB1 by \\citet{Ruci77} who noticed that its photometric period was quite close to 6$\\fp$724183, as determined by \\citet{Halli52} for its essentially circular orbit. It was classified as an RS~CVn system by \\citet{Vogt81a,Vogt81b} who found an amplitude of the $V$ light curve of 0$\\fm$43. II~Peg is among the most active RS~CVn binaries and belongs to a small subset of binaries, including V711~Tau and UX~Ari, in which H$\\alpha$ always appears in emission \\citep{NatRam81}. % The first photoelectric light curves have been published by \\citet{Chuga76}  % Long-term studies \\citep[e.g.,][]{Henry95, Rodo00} have shown dramatic changes of the photometric wave, from almost sinusoidal, to irregular or flat. Moreover, \\citet{Dona97} have clearly detected the magnetic field on this star. The so called flip-flop phenomenon (in which the dominant part of the spot activity changes the longitude every few years) has been reported in II~Peg by, e.g., \\citet{BerdyTuo98} and \\citet{Rodo00} and theoretically analyzed by \\citet{ElsKor05}.   The aim of the present work is to investigate the starspot characteristics of this two active stars with the same temperature but different gravity and activity level, as well as to study the location of the excess H$\\alpha$ emission and the degree of spatial superposition between surface inhomogeneities at different atmospheric levels. ",
        "conclusions": "We have presented a study of the surface inhomogeneities at both  photospheric and chromospheric levels based on a contemporaneous spectroscopic and photometric monitoring of the two active RS~CVn stars $\\lambda$ Andromedae and II~Pegasi.  From the same data set of medium-resolution optical spectra, we have obtained information about the chromospheric and photospheric surface inhomogeneities by using the H$\\alpha$ emission and the photospheric temperature (from line depth ratios), respectively. Additional information coming from the light curves, together with the temperature modulations, allowed us to disentangle the effects of starspot temperature and area and deduce these parameters in a unique way. A very important indication from this work is that the the starspots of $\\lambda$~And are considerably smaller and warmer than those of II~Peg, notwithstanding the nearly equal photospheric temperature. At present we cannot say if this is due to the different gravity of the two active stars, being $\\lambda$~And a giant with $\\log g\\simeq$\\,2.5 and II~Peg a subgiant ($\\log g\\simeq$\\,3.2), or if it is simply the effect of the higher activity level of II~Peg compared to  $\\lambda$~And. More active stars with different gravity and activity level must be investigated to settle this point. However, by using the values of temperature difference between photosphere and spot, $\\Delta T$, for other stars from our previous works and from the literature, we find  an increasing trend of $\\Delta T$ versus $\\log g$ that  could be explained by the magnetostatic equilibrium between gas and magnetic pressure.  We have found, for all the stars observed, a tight anti-correlation between the H$\\alpha$ emission and the photospheric temperature modulations, that indicates a close spatial association between photospheric spots and chromospheric plages.  The largest longitude shifts between plages and spots of about $\\pm 16\\degr$, have been found for $\\lambda$~And. Moreover, the area ratio of plages to spots seems to decrease when the spots get bigger.  Furthermore, in II~Peg a strong flare affecting both the H$\\alpha$ and \\ion{He}{i} lines has been observed. The energy losses in these lines have been also evaluated. A possible flare event, with a much smaller energy budget, seems to have occurred in  $\\lambda$~And. In the present work we have shown the great power of a coordinated photometric and spectroscopic monitoring of active stars for the study of the main properties of their active regions."
    },
    "0711/0711.3925_arXiv.txt": {
        "abstract": "Based on optical, IR and X-ray studies of Cas A, we propose a geometry for the remnant based on a ``jet-induced\" scenario with significant systematic departures from axial symmetry. In this model, the main jet axis is oriented in the direction of strong blue-shifted motion at an angle of 110 - 120\\degree\\ East of North and about 40 - 50\\degree\\ to the East of the line of sight. Normal to this axis would be an expanding torus as predicted by jet-induced models. In the proposed geometry, iron-peak elements in the main jet-like flow could appear ``beyond\" the portions of the remnant rich in silicon by projection effects, not the effect of mixing. In the context of the proposed geometry, the displacement of the compact object from the kinematic center of the remnant at a position angle of $\\sim$169\\degree\\ can be accommodated if the motion of the compact object is near to, but slightly off from, the direction of the main ``jet\" axis by of order 30\\degree. In this model, the classical NE ``jet,\" the SW ``counter-jet\" and other protrusions, particularly the ``hole\" in the North, are non-asymmetric flows approximately in the equatorial plane, e.g., out through the perimeter of the expanding torus, rather than being associated with the main jet. We explore the spoke-like flow in the equatorial plane in terms of Rayleigh-Taylor, Richtmyer-Meshkov and Kelvin-Helmholz instabilities and illustrate these instabilities with a jet-induced simulation. ",
        "introduction": "The star we now recognize as the supernova remnant Cassiopeia A probably exploded in 1680 \\citep{thors01,fesen07}. It has been an astrophysical mystery ever since its re-discovery by radio astronomers \\citep{ryle48} and has been the target of increasingly sophisticated studies over a range of wavelengths in the radio \\citep{and91,keo96,liszt99}, optical \\citep{fesen96,fesen01,fesenetal01,fesen06b,fesen06c,morse04}, IR \\citep{krause05,ennis06}, and X-ray \\citep{markert83,hughes00,willingale02, willingale03,hwang00,hwang01,hwang04,delaney04,laming06}. The projected image of Cas A shows a prominent jet-like structure in the NE and a ``counter-jet\" diametrically opposed in the SW. Other protrusions and interruptions in the images are notable in images at the various wavelengths. The first image taken with the {\\it Chandra Observatory} dramatically revealed the long-sought compact object. This dim object has yet to be given a firm physical identification as a neutron star or black hole. The compact object is displaced from the kinematic center of the remnant \\citep{fesen06b} to the South, nearly perpendicular to the locus of the ``jet/counter-jet\" structure. X-ray studies of the remnant have led to the discovery that iron-rich regions lie beyond silicon-rich regions \\citep{hughes00}, suggesting some sort of inversion of the expected ``onion-skin\" structure of the progenitor star. Fine-scale turbulence might mix the originally stratified composition and lead to Fe at larger radii than Si, but it is not clear that such a process can lead to a large scale segregation of the compositions, as the observations suggest.  \\citet{laming06} consider the possibility that the NE/SW ``jets\" represent a spherical explosion into an inhomogeneous CSM with ``cavities\" so that expansion is faster into those cavities, as suggested by \\citet{blondin96}. They find that such a model cannot give a sufficiently high density of plasma at high enough temperature. The ejecta expand so rapidly that the density is too low either for appreciable electron-ion equilibration to raise the electron temperature or to sufficiently ionize the plasma as observed for the NE jet.  Laming et al. consider a jet-induced model based on the calculations of \\citet{kho99} and find that such a model can provide the required conditions of ionization age and electron temperature for the collisional ionization equilibrium knots at the jet tip and the non-ionization equilibrium further back in the jet stream. Laming et al. argue that there should be a substantial amount of cold plasma at the ``jet\" head that has been cooled by radiative and adiabatic losses. This material could be composed of Fe, but is not necessarily so.  Laming et al. note that the blast wave is not seen in the direction of the NE ``jet.\" The NE jet is thus not overdense compared to the stellar envelope, as it was for the jet-induced supernova models of \\citet{kho99} \\citep[see also][] {hof01,kho01}. Laming et al. find that the kinetic energy in the NE ``jet\" is a relatively modest $10^{50}$ ergs, perhaps insufficient to represent the major power source that caused the supernova explosion.  This leaves open the basic physical mechanism of the explosion.  \\citet{hwang03} argue for an ejecta mass of about 2 \\m\\ and a total mass of the star at the time of the explosion of about 3 \\m. \\citet{chev03} have made the case for a clumpy wind that may account for the quasi-stationary flocculi and that the explosion might have corresponded to a Type IIn or Type IIb supernova, in current typology.  \\citet{young06} argue that the progenitor was a star of 15 - 20 \\m\\ on the main sequence that lost its hydrogen envelope to a binary companion and exploded with a mass of 4 - 6 \\m.  Among the elements that must be explained to fully account for the shape of Cas A are: \\begin{itemize} \\item A basic filametary morphology (Fig. 1a) \\item Quasi-stationary flocculi \\citep{fesen01} \\item Fast-moving knots of oxygen and their concentration at certain angles \\citep{fesen06c} \\item The appearance of ``rings\" at various parts of the perimeter (Fig. 1b) \\item The asymmetric distribution of mass \\citep{willingale03} \\item The distribution of silicon and other intermediate mass elements \\citep{hwang04} \\item The concentration of iron to the SE and NW, especially its appearance at larger projected radii toward the SE than silicon and its apparent absence in the NE ``jet\" \\citep{hughes00,hwang04} \\item A prominent blue-shifted region to the SE \\citep{markert83, hwang01,willingale02} \\item The prominent ``jet\" and ``counter-jet\" structure in the NE/SW \\citep{fesen01,hwang04} \\item The location and direction of motion of the compact object \\citep{fesen06b} \\item The lack of a prominent pulsar wind nebula around the compact object \\citep{fesen06a} \\item The expansion velocities and energetics as a function of distance, angle, and composition \\citep{willingale02,hwang04,fesen06b,laming06} \\item The presence or absence of a companion star \\citep{fesen06a} \\end{itemize}  As a step toward accounting for these properties we consider the physics and morphology implied by a ``jet-induced\" supernova model. The motivation for such a model has arisen in the context of recent developments in the study of the spectropolarimetry of core-collapse supernovae \\citep{wan03,maund07a,maund07b,maund07c,WW08} that show that core core collapse supernovae are routinely aspherical and frequently display a prominent axisymmetry. One obvious way to induce such an asymmetry is to explode the supernova with bi-polar jets \\citep{kho99}. \\citet{WMA07} point out that while bi-polar ``jet-like\" flow is common in core collapse supernovae, evidence for non-axisymmetric structure is also prominent. The presence of non-axisymmetric morphology is manifested as ``loops\" in the plane of the Stokes parameters Q and U.  This non-axisymmetric structure was prominent in SN~1987A \\citep{crop88} and is now recognized to be composition-dependent, with different species ejected in different directions.  Deeper study of Cas A may help to inform our interpretation of the spectropolarimetric data and vice versa. A basic picture of a bi-polar explosion with significant non-axisymmetric flow may apply in general to core-collapse explosions. Here we explore whether such a picture can help to illuminate the morphology of Cas A. ",
        "conclusions": "Despite attempts to deconvolve the three-dimensional structure, the interpretation of the kinematics, dynamics, and morphology of Cas A remains difficult. We adopt a jet-induced structure as an organizing principle to frame the analysis of observations and to provide a model to test. We suggest that Cas A exploded from a jet-like mechanism, but that, unlike many interpretations, the main jet is not the well-known NE ``jet.\" We argue that a structure to the SE with a position angle of $\\sim 125$\\degree\\ represents the main jet. This feature is less immediately discernable because of projection effects, but is suggested in Doppler maps. It is also iron rich. In the jet-induced model, the bulk of the ejecta, the intermediate-mass elements, should be ejected in an equatorial torus (for jets of equal and opposite momentum) that results from the convergence of blast waves on the equator. We interpret the NE ``jet\" and other structure as flows that are roughly normal to the main SE jet that result from non-axisymmetric structure in the equatorial plane of the jet. This structure may resolve one of the current mysteries of Cas A, that the proper motion of the compact object seems to be nearly normal to the principle axis of the dynamics. In our proposed morphology, the compact object has been ``kicked\" at an angle $\\sim 30$\\degree\\ to that of the main jet. The proposed picture also opens the way for a new interpretation of chemical structure, especially the fact that iron is seen beyond silicon: this observation may be significantly affected by projection effects, rather than resulting from true ``mixing\" of the ejecta.  There have been hints of the geometry we propose here in previous discussions of Cas A. \\citet{markert83} proposed an inclined ring model (their Figure 5), but their solution was a ring with its axis virtually along the line of sight. To account for the SE feature, they invoked non-uniform emissivity around the perimeter of the ring. The ring may have non-uniform emissivity, but their particular solution does not account for why this SE feature is iron-rich, and does not seem to be quantitatively consistent with subsequent work. \\citet{willingale02} used {\\it XMM Newton} composition and Doppler data to do a 3D deconvolution, presenting a proposed side view from the East (their Figure 11). They remark that the data is well characterized by the doughnut shape suggested by \\citet{markert83}, but their side view seems to show the axis of the torus tilted at about 45\\degree\\ to the line of sight, rather than the 87\\degree\\ advocated by \\citet{markert83}. With the spatial resolution available, the projection of \\citet{willingale02} shows no particular evidence of the NE jet. \\citet{willingale03} propose a clumpy toroidal geometry in which the axis of the toroid is oriented 40\\degree\\ East of the line of sight and 125\\degree\\ (counterclockwise) to the South of North, essentially the same as we have chosen here. Along this torus they identify concentrations of mass and energy on opposing sides to the ``North\" and ``South\" that are near to, but not quite the same as North and South on the plane of the sky (see their Figure 4 for the coordinate system they define). They refer to these concentrations of mass and energy as ``jets\" but do not relate them to the Doppler feature in the SE discussed by \\citet{willingale02}, and they are distinctly not the directions of the classic NE and SW jets. It is also not quite clear how this torus relates to the deconvolution given by \\citet{willingale02}. \\citet{willingale03} suggest that the torus they identify may be the shell of matter ejected by a progenitor Wolf-Rayet star. Here we raise the possibility that the torus is that produced by a jet-induced explosion.  The mass map presented by \\citet{willingale03} does not show any special evidence for the iron-rich mass in the SE that we identify as the main direction of the jet axis. This may be related to the fact that the hot iron moving along the jet axis has relatively little mass.  Like \\citet{willingale03}, we have no immediate explanation for the asymmetric distribution of mass around the perimeter of this torus, but propose that this is related to the instabilities and non-axisymmetric flows normal to the jet axis that we have discussed here, rather than ``jets\" {\\it per se}. \\citet{willingale03} point out that the apparent asymmetry of the explosion might suggest shear and hence turbulence and clumpiness intrinsic to the ejecta without needing a collision with an external medium. We have illustrated just this sort of effect in Figure \\ref{jets}.  \\citet{krause05} have provided another interesting and mysterious aspect concerning the nature of Cas A. They have identified an apparent bi-polar flow and associated infrared echoes. Their preferred orientation is a position angle of about 26\\degree\\ at an angle of about 82\\degree\\ to the line of sight, nearly in the plane of the sky. This angle differs from every other geometric feature discussed in this paper. \\citet{krause05} suggest that Cas A contains a magnetar that had a soft gamma-ray repeater outburst circa 1952. There is no sign of a Crab-like synchrotron nebula, which is also an issue for our suggestion that a pulsar wind might play a role in shaping the ejecta. The energy injected by a soft gamma-ray repeater burst is rather small compared to a supernova kinetic energy, but such a burst of energy might affect portions of the ejecta, enhancing previous irregularities. It would be interesting to look for other possible indications of such a magnetar burst. We note that the ionization ages identified by \\citet{willingale03} are on the scale of a century or less.  \\citet{burrows05} proposed that the main axis of the explosion of Cas A is that associated with the ``iron-rich, mass-rich, energy-rich\" SE direction and that this direction be a rotation axis. They speculate that the rotation axis precessed from the generation of the main jet to the generation of the later, secondary flow associated with the NE jet, or that the latter could be a magnetic dipole axis (although a magnetic axis would be expected to rotate in direction at the pulsar period if tilted with respect to the rotation axis). \\citet{burrows05} argue that the inferred direction of motion of the central object is along the proposed axis of explosion with an inferred blue-shift. We have made the specific proposal of a displacment of the line of motion of the compact object by $\\sim$30\\degree\\ from the axis of the main jet. \\citet{burrows05} suggest that the motion of the compact object could be due to a slightly top-bottom asymmetric bi-polar flow. \\citet{janka05} also suggest that the NE jet is induced after the main explosion, perhaps due to accretion onto the neutron star. Here we propose that there is an explicit jet/torus structure of to Cas A, and agree with the general suggestions of \\citet{burrows05} and \\citet{janka05} that the NE jet is some form of secondary flow, caused by the complex interaction of the jets or by the late-time effect of a pulsar wind, or perhaps repeated soft gamma-ray repeater outbursts.  To produce a robust explosion, the main explosion engine should involve $\\sim 10^{51}$ ergs. Even if this energy begins in jets, it is spread through a large solid angle by subsequent dynamical interaction with the surrounding progenitor star. Thus it is not clear how much energy should be directed along the main axis once free expansion is attained. This can be quantified by suitable models. It remains to be seen whether an analyis of the SE structure similar to that done by \\citet{laming06} will reveal a sufficiently large kinetic energy in the SE to be consistent with a jet in that direction being the primary origin of the explosion. The evidence that only $\\sim 10^{50}$ ergs is involved in the NE jet is consistent with that being only a secondary effect.  It is also important to determine the extent to which projection of an iron-rich SE jet on a silicon-rich equatorial flow can quantitatively account for the observed apparent composition inversion.  An important aspect of this model is that the non-axisymmetric structure arises from instabilities in or near the equatorial plane. We have sketched possible instabilities, including the effects of a pulsar wind. Three-dimensional hydrodynamics calculations or even pehaps MHD calculations \\citep{stone07} are required to determine whether or not the structure of Cas A can arise in a natural and physically self-consistent way from a jet-induced model. Such a model would have to yield a plausible explanation of transverse velocities up to 14,000 \\kms\\ in the direction of the NE jet \\citep{fesen06c} in contrast to transverse velocities of $\\sim$ 8,000 \\kms\\ along the SE structure (the radial velocities in both of these directions are more modest; \\citet{willingale02}). We note that this high velocity in the NE direction may be an indication that this is a secondary flow, perhaps driven by a pulsar wind, invoking once again the notion that the original jet energy must be dissipated throughout the ejecta to cause the explosion. A separate injection of energy may be necessary to blow holes in the original ejecta, accelerate some regions to high velocity and produce fast moving knots. \\citet{fesen06c} note that the NE jet is like a ``spray\" of ejecta knots rather than a narrow jet. That might be an important clue to the underlying process.  There is no model today that accounts for the nature of Cas A. The positive aspects of the model we present here are that it roughly accounts for the iron-rich SE structure, the extension of iron beyond silicon in that direction, and the motion of the compact object in nearly that direction. On the negative side, there is much structure in Cas A that remains unexplained. Our model gives no natural explanation for why the NE jet and SE counter-jet seem so nicely and oppositely aligned, a problem for any model in which these are secondary flows. It may also be the case that the structure seen in Cas A might be generated in the absence of a jet-induced origin. This has certainly not been demonstrated. It may also be true that the NE jet defines a symmetry axis and the compact object has been kicked at some oblique angle due to equatorial instabilities, but such a theory would also require elaboration.  An important aspect going forward is to compare models with observations. For our model we predict: a) there is enough mass and kinetic energy in the SE and NW directions to account for the explosion (some of this matter may be fast, cool and difficult to directly observe; much of the originally directed energy has spread throughout the ejecta); b) there is high-velocity iron extending to large distances in the SE (in front) and NW (in back); c) the 3D structure and kinematics should show an expanding toroidal structure, roughly normal to the SE direction we have defined here, consistent with the energy deposited in the progenitor star by the jet-induced blast waves; d) the intermediate mass elements are concentrated in this torus; e) the apparent ``overturn\" that places iron beyond silicon in the plane of the observer is largely due to differentiated, composition-dependent, bulk flow, not mixing due to turbulence; f) the compact object is moving toward us; g) the ``holes\" in the remnant caused by secondary flow, including the NE and SW jets, are the locus of faster moving material; h) the walls defining the holes are in transverse motion; the holes should be getting larger.  We urge those analyzing the data on Cas A to put these predictions to the test in order to more deeply understand this remarkable explosion. The structure to the SE (and NW) and the torus defined by \\citet{willingale03}, in particular, warrant much more detailed study. We need to better understand the kinematic and dynamical relation between the SE feature, the NE/SW ``jets,\" and this torus and to know whether the torus is more consistent with a pre-existing shell or a structure created by the supernova."
    },
    "0711/0711.4358_arXiv.txt": {
        "abstract": "The transfer of turbulent energy through an inertial range from the driving scale to dissipative scales in a kinetic plasma followed by the conversion of this energy into heat is a fundamental plasma physics process. A theoretical foundation for the study of this process is constructed, but the details of the kinetic cascade are not well understood.  Several important properties are identified: (a) the conservation of a generalized energy by the cascade; (b) the need for collisions to increase entropy and realize irreversible plasma heating; and (c) the key role played by the entropy cascade---a dual cascade of energy to small scales in both physical and velocity space---to convert ultimately the turbulent energy into heat. A strategy for nonlinear numerical simulations of kinetic turbulence is outlined. Initial numerical results are consistent with the operation of the entropy cascade. Inertial range turbulence arises in a broad range of space and astrophysical plasmas and may play an important role in the thermalization of fusion energy in burning plasmas. ",
        "introduction": "\\label{}  In a warm, magnetized kinetic plasma, the dissipation of low-frequency turbulent fluctuations by kinetic mechanisms occurs most strongly at length scales of order the plasma particle Larmor radii. If turbulence is driven in the plasma at a much larger scale, an inertial range develops in which the injected energy is cascaded by nonlinear interactions to the small scales at which the turbulence can be effectively dissipated. This transfer of turbulent energy through an inertial range from the driving scale to dissipative scales followed by the conversion of this energy into heat is a fundamental kinetic plasma physics process. In this paper, I lay out a theoretical foundation on which to build our knowledge of \\emph{inertial range turbulence in kinetic plasmas} and outline a practical strategy for the nonlinear numerical modeling of this process.  The study of turbulence in kinetic plasmas is not a new endeavor; in magnetic fusion community, understanding the effect of turbulence on the transport properties of the plasma is one of the primary goals of the numerical modeling effort. But the turbulence most often investigated in laboratory plasmas is driven by microinstabilities arising from gradients in the background plasma---for example, ion temperature gradient (ITG) and electron temperature gradient (ETG) instabilities \\cite{Dimits:1996,Dorland:2000,Jenko:2000,Jenko:2001a,Rogers:2000, Jenko:2001b,Jenko:2002,Candy:2004,Parker:2004}. These instabilities inject energy into the plasma at the scale of the particle Larmor radius; the resulting turbulence does not develop an inertial range because the driving and kinetic dissipation scales are of roughly the same order. Inertial range turbulence in kinetic plasmas, therefore, represents a fundamentally different process than the kinetic turbulence traditionally studied by the fusion community.  The turbulence in most astrophysical contexts, on the other hand, is typically driven at scales much larger than the Larmor radius and gives rise to a turbulent cascade through an inertial range to the dissipative scales at which the energy of turbulent fluctuations is ultimately converted into particle thermal energy.  Before constructing the detailed theoretical foundation for the study of this process, presented in \\secref{sec:theory}, I will provide here a simplified blueprint of the physical mechanisms guiding the flow of energy. It is important to emphasize that many of the mechanisms described below are not well understood.  Here I aim only to paint the overall picture in broad strokes, deferring discussion of the many uncertainties to \\secref{sec:theory}.  At a scale larger than the particle mean free path---a collisional scale for which magnetohydrodynamics (MHD) provides an adequate description---the plasma is stirred by some external mechanism, driving an assortment of MHD Alfv\\'en, fast, slow, and entropy mode fluctuations in the plasma.  At this driving scale, these modes are undamped; thus, a turbulent cascade develops nonlinearly to transfer the fluctuation energy to smaller scales. The compressive modes become damped as the cascade reaches scales of order or smaller than the collisional mean free path, but the Alfv\\'enic cascade continues undamped down to the scale of the ion Larmor radius.  At this kinetic scale, the electromagnetic fluctuations may be damped collisionlessly by the Landau resonance with the ions. In the absence of collisions, this process conserves a generalized energy---the free energy removed from the electromagnetic fluctuations generates nonthermal structure in velocity space of the ion distribution function. The remaining electromagnetic fluctuation energy continues to cascade below the scale of the ion Larmor radius as a kinetic \\Alfven wave cascade.  Upon reaching the scale of the electron Larmor radius, the electromagnetic fluctuations of the kinetic \\Alfven wave cascade are completely damped via the Landau resonance with the electrons; again, a generalized energy is conserved in this process, leading to the creation of nonthermal structure in velocity space of the electrons. But the damping of the electromagnetic fluctuations and consequent generation of structure in velocity space does not correspond to heating; irreversible heating requires an increase in entropy that can only be achieved by collisions. The thermalization of the turbulent energy by collisions is ultimately achieved thanks to a cascade to small scales in velocity space of the particle distribution functions---an \\emph{entropy cascade}.  The entropy cascade drives the distribution function structure in velocity space to scales small enough that even weak collisions are sufficient to smooth out that structure towards the Maxwellian, causing entropy to increase---this is the final step in the conversion of the energy of the turbulent fluctuations to thermal energy of the plasma particles.  The entire process described above is the kinetic generalization of the familiar cascade of energy in a fluid turbulent system; this fundamental kinetic plasma physics process, encompassing the dual cascade in both physical and velocity space, is referred to as the \\emph{kinetic cascade} of the generalized energy \\cite{Schekochihin:2007}. Neither the detailed interactions of the kinetic cascade nor its implications for the flow of energy in turbulent systems are well understood; a theoretical foundation for its study is layed out in \\secref{sec:theory}. Nonlinear numerical simulations of the kinetic cascade occurring in inertial range turbulence are expected to play a leading role in shedding light on the complex physical mechanisms involved---a strategy for the numerical approach to this problem is presented in \\secref{sec:numerics} and initial numerical results are discussed in \\secref{sec:results}. Astrophysical and laboratory environments in which the kinetic cascade of inertial range turbulence may play an important role are identified in \\secref{sec:app}. ",
        "conclusions": "\\label{sec:conc}  In this paper, I have presented a general picture of a fundamental plasma physics process: the transfer of turbulent energy through an inertial range from the driving scale to dissipative scales followed by the conversion of this energy into heat. The details of the kinetic cascade of the generalized energy are not yet well understood. Numerical simulations of kinetic turbulence will play a crucial role in illuminating the key elements of the cascade. Here I emphasize several of these important elements: \\begin{enumerate} \\item The nonlinear cascade of energy in kinetic plasmas preserves a generalized energy given by \\eqref{eq:consw}. \\item Collisions are required to increase the entropy and thus achieve irreversible heating of the plasma. \\item The ultimate thermalization of the turbulent energy is achieved in weakly collisional plasmas through an entropy cascade, a dual cascade of energy to small scales in both physical and velocity space. \\end{enumerate}  I have outlined the numerical approach chosen to investigate turbulence in the transition regime from the \\Alfven wave to the kinetic \\Alfven wave cascade. The computationally demanding simulations use a Langevin antenna to model driving from larger scales in the inertial range and a hypercollisional operator to remove energy at the smallest resolved scales. Results from a simulation with high velocity-space resolution are consistent with the development of $v_\\perp$ structure due to the action of the entropy cascade.  Inertial range turbulence arises in kinetic plasmas across a broad range of space and astrophysical plasmas---including the solar wind, the interstellar medium, and accretion disks around compact objects---and may play an important role in the thermalization of fusion $\\alpha$-particle energy in next-generation burning plasmas.  A program of nonlinear gyrokinetic simulations of kinetic, inertial range turbulence is necessary to explore the key elements of the kinetic cascade described here. The rich plasma physics involved in the inherently kinetic and nonlinear nature of the predicted entropy cascade, along with the computational intensity of the numerical work, make this important physics problem both fun and challenging. Additional aspects of kinetic turbulence must also be examined, such as the effect of strong fast wave turbulence on the Alfv\\'en, slow, and entropy modes; this requires a treatment beyond the scope of gyrokinetics."
    },
    "0711/0711.4391_arXiv.txt": {
        "abstract": "The dynamics and structure of accretion disks, which accumulate the vertical magnetic field in the centers, are investigated using two- and three-dimensional MHD simulations. The central field can be built up to the equipartition level and disrupts a nearly axisymmetric outer accretion disk inside a magnetospheric radius, forming a magnetically arrested disk (MAD). In the MAD, the mass accretes in a form of irregular dense spiral streams and the vertical field, split into separate bundles, penetrates through the disk plane in low-density magnetic islands. The accreting mass, when spiraling inward, drags the field and twists it around the axis of rotation, resulting in collimated Poynting jets in the polar directions. These jets are powered by the accretion flow with the efficiency up to $\\sim 1.5\\%$ (in units $\\dot{M}_{in}c^2$). The spiral flow pattern in the MAD is dominated by modes with low azimuthal wavenumbers $m\\sim 1-5$ and can be a source of quasi-periodic oscillations in the outgoing radiation. The formation of MAD and Poynting jets can naturally explain the observed changes of spectral states in Galactic black hole binaries. Our study is focused on black hole accretion flows; however, the results can also be applicable to accretion disks around nonrelativistic objects, such as young stellar objects and stars in binary systems. ",
        "introduction": "Accretion disks can carry small- and large-scale magnetic fields. The small-scale field ($\\ell\\la R$, where $\\ell$ is the field scale length and $R$ measures the radial distance from the disk center) can be locally generated by the MHD dynamo (Brandenburg et al. 1995; Stone et al. 1996) supported by the turbulence, which results from the magneto-rotational instability (MRI, Balbus \\& Hawley 1991). This field can provide the outward transport of angular momentum in the bulk of the disk with the help of local Maxwell stresses (Shakura \\& Sunyaev 1973; Hawley, Gammie, \\& Balbus 1996). The large-scale field ($\\ell > R$) is unlikely produced in accretion disks (however, see Tout \\& Pringle 1996), and can either be captured from the environment and dragged inward by an accretion flow (Bisnovatyi-Kogan \\& Ruzmaikin 1974, 1976), or inherited from the past evolution (see \\S 4). The large-scale field can remove the angular momentum from accretion disks by global Maxwell stresses through the magnetized disk corona (K\\\"onigl 1989).  A large-scale bipolar field, unlike a small-scale field, can not dissipate locally due to the magnetic diffusivity and can not be absorbed by the central black hole. In the case of inefficient outward diffusion of the bipolar field through the disk (see Narayan, Igumenshchev, \\& Abramowicz 2003; Spruit \\& Uzdensky 2005; Bisnovatyi-Kogan \\& Lovelace 2007; and for other possibility, see van Ballegooijen 1989; Lubow, Papaloizou, \\& Pringle 1994; Lovelace, Romanova, \\& Newman 1994; Heyvaerts, Priest, \\& Bardou 1996; Agapitou \\& Papaloizou 1996; Livio, Ogilvie, \\& Pringle 1999), this field is accumulated in the innermost region of an accretion disk and forms a ``magnetically arrested disk,'' or MAD (Narayan et al. 2003). The MAD consists of two parts: the outer, almost axisymmetric, Keplerian accretion disk and the inner magnetically dominated region, in which the accumulated vertical field disrupts the outer disk at the magnetospheric radius $R_{\\rm m}\\sim 8\\pi GM\\rho/B^2$, where $M$ is the central mass, $\\rho$ is the accretion mass density, and $B$ is the magnetic induction.  It is believed that the large-scale bipolar field in accretion disks is responsible for the formation of jets observed in a large variety of astrophysical objects (e.g., Livio, Pringle, \\& King 2003). The magnetically driven jets can be of two types, basically depending on a mass load by the disk matter (e.g., Lovelace, Gandhi, \\& Romanova 2005): Poynting jets and hydromagnetic jets, which have, respectively, small and large mass loads. The hydromagnetic jets can be formed by two mechanisms: the magneto-centrifugal mechanism (Blandford \\& Payne 1982; K\\\"onigl \\& Pudritz 2000) and the toroidal-field pressure generated by the disk rotation (Lynden-Bell 2003; Kato, Mineshige, \\& Shibata 2004). The magneto-centrifugal mechanism produces relatively wide outflows and, to be consistent with observations of the collimated jets, requires an additional focusing mechanism. Jets driven by the toroidal-field pressure can have a high degree of collimation, but these jets are known to be kink unstable (Eichler 1993; Appl 1996; Spruit, Foglizzo, \\& Stehle 1997). Poynting jets  are naturally self-collimated and expected to be marginally kink stable (Li 2000; Tomimatsu, Matsuoka, \\& Takahashi 2001). These jets can originate in the innermost region of accretion disks and powered either by the disks themselves (Lovelace, Wang, \\& Sulkanen 1987; Lovelace et al. 2002) or by rotating black holes (Blandford \\& Znajek 1977; Punsly 2001; also see, Takahashi et al. 1990; Komissarov 2005; Hawley \\& Krolik 2006; McKinney 2006).  Our study is based on two- and three-dimensional (2D and 3D, respectively) MHD simulations and has two main goals. First, we investigate the dynamics and structure of MAD's. We show that MAD's are formed in the accretion flows, which carry inward large-scale poloidal magnetic fields. Inside the magnetospheric radius $R_{\\rm m}$, the matter accretes as discrete streams and blobs, fighting its way through the strong vertical magnetic field fragmented in separate bundles. Because of rotation, the streams take spiral shapes. Second, we demonstrate a link between the existence of MAD's and production of powerful Poynting jets. These jets should always be generated in MAD's because of the interaction of the spiraling accretion flow with the vertical magnetic bundles, which, as the result, are twisted around the axis of rotation.  This paper extends the work of Igumenshchev, Narayan, \\& Abramowicz (2003) by studying in more detail the radiatively inefficient accretion disks with poloidal magnetic fields. We employ a new version of our 3D MHD code, which can be utilized in multi-processor simulations. The paper is organized as follows: We describe the solved equations, the numerical method used, and initial and boundary conditions in \\S 2. We present our numerical results in \\S 3, and discuss and summarize them in \\S 4. ",
        "conclusions": "We have performed a numerical study of the formation and evolution of quasi-stationary MAD, which is characterized by a strong vertical magnetic field accumulated at the disk center. We employ the simulation design, in which the poloidal magnetic field of one sign is permanently injected into the computational domain at the outer numerical boundary and the unipolar vertical field is transported inward by the accretion flow. The accumulated field has a significant impact on the inner flow structure and dynamics in both 2D and 3D simulations. In the axisymmetric 2D simulations, the field pressure can temporarily halt the mass falling into the black hole, resulting in the cycle accretion, in which the longer periods of accumulation of the mass at the magnetospheric radius $R_m$ are followed by the short periods of accretion. The 3D simulations have shown, however, that the axisymmetric cycle accretion is not realized. Instead, the accumulated field causes the mass to accrete quasi-regularly in the form of non-axisymmetric spiral streams and blobs. We have demonstrated that 3D MAD's can be efficient sources of collimated, bipolar Poynting jets, which originate in the vicinity of the central black hole. These jets develop due to and are powered by the interaction of the spiral mass inflows with the central field split into separate magnetic bundles. The efficiency of conversion of the accretion energy $\\dot{M}_{in} c^2$ into the Poynting jet energy is up to 1.5\\% in our simulations. This estimate may not be accurate (we believe, underestimated) because of the use of the pseudo-Newtonian approximation (see \\S 2). The better estimate of the efficiency can be obtained using general relativistic MHD simulations.  We have presented the simulation results from three models of radiatively inefficient accretion disks, which differ by the strength of the injected field. In accordance with the previous studies (e.g., Stone \\& Pringle 2001), the structure of the outer disks in these models is determined by the field strength. In Model~A, which has the largest injected field, the MRI is suppressed and the accretion flow is driven by global Maxwell stresses, which transport the excessive angular momentum outward from the disk through the disk corona. In Models~B and C, which have the smaller injected fields, the MRI and turbulent motions are developed. The turbulence in these models is axisymmetric and partially supported by the efficient convection motions resulting from dissipation of toroidal and small-scale, poloidal magnetic fields. These motions cause the increase of the disk thicknesses in comparison with non-turbulent Model~A. The 3D simulations of Models~A and B have demonstrated that non-axisymmetric motions are not important in the outer parts of the disks, outside the inner magnetically dominated region (the disks remain almost axisymmetric), but very important inside this region, resulting in the development of the spiral accretion flows and bi-polar Poynting jets.  Numerical magnetic dissipations and reconnections result in a magnetic diffusivity, which influences the structure and dynamics of our models. The spatial scale, on which the diffusivity occurs in our simulations, is determined by the gridsize, which greatly exceeds the scales of various resistive mechanisms (e.g., Coulomb collisions, dissipative plasma instabilities) in the relevant astrophysical conditions. Therefore, the magnetic diffusivity is significantly overestimated in our models and results in an excessive slippage of an accretion flow through magnetic field. This slippage reduces the ability of the flow to drug inward the vertical field, but, however, the numerical diffusivity is not efficient enough to totally prevent the field accumulation at the disk center. The numerical diffusivity suppresses the MRI on the scales of the gridsize, and, therefore, prevents the development of turbulence in the outer regions of our models, where the gridsize is increased. Other effect of the numerical magnetic diffusivity is the enhancement of the ablation of magnetic islands, which are found in the 3D simulations (see \\S 3.1). To test the sensitivity of our models to magnetic dissipations, we have performed a 2D simulation of the model, which is similar to Model~A, but has the double number of grid points in the $R$- and $\\theta$-directions. The simulation has demonstrated the qualitative similarity of axisymmetric evolution of the high resolution model and Model~A: both models show the formation of accretion disks, accumulation of the vertical field in the disk centers, and development of the cycle accretion. The high resolution model forms a thiner laminar accretion disk. Unfortunately, a more detailed quantitative comparison of these models meets some difficulties because of the different properties of the mass and field injection region (see \\S 2), which are changed with the change of the resolution.  The main results of our study, the formation of MAD's and Poynting jets, have been obtained under the assumption of radiatively inefficient flows, but, we believe that these results can also be applied to the radiatively efficient, dense accretion disks (e.g., Kato, Fukue, \\& Mineshige 1998). The formation of MAD's should not be affected by the radiative cooling as soon as the central field satisfies the equipartition condition, \\begin{equation} {B^2\\over 8\\pi} \\sim {GM\\rho\\over R_g}, \\end{equation} where $\\rho$ is the mass density in the innermost region. The radiative losses results in the higher $\\rho$ and, therefore, the larger B is necessary to obtain radiative MAD's. We expect that the qualitatively similar spiral structure of the inner magnetically dominated region, to that found in our simulations, can be developed in the case of the radiative disks. Poynting jets should be a necessary attribute of the radiative MAD's as well.  The formation of MAD's in the radiative disks can be used to explain the observations of the low/hard state in black hole binaries (for a review, see Remillard \\& McClintock 2006). Here, we briefly discuss basic moments of this application of MAD's and leave more quantitative considerations for future works. We assume, for example, the development of the MAD in the radiation pressure dominated accretion disk at the subcritical regime (see Shakura \\& Sunyaev 1973). In such a disk, the radiation diffusion time scale $t_{rad}\\simeq H^2\\sigma_T\\rho/c$ can significantly exceed the Keplerian time $t_{K}= 2\\pi R^{3/2}/\\sqrt{GM}$ at small values of the $\\alpha$-parameter, $\\alpha \\la 0.1$, in virtu of the relation \\begin{equation} {t_{K}\\over t_{rad}} \\simeq 3.4\\alpha, \\end{equation} which follows from the Shakura-Sunyaev solution. Here, we denote $H$ to be the disk half-thickness and $\\sigma_T=0.4$ cm$^2$/g to be the Thomson scattering cross-section. As soon as $t_{rad}\\gg t_{K}$ in the outer Shakura-Sunyaev disk, $t_{rad}$ can also significantly exceed the accretion velocity $t_{accr}\\sim t_{K}$ in the inner spiral flow in the MAD, in virtu of the relation \\begin{equation} {t_{accr}\\over t_{rad}}\\propto R, \\end{equation} which is satisfied in accretion flows with the scaling law of the accretion velocity $v_{accr}\\propto R^{-1/2}$ and $H\\propto R$. Having $t_{rad}\\gg t_{accr}$, one concludes that the radiation is traped inside the spiral flow on the accretion time scale and, therefore, this flow is radiatively inefficient. From the point of view of an observer, which detects the softer part of the spectrum of outgoing radiation (below $\\sim$ eVs), the MAD will look like a Shakura-Sunyaev disk truncated at the inner radius $R_{tr}$, which coinsides with the transition radius between the inner magnetically dominated region and outer axisymmetric accretion disk. Typically, in observations, $R_{tr}$ is in the interval from a few tens to hundreds of $R_g$ (e.g., in Cyg X-1, see Done \\& Zycki 1999), which is consistent with that obtained in our Models~A and B. The observed specta of black hole binaries in the low/hard state are dominated by the hard x-ray component (e.g., Done, Gierli\\'nski, \\& Kubota 2007) and this can be explained by the radiation from the hot, optically thin magnetized medium, which surrounds the accreting spiral flows and in which the binding energy of these flows is released. The synchrotron radiation from the magnetized medium and Poynting jets in the MAD (see Goldston, Quataert, \\& Igumenshchev 2005) can be used to explain the observed radio luminosity in the low/hard state; this luminosity is believed to be due to steady jets (e.g., Corbel et al. 2003; Gallo, Fender, \\& Pooley 2003).  Our simulations assume accretion disks around black holes and can be relevant to objects with relativistic jets containing accreting stellar-mass black holes (e.g., in micro-quasars; black holes resulted from type Ib/c supernova explosions and mergers of two compact objects) and supermassive black holes (in galactic centers). However, we believe that our main results can also be relevant to nonrelativistic astrophysical objects, in which accretion disks and jets are observed. These objects include, for example, young stellar objects and accreting stars (e.g., white dwarfs) in binary systems. Qualitatively, we expect that the structure and dynamics of MAD's and Poynting jets are similar in the both relativistic and nonrelativistic cases. We expect, however, large quantitative differences in these two cases because of the different energy-density scales involved in the regions of jet formation. Black holes, which are capable of launching jets almost from the event horizon, can produce ultra relativistic jets (e.g., McKinney 2006). Jets from nonrelativistic objects, which have the surface radius $R_*\\gg R_g$, are limited by the velocities $v\\sim\\sqrt{GM/R_*}\\ll c$. For example, the latter formula gives the upper estimate of the jet velocity $\\sim 400\\,km/s$ from the solar-type stars.  The problem of inward transport and amplification of the vertical field in turbulent accretion disks was intensively discussed in past and recent years (see, e.g., Spruit \\& Uzdensky 2005). The solution of this problem can help to discriminate models of accretion disks, which are consistent with observations (e.g., Meier \\& Nakamura 2006; Schild, Leiter, \\& Robertson 2006). Our simulation results show that the vertical field is transported inward and amplified independent of the disk structure, either laminar or turbulent. It is worth noting, however, that magnetic fields in our models are imposed and relatively large. The assumed strength of these fields exceeds the possible strength of the self-sustained magnetic fields that could be developed due to the MRI (Sano et al. 2004). Therefore, these results should be considered with some caution, because they do not represent the case of weak vertical magnetic fields.  The strong vertical field in the center of accretion disks can be, in principal, a relic field that is inherited from the previous evolution. This field can appear, for example, in the merger scenario (merger of two magnetized neutron stars or a black hole with a magnetized neutron star, e.g., Berger et al. 2005) or in the course of the gravitational collapse of an extended ``proto\" object (e.g., proto-stellar cloud, supernova progenitor), which produces a significantly more compact object (protostar, black hole). In the latter case, the proto object may contain some amount of the poloidal field, which will be amplified and accumulated at the center during the collapse. After the formation of the compact object, the remained noncollapsed mass can still move inward, forming an accretion disk and confining the field in the vicinity of the object. Depending on the relative strength of this relic field and the mass accretion rate, MAD's and Poynting jets can be developed. The considered scenario can be applied to young stellar objects (T-Tauri stars, e.g., Donati et al. 2007) and the hyper-accretion model for gamma-ray bursts (Woosley 1993; Paczy\\'nski 1998).  The spiral-flow pattern in MAD's rotates with about the same angular velocity at all radial distances; i.e., it rotates almost as a rigid body. Such a rotation can result in quasi-periodic oscillations (QPO's) in the emitted radiation, if the disk's axis is inclined to the line of view of an observer (e.g., Alpar \\& Shaham 1985; Lamb et al. 1985; Strohmayer et al. 1996; Lamb \\& Miller 2001; Titarchuk 2003). The frequency of these QPO's should be related to the rotation of the spiral pattern, which angular velocity is defined by the radius of the magnetically dominated region and can be a fraction ($\\sim 0.5-1$) of the orbital velocity at this radius. More investigations are required to make quantitative predictions about QPO's from MAD's."
    },
    "0711/0711.2504_arXiv.txt": {
        "abstract": "We review the classical observational results for isotopic abundance variations for several elements in CP stars.  We concentrate on the ``newest'' anomaly, in calcium.  The cosmically very rare isotope, $^{48}$Ca can rival and even dominate the more common, alpha nuclide, $^{40}$Ca.  Relevant examples are found in the hot, non-magnetic HgMn stars, and the field horizontal-branch star, Feige 86.   The calcium anomaly is also present in cool, magnetic stars, including the notorious HD 101065, Przybylski's star. ",
        "introduction": "\\label{intr}  The earliest observations of isotopic effects in stellar spectra were for diatomic molecules. The Michigan astronomer W. C. Rufus (1916) noted systematic absorptions in the spectra of R stars that we now know were due to differing isotopic species in the Swan bands of C$_2$. At the time, the Swan bands were known to be due to carbon, but the identification with the C$_2$ molecule was made some ten years later (see Sanford 1932). By the late 1940's, Andrew McKellar (1948) was able to make a systematic survey of the $^{12}$C/$^{13}$C ratio in carbon stars by a study of the molecular bands.  The deuterium analog of the famous 21 cm radio line of neutral hydrogen was observed in the 1960's (Weinreb 1962). The most extensive observations of isotopic spectra are for interstellar molecules.  The field was already mature several decades ago (cf. Wanneir 1980; Wilson and Rood 1994). ",
        "conclusions": ""
    },
    "0711/0711.4996.txt": {
        "abstract": "An extraordinarily rich zoo of non-baryonic Dark Matter candidates has been proposed over the last three decades. Here we present a 10-point test that a new particle has to pass, in order to be considered a viable DM candidate: I.) Does it match the appropriate relic density? II.) Is it {\\it cold}? III.) Is it neutral? IV.) Is it consistent with BBN? V.) Does it leave stellar evolution unchanged? VI.) Is it compatible with constraints on self-interactions? VII.) Is it consistent with {\\it direct} DM searches? VIII.) Is it compatible with gamma-ray constraints? IX.) Is it compatible with other astrophysical bounds? X.) Can it be probed experimentally? ",
        "introduction": "\\label{sec:chapter one}   The analysis of the Cosmic Microwave Background (CMB) anisotropies is a powerful tool to test cosmological models, and to extract the corresponding cosmological parameters. For instance, the angular position of the peaks in the power spectrum of temperature anisotropies is a sensitive probe of the curvature of the Universe (see e.g. \\cite{Scott Smoot, Dodelson} for a review and a more extended discussion). The power spectrum of CMB anisotropies is fitted within the Standard Cosmological Model with a number or free parameters that depends onto the priors.\\\\ The best fit of the three years WMAP data, with a 6 parameters flat $\\Lambda$CDM model and a power-law spectrum of primordial fluctuations, gives \\cite{Spergel} $$ \\Omega_{b}h^2=0.0223^{+0.0007}_{-0.0009}  \\mbox{    } \\Omega_{M}h^2=0.127^{+0.007}_{-0.013}   $$ for the abundance of baryons and matter, respectively. The normalized abundance $\\Omega_i$ is defined as $\\Omega_{i}=\\rho_{i}/\\rho_{c},$ where $\\rho_c$ is the critical density, and the scaled Hubble parameter $h$ is defined as $H_0 \\equiv 100h \\mbox{ km} \\mbox{ s}^{-1} \\mbox{ Mpc}^{-1}.$ A joint-likelihood analysis on a larger data-sets including, besides WMAP3, also small scale CMB experiments (BOOMERang, ACBAR, CBI and VSA), Large-Scale Structures (SDSS, 2dFGRS) and SuperNova (HST/GOODS, SNLS), further strengthens the constraints to \\cite{http://lambda.gsfc.nasa.gov} $$ \\Omega_{b}h^2=0.0220^{+0.0006}_{-0.0008}  \\mbox{    } \\Omega_{M}h^2=0.131^{+0.004}_{-0.010}.   $$ Note that the baryonic density is consistent with the determination from big bang nucleosynthesis \\cite{Sarkar? vedi PDG} $$ 0.017<\\Omega_b h^2 <0.024 \\mbox{ (95 \\% CL)}.$$  For a new particle to be considered a good DM candidate, a production mechanism that reproduce the appropriate value of the relic density must exist. Moreover, to guarantee its stability, its lifetime must exceed the present age of the Universe. Taking in account the estimates of the Hubble Space Telescope Key Project \\cite{Hubble constant} and in agreement with the result derived by WMAP, $H_0 = 72 \\pm 3 \\mbox{ (statistical)}\\pm 7 \\mbox{ (systematic)} \\mbox{ km} \\mbox{ s}^{-1} \\mbox{Mpc}^{-1}, $ we require a lifetime $\\tau \\gtrsim 4.3 \\times 10^{17} \\mbox{ s}.$  In many proposed extensions of the Standard Model of particle physics, the stability of the DM particle is ensured by imposing a symmetry that forbids the decay of DM into Standard Model particles. For example, R-parity in Supersymmetry models (SUSY) \\cite{Rparity 1,Rparity 2}, K-parity in Universal Extra Dimensions Models (UED) \\cite{K parity}, and T-parity in Little Higgs Models \\cite{t parity}, prevent the lightest {\\it new} particle in the respective theories from decay (see for example Ref.\\cite{Jungman Kamionkowki Griest} for a detailed discussion on SUSY DM and Ref.\\cite{Profumo Review} for a review on UED DM).\\\\   \\subsection*{Thermal relics}  Among the best DM candidates, there is a class of particles called WIMPs (for weakly interacting massive particles), that are thermal relics and naturally achieve the appropriate relic density.  The scenario goes as follows: the WIMP is in thermodynamic equilibrium with the plasma in the early Universe, and it decouples when its interaction rate drops below the expansion rate of the Universe. For a non-relativistic particle at decoupling, the number density over the entropy density remains frozen, i.e. the thermal relic {\\it freezes-out}. The evolution of the number density of a generic species $\\chi$ in the Universe, is described by the following Boltzmann equation: $$   \\dot{n}_{eq} + 3Hn = -\\langle \\sigma_{ann}v \\rangle \\left[ n^2 -n_{eq}^2 \\right].$$ The second term in the l.h.s of the equation takes into account the dilution of the number density due to the expansion of the Universe. $\\langle \\sigma_{ann} v \\rangle $  is the thermal average of the annihilation cross section times velocity and it is parametrized with a non-relativistic expansion in powers of $v^2$, as: $ \\langle \\sigma_{ann} v \\rangle = a + b \\langle v^2\\rangle +\\mathcal{O}(\\langle v^4 \\rangle) \\simeq a + 6b/x,$ with $x\\equiv m/T.$  $n_{eq}$ is the equilibrium density of WIMPs in the plasma at temperature T and for a non-relativistic specie is given by $n_{eq } = g(\\frac{mT}{2\\pi})^{3/2} e^{-\\frac{m_{\\chi}}{T}} $, where $g$ denotes the number of degrees of freedom of $\\chi$ and $m_{\\chi}$ is the WIMP mass.  The Boltzmann equation can be solved integrating it in two extreme regions, long before and long after the WIMP freeze-out (e.g. WIMP decoupling), and matching then the solutions. Skipping the calculation details, that can be reviewed e.g. in \\cite{Kolb Turner}, the relic density today for a generic WIMP $\\chi$  is \\cite{bertone review}: \\begin{eqnarray} \\Omega_{\\chi} h^2 & \\thickapprox &\\frac{1.07 \\times 10^{9} \\mbox{ GeV}^{-1} }{M_{Pl} }\\frac{x_f}{\\sqrt{g_{*f}}} \\frac{1}{(a+3b/x_f)} \\nonumber \\\\ & \\thickapprox & \\frac{3 \\times 10^{-27} \\mbox{ cm}^3\\mbox{s}^{-1}}{\\langle \\sigma_{ann} v\\rangle}. \\label{eqn: Omega} \\end{eqnarray} where $g_{*f}$ counts the relativistic degrees of freedom at the decoupling, $M_{Pl}$ is the Planck mass and $x_f \\equiv m_{\\chi}/T_f$ with $T_f$ the freeze-out temperature. The last line is an order of magnitude estimate and it shows that the relic abundance of a non relativistic decoupled specie strictly depends on the annihilation cross section at freeze-out \\cite{Jungman Kamionkowki Griest}. Furthermore, it has to be noticed that the annihilation cross section, for a particle of given mass, has a maximum, imposed by the partial wave unitarity of the S matrix, $\\langle \\sigma_{ann} v\\rangle_{max} \\sim 1/m_{\\chi}^2 $ \\cite{Unitarity bound, Hui Ubound}. Thus, with the use of Eq. \\ref{eqn: Omega}, the requirement $\\Omega_{M} h^2 \\lesssim 1$ implies the following constraint on the mass of the DM particle, also called \"unitarity bound\" \\cite{Unitarity bound} $$m_{DM} \\lesssim 340 \\mbox{ TeV}.$$ Applying the more stringent constraint, obtained by WMAP, the upper bound on $m_{DM}$ becomes: $$ m_{DM} \\lesssim 120 \\mbox{ TeV}.$$  \\begin{figure*}[tb] \\includegraphics[width=6.5cm]{Coannihleptons.eps} \\includegraphics[width=6.5cm]{Coannquarks.eps} \\caption{Left: Relic Abundance of $B^{(1)}$ in the UED model as a function of its mass after including no coannihilation (black line), coannihilation with all leptons (blue) and all electroweak particles (red). For the cases with coannihilation, the solid and dashed lines are computed with a mass splitting $\\delta = 0.01$ and $0.05$ respectively. Right: The same as in the left panel but accounting for coannihilation of $B^{(1)}$ with all electroweak particles and quarks (blue line), and all level-one KK particles, including KK gluons (red line). Solid and dashed lines are for a mass splitting $\\delta = 0.01$ and $0.05$ respectively. From Ref.\\cite{Burnell}.} \\label{fig:coannihilation} \\end{figure*}  However, this constraint was derived under the assumption that particles were in thermal equilibrium in the early universe, thus it applies only to thermal relics, and can be evaded by species which are non-thermally produced.  The standard computation of the thermal relic abundance discussed above presents three important exceptions, as it has been shown, following previous ideas \\cite{prima griest}, by Griest and Seckel \\cite{Griest Seckel}. They take place for WIMPs lying near a mass threshold, for annihilations near to a pole in the cross section, or in presence of coannihilations. The last effect occur when a particle that shares a quantum number with the WIMP, is nearly degenerate in mass with it. If the mass gap is low enough (roughly $\\lesssim 10 \\%$) the coannihilation reactions, involving WIMP particles, can control the WIMP abundance and lower or enhance it.\\\\ Full relic density calculations, including all coannihilations, have been performed e.g. for the supersymmetric neutralino, for which numerical codes such as DarkSUSY \\cite{darksusy} and micrOMEGAs \\cite{micromega} are publicly available. Coannihilations have a dramatic effect on the relic density, and they can lower it by a factor of up to several hundreds (see e.g. \\cite{Edsio}).  %The computation of coannihilation effects in a more general Minimal %Supersymmetric Standard Model contest  Coannihilations are also important in UED models, where the relic density of the first excited state of the B, which may be the lightest Kaluza-Klein particle (LKP) and a viable DM candidate, may be enhanced or lowered depending on the coannihilation channel \\cite{Servant Tait, Burnell, Matchev}. See Fig.\\ref{fig:coannihilation}.\\\\  Deviations from Standard Cosmology can substantially change the picture. For example, due to the presence of scalar fields, the universe may undergo a period of much higher expansion rate and the relic abundance of a WIMP may result increased by several orders of magnitude \\cite{Salati, Masiero}. Furthermore, the production of entropy in the Universe after the WIMP decoupling may dilute its abundance, e.g. due to out-of-equilibrium decays of non-relativistic particles or to first-order phase transitions \\cite{Kawasaki Entropy,Abazajan Koushiappas,Asaka Kusenko Shapo,Ellis Nanopoulos Entropy}.\\\\  It is also possible that DM particles did not experience the thermal history depicted above, and that they have inherited the appropriate relic density through the decay of a more massive species, that has earlier decoupled from the thermal bath. This is e.g. the case for SuperWeakly Interacting Massive Particles (SWIMP), such as the LSP gravitino in SUSY and the first excitation of the graviton in UED, which are produced by late decays of the next-to lightest particles (NLSP/NLKP) in the respective theories \\cite{Feng Rajaraman03, EllisOliveSantoso04, Feng RefPRD03, FengRajaTakaPRD03} and whose relic abundances are simply the rescaled thermal relic densities of the NLPs: $$\\Omega_{\\rm{SWIMP}} = \\frac{m_{\\rm{SWIMP}}}{m_{\\rm{NLP}}} \\Omega_{\\rm{NLP}}.$$  Other production mechanisms may actually be concomitant for such candidates, such as the production at reheating after the end of the inflationary era (see e.g. \\cite{Feng RefPRD03, Boltz Brandeburg Bucch}. See also below for a brief discussion of gravitino production).\\\\  \\subsection*{Other production mechanisms}  Very heavy DM candidates, such as the so-called wimpzillas, have been proposed, with masses as large as $10^{15} \\mbox{ GeV}$, i.e. well above the unitarity limit (see e.g. \\cite{Wimpzillas Review} for a review). For mechanisms that produce these super-massive particles with $\\Omega_{DM}\\sim 1,$ departure from thermal equilibrium is automatic \\cite{Wimpzillas Review}, and the challenge is not to overproduce them. Several mechanisms have been proposed: for instance they could be created during reheating after inflation, with masses a factor $10^3$ larger than the reheating temperature \\cite{Reheating Wimpzillas}, or during a pre-heating stage, with masses up to the Grand Unification scale ($10^{15}$ GeV) \\cite{Preheating Wimpzillas GUT} or even to the Planck Scale \\cite{Preheating Wimpzillas Planck}, or again from bubble collisions if the inflation exit is realized by a first-order transition \\cite{Bubble Wimpzillas}. Another very interesting mechanism is of gravitational nature: wimpzillas may be created by amplification of quantum fluctuations in the transition between the inflationary regime and the matter (radiation) dominated one, due to the nonadiabatic expansion of the spacetime \\cite{Riotto Chung Kolb,Kolb Starobinski}. This scenario can produce particles with mass of the order of the inflaton mass and do not require couplings of wimpzillas with inflaton or other particles.  These particles can be accomodated in existing theoretical frameworks. For instance, stable or metastable bound states called cryptons arise in M-theory, and other possibilities are contemplated in string theories \\cite{pyrons}. Furthermore, messenger bosons in soft supersymmetry breaking models may be very massive and in presence of accidental symmetries in the messengers sectors, might be stable \\cite{yanagida}.  Although wimpzillas have been invoked in top-down scenarios that seek to explain the origin of Ultra High Energy Cosmic Rays \\cite{Kolb Starobinski,Sarkar,Ziaeepour}, this interpretation is today problematic because it predicts a large photon component in the UEHCRs spectrum, in disagreement with the recent results of the Auger experiment \\cite{Semikoz}.  Several production mechanisms can act together to produce a given species, and its relic abundance receives contributions from each of them. The calculation depends on the details of the particle physics and cosmological models adopted. In the case of axions, i.e. light pseudoscalar particles introduced to solve the Strong CP Problem, the production mechanisms in the early Universe are scattering in the hot thermal plasma and possibly radiation by topological defects like axion-strings. Another relevant production mechanism is the so-called {\\it misalignment}: the axion field rolls towards its minimum, near the QCD epoch, and it ends with coherent oscillations that produce a Cold Dark Matter condensate. A lower bound on the axion mass can be inferred requiring that they do not overclose the Universe, but the uncertainties in the calculation of their production make the constraint rather weak (for recent reviews of axions see \\cite{Sikivie,Raffelt Axions}).\\\\  As mentioned before, gravitinos can be copiously emitted by the decay of the NLP in SUSY but they can also be produced, during reheating, by inelastic 2$\\rightarrow$2 scattering processes off particles in the thermal bath and in some scenarios they can act as Cold Dark Matter candidates (e.g. \\cite{graviCDM1, graviCDM2, graviCDM3, Moroi TRehe, graviCDM5, graviCDM6}, see e.g. \\cite{SteffenRevGrav, Buchmuller gravitino Leptog} for more details and references on gravitino DM models). The efficiency of the production depends on the reheating temperature $T_R$ so the bound on $\\Omega_{DM}$ translates into an upper limit on $T_R$ \\cite{Boltz Brandeburg Bucch, Moroi TRehe, Pradler}. In addition to thermal production and late decays of the NLSP, other non thermal and inflation model dependent contributions can arise and change considerably the predictions \\cite{Giudice Riotto, NillesPelosoSorbo}.  Sterile neutrinos, which arise naturally in theoretical frameworks \\cite{GUT Sterile,Mirror Sterile,String Sterile,ExtraDimension Sterile} or in the phenomenological $\\nu$MSM \\cite{nuMSMM}, have been proposed as a solution of the LSND anomaly \\cite{LSND}, as explanation of the high pulsar velocities \\cite{Pulsar velocities} and as Dark Matter candidates. Recently, the MiniBoone collaboration has reported its first results, excluding at 98\\% C.L. the two-neutrino appearance oscillation scheme obtained from LSND data \\cite{Aguilar MiniBoone}. The (3+1) scheme, involving one sterile neutrino specie, is excluded and also models with two or three sterile species are not viable because of the tension between appearance and disappearance data \\cite{Maltoni Schwtz}.  Sterile neutrinos may be produced in the early Universe from collision oscillation conversions of active thermal neutrinos. Their momentum distribution is significantly distorted with respect to a thermal spectrum due to the effects of quark-hadron transition, the modification of the neutrino thermal potential caused by the presence of thermal leptons and the heating of the coupled species (see e.g. \\cite{Abazajan Production} for precise computation of relic abundance). Moreover it has been proposed an enhanced resonant production, in presence of a lepton asymmetry in the early universe significantly higher than the baryonic one \\cite{Abazajian Fuller,Resonant Sterile}. ",
        "conclusions": "In conclusion, we have presented a set of requirements that a particle has to fulfil in order to be considered a viable DM candidate. The requirements are presented in the form of a ten-point test, and we have discussed each of them in a dedicated section that describes the nature of the requirement and guides the reader through the relevant literature.  The test performance of a small subset of DM candidates proposed over the years is shown in Tab.\\ref{tab:tabella}. The \\checkmark symbol is used when the candidates satisfy the corresponding requirement, and it is accompanied by a ! symbol, in the case that present and upcoming experiment will soon probe a significant portion of the candidate's parameter space. If the requirement can be satisfied only in less natural, or non-standard scenarios, or in the case of tension with observational data, the symbol $\\sim$  is used instead. Candidates with a $\\sim$ symbol in the last column, where the final result is shown, should still be considered viable. If one of the requirements is not satisfied, then the symbol $\\times$ is used, and since these requirements are {\\it necessary} conditions, the presence of a single $\\times$ is sufficient to rule out the particle as a viable DM candidate.  \\begin{table*} %[htbp] \\begin{center} \\begin{tabular}{|p{5 cm}||*{11}{c|}|} \\hline & {\\bf I.} & {\\bf II.} & {\\bf III.} & {\\bf IV.} & {\\bf V.} & {\\bf VI.} & {\\bf VII.} & {\\bf VIII.} & {\\bf IX.} & {\\bf X.}& {\\bf Result}\\\\ {\\bf {\\it DM candidate}}  &   $\\Omega h^2$   & Cold & Neutral & BBN & Stars & Self & Direct & $\\gamma$-rays & Astro & Probed &  \\\\ \\hline \\hline  SM Neutrinos & $\\times$ & $\\times$ & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark & -- & -- & \\checkmark &$\\times$ \\\\ \\hline {\\bf Sterile Neutrinos} & $\\sim$ & $\\sim$ & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark ! & \\checkmark & $\\sim$ \\\\ \\hline {\\bf Neutralino} & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark ! & \\checkmark ! & \\checkmark ! & \\checkmark & \\checkmark \\\\ \\hline {\\bf Gravitino} & \\checkmark & \\checkmark & \\checkmark & $\\sim$ & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark & $\\sim$ \\\\ \\hline {\\bf Gravitino (broken R-parity)} & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark \\\\ \\hline Sneutrino $\\tilde{\\nu}_{L}$ & $\\sim$ & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark & $\\times$ &\\checkmark  ! & \\checkmark ! & \\checkmark & $\\times$ \\\\ \\hline {\\bf Sneutrino} $\\tilde{\\nu}_{R}$ & \\checkmark & \\checkmark & \\checkmark& \\checkmark & \\checkmark & \\checkmark & \\checkmark ! & \\checkmark ! &\\checkmark  ! & \\checkmark & \\checkmark \\\\ \\hline {\\bf Axino} & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark \\\\ \\hline {\\bf SUSY Q-balls} & \\checkmark & \\checkmark & \\checkmark & \\checkmark & $\\sim$& -- & \\checkmark ! & \\checkmark & \\checkmark & \\checkmark  & $\\sim$\\\\ \\hline {\\bf $B^1$ UED} & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark ! & \\checkmark !& \\checkmark ! & \\checkmark & \\checkmark \\\\ \\hline First level graviton UED & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark & $\\times$ & $\\times$ & \\checkmark & $\\times$$^a$ \\\\ \\hline {\\bf Axion} & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark ! & \\checkmark & \\checkmark & \\checkmark & \\checkmark \\\\ \\hline {\\bf Heavy photon (Little Higgs)} & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark !& \\checkmark ! & \\checkmark & \\checkmark \\\\ \\hline {\\bf Inert Higgs model} & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark ! & -- & \\checkmark & \\checkmark \\\\ \\hline Champs & \\checkmark & \\checkmark & $\\times$ & \\checkmark & $\\times$ & -- & -- & -- & -- & \\checkmark & $\\times$ \\\\ \\hline {\\bf Wimpzillas} & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark & \\checkmark & $\\sim$& $\\sim$ \\\\ \\hline \\end{tabular} \\end{center} \\caption{Test performance of selected DM candidates. The \\checkmark symbol is used when the candidates satisfy the corresponding requirement, and it is accompanied by a ! symbol, in the case that present and upcoming experiment will soon probe a significant portion of the candidate's parameter space. If the requirement can be satisfied only in less natural, or non-standard scenarios, or in the case of tension with observational data, the symbol $\\sim$  is used instead. Candidates with a $\\sim$ symbol in the last column, where the final result is shown, should still be considered viable. If one of the requirements is not satisfied, then the symbol $\\times$ is used, and since these requirements are {\\it necessary} conditions, the presence of a single $\\times$ is sufficent to rule out the particle as a viable DM candidate. Footnotes: $^a$ It is possible to reconcile a graviton LKP scenario with CMB and diffuse photon background measurements, if the minimal UED model is extended with right-handed neutrinos, Ref.\\cite{MatsumotoSato}.} \\label{tab:tabella} \\end{table*}  %\\footnotetext{It is possible to reconcile a graviton LKP scenario %with CMB and diffuse photon background measurements and BBN %constraints, if the minimal UED model is extended with right-handed %neutrinos, Ref. }"
    },
    "0711/0711.3734_arXiv.txt": {
        "abstract": "{} {We investigate mid-infrared and X-ray properties of the dusty torus invoked in the unification scenario for active galactic nuclei by using the relation between mid IR and hard X-ray luminosities to constrain the geometry and physical state of the dusty torus.} {We present new VISIR observations of 17 nearby AGN and combined these with our earlier VISIR sample of 8 Seyfert galaxies. After combining these observations with X-ray data from the literature, we studied the correlation between their mid IR and hard X-ray luminosities.} {A statistically highly significant correlation is found between the rest frame $12.3 \\, \\mu$m ($L_{\\mathrm{MIR}}$) and 2-10 keV ($L_{\\mathrm{X}}$) luminosities. Furthermore, with a probability of 97 \\%, we find that Sy 1 and Sy 2 nuclei have the same distribution of $L_{\\mathrm{MIR}}$ over $L_{\\mathrm{X}}$.} {The high resolution of our MIR imaging allows us to exclude any significant non-torus contribution to the AGN mid IR continuum, thereby implying that the similarity in the $L_{\\mathrm{MIR}}$ / $L_{\\mathrm{X}}$ ratio between Sy 1s and Sy 2s is intrinsic to AGN. We argue that this is best explained by clumpy torus models. The slope of the correlation is in good agreement with the expectations from the unified scenario and indicates little to no change in the torus geometry with luminosity. In addition, we demonstrate that the high angular resolution is crucial for AGN studies in the IR regime.} ",
        "introduction": "The unification model for active galactic nuclei (AGN) interprets the different appearance of Seyfert 1 (Sy1) and Seyfert 2 (Sy2) galaxies uniquely as the result of an orientation effect \\citep{anton93,barthel94,urry96}. The central engine is considered to be surrounded by an optically and geometrically thick molecular torus. Associated with this torus are large masses of dust that supposedly reprocess the X-ray and UV emission from the accretion disk and re-emit it in the mid infrared (MIR) regime \\citep{pier93}.  It is thus very attractive to search for correlations between IR continuum and hard X-ray emission in order to test the unification scenario for AGN. A tight correlation between the $10.5 \\, \\mu$m continuum and the absorption-corrected 2-10 keV luminosities for 8 nearby Seyfert galaxies was reported by \\citet{krabbe01} using $1\\farcs 2$ resolution MIR imaging. More recently, \\citet{lutz04} found a correlation between the rest frame $6 \\, \\mu$m luminosity and the absorption-corrected hard X-ray luminosity for a sample of $71$ AGN. This sample was comprised of objects for which $24 \\arcsec$ angular resolution ISOPHOT spectra and hard X-ray observations were available; in particular it does not contain Compton-thick objects. However, the authors reported two problems that the unification scenario faces when compared  with their observations. (I) The scatter of the relation is about an order of magnitude larger than expected from the results of \\citet{krabbe01}. (II) There is no significant difference between type 1 Seyferts (Sy types 1 to 1.5) and type 2 Seyferts (Sy types 1.8 to 2) objects in the average ratio of mid-infrared to hard X-ray emission, as would be expected from an optically and geometrically thick as well as smooth torus dominating the mid IR AGN continuum. This was also supported by \\citet{alonso02} who investigated the correlation between the $10 \\, \\mu$m luminosity and the black hole mass of AGN. In an earlier paper \\citep[][hereafter Paper~I]{horst06}, \\defcitealias{horst06}{Paper~I} we presented VISIR observations of 8 nearby Seyfert nuclei. With our angular resolution of $0 \\farcs 35$ at FWHM we then improved upon \\citet{krabbe01} and \\citet{lutz04} by a factor of 3 and 80, respectively, thus minimising contributions from extra-nuclear emission. We found a strong correlation between the rest frame $12.3 \\, \\mu$m and 2-10 keV luminosities with type 1 and type 2 Seyferts following the same correlation of $L_{\\mathrm{MIR}} \\propto L_{\\mathrm{X}}^{1.60 \\pm 0.22}$. We interpreted this as a strong indication for the obscuring medium to be clumpy rather than smoothly distributed and thus appearing as optically thin in the mid infrared. Moreover, we found the slope to be in good agreement with the theoretical predictions by \\citet{beckert04}.  Here we present results from our enlarged AGN sample which contains 21 newly observed objects, out of which 17 were detected with VISIR. Thus, we can now analyse the mid IR and hard X-ray properties of 25 objects.  Throughout this paper we assume $H_{0} = 73$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\Lambda} = 0.72$ and $\\Omega_{\\mathrm{m}} = 0.24$ \\citep{spergel06}. ",
        "conclusions": "\\label{conclusion}  We have presented new high resolution mid infrared photometry of 17 AGN, increasing our total sample to 25 detected and 4 non-detected objects. We have found angular resolution to be a crucial requirement for our study as poorly resolved sources can be heavily contaminated by non-AGN emission in the mid IR band.  The rest frame $12.3 \\, \\mu$m luminosities $L_{\\mathrm{MIR}}$ of our well resolved sources strongly correlate with their rest frame 2-10 keV luminosities $L_{\\mathrm{X}}$, matching the expectations from the unified scenario for AGN. With a probability of 97 \\%, type 1 and type 2 Seyferts have the same distribution of $L_{\\mathrm{MIR}} / L_{\\mathrm{X}}$. This similarity is most probably intrinsic to the AGN and not caused by extra-nuclear emission contaminating the MIR flux. These results have two important implications: \\begin{itemize} \\item As we do not find the offset in $L_{\\mathrm{MIR}} / L_{\\mathrm{X}}$ between type 1s and type 2s, predicted by smooth torus models, AGN tori are likely to be clumpy. The volume filling factor of these clumps has to be small in order to reproduce our results. \\item At a $7 \\sigma$ level, the slope of the correlation is not compatible with the predictions of the receding torus model by \\citet{simpson05}. We suggest that the break luminosity for this model is higher than hitherto assumed. A population of Compton thick AGN at moderate luminosities being missed in current X-ray surveys could resolve this problem. \\end{itemize} The question, whether AGN at low luminosities, with $L_{\\mathrm{bol}} < 10^{42}$ erg s$^{-1}$ have a qualitatively different appearance than objects of higher luminosities, remains open. Fortunately, we have recently been awarded time to observe a sample of 14 heavily obscured ($N_{\\mathrm{H}} = 10^{24} \\sim 10^{25}$ cm$^{-2}$), mostly low luminosity AGN with VISIR which may help to shed light on this issue."
    },
    "0711/0711.3819_arXiv.txt": {
        "abstract": "{The \\gamdorr\\ stars are a recent class of variable main sequence F-type stars located on the red edge of the Cepheid instability strip. They pulsate in gravity modes, and this makes them particularly interesting for detailed asteroseismic analysis, which can provide fundamental knowledge of properties near the convective cores of intermediate-mass main sequence stars.}  % {To improve current understanding of \\gamdor\\ stars through theoretical modelling, additional constraints are needed. Our aim is to estimate the fundamental atmospheric parameters and determine the chemical composition of these stars. Detailed analyses of single stars have previously suggested links to Am and \\lamboo\\ stars, so we wish to explore this interesting connection between chemical peculiarity and pulsation.} % {We have analysed a sample of \\gamdor\\ stars for the first time, including nine bona fide and three candidate members of the class. We determined the fundamental atmospheric parameters and compared the abundance pattern with other A-type stars. We used the semi-automatic software package \\vwa\\ for the analysis. This code relies on the calculation of synthetic spectra and thus takes line-blending into account. This is important because of the fast rotation in some of the sample stars, and we made a thorough analysis of how \\vwa\\ performs when increasing \\vsini. We obtained good results in agreement with previously derived fundamental parameters and abundances in a few selected reference stars with properties similar to the \\gamdor\\ stars. }  % {We find that the abundance pattern in the \\gamdor\\ stars is not distinct from the constant A- and F-type stars we analysed.} % {} % ",
        "introduction": "The members of the \\gamdor\\ class of variable stars are found near the main sequence on the cool edge of the Cepheid instability strip with spectral types A7--F5. They thus share properties with \\dss\\ star variables, but the \\gamdor\\ periods are an order of magnitude longer, indicative of $g$~mode pulsation. The \\gamdor\\ phenomenon was first identified by \\cite{balona94}, and \\cite{kris95} presented the first list of six candidates. \\cite{henry07} presents a list of 66 \\gamdor\\ stars, and the group continues to grow as new members are discovered both among field stars \\citep{mathias04, henry05, decat06} and in open clusters \\citep{arentoft07}. Only a few stars show pulsations characteristic of \\gamdor\\ and \\dss\\ stars simultaneously \\citep{henry05,rowe06,king07}.  The \\gamdor\\ stars have given new hope for a deeper understanding of main sequence stars with masses around 2\\,\\msun\\ through asteroseismic analyses. Several \\dss\\ stars have been studied extensively through both photometry and spectroscopy, and dozens of individual modes are now known in a few field stars \\citep{breger05} and also in members of open clusters \\citep{bruntt07}. Whilst observational work has been very successful, comparison with theoretical models has so far not been able to provide a fully adequate description of all the observations (see \\citealt{zima06} for recent developments). While the driving in \\dss\\ stars is well understood in terms of the opacity or $\\kappa$ mechanism, the link to predicting observed mode amplitudes is weak. Furthermore, theoretical models of \\dss\\ stars show that even moderate rotation leads to significant shifts in the mode frequencies \\citep{suarez06a,suarez06b}, which complicates the confrontation of observations and models.  The theoretical framework for interpreting the observed pulsation in \\gamdor\\ stars is well under way. Pulsations are thought to be driven by a flux blocking mechanism near the base of their convective envelopes \\citep{dupret04, dupret06}. \\cite{moya05} investigated a method of constraining the models of \\gamdor\\ using frequency ratios. This method has been attempted on individual \\gamdor\\ stars \\citep{moya05,rodriguez06}. % While the method is indeed very useful for providing constraints, no unique models that fit all the observations were found. An improvement would be to better constrain the fundamental atmospheric parameters including metallicity. The star studied by \\cite{moya05}, HD~12901, is included in our sample.  In the current study we carry out a detailed abundance analysis of a sample of \\gamdor\\ stars described by (\\citealt{decat06}; hereafter \\paper) Thus, the current work is the second part of our detailed spectroscopic analysis of a sample of southern candidate \\gamdor\\ stars. In \\paper\\ we made a detailed analysis of the spectra to study binarity and the pulsation properties. We identified 10 new bona fide   % \\gamdor\\ stars of which 40\\% are binary stars.  Detailed abundance analyses of \\gamdor\\ stars have only been done for a few individual stars. \\cite{bruntt02} analysed the \\gamdor\\ star HD~49434 and found a metallicity slightly below solar, but the analysis was hampered by the high \\vsini\\,$=85$\\,\\kms. \\cite{sadakane06} analysed HD~218396 and found solar abundance of C and O (but not S) and abundances of iron peak elements of $-0.5$~dex, thus suggesting a \\lamboo\\ nature for this star. \\cite{henry05} found evidence that the \\gamdor\\ star HD~8801 is an Am star based on the strength of the Ca~K line. Our aim is to shed light on these intriguing links that have been suggested between the \\gamdor\\ variables and the chemically peculiar \\lamboo\\ and Am-type stars \\citep{gray99,sadakane06}. ",
        "conclusions": "We have presented a detailed abundance analysis of a group of bona fide and candidate \\gamdor\\ stars. In addition we analysed a number of constant stars with similar stellar parameters. There seems to be larger scatter in the abundances for the \\gamdor\\ stars, but we find no strong evidence that the overall abundance pattern is different from other A- and F-type stars. Furthermore, the metallicity is quite close to the solar value in all cases. We have constrained the fundamental parameters of 18 single field stars from \\paper, of which about half are potential \\gamdor\\ stars.  We also analysed a few reference stars in order to thoroughly test the performance of the \\vwa\\ software package. The software gives reliable results for the value of \\teff\\ and \\logg\\ for our primary and secondary reference stars, \\ie\\ the Sun and the astrometric binary HD~110379; the latter has a well-determined \\logg\\ but poorly determined \\teff. Our results also agree well with previous analyses of three tertiary reference A- and F-type stars, although these single field stars do not have well-determined values of \\teff\\ and \\logg. Our analysis of synthetic spectra with increasing rotational velocity shows that, for stars with $v \\sin i > 50$\\,\\kms, our method cannot be used to constrain the microturbulence, \\teff, and \\logg\\ simultaneously. For the slowly rotating stars with $v \\sin i < 25$\\,\\kms, we can constrain \\teff, \\logg, and \\feh\\ to about 120~K, 0.13~dex, and 0.09~dex including estimated uncertainties of the applied model atmospheres. These results are certainly an improvement over photometric uncertainties, which are typically at least twice as large. We expect that our results will be useful in future asteroseismic studies of \\gamdor\\ stars."
    },
    "0711/0711.1394_arXiv.txt": {
        "abstract": "We present the chemical compositions for eight bright giants in the globular cluster NGC 1851. Our analysis reveals large star-to-star abundance variations and correlations of the light elements O, Na, and Al, a feature found in every well studied globular cluster. However, NGC 1851 also exhibits large star-to-star abundance variations of the $s$-process elements Zr and La. These $s$-process elements are correlated with Al, and anticorrelated with O. Furthermore, the Zr and La abundances appear to cluster around two distinct values. A recent study revealed a double subgiant branch in NGC 1851. Our data reinforce the notion that there are two stellar populations in NGC 1851 and indicate that this cluster has experienced a complicated formation history with similarities to $\\omega$ Centauri. ",
        "introduction": "\\label{sec:intro}  Recent studies have revealed that the Galactic globular clusters $\\omega$ Centauri \\citep{bedin04}, NGC 2808 \\citep{dantona05,piotto07}, and M 54 \\citep{siegel07} contain multiple stellar populations. The evidence consists of HST color-magnitude diagrams (CMDs) in which multiple main sequences have been identified. In all cases, the most plausible explanation is that these globular clusters contain populations with distinct compositions and/or ages, in contrast to the classical concept that globular clusters are monometallic, coeval, simple stellar populations. The most recent addition to this intriguing collection is NGC 1851, a globular cluster with a bimodal horizontal branch (HB), whose CMD displays a double subgiant branch \\citep{milone07}.  Both $\\omega$ Centauri and M 54 possess a large metallicity spread and are the two most massive Galactic globular clusters. $\\omega$ Centauri is widely regarded as the nucleus of an accreted dwarf galaxy (e.g., \\citealt{smith00}) and M 54 is suspected to be the nucleus of the Sagittarius dwarf spheroidal galaxy (e.g., \\citealt{layden00}). However, not all clusters with multiple main sequences show a metallicity spread and although there is speculation that all globular clusters may have been the nuclei of ancient dwarf galaxies \\citep{bekki06}, currently there is little direct observational evidence. For NGC 2808, the multiple main sequences and complex HB morphology can be attributed to a large He abundance variation \\citep{dantona04,dantona05,piotto07}, but the Fe abundances and ages show little, if any, variation among the populations. Although there are large star-to-star abundance variations of O, Na, and Al in NGC 2808 \\citep{carretta06}, such patterns are ubiquitous in Galactic globular clusters (e.g., \\citealt{kraft94} and \\citealt{gratton04}).  The situation regarding the chemical compositions for stars in the globular cluster NGC 1851 is less clear due to the lack of data. The sole spectroscopic abundance analysis was conducted by \\citet{hesser82}. Low resolution spectra revealed extremely strong CN bands for three out of eight bright red giants. \\citet{hesser82} also found that the Sr and Ba lines may be enhanced in CN-strong stars relative to other giants in this cluster. Such abundance patterns, if confirmed, would bear a striking resemblance to $\\omega$ Centauri. A study of the stellar chemical compositions in NGC 1851 is of great interest in order to provide constraints upon the interpretation of the double subgiant branch and to explore the chemical evolution of this cluster. ",
        "conclusions": "\\label{sec:discussion}  \\citet{milone07} explored various possibilities for explaining the double subgiant branch in NGC 1851. If age is the sole parameter, a difference of 1 Gyr would be required. If composition is the sole parameter, any abundance variation invoked to account for the double subgiant branch is heavily constrained by the width of the main sequence (and giant branch). \\citet{milone07} estimate that the maximum possible abundance variation on the main sequence would be $\\Delta$[Fe/H] = 0.1 dex or a helium abundance $\\Delta Y$ = 0.026. They argue that the He abundance alone cannot explain the observed subgiant branch. They suggest that a population with a 0.2 dex metallicity increase could explain the subgiant branch, but that such an abundance spread is precluded by the width of the main sequence and giant branch. However, a combination of increased [Fe/H] by 0.2 dex and helium from Y = 0.247 to Y = 0.30 could reproduce the observed subgiant branch behavior as well as ensuring a sufficiently narrow main sequence and giant branch.  Available measurements indicate that the abundances of Fe, O, Na, Al, and neutron capture elements do not change as a function of evolutionary status from the main sequence to the giant branch in globular clusters (e.g., \\citealt{gratton01}, \\citealt{james04} and \\citealt{cohen05}). If we assume that the abundances for giant stars are representative of subgiant stars in NGC 1851, then our measured abundances eliminate the possibility that the subgiant branch split is due to an abundance difference of 0.2 dex for Fe and a shift of $Y$ from 0.247 to 0.30. While our sample size is limited, there is no hint of a bimodal distribution of Fe abundances separated by 0.2 dex, as required for the chemical composition explanation for the double subgiant branch. Indeed, \\citet{milone07} noted that such a possibility is also precluded by the magnitudes of blue HB stars. Interestingly, the dispersion in our Fe abundances ($\\sigma$ = 0.09 dex) is within the 0.1 dex ``limit'' imposed by the width of the main sequence.  The abundances of additional elements add further complexity to NGC 1851. The star-to-star abundance variation of O, Na, and Al and the anticorrelation of O and Na exist in NGC 1851 and indeed in all globular clusters. While our sample size is small, the amplitude of these abundance variations is comparable to clusters at a similar metallicity. However, the light $s$-process element Zr and the heavy $s$-process element La exhibit a large star-to-star abundance variation. (Our spectra also indicate a large range in Ba line strengths, and presumably Ba abundance, but we are unable to quantify the range.) Furthermore, the Zr and La abundances appear to have a bimodal distribution and reinforce the idea that there are two stellar populations in this cluster. (The ratio of heavy to light $s$-elements, [La/Zr] is similar for the ``two populations''.) La and Al (and Zr and Al) are correlated and so the source of the Al variation is likely to be the source of the La variation. The hint of a correlation between Al and Zr has also been found in NGC 6752 \\citep{yong05}. Asymptotic giant branch (AGB) stars may produce large enhancements of Al and La, albeit via different processes in AGB stars of different masses (e.g., \\citealt{busso99}, \\citealt{karakas03}). The He, C, N, and O produced by such AGB stars may play a role on the age differences for the subgiant branch populations.  \\citet{milone07} raise the tantalizing prospect that the relative frequency of stars on the fainter/brighter subgiant branches (45\\% vs.\\ 55\\%) roughly matches the relative frequency of HB stars bluer/redder than the instability strip (37\\% vs.\\ 63\\%). The fraction of CN-strong, Ba-strong, Sr-strong stars found by \\citet{hesser82} was $\\sim$ 40\\%. In this study, the fraction of Zr-strong, La-strong stars was also $\\sim$ 40\\%. Based on the relative numbers, we speculate that the brighter subgiant branch stars have ``normal'' CN, Sr, Zr, Ba, and La abundances and populate the red HB. We also speculate that the fainter subgiant branch stars are CN-strong, enriched in Sr, Zr, Ba, and La, and populate the blue HB. Indeed if the CN, Zr, and La enriched stars are also He-rich, they are expected to populate the blue HB \\citep{sweigart97}. However, we caution that the relative numbers of the two populations will not be the same at all evolutionary phases. Nevertheless, the abundances in subgiant branch, giant, and red HB stars can test this hypothesis.  Since the fainter subgiant branch stars are older than the brighter subgiant branch stars (assuming age is the only parameter controlling the double subgiant branch), then our tentative speculation would require that the Zr-rich, La-rich stars are younger than the Zr-normal, La-normal stars. While this unpalatable scenario would be a challenge to explain, \\citet{villanova07} find that the most metal-rich subgiant stars in $\\omega$ Centauri are among the oldest stars in the cluster. Clearly, the sequence of events that led to the formation of the different stellar populations in NGC 1851 and $\\omega$ Centauri requires some imagination. On the other hand, an analysis of a larger sample of giant stars may reveal that the CN-strong, Zr-rich, La-rich stars are associated with the majority population of brighter subgiant branch stars. In this case, if the CN-strong, Zr-rich, La-rich stars are the progeny of the CN-normal, Zr-normal, La-normal stars, the assumed 1 Gyr age difference would then place constraints upon the mass range of the AGB stars that produced the CN, Zr, and La excess. At the metallicity of NGC 1851, Karakas (2007 priv.\\ comm.) suggests that the minimum mass (ZAMS) of the AGB stars would be approximately 2 M$_{\\odot}$. Further, the AGB stars that produce La have M $\\gtrsim$ 1 M$_\\odot$ and the AGB stars that produce Al have M $\\gtrsim$ 4 M$_\\odot$.  Finally, we note that the only other globular cluster that exhibits a large range in C, N, O, Na, Al, Sr, Zr, Ba, and La abundances is $\\omega$ Centauri. Both NGC 1851 and $\\omega$ Centauri display multiple subgiant branches which are presumably due to stellar populations with different ages and compositions. Further, NGC 1851 and $\\omega$ Centauri are the only clusters that display a large variation in the Str{\\\" om}gren $m_1$ index, traditionally used as a metallicity indicator. NGC 1851 appears to exhibit a bimodal $m_1$ distribution on the giant branch suggesting an extreme variation in C and/or N abundances (Grundahl 2007 in preparation). In the context of the formation of the Galactic globular cluster system, we speculate that NGC 1851 may represent a ``bridge'' between $\\omega$ Centauri (large variation of all elements) and NGC 6752-like clusters (constant Fe but large variations of light elements C-Al). A more detailed analysis of additional elements and additional stars in NGC 1851 is of great interest."
    },
    "0711/0711.1677_arXiv.txt": {
        "abstract": "If a class of stars orbits the central black hole in our galaxy in short period ($\\sim 0.1$ year), high eccentricity ($\\sim 0.9$) orbits, they will experience precessions of their orbital planes induced by both relativistic frame-dragging and the quadrupolar gravity of the hole, at levels that could be as large as 10 $\\mu$arcseconds per year, if the black hole is rotating faster than 1/2 of its maximum rotation rate.  Astrometric observations of the orbits of at least two such stars can in principle lead to a determination of the angular momentum vector ${\\bf J}$ of the black hole and its quadrupole moment $Q_2$.  This could lead to a test of the general relativistic no-hair theorems, which demand that $Q_2 = - J^2/M$.   Future high-precision adaptive infrared optics instruments may make such a fundamental test of the black-hole paradigm possible. ",
        "introduction": "\\label{secI}  The discovery, using infrared telescopes, of stars orbiting within an arcsecond of the central object SgrA$^*$ in our galaxy, together with accurate determinations of their orbits, has provided strong evidence for the existence there of a massive black hole (MBH) of around $3.6 \\times 10^6 \\,M_\\odot$ (see \\cite{alexander05} for a review). In addition to opening a window on the innermost region of the galactic center, the discovery of these stars has made it possible to contemplate using orbital dynamics to probe the curved spacetime of a rotating black hole.  The orbital periods of these stars are on the scale of tens of years, and thus most relativistic effects, such as the advance of the pericenter, are too small to be observed at present (see, however \\cite{zucker}). Nevertheless, there seems to be every expectation that, with improved observing capabilities,  a population of stars significantly closer to the hole will eventually be discovered, making orbital relativistic effects detectable \\citep{jaroszynski,fragile,rubilar,weinberg,kraniotis}.  Furthermore, plans are being developed to achieve infrared astrometry on such objects at the level of 10 $\\mu$arcseconds \\citep{eisenhauer}.  High-precision Doppler measurements may also be possible \\citep{zucker}.  This makes it possible to consider doing more than merely detect relativistic effects, but rather to provide the first test of the black hole ``no-hair'' or uniqueness theorems of general relativity. According to those theorems, an electrically neutral black hole is completely characterized by its mass $M$ and angular momentum $J$.  As a consequence, all the multipole moments of its external spacetime are functions of $M$ and $J$, specifically, the quadrupole moment $Q_2 = -J^2/M$ (in units where $G=c=1$).  If the black hole were non-rotating ($J=0$), then its exterior metric would be that of Schwarzschild, and the most important relativistic effect would be the advance of the pericenter. If it is rotating, then two new phenomena occur, the dragging of inertial frames and the effects of the hole's quadrupole moment, leading not only to an additional pericenter precession, but also to a precession of the orbital plane of the star.  These precessions are smaller than the Schwarzschild effect in magnitude because they depend on the dimensionless angular momentum parameter $\\chi \\equiv J/M^2$, which is always less than one, and because they fall off faster with distance from the black hole. However, accumulating evidence suggests that MBH should be rather rapidly rotating, with $\\chi$ larger than 0.5 and possibly as large as 0.9, so these effects could be significant.  The purpose of this paper is to point out that, if a class of stars were to be found with orbital periods of fractions of a year, and with sufficiently large orbital eccentricities, then the quadrupole-induced precessions could be as large as 10 $\\mu$as per year.   Figure \\ref{fig1} illustrates this: assuming a black hole with $\\chi = 0.7$, it shows the orbital period required as a function of eccentricity, for the rates of precessions due to Schwarzschild (S), frame-dragging ($J$) and quadrupole ($Q_2$) terms to be as large as 10, 5, and 1 $\\mu$as per year.   \\begin{figure}[t] \\begin{center} \\epsfig{file=f1.eps,width=3.0in} \\caption{Orbital periods vs. eccentricity required to give measurable relativistic precession rates.  Dotted curves show minimum periods vs. $e$ that avoid tidal disruption, for various stellar masses. \\label{fig1}} \\end{center} \\end{figure}  \\begin{figure}[t] \\begin{center} \\epsfig{file=f2.eps,width=3.0in} \\caption{Relativistic precession amplitudes vs. black hole spin parameter $\\chi$. \\label{fig2}} \\end{center} \\end{figure}   Figure \\ref{fig2} shows the effect of black hole spin on the amplitudes of the relativistic effects.  For orbits with eccentricity 0.9 and periods of one year and 0.1 years, the amplitudes of the three effects are plotted in $\\mu$as per year.  The precession of the orbital plane is the most important effect here, because it depends only on $J$ and $Q_2$; the Schwarzschild part of the metric affects only the pericenter advance. The orbital plane is determined by its inclination angle $i$ relative to the plane of the sky and by the angle of nodes $\\Omega$ between a reference direction and the intersection of the two planes.  Standard astrometric and Doppler observations can determine $\\Omega$, $i$, the pericenter angle $\\omega$, the semimajor axis $a$ and the orbital eccentricity $e$, and, given sufficient observation time, the secular rates of change $d\\Omega/dt$, $di/dt$, and $d\\omega/dt$.  However, in order to test the no-hair theorems, one must determine five parameters: the mass of the black hole, the magnitude and two angles of its spin, and the value of the quadrupole moment. The ``Kepler-measured'' mass is determined from the orbital periods of stars, but may require data from a number of stars to fix it separately from any extended distribution of mass. Then, to measure $\\bf J$ and $Q_2$, it is necessary and sufficient to measure $d\\Omega/dt$ and  $di/dt$ for two stars in non-degenerate orbits. A test of the no-hairness of the central object in our galaxy would be compelling evidence that it is truly a black hole of general relativity. ",
        "conclusions": "\\label{secV}  We have shown that a class of stars orbiting a rotating central black hole in our galaxy in short period, high eccentricity orbits, will experience precessions of their orbital planes induced by both frame dragging and the quadrupolar gravity of the hole, at levels that could be as large as 10 $\\mu$arcseconds per year.  Observations of the orbits of at least two such stars can in principle lead to a determination of the angular momentum vector ${\\bf J}$ and quadrupole moment $Q_2$ of the black hole, and could provide a test of the no-hair theorems of general relativity.  Alternative possibilities for no-hair tests involve timing measurements of pulsars orbiting black-hole companions \\citep{kopeikin}, gravitational-wave measurements of compact objects spiralling into massive black holes \\citep{ryan,glampedakis,hughes}, or detection of quasi-normal ``ringdown'' gravitational radiation of perturbed black holes \\citep{dreyer,bcw}.  Detecting such stars so close to the black hole, and carrying out infrared astrometry to $10 \\, \\mu$arcsecond accuracy will be a challenge.  However, if this challenge can be met with future improved adaptive optics systems currently under study, such as GRAVITY \\citep{eisenhauer}, it could lead to a powerful test of the black-hole paradigm.  In future work, we plan to study in detail such complicating effects as second-post-Newtonian (2PN) corrections to the Schwarzschild part of the pericenter advance, tidal effects, effects of unseen mass distributions within the observed stellar orbits, and light deflection and Shapiro time delay effects \\citep{rubilar,weinberg}. For example, a torus of matter of mass $m$ orbiting the black hole at a distance $R$ will induce fractional changes in the apparent angular momentum and quadrupole moment of order $\\delta J/J \\sim (m/M)(R/M)^{1/2}(1/\\chi)$ and $\\delta Q/Q \\sim (m/M)(R/M)^2 (1/\\chi)^2$, so only  a very massive and/or very distant torus will be relevant. We also plan to carry out covariance analyses to obtain more realistic estimates of the accuracies that might be obtained for the no-hair test for given raw astrometric accuracies, and for a range of observing schedules."
    },
    "0711/0711.3358_arXiv.txt": {
        "abstract": "Intergalactic space is filled with a pervasive medium of ionized gas, the Intergalactic Medium (IGM). A residual neutral fraction is detected in the spectra of Quasi-Stellar Objects at both low and high redshifts, revealing a highly fluctuating medium with temperatures characteristic of photoionized gas. The statistics of the fluctuations are well-reproduced by numerical gravity-hydrodynamics simulations within the context of standard cosmological structure formation scenarios. As such, the study of the IGM offers an opportunity to probe the nature of the primordial density fluctuations on scales unavailable to other methods. The simulations also suggest the IGM is the dominant reservoir of baryons produced by the Big Bang, and so the principal source of the matter from which galaxies formed. The detection of metal systems within the IGM shows that it was enriched by evolved stars early in its history, demonstrating an intimate connection between galaxy formation and the IGM. The author presents a comprehensive review of the current understanding of the structure and physical properties of the IGM and its relation to galaxies, concluding with comments on prospects for furthering the study of the IGM using future ground-based facilities and space-based experiments. ",
        "introduction": "\\label{sec:intro}  According to the Big Bang theory, the primordial hydrogen and helium created in the Universe first materialized in the form of an extremely hot ionized gas. By the time the Universe was three hundred thousand years old, adiabatic expansion cooling brought the temperature of the primordial plasma down until the hydrogen and helium recombined. The radiation last scattered at this time appears today as the Cosmic Microwave Background (CMB). The search for the IGM began well before the discovery of the CMB with an attempt by \\textcite{1959ApJ...129..525F} to detect the hyperfine 21cm absorption signature from hydrogen along the line of sight to the extragalactic radio galaxy Cygnus A. Although no detection was made, combining the optical depth with the measured temperature of the CMB discovered in 1965 would have been sufficient to exclude the possibility that the Universe was closed by baryons, with an upper limit on the baryon density of only 20\\% of the closure density required for an Einstein-deSitter (flat) universe. \\footnote{This conclusion requires making the (at the time) reasonable assumptions that the hydrogen was all neutral and that the hyperfine structure levels of the hydrogen were in thermal equilibrium with the CMB.}  Nearly coincident with the discovery of the CMB, however, a considerably improved measurement of the density of intergalactic neutral hydrogen was made. Soon after the discovery of the first Quasi-Stellar Object (QSO) \\cite{1963Natur.197.1037H, 1963Natur.197.1040S}, \\textcite{1965ApJ...142.1633G} reported a small decrement in a QSO spectrum shortward of its \\Lya emission line. Attributing the decrement to the \\Lya resonance scattering of radiation from the QSO by intergalactic neutral hydrogen, Gunn and Peterson demonstrated that the cosmic mass density of neutral hydrogen was exceedingly smaller than the Einstein-deSitter density. In fact, it was far smaller than the spatially averaged hydrogen of all the stars in the Universe. If the Big Bang theory was correct, it meant either that galaxy formation was an extraordinarily efficient process, sweeping up all but a tiny residue of the primordial hydrogen, or that the gas was reionized.  These two themes, the detection of intergalactic gas through the 21cm signature in the radio or through Lyman resonance scattering lines in the optical or ultraviolet (UV), continue to dominate studies of the Intergalactic Medium (IGM). To date, almost all that is known about the structure of the IGM has been derived from optical and UV data. This situation is expected to change dramatically in the near future with the development of low frequency radio arrays like the LOw Frequency ARray (LOFAR)\\footnote{http://www.lofar.org}, the Murchison Widefield Array (MWA)\\footnote{http://www.haystack.mit.edu/ast/arrays/mwa/}, the Primeval Structure Telescope (PaST/21CMA) \\footnote{http://web.phys.cmu.edu/$\\sim$past/}, the Precision Array to Probe the Epoch of Reionization (PAPER) \\footnote{http://astro.berkeley.edu/$\\sim$dbacker/eor/}, and a possible Square Kilometre Array (SKA). \\footnote{http://www.skatelescope.org} A primary science driver of these instruments is the direct imaging of the IGM prior to the completion of the Epoch of Reionization \\cite{MMR97}. Most of this review focuses on the current understanding of the state of the IGM as determined from optical and UV measurements. The observations have relied almost exclusively on the spectra of QSOs, although IGM absorption features have also been detected in the spectra of Gamma-Ray Bursts (GRBs) \\cite{2006PASJ...58..485T}, and indeed played a key role in establishing the extragalactic character of some bursts \\cite{1997Natur.387..878M}.  Almost immediately after Gunn and Peterson publicized their finding, it was recognized that individual \\Lya absorption features should appear from neutral hydrogen concentrated into cosmological structures \\cite{1965ApJ...142.1677B, 1967ApJ...149..465W}. Absorption features had in fact been detected in higher resolution QSO spectra, but these were identified with intervening ionized metal absorption systems \\cite{1968ApJ...153..689B}, as was expected if galaxies had hot halos of ionized gas:\\ the lines of sight to background QSOs were expected to pass through such hot halos and intercept any ionized gas clouds within them \\cite{1969ApJ...156L..63B}. The features, however, were uncomfortably common, hinting at a class of unknown structures not associated with galaxies. The \\Lya resonance line features continued to prove elusive until 1971, when \\textcite{1971ApJ...164L..73L} recognized several absorption features shortward of the \\Lya emission line of a QSO as \\Lya lines\\footnote{These are not true absorption features involving the net destruction of a photon, but the scattering out of the line of sight of resonance line photons.} arising in a population of discrete absorption systems also showing metal features. The \\Lya lines form a plethora of distinct absorption features in the spectra of high redshift QSOs; they are collectively known today as the \\Lya forest.  The properties of the \\Lya forest came under increasing scrutiny, with the first systematic survey conducted several years later by \\textcite{1980ApJS...42...41S}, convincingly demonstrating through the homogeneity of the observed properties of the absorbers their intergalactic origin, as opposed to clouds ejected by the QSOs observed. Although limited by the resolution of the spectrograph, the measured velocity widths of the \\Lya features corresponded to gas temperatures of a few to several $10^4$~K, characteristic photoionization temperatures for gas of a primordial composition, for which there is no significant cooling by metals. The number of features per comoving length was shown to increase with redshift, demonstrating that the systems were evolving \\cite{1982ApJ...252...10Y}.  The past decade has witnessed a dramatic improvement in precision studies of the \\Lya forest with the advent of 8--10~m class telescopes, particularly the Keck 10-m and the 8.2m Very Large Telescope (VLT). For the first time, the individual absorption features in the \\Lya forest were spectroscopically resolved over their full range. Velocity widths of $\\sim25\\kms$ are typical. The neutral hydrogen column densities of the absorbers range from roughly $10^{12}-10^{22}\\,{\\rm cm^{-2}}$. The highest column density systems, the Damped \\Lya Absorbers (DLAs), are of particular interest for galaxy formation, as they are suspected of containing the neutral gas that formed the bulk of the stars in present day galaxies.  As the number per length of absorption systems increases along a line of sight with increasing redshift, so does the mean flux decrement in a background QSO spectrum due to \\Lya scattering. The QSO spectra measured as part of the Sloan Digital Sky Survey (SDSS) \\footnote{http://www.sdss.org} show a rapid rise in the flux decrement at $z\\gta5.5$, suggesting that the epoch of \\HI\\ reionization may lie not far beyond $z\\simeq6$. Many of the hydrogen features also show absorption lines from metals, including carbon, silicon, nitrogen, oxygen, magnesium, iron and others. The abundances of the metals, however, are at most about 10\\% of solar at low redshifts, and as low as 1\\% at high redshifts, indicating that the absorption systems are comprised largely of primordial material. The primordial nature of the gas received further important confirmation in 1994 with the discovery by \\textcite{1994Natur.370...35J} of intergalactic helium using the {\\it Hubble Space Telescope} ({\\it HST}) through the detection of \\HeII\\ \\Lya absorption.  Because the baryons in the IGM are detected only though their absorption signatures, the physical structures that give rise to the features must be modeled. Early models characterized the systems as discrete clouds of gas, with most of the focus on either clouds pressure-confined by a hot medium \\cite{1980ApJS...42...41S}, or gravitationally confined in a dark matter halo \\cite{1986Ap&SS.118..509I,1986MNRAS.218P..25R}. At the time it was believed that the absorption systems accounted for only a few percent of the baryons produced in the Big Bang, much like galaxies, their visible counterparts. A paradigm shift in the models occurred in the mid 1990s. Measurements of coincident absorption features along neighboring lines of sight suggested sizes of tens to hundreds of kiloparsecs for the absorbers \\cite{1992ApJ...389...39S, 1994ApJ...437L..83B, 1997ApJ...491...45D}, much larger than expected for clouds confined by pressure or dark matter halos. A radical transformation in the understanding of the nature and structure of the IGM was initiated by numerical simulations of cosmological structure formation. Today essentially all of the baryons produced in the Big Bang are believed to have undergone the same gravitational instability process initiated by primordial density fluctuations that was responsible for the formation of galaxies. The computation of the structure of the IGM has been converted into an initial value problem similar to that of the CMB fluctuations. Fluctuations in the CMB are solved for by following the gravitational instability growth of a postulated spectrum of primordial dark matter density fluctuations. The growth of structure in the IGM is now treated as the nonlinear extension of these computations. The result is a network of filamentary structures, the so-called ``cosmic web'' \\cite{1996Natur.380..603B}. The \\Lya forest is believed to be a signature of the cosmic web. Early simulations broadly reproduced the statistics of the \\Lya forest spectacularly well \\cite{Cen94, 1995A&A...295L...9P, ZAN95, 1996ApJ...457L..51H, 1997ApJ...485..496Z, Theuns98}. An immediate conclusion was that at $z\\gta1.5$, some 90\\% of the baryons produced in the Big Bang are contained within the IGM, with only 10\\% in galaxies, galaxy clusters or possibly locked up in an early generation of compact stars.  Soon after the discovery of intervening absorption features, it was recognized that they provided potentially powerful tests of fundamental properties of the Universe. The split in the fine structure lines of the metals was used to set constraints on the variability of the fine structure constant \\cite{1967ApJ...149L..11B}. A bunching of absorption features near $z\\simeq1.9$ (now known to be fortuitous), gave rise to the (re)introduction of a cosmological constant to account for the numerous features as multiple images due to lensing \\cite{1967ApJ...147.1222P}. The expected primordial composition of the IGM offered the potential of placing constraints on the photon-to-baryon ratio of the Universe through measurements of the intergalactic deuterium abundance. More recently, the success of the models has inspired attempts to exploit the \\Lya forest as a new means to constrain cosmological structure formation models and obtain stringent constraints on the cosmological parameters.  The description of the IGM by the simulations, however, is far from complete. There remain many unsolved problems. The current simulations do not reproduce all the observed properties of the IGM. The absorption lines are predicted to be substantially narrower than measured. This likely stems from the principal outstanding missing piece of physics, the reionization of the IGM. Not only must hydrogen be ionized, but helium as well. The ionization heats the gas through the photoelectric effect. Detailed radiative transfer computations are required to recover the temperatures, for which there is still limited success. The sources of the reionization and the epochs of reionization, both of hydrogen and helium, are still not firmly established. The origin of the metal absorption systems in the diffuse IGM is still unknown, although it is widely expected they were deposited by winds from galaxies, possibly driven by intense episodes of star formation. As such, the metal absorption lines in principle offer an important means of studying the history of cosmic star formation. Most fundamentally, the relation of the IGM to the galaxies that form from it is still mostly unknown, but offers perhaps the most exciting prospects for new research directions.  The purpose of this review is to describe the progress made in the understanding of the origin of structure within the IGM, with a view to presenting the underlying physics that determines the structure. An understanding of the physics is necessary for future progress. The past decade has revealed the IGM to be a complex dynamical arena involving interactions between the IGM gas, galaxies and QSOs. It is becoming increasingly apparent that the separation of these systems into distinct and isolated entities is an artificial construct. Galaxies and QSOs originated from the IGM, and their radiation and outflows impacted upon it. Any complete understanding of the origin of these systems requires a unified treatment. In this way, the IGM resembles the interstellar medium of disk galaxies in which the gaseous component is intimately linked to the stars and their evolution and impact. Interpreting the increasingly refined observations requires detailed modeling, which relies on large-scale numerical computations involving gas, radiative processes, and gravity. The physics involved is intermediate in complexity between that of the CMB and galaxy formation, rendering the IGM a bridge between these extreme scales of cosmological structure formation. Unraveling the processes that led to the formation and structure of the IGM may thus serve as a crucial step in the solution of the much more involved problem of galaxy formation.  \\begin{figure} \\includegraphics[width=8cm, height=6.5cm]{meiksin_fig01.eps} \\caption{Number of refereed papers on the intergalactic medium and QSO intervening absorption systems published from 1965--2007, as provided by the ADS abstract service (adsabs.harvard.edu). The sharp rises follow key developments in astronomical instrumentation and observations. The lower histogram shows the rise of papers on the reionization and subsequent ionization state of the IGM. } \\label{fig:citations} \\end{figure}  A search of the literature for papers on the IGM since the Gunn \\& Peterson (1965) measurement produces close to 6000 references.\\footnote{Based on a boolean search of the Astrophysics Data System abstract service (adsabs.harvard.edu) for all refereed papers with abstracts or keywords containing ``(intergalactic and medium) or (quasar and absorption and line).''  Reionization papers are selected as the subset discussing the reionization of the IGM or the subsequent ionization structure of the IGM.} While this review does not have the space to describe the observational methods used to measure the IGM, it should be recognized that progress in the understanding of the properties of the IGM is indebted to advances in observational techniques. This is well illustrated by a plot of the number of refereed publications in the field against time. Periods of relatively steady output are punctuated by four leaps. The first occurs at the end of the 1960s and beginning of the 1970s with the introduction of image tube spectrographs coupled with integrating television systems for photon counting \\cite{1972ailt.conf..295B, 1972ailt.conf..317M}, which greatly facilitated the taking of spectra. The next occurs in the mid-1970s with the development of x-ray astronomy following the launch of the {\\it UHURU} satellite in 1970 and the recognition that galaxy clusters contain an extended and pervasive medium of hot, radiating gas \\cite{1973ApJ...185L..13K, 1973ApJ...184L.105L}. Another sharp rise occurs in the early 1990s following the launch of the {\\it Hubble Space Telescope} in 1990, the installation of EMMI and its echelle spectroscope on the New Technology Telescope (NTT) \\cite{1990Msngr..61...51D}, and the delivery of the HIRES spectrometer \\cite{1994SPIE.2198..362V} to the Keck telescope. The fourth occurred in 2000 with the introduction of the UV and Visible Echelle Spectrograph (UVES) to the Very Large Telescope (VLT) \\cite{2000SPIE.4005..121D}, the launch of the {\\it Far Ultraviolet Spectroscopic Explorer} ({\\it FUSE}) \\cite{2000ApJ...538L...1M}, and the beginning of operations of the Sloan Digital Sky Survey \\cite{2000AJ....120.1579Y}. The latter in particular triggered a surge of activity in reionization studies following the discovery of high redshift QSOs ($z>6$) with spectra indicating a rapid rise in the effective optical depth of the IGM to \\Lya photons, hinting that the Epoch of Reionization was being approached \\cite{2001AJ....122.2850B, 2002AJ....123.1247F}. This is indicated by a rise in IGM reionization papers in Fig.~\\ref{fig:citations}, a trend which continues today, fueled in part by the growing interest in the influence of reionization on the CMB fluctuations measured by the {\\it Wilkinson Microwave Anisotropy Probe} ({\\it WMAP}) \\cite{Kogut03}. The next major leap may well come from the anticipated radio measurements.  The rapidly rising tide of IGM studies has brought along with it several reviews. This is fortunate, as it is impossible in a single review to cover all areas thoroughly today. Early reviews of the QSO absorption line literature, now largely historical, were provided by \\textcite{1976ARA&A..14..307S} and \\textcite{1981ARA&A..19...41W}. \\textcite{1998ARA&A..36..267R} and \\textcite{Bechtold03} provide reviews of the \\Lya forest that are still largely up-to-date in the sense that most of the topics currently engaging the community are treated, with developments since mostly improvements in accuracy and in the details of the numerical models. Reviews of the low redshift IGM are provided by \\textcite{2003ASSL..281....1S} and \\textcite{2006pces.conf..111S}. \\textcite{2005ARA&A..43..861W} have reviewed the current understanding of Damped \\Lya Absorbers, an absorber class of special concern as it represents the closest link to the gaseous component forming present day galaxies. A new series of reviews followed the recent explosion of activity on the reionization of the IGM in anticipation of the detection of the Epoch of Reionization (EoR) through its high redshift 21cm signature. An early review in this area is by \\textcite{2001ARA&A..39...19L}. Since then, current observational and theoretical aspects of reionization have been exhaustively covered by \\textcite{2006ARA&A..44..415F}, \\textcite{2006PhR...433..181F}, and \\textcite{2007RPPh...70..627B}.  Rather than repeating the wide range of IGM phenomenology already covered by previous reviewers, I concentrate here on the physics underlying the structure of the IGM. One aim is to describe the physical underpinnings of current numerical simulations as required for future simulations to further progress. As observations are crucial for constraining any models, I begin by giving a broad overview of the current observational situation. The next section describes the physics of ionization equilibrium, followed by a discussion of the UV metagalactic background that maintains the ionization of the IGM. A brief review of early models of the \\Lya forest absorbers is then presented, followed by a discussion of current numerical models. The reionization of the IGM is then summarized along with means for its detection. This is followed by a discussion of the connection between galaxies and the IGM before concluding with observational and theoretical prospects for the future.  Unless stated otherwise, the cosmological parameters $\\Omega_m=0.3$ and $\\Omega_v=0.7$ are assumed for, respectively, the ratios of the present day matter density and the present day vacuum energy density to the current Einstein-deSitter density $\\rho_{\\rm crit}(0)=3H_0^2/8\\pi G$, a Hubble Constant of $H_0=100h\\,{\\rm km\\,s^{-1}\\,Mpc^{-1}}$ with $h=0.7$ (${\\rm Mpc} = {\\rm Megaparsec}\\simeq 3.0856\\times10^{22}\\,{\\rm m}$), and a baryon density of $\\Omega_bh^2=0.022$. These values are consistent within the errors with the current best estimates for a flat Universe based on CMB measurements using Year-5 {\\it WMAP} data of $\\Omega_m=0.279\\pm0.014$, $\\Omega_v=0.721\\pm0.015$, $h=0.701\\pm0.013$, and $\\Omega_bh^2=0.02265\\pm0.00059$ \\cite{2008arXiv0803.0547K}, or intergalactic $D/H$ measurements, giving $\\Omega_bh^2=0.021\\pm0.001$ \\cite{2006ApJ...649L..61O}. ",
        "conclusions": ""
    },
    "0711/0711.3499_arXiv.txt": {
        "abstract": "{Luminous quasars are known to display a sharp steepening of the continuum near 1100\\,\\AA. This spectral feature is not well fitted by current accretion disk models, unless comptonization of the disk emission is invoked. Absorption by carbon crystalline dust has been proposed to account for this feature.} {\\ton\\ ($\\zq=1.928$) exhibits the steepest far-UV decline ($\\Fnu \\propto \\nu^{-5.3}$) among the 183 quasar HST-FOS spectra analyzed by Telfer et\\,al. It is an ideal object to test the crystalline dust hypothesis as well as alternative interpretations of the UV break.}{We reconstruct the UV spectral energy distribution of \\ton\\ by combining HST, IUE and Palomar spectra.} {The far-UV continuum shows a  very deep continuum trough, which is bounded by a steep far-UV rise. We fit the trough assuming nanodiamond dust grains.}{Extinction by carbon crystalline dust reproduces the deep absorption trough of \\ton\\ reasonably well, but not the observed steep rise in the extreme UV. We also study the possibility of an intrinsic continuum rollover. The dust might be part of a high velocity outflow ($\\simeq 13000$\\,\\kms), which is observed in absorption in the lines of \\CIVw, \\OVIw, \\NVw\\ and \\Lya.} ",
        "introduction": "\\label{sec:intro}  The spectral energy distribution (\\sed) of  quasar and Seyfert\\,I galaxies is composed of emission lines superimposed on various continuum emission components. The near-infrared to visible domain is reasonably well reproduced by a powerlaw. The ultraviolet spectral region is characterized by a broad continuum excess, which is referred to as the big blue bump (BBB). According to general belief, it corresponds to the thermal signature from a hot accretion disk orbiting a supermassive black hole. The extension in the extreme UV of the BBB is expected to provide the ionizing flux that powers most of the emission lines. A serious problem with this picture, however, is that the BBB appears to be too soft\\footnote{There is an apparent inconsistency between the \\seds\\ typically observed and the much harder ones preferred in photoionization BELR models (e.g. Baldwin et\\,al. 1995; Casebeer, Leighly \\& Baron 2006).} to account for the high excitation lines from the broad emission line region (BELR) (Korista et\\,al. 1997). In effect, a marked continuum decline (i.e. steepening) takes place shortward of $\\simeq 1100$\\,\\AA\\ (rest-frame), which we hereafter refer to as the far-UV break\\footnote{We will also refer to the wavelength domain longward and shortward of the \\sed\\ steepening (or break) as the near- and far-UV regions, respectively.}. A possibility might be that this break is more akin to a localized continuum trough, followed by a marked recovery in the extreme UV, which is the energy region responsible for producing the high-excitation emission lines. Various mechanisms that could generate such a trough are summarized by Binette et\\,al. (2007).   Our aim is to probe the nature of the far-UV break and to test alternative interpretations of it by focussing on the more extreme cases. For instance, many quasar spectra reveal an \\sed\\ significantly steeper than the `average' \\sed\\ derived by Telfer et\\,al.( 2002; hereafter TZ02), which behaves as $\\nu^{-1.76}$ (in $F_{\\nu}$) shortward of the break. Among the 77 far-UV indices measured by TZ02, there were 3 objects with an ionizing continuum steeper than $\\nu^{-3}$. In this respect, \\ton\\ (alternatively named PG\\,1017+280 or J1019+2745) at redshift $\\zq=1.928$ is the most extreme case, with a powerlaw index as steep as $\\nu^{-5.3}$ (TZ02). For this reason, we consider it an ideal object to test competing models of the physical origin of the 1100\\,\\AA\\ break. In this paper, we used different bibliographical sources to build an \\sed\\ as complete as possible of \\ton. As the data are of limited quality, we intend to obtain higher quality data that would cover the X-ray domain down to the optical UV.  Our initial objective was to verify whether the hypothesis of absorption by crystalline carbon dust grains of Binette et\\,al. (2005; hereafter B05) would survive the test of modeling the extreme \\sed\\ of \\ton. Using two flavors of nanodiamonds, B05 could successfully reproduce the position and detailed shape of the far-UV break in 50 quasars, out of a total sample of 61 objects from HST-FOS archives\\footnote{The analysis of B05 included only  {\\it multigrating} HST spectra that extended down to at least 900\\,\\AA\\ (rest-frame), i.e. 61 quasars in total.  \\ton\\ was not included, since only a single grating HST spectrum exists.}. More recently, Haro-Corzo et\\,al. (2007, hereafter H07) reduced the dust model to a single flavor (that excludes the meteoritic type), since a particular emission feature expected near 3.5\\,$\\mu$m with the meteoritic case is absent from the mid-IR spectrum of 3C298 (de Diego et\\,al. 2007).  In this paper, we discuss the merits of the dust absorption model as well as its limitations in reproducing the deep trough that characterize the \\ton\\ \\sed. In Appendix\\,\\ref{app:a}, we compare different extinction curves and illustrate how the far-UV extinction is affected. In a companion Paper (Binette \\& Krongold 2007), we report on some of the peculiarities of the emission line spectrum of this unusual quasar.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{binette_fig1.eps}} \\caption{Spectral segments of \\ton\\ extracted from different archival or bibliographical sources. Top panel: blue and red optical spectra (in arbitrary units) from Sargent, Boksenberg \\& Steidel (1988); lower right panel: HST-FOS spectrum using grating G270H, lower left panel: IUE spectra (LWP and SWP) (see Sect.\\,\\ref{sec:spc}). \\label{fig:spc}} \\end{figure} ",
        "conclusions": "\\label{sec:disc}  The \\sed\\ observed in  \\ton\\ is rather unique, due to a striking lack of soft ionizing photons. The far-UV break is definitely more extreme than in the other quasars of the TZ02 sample (B05). There is no consensus on the physical origin of the quasar far-UV break. Binette et\\,al. (2007) have recently presented a review of alternative interpretations, some of which, however, lack detailed calculations that prevent making a direct comparison with observations. Among the promising interpretations lies the possibility of continuum reprocessing by a  wind arising from the accretion disk or absorption from gas progressively accelerated up to quasi relativistic velocities (Eastman, MacAlpine \\& Richstone 1983). In this paper, we explore the dust absorption interpretation, for which we dispose of a detailed model, but it is certainly premature to exclude  other interpretations at this stage.  The nanodiamond dust absorption hypothesis is consistent with the \\ton\\ \\sed, since it produces a broad absorption feature that is compatible with the position, depth and shape of the observed trough. However, the far UV-rise takes place at slightly shorter wavelengths than predicted by the simplest model (Model\\,\\I). If not due to limitations in the S/N or systematics in the instrumental system of IUE, this can be accounted for by assuming a roll-over near 670\\,\\AA\\ or photoelectric absorption by \\HeI. Alternatively, the shortcoming might result from the limitations of the dust grain model. Improved quality and higher resolution spectra will be essential to determine the true continuum level at the bottom of the trough and to confirm whether  a very steep rise occurs in the extreme UV.  Reproducing the observed  continuum trough, in any case, is not sufficient to vindicate the dust model. For instance, we need to find evidence in the mid-IR of re-radiation  of the energy absorbed by the dust. Using Spitzer data on 3C298, de\\,Diego et\\,al. (2007) ruled out the presence of emission bands at 3.43 and 3.53\\mic\\ as is expected in the case of meteoritic nanodiamonds [i.e. nanodiamonds with surface adsorbates that produce C$-$H stretch emission (Andersen et\\,al. 1998; Jones \\& d'Hendecourt 2000)]. The model used in this paper is not affected, since it is based on cubic diamonds, which may emit either via the one-phonon mode due to internal bulk impurities within the crystals (Andersen 1999)  or, if made of pure crystals, via multi-phonon modes, which are unfortunately quite inefficient (Braatz et\\,al. 2000; Edwards 1985; Jones \\& d'Hendecourt 2000). However, testing these two remaining possibilities still remains a task to be done.  The absence of strong bound-free absorption of \\HI\\ shortward of the Lyman limit implies a high ionization parameter ($\\up \\ga 0.05$) for the dust screen, if we assume photoionization  (or a temperature above 63\\,000\\,K for the thermally ionized case). In such an environment, crystallite dust would not lie in thermal equilibrium with the ambient gas and therefore must continually be replenished at a rate near that at which it is being destroyed. A dust screen within a thin funnel-shaped wind structure as contemplated for the ionized gas by Elvis (2000) and as inferred from the warm absorbers in NGC\\,4051 (Krongold et\\,al. 2007) might resolve this problem. Two conditions would have to be met, however: the wind must be launched from a cool region of the accretion disk where the dust can be formed (see Konigl \\& Kartje 1994; Everett 2005) and the wind must cross our line-of-sight to the UV emission region of the disk.  The possibility of such a wind is particularly relevant to the problem of detecting the expected absorption lines (such as \\CIV\\ or \\OVI)  associated with the dust. We do find such lines in an absorption system that is blueshifted by 13200\\,\\kms. To validate this possibility, it is necessary to confirm the presence of similar absorption systems in other quasars that possess a prominent far-UV break.  To conclude, despite the success of the  crystalline dust model in reproducing reasonably well the  far-UV trough observed in \\ton\\ and in other quasars (BI05), there arise various objections to such models that we have partially addressed in this paper. Further studies are needed to reach firm conclusions in favor of or against the possibility of crystalline dust as an explanation of the far-UV break in quasars."
    },
    "0711/0711.0515_arXiv.txt": {
        "abstract": "We present a spatially resolved spectroscopic study of the thermal composite supernova remnant \\snr\\ with \\Chandra. The X-ray spectrum of \\snr\\ is characterized by K~lines from Mg, Si, S, Ar, and Ca. The X-ray--emitting gas is found to be enriched in sulfur and calcium. The broadband and tricolor images show two incomplete shell-like features in the northeastern half and brightness fading with increasing radius to the southwest. There are over 30 unresolved sources within the remnant.  None shows characteristics typical of a young neutron star. The maximum diffuse X-ray intensity coincides with a radio-bright region along the eastern border. In general, gas in the inner region is at higher temperature, and the emission is brighter, than that in the outer region. The gas in the remnant appears to be near ionization equilibrium. The overall morphology can be explained by the evolution of the remnant in an ambient medium with a density enhancement from west to east. We suggest that the remnant was born in a preexisting cavity and that the bright inner emission is due to the reflection of the initial shock from the dense cavity wall. This scenario may provide a new candidate mechanism to explain the X-ray morphology of other thermal composite supernova remnants. ",
        "introduction": "\\label{sec:intro} Massive stars evolve rapidly and can explode not far from the dense clouds that were their birthplaces.  The resulting supernova remnants (SNRs) thus are expected to be in the vicinity of, and to interact with, molecular and/or \\HI\\ clouds. Such interactions may account for the characteristics of the thermal composite (or mixed-morphology) SNRs, which radiate bright thermal X-ray emission interior to their radio shells and have faint X-ray rims (Green \\etal\\ 1997; Rho \\& Petre 1998; Yusef-Zadeh et al.\\ 2003). Several mechanisms have been proposed to interpret the X-ray morphology of the thermal composites. These include (1) a radiatively cooled rim and a hot interior (e.g., Harrus et al. 1997; Rho \\& Petre 1998), (2) a hot interior with density increased by thermal conduction (Cox et al.\\ 1999; Shelton et al.\\ 1999), (3) an increase in internal gas density due to evaporation of the engulfed cloudlets (White \\& Long 1991), and (4) shock interaction at the edge of a cloud but seen brightened in the projected interior (Petruk 2001; as summarized in, e.g., Chen et al.\\ 2004). Recently, Shelton et al. (2004) have explained the X-ray properties of SNR~W44 by invoking thermal conduction and bright metal emission due to dust destruction and ejecta enrichment in the interior.  Also, a study based on ionization states of hot interior plasma suggests that the thermal composite (or mixed-morphology) phenomenon may be an evolutionary state of SNRs due to thermal conduction (Kawasaki et al.\\ 2005). The sample comprised six thermal composites including \\snr, which is the subject of this paper.  In none of these mechanisms has the influence of the supernova progenitor on the surrounding medium been considered. Massive stars may sculpt a cavity with their energetic stellar winds and ionizing radiation before they explode as core-collapse supernovae. The observational effect of the SNR's blast wave impacting the cavity wall should be of particular interest. After almost free expansion in the cavity, the blast wave will ``reflect'' from the cavity wall and a strong reverse shock will heat the interior material. Thus, a particularly X-ray bright interior may be expected in this scenario, as we will outline for the case of \\snr.  Kesteven~27 (G327.4$+$0.4) is one of the archetypical thermal composite SNRs (Rho \\& Petre 1998; Enoguchi et al.\\ 2002). As seen in X-rays with \\Einstein\\ (Seward 1990), \\ROSAT\\ (Seward et al.\\ 1996), and \\ASCA\\ (Enoguchi et al.\\ 2002), it appears prominently brightened in the center, especially along an inner, broken ring. This is different from most other thermal composites, which are centrally-peaked in X-rays. In the radio band (Kesteven \\& Caswell 1987; Whiteoak \\& Green 1996) there is a clear-cut rim, except in the southwest. The radio emission is brightest in the east, with the surface brightness peaking near the southeastern border, where the blast wave may be striking a dense cloud. The radio emission gradually fades toward the southwest, where the remnant seems to break out to a lower density region. It was noted that the radio image displays multiple arcs centered on the eastern half, from which a series of arclike shells may have emanated (Milne et al.\\ 1989).  Although OH masers, which usually mark shock interaction with molecular clouds, are not detected in \\snr\\ (unlike other thermal composites such as 3C~391, W28, and W44) (Green et al.\\ 1997), this SNR is found to be embedded in an \\HI\\ cloud complex at a local standard of rest velocity of $70\\km\\ps$. There is an \\HI\\ ridge or shell just exterior to the rim, delineated by the radio continuum contours, except on the west and southwest sides (McClure-Griffiths et al.\\ 2001). McClure-Griffiths et al.\\ suggest that this confirms that the remnant shock is impacting a density enhancement from the inner void. The \\HI\\ observation provides a dynamical distance estimate of $\\sim4.3\\kpc$, on the far side of the Scutum-Crux arm. In this paper, we take this value as a reference for the distance: $d=4.3\\du\\kpc$.  We present a spatially resolved X-ray spectroscopic study of the remnant using a \\Chandra\\ ACIS-I observation. In \\S~2, we briefly describe the observation and data calibration and present our analysis and results. The physical properties of the thermal X-ray emission are discussed in \\S~3. We summarize our results in \\S~4. Statistical errors are all presented at the 90\\% confidence level.  \\section {Observations and Data Analysis}  \\snr\\ was observed with the Advanced CCD Imaging Spectrometer (ACIS) on board the \\Chandra\\ X-Ray Observatory on 2003 June 21 (ObsID 3852) for 39 ks. The target center ($\\RA{15}{48}{31.5}$, $\\decl{-53}{46}{20}$; J2000) was placed for optimal coverage of the X-ray--bright region of the remnant (see Fig.~\\ref{f:bband}) on the four ACIS-I CCD chips. We reprocessed the event files (from level 1 to level 2) using the \\Chandra\\ Interactive Analysis of Observations ({\\small CIAO}) data processing software (ver.\\ 3.2.2) \\footnote{http://cxc.harvard.edu/ciao} to remove pixel randomization and to correct for CCD charge transfer inefficiencies. After removing flares with count rates greater than 1.2 times the mean light-curve value, a net exposure of 37 ks remained and was used for analysis.  \\begin{figure}[tbh!] \\centerline{ {\\hfil\\hfil \\psfig{figure=f1.eps,height=3.in,angle=270, clip=} \\hfil\\hfil}} \\caption{ Broadband (0.3--7 keV) diffuse X-ray image (on a square root scale between $64$ and $480$ photons$\\cm^{-2}\\ps\\,\\mbox{sr}^{-1}$) overlaid with MOST 843~MHz radio contours (at square root scale levels 0.01, 0.02, 0.05, 0.10, 0.17, 0.25, and 0.36 ${\\rm~Jy~beam^{-1}}$; from Whiteoak \\& Green 1996). The X-ray image is exposure-corrected and was adaptively smoothed to achieve a S/N of 4 (using the CIAO program {\\em csmooth}). The dashed box denotes the field of view of ACIS-I. } \\label{f:bband} \\end{figure}  \\subsection{Source Identification and Optical Counterparts}  A major goal of this observation was to search for an internal compact object.  There is obviously no bright Crab-like pulsar or pulsar wind nebula (PWN) within Kes 27.  However, less energetic PSR/PWN examples have been observed in many other remnants, for example, IC~443 (1.5 kpc distant) and the Vela SNR (0.25 kpc distant) (Olbert et al. 2001; Helfand et al. 2001).  Upper limits can be set for these types of PWNe.  If moved to the more distant Kes 27, the IC 443 PWN ($L_x = 3 \\times 10^{33}$ ergs s$^{-1}$, $2^{\\prime}\\times 1^{\\prime}$ in extent) would appear as a $0.7^{\\prime} \\times 0.35^{\\prime}$ patch in the 2--8 keV X-ray band, with strength $\\approx 20$ counts ks$^{-1}$.  We can set an upper limit of 1 count ks$^{-1}$ for an object of this size in the interior of Kes 27, so $L_x \\leq 1.5 \\times 10^{32}$ ergs s$^{-1}$.  Similar reasoning for the Vela PWN ($1.6^{\\prime}$ in diameter, $L_x = 7 \\times 10^{32}\\ergs\\ps$) predicts 7 counts ks$^{-1}$ from a $6^{\\prime\\prime}$ region.  In this case we can set an upper limit of 0.5 counts ks$^{-1}$, corresponding to $L_x \\leq 0.5 \\times 10^{32}$ ergs s$^{-1}$.  Note that an upper limit depends strongly on the size of the object. Leahy (2004) has suggested that the PSR/PWN within the IC 443 boundary is really associated with the nearer remnant G189.6+3.3.  Although this would reduce its luminosity to about one-quarter of that quoted above, our upper limit does not change.  The thermal composite remnant Kes 79 (7 kpc distant) contains an unresolved central compact object (CCO) with $L_x = 7 \\times 10^{33}$ ergs s$^{-1}$ (Seward et al.\\ 2003). The spectrum is soft, but since absorption is comparable in the two remnants, if this object were in Kes 27, we would expect 40 counts ks$^{-1}$, twice as bright as the brightest observed unresolved source.  The hardness ratio (defined in the notes to Table~\\ref{T:sources}) should be $\\approx 0.5$, and no variation or optical counterpart is expected.  As described below, we searched for such an object and did not find one.  An upper limit of 0.8 counts ks$^{-1}$ yields $L_x \\leq 1.4 \\times 10^{32}$ ergs s$^{-1}$ for any CCO in Kes 27.  \\begin{center} \\begin{deluxetable*}{ccrcccccccc} \\tabletypesize{\\footnotesize} \\tablecaption{{\\sl Chandra} Source List \\label{acis_source_list}} \\tablewidth{0pt} \\tablehead{ \\colhead{Source} & \\colhead{CXOU Name} & \\colhead{Count rate $({\\rm~cts~ks}^{-1})$} & \\colhead{HR} & \\colhead{$J$} & \\colhead{$H$} & \\colhead{$K$} & \\colhead{$B$} & \\colhead{$R$} & \\colhead{$I$} & \\colhead{Comments}\\\\ \\colhead{(1)} & \\colhead{(2)} & \\colhead{(3)} & \\colhead{(4)} & \\colhead{(5)} & \\colhead{(6)} & \\colhead{(7)} & \\colhead{(8)} & \\colhead{(9)} & \\colhead{(10)} & \\colhead{(11)} } \\startdata  1 & J154722.45--534552.6 & 0.81$\\pm$0.18 &  0.80$\\pm$0.18 & 14.07 & 13.54 & 13.45 & 17.42 & 16.05 & 15.27 & ? \\\\ % 2 & J154737.93--534639.6 & 2.55$\\pm$0.29 & -0.10$\\pm$0.12 & \\nodata & \\nodata & \\nodata & \\nodata & \\nodata & \\nodata & star?, varies \\\\ % 3 & J154802.08--534549.3 & 1.53$\\pm$0.23 & -0.19$\\pm$0.16 & 12.16 & 11.44 & 11.17 & 16.75 & 14.77 & 14.02 & star \\\\ % 4 & J154805.52--534050.5 & 0.75$\\pm$0.17 &  -0.64$\\pm$0.21 & 11.50 & 11.23 & 11.17 & \\nodata & \\nodata & \\nodata & star \\\\ % 5 & J154805.67--534018.7 & 0.67$\\pm$0.16 &  -0.44$\\pm$0.25 & 13.23 & 12.77 & 12.71 & 16.01 & 14.63 & 14.28 & star \\\\ % 6 & J154808.11--534128.1 & 1.07$\\pm$0.20 & -0.70$\\pm$0.16 & 10.84 & 10.67 & 10.61 & 11.69 & 11.59 & 11.56 & star \\\\ % 7 & J154816.79--534125.5 & 7.03$\\pm$0.46 &  0.95$\\pm$0.03 & \\nodata & \\nodata & \\nodata & \\nodata & \\nodata & \\nodata & AGN \\\\ % 8 & J154817.90--535654.4 & 0.86$\\pm$0.18 &  -0.44$\\pm$0.21 & 13.30 & 11.66 & 10.81 & 15.58 & 14.26 & 13.99 & star\\\\ % 9 & J154819.85--534618.3 & 2.17$\\pm$0.27 &  0.28$\\pm$0.13 & 10.83 & 9.74 & 9.31 & 19.49 & 15.63 & 13.95 &  star \\\\ % 10 & J154820.19--535447.9 & 1.05$\\pm$0.20 & -0.54$\\pm$0.18 & 14.86 & 14.11 & 13.95 & 18.08 & 16.43 & 15.37 & star \\\\ % 11 & J154820.78--534222.4 & 1.07$\\pm$0.20 &  0.15$\\pm$0.20 & 14.01 & 13.04 & 12.65 & 20.61 & 18.54 & 17.11 & star \\\\ % 12 & J154822.10--535548.0 & 0.83$\\pm$0.18 &  -0.29$\\pm$0.22 & 12.88 & 12.34 & 12.21 & 16.02 & 14.64 & 14.05 & star \\\\ % 13 & J154826.38--534940.0 & 2.63$\\pm$0.29 & -0.37$\\pm$0.11 & 13.42 & 12.70 & 12.47 & 17.56 & 14.87 & 14.28 & star \\\\ % 14 & J154829.60--534627.1 & 3.57$\\pm$0.34 & -0.53$\\pm$0.09 & 13.24 & 12.66 & 12.44 & \\nodata & \\nodata & \\nodata & star, flare x10 \\\\ % 15 & J154832.21--533736.3 & 1.80$\\pm$0.25 &  0.64$\\pm$0.12 & \\nodata & \\nodata & \\nodata & \\nodata & \\nodata & \\nodata & AGN? \\\\ % 16 & J154832.24--534028.0 & 0.97$\\pm$0.19 &  0.89$\\pm$0.14 & \\nodata & \\nodata & \\nodata & \\nodata & \\nodata & \\nodata & AGN? \\\\ % 17 & J154832.37--533916.8 &57.55$\\pm$1.27 & -0.55$\\pm$0.02 & 10.24 & 9.86 & 9.75 & 12.20 & 11.28 & 10.90 & star, varies x2  \\\\ % 18 & J154832.96--535249.7 & 0.78$\\pm$0.17 &  -0.52$\\pm$0.22 & \\nodata & \\nodata & \\nodata & 20.02 & 18.70 & 16.68 & star, varies x3 \\\\ % 19 & J154837.83--534955.7 & 2.47$\\pm$0.29 & -0.52$\\pm$0.11 & 10.59 & 9.93 & 9.67 & 14.18 & 12.77 & 11.36 & star \\\\ % 20 & J154839.05--534316.1 & 0.72$\\pm$0.17 &  -0.19$\\pm$0.25 & 10.77 & 9.99 & 9.66 & 17.01 & 14.36 & 12.92 & star \\\\ % 21 & J154839.32--534527.1 & 0.97$\\pm$0.19 &  -0.28$\\pm$0.20 & 11.57 & 11.30 & 11.20 & 13.43 & 12.41 & 11.09 & star \\\\ % 22 & J154840.21--534043.6 &13.20$\\pm$0.62 & -0.71$\\pm$0.0 & 11.50 & 10.87 & 10.64 & \\nodata & \\nodata & \\nodata & star, varies x3  \\\\ % 23 & J154840.68--534635.9 & 0.67$\\pm$0.16 &  -0.44$\\pm$0.25 & 12.52 & 12.16 & 12.08 & 14.57 & 14.06 & 13.69 & star \\\\ % 24 & J154841.14--534005.2 & 0.78$\\pm$0.17 &  -0.03$\\pm$0.24 & 13.74 & 13.22 & 13.01 & 17.27 & 15.55 & 14.91 & star \\\\ % 25 & J154842.55--535610.1 & 0.83$\\pm$0.18 &  0.81$\\pm$0.18 & 16.14 & 15.13 & 14.56 & \\nodata & \\nodata & \\nodata & ? \\\\ % 26 & J154842.77--535213.4 & 0.83$\\pm$0.18 &  -0.36$\\pm$0.22 & 13.71 & 13.27 & 13.16 & 16.88 & 15.32 & 14.96 & star\\\\ % 27 & J154853.11--534806.2 & 1.50$\\pm$0.23 & -0.43$\\pm$0.15 & 13.76 & 13.37 & 13.30 & 16.79 & 15.31 & 14.82 & star\\\\ % 28 & J154915.48--534407.8 & 1.34$\\pm$0.22 &  0.72$\\pm$0.14 & 15.91 & 14.97 & 14.34 & \\nodata & \\nodata & \\nodata & ?\\\\ % 29 & J154916.18--534244.9 & 1.34$\\pm$0.22 &  0.80$\\pm$0.13 & \\nodata & \\nodata & \\nodata & \\nodata & \\nodata & \\nodata & AGN? \\\\ % 30 & J154932.56--535026.4 & 0.59$\\pm$0.16 &  0.00$\\pm$0.29 & 11.55 & 10.80 & 10.49 & 16.40 & 14.86 & 13.89 & star \\\\ % \\enddata \\tablecomments{ Col.~(1): Generic source number. Col.~(2): {\\sl Chandra} X-Ray Observatory (unregistered) source name, following the {\\sl Chandra} naming convention and the IAU Recommendation for Nomenclature (e.g., http://cdsweb.u-strasbg.fr/iau-spec.html). Col.~(3): On-axis source broad-band count rate. Col.~(4): The hardness ratio defined as ${\\rm HR}=({\\cal H}-{\\cal S})/({\\cal H}+{\\cal S})$, where ${\\cal S}$ and ${\\cal H}$ are the source count rates in the 0.3--1.5 and 1.5--7~keV bands, respectively. This table only lists sources with individual signal-to-noise ratios greater than 4 in the broad band (${\\cal S}+{\\cal H}$), with an exception made for source 30, which has a S/N of 3.8 but is obviously discernible in the tricolor image (Fig.3). Cols.~(5)--(7): $JHK$ magnitudes from the 2MASS catalog (Skrutskie et al.\\ 2006). Cols.~(8)--(10): $BRI$ magnitudes from the USNO-B1.0 catalog (Monet et al.\\ 2003). } \\label{T:sources} \\end{deluxetable*} \\end{center}  We first searched for pointlike sources in the broad band (0.3--7.0 keV) with a wavelet source detection algorithm. Thirty sources with a signal-to-noise ratio S/N$\\gsim 4$ in the field of view (the range of the four ACIS-I chips) are listed in Table~\\ref{T:sources} and identified in Figure~\\ref{f:sources}. The spectrum of each source is characterized by a hardness ratio. This is useful in identifying red stars and active galactic nuclei (AGNs). The count rates listed were taken from \\Chandra\\ automatic processing. Optical data were taken from the Two Micron All Sky Survey (2MASS) catalog (Skrutskie et al.\\ 2006; \\footnote{See http://www.ipac.caltech.edu/2mass.} $JHK$) and the USNO-B1.0 catalog (Monet et al.\\ 2003; \\footnote{Available at http://www.nofs.navy.mil/data/fchpix.} $BRI$). If no cataloged source was within $1''$ of the X-ray position, no optical data are listed.  Most Table~\\ref{T:sources} sources have soft X-ray spectra and optical counterparts with $0.0 < J-H < 0.8$ and $0.0 < H-K < 0.4$. These are identified as stars and are probably of spectral type K or M (Tokunaga 2000).  The X-ray spectra of these clearly show less absorption than that of Kes 27, so these are certainly foreground objects --- as well they should be. Sources 9 and 11, with harder spectra, are probably highly reddened stars. Even though not many photons were collected, the X-ray flux from some sources is quite variable. The approximate variation is listed.  \\begin{figure}[tbh!] \\centerline{ {\\hfil\\hfil \\psfig{figure=f2.ps,height=3.5in,angle=0, clip=} \\hfil\\hfil}} \\caption{Identification of the unresolved sources listed in Table~\\ref{T:sources}, showing the adaptively smoothed 0.3--7~keV emission in logarithmic gray scale. } \\label{f:sources} \\end{figure}  The bright source No.7 is quite hard, is strongly absorbed, and has no counterpart.  This is probably a background AGN.  Sources 15, 16, and 29 are also hard with no counterparts and are likely to be AGNs. The question marks in Table~1 for these sources indicate classification based on hardness ratio.  There were not enough events to determine a spectral shape.  A typical AGN with strength of 1 count ks$^{-1}$ and with extinction corresponding to at least $2 \\times 10^{22}\\cm^{-2}$, the column density along the path to Kes~27, would have a hardness ratio $\\geq 0.7$ with $B \\sim 27$, $R \\sim 21$, and  $H \\sim 13$~mag (Seward 2000; NASA/IPAC Extragalactic Database \\footnote{See http://nedwww.ipac.caltech.edu.}). There might be a counterpart in the infrared ($JHK$), but not in the optical bands ($BR$). Sources 1, 25, and 28 are hard, with apparent counterparts.  Since these three are far from the focal point, the X-ray location is not as accurate as for the more central sources, and there is a higher probability that the counterparts listed are accidental. Sources 25 and 28, with infrared counterparts only, could be AGNs. Sources 2 and 18 are soft, with no 2MASS counterpart.  They are classified as stars because they were observed to vary.  The brightest unresolved source, No.17, produced enough counts to derive spectral parameters.  The spectrum can best be fitted with a thermal gas plus power law model (see Fig.~\\ref{f:nspot1}). This source is weakly absorbed [$\\NH\\sim(6\\pm4)\\E{20}\\cm^{-2}\\ps$], with power-law photon index $\\Gamma\\sim2.9\\pm0.2$, gas temperature $kT\\sim0.77^{+0.07}_{-0.13}\\keV$, probably overabundant Ne ([Ne/H]$\\sim3.4^{+3.0}_{-2.3}$), and an unabsorbed 0.5--10~keV flux $F_x\\sim5.9\\E{-13}\\ergs\\cm^{-2}\\ps$ ($\\chi^2=63.3$ with 73 degrees of freedom [dof]). The spectrum of the second brightest unresolved source, No.22, can be fitted with a double thermal gas model [$\\NH\\sim(6\\pm1)\\E{21}\\cm^{-2}$, $kT_1\\sim0.2^{+0.05}_{-0.02}\\keV$, $kT_2\\sim1.3^{+1.4}_{-0.6}\\keV$, $F_x\\sim1.7\\E{-12}\\ergs\\cm^{-2}\\ps$, $\\chi^2/{\\rm dof}=14.3/17$].  These spectra and column densities are quite consistent with the identification as foreground stars.  \\begin{figure}[tbh!] \\centerline{ {\\hfil\\hfil \\psfig{figure=f3.ps,height=2.3in,angle=270, clip=} \\hfil\\hfil}} \\caption{Spectrum of the northern hot spot (source 17), fitted with a {\\em vmekal}+{\\em powerlaw} model. The source spectrum was extracted from within a 3~$\\sigma$ region, and an area of twice the 3~$\\sigma$ size surrounding the source region was used as the background. The net count rate is $(5.63\\pm0.12)\\E{-2}\\ps$. The spectrum has been adaptively binned to achieve a background-subtracted S/N of 4. } \\label{f:nspot1} \\end{figure}  Photons were sparse for most Table 1 sources. A few sources were not easy to classify. Source 2 is soft and has no optical counterpart, but it does not exhibit enough interstellar absorption to be at the distance of the remnant. Source 15 has no counterpart, and the hardness ratio is close to the 0.5 expected from a CCO.  Since it is at the limb of the remnant, an AGN identification is more likely.  Two interior sources in the previous \\ROSAT\\ image, Nos.\\ 2 and 4, were identified as possible PSR/PWN candidates (Seward et al.\\ 1996). These correspond respectively to \\Chandra\\ sources 14 and 19, which are identified here as foreground stars.  \\subsection{Spatial Analysis} \\label{sec:spatial}  Figure~\\ref{f:bband} presents the \\Chandra\\ image of diffuse emission in the $0.3--7.0\\keV$ broad band. Point sources were subtracted, and the pixel values in source regions were replaced with values interpolated from the surrounding area. The X-ray image was then exposure-corrected and adaptively smoothed to achieve a S/N of 4 (using the CIAO program {\\em csmooth}). Contours of the Molonglo Observatory Synthesis Telescope (MOST) 843 MHz radio emission observation of the remnant (Whiteoak \\& Green 1996) are superposed on the X-ray image.  To demonstrate the energy dependence of the SNR morphology and the point sources, we also created a tricolor image in three energy bands, as shown in Figure~\\ref{f:3color}: 0.3--1.5 keV ({\\em red}), 1.5--2.2 keV ({\\em green}), and 2.2--7.0 keV ({\\em blue}). These three bands respectively contain the Mg, Si, and S lines, which dominate the remnant's thermal emission (see \\S\\ref{sec:spec}), and have comparable photon fluxes. We first produced exposure maps in the three bands and used these for flat-fielding, accounting for bad-pixel removal, correcting for telescope vignetting, and correcting for variations in quantum efficiency across the detector. The images in the three bands were then smoothed with an adaptive filter (using program {\\em csmooth}) to achieve a broadband (0.3--7 keV) S/N of 4.  \\begin{figure}[tbh!] \\centerline{ {\\hfil\\hfil \\psfig{figure=f4.eps,height=3.in,angle=0, clip=} \\hfil\\hfil}} \\caption{ Tricolor ACIS-I image of \\snr.  The X-ray intensities in the 0.3--1.5, 1.5--2.2, and 2.2--7.0 keV bands are coded red, green, and blue, respectively, and are scaled logarithmically in the ranges 19--170, 19--290, and 27--270 photons$\\cm^{-2}\\ps\\,\\mbox{sr}^{-1}$. The exposure-corrected X-ray maps were adaptively smoothed to achieve a broadband (0.3--7 keV) S/N of 4 (using the CIAO program {\\em csmooth}). } \\label{f:3color} \\end{figure}  Figures~\\ref{f:bband} and \\ref{f:3color} show that the X-ray--emitting gas occupies a volume of linear size $\\ge17'$ (the size of the ACIS-I CCD chips) and that the centroid of the X-ray emission lies in the eastern half of the remnant. The bright X-ray interior previously seen is now resolved into two semicircular arcs in the northeastern half of the remnant. The emission from the inner arc is harder than that from the outer arc (see Fig.~\\ref{f:3color}). Apart from this bright, hard arc, there is an incomplete, soft arc along the northeastern outskirt. This outer arc is interior to the radio border. The X-ray surface brightness peaks near the southeastern border, where there is a bright X-ray knot (region E1 in Fig.~\\ref{f:reg}) just inside the peak radio emission. This phenomenon is very similar to that seen on the western rim of SNR~3C~391, where the blast wave has been suggested to be propagating into a small, dense region, causing drastic shock deceleration or magnetic field compression and amplification (Chen et al.\\ 2004). The soft, bright, bar-like part of the outer arc in the north appears coincident with a radio arc (Fig.~\\ref{f:bband}). The surface brightness radial profiles are plotted in Figure~\\ref{f:rp}. The two peaks in the northeast profile correspond to the two X-ray arcs, with mean radii of $2.5'$ and $6'$. The southwest profile shows that the X-ray emission fades with increasing radius, and no shell-like structure can be discerned.  \\begin{figure}[tbh!] \\centerline{ {\\hfil\\hfil \\psfig{figure=f5.eps,height=2.5in,angle=0, clip=} \\hfil\\hfil}} \\caption{ Raw image of the ACIS observation (with point-like sources removed). The labeled regions are used for spectrum extraction. } \\label{f:reg} \\end{figure}  \\begin{figure}[tbh!] \\centerline{ {\\hfil\\hfil \\psfig{figure=f6.eps,height=2.7in,angle=270, clip=} \\hfil\\hfil}} \\caption{ Exposure-corrected radial X-ray (0.3--7 keV) surface profiles for the northeast and southwest parts, centered at R.A.=$\\RA{15}{48}{37.9}$, decl.=$\\decl{-53}{48}{01}.9$. The NE part is from position angle $-46.8^{\\circ}$ to $133.2^{\\circ}$ and the SW part is from $133.2^{\\circ}$ to $313.2^{\\circ}$. A blank-sky exposure on the ACIS-I chips has been used for background. The dotted lines are plotted according to the White \\& Long (1991) model, relatively scaled to match the maximum brightness, with parameters $(\\tau, C)=(10,50)$ for the left panel and $(\\tau, C)=(20,56)$ for the right (see text). } \\label{f:rp} \\end{figure}  \\subsection{Spectral Analysis of the Diffuse Emission}\\label{sec:spec}  After removal of point sources, we extracted the overall spectrum of the diffuse emission from the whole area of the four ACIS-I chips (excluding the peripheral edge) and other on-SNR spectra from 10 substructures marked in the raw-counts image shown in Figure~\\ref{f:reg}. Because the field of view of the four ACIS-I chips is almost entirely covered by the remnant, a local background cannot be obtained from these chips, and hence a double background subtraction method was used. We extracted an off-SNR source-free spectrum from the ACIS-S2 chip (excluding the peripheral edge again; Fig.~\\ref{f:reg}). The respective blank-sky background \\footnote{See http://cxc.harvard.edu/contrib/maxim/acisbg, period D.} contributions estimated from the same regions and scaled by the flux at 9.5--12~keV \\footnote{A high-energy band free of sky emission; see also http://cxc.harvard.edu/contrib/maxim/acisbg.} were subtracted from the on- and off-SNR spectra. Individual on-SNR spectra were adaptively binned to achieve a background-subtracted signal-to-noise ratio of 4 and the off-SNR spectrum was binned using a ratio of 3. Each on-SNR spectrum was jointly fitted together with the off-SNR spectrum (Fig.~\\ref{f:spectra}). The on-SNR sky background was determined by scaling the off-SNR emission according to the region sizes (Fig.~\\ref{f:spectra}, {\\em green lines}) and is phenomenologically well described by a power law with photon index $\\sim0.7$. Because of the poor signal-to-noise ratio of the off-SNR spectrum below 1.2~keV, the counts below 1.2~keV in all the on-SNR spectra were ignored so as to match the off-SNR one. The {\\small XSPEC} spectral fitting package (ver.\\ 11.3) \\footnote{See http://heasarc.gsfc.nasa.gov/docs/software/lheasoft/xanadu/ xspec/xspec11/index.html.} was used throughout.  For the foreground absorption, the cross sections from Morrison \\& McCammon (1983) were used, and solar abundances were assumed.  \\begin{figure*}[tbh!] \\centerline{ {\\hfil\\hfil \\psfig{figure=f7.eps,height=9in,angle=270, clip=} \\hfil\\hfil}} \\caption{ ACIS spectra of the entire remnant and the small regions shown in Fig.~\\ref{f:reg}. The on-SNR spectra (black solid lines) were rebinned to achieve a background-subtracted S/N of 4, and the off-SNR one (red) was rebinned using a ratio of 3. The net on-SNR spectra (blue) are fitted with an absorbed {\\em vpshock} model, except for the spectrum of the entire remnant, for which an absorbed {\\em vsedov} model is used. The sky backgrounds (green) are scaled from the off-SNR spectrum, which is mimicked by a {\\em power-law} model. The dot-dashed lines around 3.12 keV for the spectral fits of the entire remnant and region C1 are Gaussians accounting for the Ar~XVII line. } \\label{f:spectra} \\end{figure*}  \\subsubsection{Overall Spectral Properties} \\label{sec:glo} The spectrum of the entire remnant (Fig.~\\ref{f:spectra}, top left) shows distinct K$\\alpha$ emission line features of metal elements: Mg~XI ($\\sim1.34\\keV$), Si~XIII ($\\sim1.84\\keV$), S~XV ($\\sim2.43\\keV$), Ar~XVII ($\\sim3.12\\keV$), and Ca~XIX ($\\sim3.86\\keV$), confirming the thermal origin of the emission. We fitted the net on-SNR spectrum with two absorbed non-equilibrium ionization (NEI) thermal plasma models, {\\em vpshock} and {\\em vsedov} (Borkowski et al.\\ 2001; see Table~\\ref{T:ent_spec}). The former model characterizes the plasma parameters of a plane-parallel shock, typified by a constant electron temperature and the shock ionization timescale (the product of the electron density and the time since passage of the shock). The latter is based on Sedov (1959) dynamics, typified by the mean and the electron temperatures immediately behind the shock and by the ionization timescale (the product of the electron density immediately behind the shock and the remnant's age). Because the NEI models do not include emission from argon, we added a Gaussian at 3.12~keV for the Ar line. Since this spectrum comprises contributions from various regions with different physical properties, it is not expected to be well fitted by a model with a single thermal component. However, it helps to sketch the overall properties of the hot gas interior to the remnant. The two models produce similar physical parameters, but from comparison of the statistical goodness, the {\\em vsedov} model seems to better fit the spectrum than {\\em vpshock}. Both models produce elevated abundances of S ($\\sim1.9$--2.8 times solar) and Ca ($\\sim$2.8--6.8), indicating that the plasma is metal-enriched. The postblast shock temperature $kT_s\\sim0.37$--$0.51\\keV$ derived with the {\\em vsedov} model corresponds to a mean gas temperature $\\approx1.27kT_s\\sim0.47$--$0.65\\keV$ for the Sedov (1959) case, comparable to the mean temperature $\\sim0.62$--$0.66\\keV$ derived with the {\\em vpshock} model. The mean ionization timescale is $\\lsim10^{12}\\cm^{-3}\\,\\mbox{s}$, which indicates that the remnant as a whole is close to, but has not completely reached, ionization equilibrium. The unabsorbed fluxes (0.5--$10\\keV$) in the field of view inferred from the two models are (2.3--$2.7)\\E{-10}$ and (5.4--$8.6)\\E{-10}\\ergs\\cm^{-2}\\ps$, corresponding to X-ray luminosities of (5.1--$6.0)\\E{35}\\du^2$ and (1.2--$1.9)\\E{36}\\du^2\\ergs\\ps$, respectively.  \\begin{center} \\begin{deluxetable*}{c|cc} \\tabletypesize{\\footnotesize} \\tablecaption{Spectral Fitting Results for the Entire Remnant and Estimates of the Das Density} \\tablewidth{0pt} \\tablehead{ \\colhead{Parameters} \\vline & \\colhead{\\em vpshock} & \\colhead{\\em vsedov} } \\startdata $\\chi_{\\nu}^{2}$ (dof) & 1.89 (162) & 1.57 (161) \\\\ $\\NH$ ($10^{22}\\cm^{-2}$) & $2.3\\pm0.1$ & $2.3\\pm0.1$ \\\\ Temperature ($\\keV$) & $kT_x=0.63^{+0.03}_{-0.01}$$^{\\,\\,\\rm a}$ & $kT_s=0.41^{+0.10}_{-0.04}$$^{\\,\\,\\rm b}$ \\\\ &       & $kT_e=0.41^{+0.04}_{-0.03}$$^{\\,\\,\\rm b}$ \\\\ Ion.\\ Timescale ($10^{11}\\cm^{-3}\\,{\\rm s}$)$^{\\rm c}$ & $12.0^{+2.7}_{-4.8}$ & $4.19^{+1.15}_{-0.83}$ \\\\ $f\\nel\\nH V/\\du^{2}$ ($10^{58}\\cm^{-3}$)$^{\\rm d}$ & $2.01^{+0.10}_{-0.19}$ & $3.36^{+0.85}_{-0.72}$ \\\\ {[Mg/H]} & $0.85^{+0.09}_{-0.08}$ & $0.63^{+0.08}_{-0.07}$ \\\\ {[Si/H]} & $1.00^{+0.07}_{-0.06}$ & $1.13^{+0.12}_{-0.09}$ \\\\ {[S/H]} & $2.10^{+0.20}_{-0.22}$ & $2.47^{+0.36}_{-0.28}$ \\\\ {[Ca/H]} & $4.87^{+1.92}_{-1.77}$ & $4.45^{+1.84}_{-1.66}$ \\\\ Flux ($10^{-10}\\ergs\\cm^{-2}\\ps$) & $2.56^{+0.12}_{-0.25}$ & $6.88^{+1.74}_{-1.47}$ \\\\ $\\nH/f^{-1/2}\\du^{-1/2}$ ($\\cm^{-3}$)\\hspace{1mm}$^{\\rm e}$ & $0.34^{+0.01}_{-0.02}$ & $0.44^{+0.06}_{-0.05}$ \\tablecomments{The unabsorbed fluxes are in the 0.5--$10\\keV$ band. The net count rates of the on- and off-source spectra are $(8.52\\pm0.06)\\E{-1}$ and $(1.86\\pm0.14)\\E{-2}$ cts$\\ps$, respectively. Confidence ranges are at the 90\\% level. } \\enddata \\tablenotetext{a}{\\phantom{0} Mean electron temperature.} \\tablenotetext{b}{\\phantom{0} Mean shock temperature and electron temperature immediately behind the shock front, respectively.} \\tablenotetext{c}{\\phantom{0} The ionization timescale is defined as the product of the electron density and the time since the passage of the shock in the {\\em vpshock} model, and as the product of the electron density immediately behind the shock and the remnant's age in the {\\em vsedov} model. The value listed for {\\em vpshock} is the upper limit for the range of the timescale.} \\tablenotetext{d}{\\phantom{0} Where $f$ denotes the filling factor of the hot gas.} \\tablenotetext{e}{\\phantom{0} In the estimate of the densities, we assume a spherical volume of diameter $17'$ for the X-ray--emitting gas.} \\label{T:ent_spec} \\end{deluxetable*} \\end{center}  The spatial distribution of the relative strength of the S He$\\alpha$ emission is shown in the equivalent width (EW) image of this line in Figure~\\ref{f:ew}. This image was constructed using a method similar to those of Hwang et al.\\ (2000) and Park et al.\\ (2002). A major difference of our method from theirs is that we rebin the data using an adaptive mesh, with each bin in each narrow-band image including at least 10 counts (see Warren et al.\\ 2003 for similar binning). The image shows that the S~line emission is distributed over a broad region but is strongest in the bright eastern portion of the inner arc. The Ca line is very weak, and we did not produce a Ca EW image.  \\begin{figure}[tbh!] \\centerline{ {\\hfil\\hfil \\psfig{figure=f8.ps,height=3.in,angle=270, clip=} \\hfil\\hfil}} \\caption{ Equivalent width image of the S~He$\\alpha$ ($\\sim2.46\\keV$) line on a square root scale overlaid with the broadband (0.3--7 keV) diffuse X-ray contours (at seven logarithmic scale levels between 120 and 400 photons$\\cm^{-2}\\ps\\,\\mbox{sr}^{-1}$). The energy range used to extract the S~line image is 2.28--2.70 keV, while the low- and high-energy ranges around the S~line used to estimate the underlying continuum are 2.07--2.28 and 3.30--3.70 keV. The dashed box denotes the field of view of ACIS-I. } \\label{f:ew} \\end{figure}  \\subsubsection{Small-Scale Properties} The spectra of the 10 substructures all show distinct emission lines, indicating optically thin thermal plasma. They were fitted with the {\\em vpshock} model, and the spectral fit parameters are summarized in Table~\\ref{T:xspec}. A Gaussian at $\\sim3.1\\keV$ is added to account for the Ar~XVII line in the region C1 spectrum. The model fits allow the abundances of key line-producing elements to vary while leaving other abundances fixed to the solar value. Estimates of the hydrogen number densities ($\\nH$) inferred from the volume emission measures ($fn_e\\nH V$, where $f$ is the filling factor of the hot gas and $n_e\\sim1.2\\nH$ is assumed) are also given in Table~\\ref{T:xspec} (and Table~\\ref{T:ent_spec} for the entire remnant). In the derivation of the densities we have assumed that the three-dimentional shapes of the elliptical regions (E1, C1, C2, W1, W2, S, L1, and L2) are ellipsoids, those of the arclike regions (E2 and N) are projections of spherical segments, and that of the entire X-ray--emitting volume is a sphere of diameter $17'$. Deviations from these assumptions and the nonuniformity of the X-ray--emitting plasma are consolidated into the factors $f$ of the individual regions. The regions show significant spectral variation. The gas temperatures ($\\sim 1\\keV$) along the inner semicircular arc (regions C1 and C2) are higher than those ($\\sim0.4$--$0.7\\keV$) along the softer outer X-ray arc (regions E1, E2, and N). The densest gas in the eastern region E1 is consistent with the suggestion, based on the radio and X-ray brightness peaks, that the forward shock encounters a small dense region there (\\S~\\ref{sec:spatial}). The gas density of the low surface brightness regions (L1 and L2) between the two arcs is lower than that in the two bright arcs. The ionization timescales for the small-scale regions (except E1 and E2) are around $10^{12}\\cm^{-3}\\mbox{s}$ or even higher, and so they may be close to ionization equilibrium, but the ionization timescales for regions E1 and E2 (along the eastern outer shell) are relatively low [(0.6--$3.0)\\E{11}\\cm^{-3}\\mbox{s}$], indicating that the plasma is far from ionization equilibrium. The sulfur abundance in most of the small-scale regions aside from E2 is somewhat higher than solar and seems to peak in region C1, the eastern portion of the inner arc, as shown in the S He$\\alpha$ line distribution in the EW map.  \\begin{center} \\begin{deluxetable*}{cc|ccccccccc} \\tabletypesize{\\footnotesize} \\tablecaption{{\\em vpshock} Fitting Results and Estimates of the Gas Density} \\tablewidth{0pt} \\tablehead{ \\colhead{Regions} & \\colhead{Net Count Rate} \\vline & \\colhead{$\\chi^{2}_{\\nu}$ (d.o.f.)} & \\colhead{$\\NH$} & \\colhead{$kT_{x}$} & \\colhead{$n_e t$} & \\colhead{$f\\nel\\nH V/\\du^{2}$\\hspace{1mm}$^{\\rm a}$} & \\colhead{[Mg/H]} & \\colhead{[Si/H]} & \\colhead{[S/H]} & $\\nH/f^{-1/2}\\du^{-1/2}$\\hspace{1mm}$^{\\rm b}$\\\\ \\noalign{\\smallskip} \\colhead{} & \\colhead{($10^{-2}$ cts$\\ps$)} \\vline & \\colhead{} & \\colhead{($10^{22}\\cm^{-2}$)} & \\colhead{(keV)} & \\colhead{($10^{11}\\cm^{-3}\\,{\\rm s}$)} & \\colhead{($10^{56}\\cm^{-3}$)} & \\colhead{} & \\colhead{} & \\colhead{} & \\colhead{(cm$^{-3}$)} } \\startdata E1 & $4.00\\pm0.11$ & 1.07 (54) & $2.7\\pm0.2$ & $0.47\\pm0.04$ & $1.68^{+1.36}_{-0.42}$ & $36.2^{+14.3}_{-11.1}$ & $0.74^{+0.28}_{-0.20}$ & $0.86^{+0.25}_{-0.19}$ & $1.86^{+0.82}_{-0.63}$ & $2.1^{+0.4}_{-0.3}$\\\\ E2 & $6.05\\pm0.14$ & 0.92 (69) & $2.5^{+0.1}_{-0.2}$ & $0.60^{+0.07}_{-0.06}$ & $0.89^{+0.76}_{-0.34}$ & $29.1^{+10.2}_{-4.6}$ & $0.68^{+0.22}_{-0.17}$ & $0.74^{+0.21}_{-0.08}$ & $1.29^{+0.52}_{-0.45}$ & $0.73^{+0.13}_{-0.06}$\\\\ N & $7.58\\pm0.15$ & 1.02 (75) & $1.9\\pm0.2$ & $0.60^{+0.04}_{-0.02}$ & 18.9 ($>5.6$) & $16.7^{+3.4}_{-4.2}$ & $0.67^{+0.14}_{-0.21}$ & $1.06^{+0.27}_{-0.22}$ & $2.48^{+0.94}_{-0.66}$ & $0.43^{+0.04}_{-0.05}$\\\\ S  & $4.18\\pm0.12$ & 0.98 (54) & $2.3\\pm0.2$ & $0.72^{+0.07}_{-0.05}$ & 12.7 ($>3.8$) & $7.2^{+1.9}_{-1.0}$ & $1^{\\rm c}$ & $1^{\\rm c}$ & $1.73^{+0.73}_{-0.55}$ & $0.25^{+0.03}_{-0.02}$\\\\ C1  & $8.19\\pm0.16$ & 1.26 (90) & $1.8\\pm0.6$ & $1.21^{+0.45}_{-0.39}$ & $3.7^{+22.6}_{-1.6}$ & $4.7^{+6.2}_{-1.3}$ & $0.83^{+0.52}_{-0.38}$ & $1.44^{+0.43}_{-0.32}$ & $2.97^{+0.39}_{-0.55}$ & $0.43^{+0.28}_{-0.06}$\\\\ C2  & $5.13\\pm0.13$ & 0.85 (68) & $2.1^{+0.2}_{-0.3}$ & $0.76^{+0.25}_{-0.09}$ & $7.2^{+31.8}_{-5.0}$ & $6.7^{+2.4}_{-2.0}$ & $1.22^{+0.69}_{-0.45}$ & $1.10^{+0.40}_{-0.28}$ & $1.92^{+0.88}_{-0.59}$ & $0.63^{+0.11}_{-0.09}$\\\\ W1  & $4.12\\pm0.11$ & 1.00 (54) & $2.6^{+0.3}_{-0.2}$ & $0.69^{+0.05}_{-0.04}$ & $7.6^{+72.2}_{-5.2}$ & $9.8\\pm2.7$ & $0.84^{+0.47}_{-0.40}$ & $1.05^{+0.38}_{-0.29}$ & $1.79^{+0.78}_{-0.56}$ & $0.37\\pm0.05$\\\\ W2  & $2.73\\pm0.10$ & 1.15 (39) & $2.3^{+0.3}_{-0.4}$ & $0.62^{+0.05}_{-0.07}$ & 7.3 ($>2.0$) & $6.7^{+3.8}_{-2.0}$ & $1^{\\rm c}$ & $1^{\\rm c}$ & $1.92^{+1.15}_{-0.81}$ & $0.30^{+0.09}_{-0.04}$\\\\ L1+L2 & $3.21\\pm0.10$ & 1.17 (47) & $2.2^{+0.2}_{-0.3}$ & $0.89\\pm0.06$ & $5.2^{+13.6}_{-2.8}$ & $3.9\\pm1.3$ & $1^{\\rm c}$ & $1^{\\rm c}$ & $1.60^{+0.60}_{-0.54}$ & $0.26\\pm0.04$ \\tablecomments{ The net count rates of the on-source spectra are listed in the second column; that of the off-source spectrum is $(1.86\\pm0.14)\\E{-2}$ cts$\\ps$. Confidence ranges are at the 90\\% level. } \\enddata \\tablenotetext{a}{\\phantom{0} Where $f$ denotes the filling factor of the hot gas.} \\tablenotetext{b}{\\phantom{0} In estimating the densities, we assume oblate spheroids for elliptical regions E1 (with half-axes $1.50'\\times1.50'\\times1.29'$), C1 ($2.62'\\times2.62'\\times1.31'$), and C2 ($2.22'\\times2.22'\\times1.18'$), ellipsoids (with half line-of-sight size to be the outer radius of shell region N of $8'$) for regions W1 (with half-axes $2.11'\\times1.49'\\times8'$), W2 ($2.22'\\times1.50'\\times8'$), S ($2.34'\\times2.20'\\times8'$), L1 ($1.43'\\times0.82'\\times8'$), and L2 ($1.87'\\times0.84'\\times8'$), and spherical segments for regions E2 (with inner and outer radii $5.8'$ and $7.6'$ and sector angle $80^{\\circ}$) and N ($5.8'$, $8.0'$, $85^{\\circ}$).} \\tablenotetext{c}{\\phantom{0} Fixed to the solar abundance.} \\label{T:xspec} \\end{deluxetable*} \\end{center} ",
        "conclusions": "\\snr\\ has been classified as a thermal composite (or mixed-morphology) remnant. It is unusual in that the diffuse thermal X-ray emission peaks not in the center, but along two concentric semicircular arcs. The centroid of the X-ray emission lies in the eastern half of the remnant, and the emission of the inner arc is harder than that of the outer region.  The thermal properties of \\snr, as inferred from \\S~\\ref{sec:spec}, suggest that the X-ray--emitting gas is essentially close to equilibrium of ionization (except along the eastern rim) and enriched with sulfur and calcium. If we adopt $\\nH\\sim0.4f^{-1/2}\\du^{-1/2}\\cm^{-3}$ as the mean gas density, then the mass of hot gas is $M=1.4\\nH\\mH fV\\sim110f^{1/2}\\du^{5/2}\\Du^3\\Msun$, where $\\mH$ is the hydrogen atomic mass and $\\Du$ is the remnant's diameter $D$ scaled by the radio size, $\\sim20'$ (thus $D\\approx25\\Du\\du\\parsec$). The thermal energy contained in the remnant is $E_{th}=(3/2)2.3\\nH kT fV\\sim3\\E{50}f^{1/2}\\du^{5/2}\\Du^3\\ergs$, where a mean temperature $kT\\sim0.6\\keV$ for the hot gas is used. With a shock temperature $kT_s\\sim0.4\\keV$, the shock velocity would be $v_s=(16kT_s/3\\bar{\\mu}\\mH)^{1/2}\\sim580\\km\\ps$, where the mean atomic weight $\\bar{\\mu}=0.61$. The dynamical age of the remnant, $\\sim D/5v_s\\sim8.4\\E{3}\\du\\Du\\yr$, is derived with the Sedov (1959) solution, which, because of the unusual morphology and complicated gas components, is only an approximation. The age estimate ($\\sim3.5\\E{3}\\yr$) based on the \\ROSAT\\ X-ray observation used a gas temperature ($3.2\\keV$) that was too high (Seward et al.\\ 1996). The Sedov model fit yields an ionization timescale $\\tau_{\\rm Sedov}\\sim4.2\\E{11}\\cm^{-3}\\,\\mbox{s}$, which implies a dynamical age $t\\approx\\tau_{\\rm Sedov}/(1.2\\times4\\nH)\\sim7\\E{3}f^{1/2}\\du^{1/2}\\yr$ (where we have assumed that the preshock density is similar to the mean interior density $\\nH$). This is similar to the above Sedov age. On the other hand, the upper limit on the mean ionization timescale derived with the {\\em vpshock} model (Table~\\ref{T:ent_spec}) gives an age $\\sim9\\E{4}f^{1/2}\\du^{1/2}\\yr$. This is similar to the ionization age ($>8\\E{4}\\yr$) found in the \\ASCA\\ X-ray study (Enoguchi et al.\\ 2002) and is comparable to the remnant's radiative cooling timescale, $t_{\\rm cool}\\sim4\\E{4}E_{51}^{0.24}[n_0/(0.4\\cm^{-3})]^{-0.52}\\yr$, where $E_{51}$ is the SNR's explosion energy in units of $10^{51}\\ergs$ and $n_0$ is the preshock density (Falle 1981). In this case the postshock temperature would be $\\lsim1\\E{6}\\K$ and thus the gas would be X-ray--faint. This does not agree with the postshock temperatures derived from the spectral fits and the fact that the eastern rim of the SNR is X-ray--bright. A possibility to reconcile the different estimates might be for the X-ray emitting gas to have a very low filling factor, $f\\lsim1\\E{-2}\\Du^2\\du$, but this seems unlikely.   The mass of the X-ray--emitting gas, $\\sim 10^2 f^{1/2} \\Msun$, suggests that the SNR is not in the free-expansion phase and is probably dominated by the swept-up ambient medium.  Even so, Kes~27 appears to be enriched in S and Ca, and this is not a unique case. One sample of 23 thermal composite (or mixed-morphology) SNRs showed that 10 of them to be detected as metal-enhanced (Lazendic \\& Slane 2006). The contribution of supernova ejecta has been suggested to be the cause of both the elevated metal abundances and the X-ray brightening in the central regions. For example, the overabundant Mg, Si, and (possibly) S in thermal composite SNR~G290.1$-$0.8 (Slane et al.\\ 2002) and the Ne, Si, and S enhancement in W44 are ascribed to supernova ejecta (Shelton et al.\\ 2004).  Environments of nonuniform density are often seen around SNRs, especially those adjacent to dense clouds. This sometimes is consistent with systematic variation of the intervening hydrogen column density along the line of sight to the SNR, as for 3C~391 (Chen et al.\\ 2004) and 3C~397 (Safi-Harb et al.\\ 2005). \\snr's HI density enhancement is found to the east of the remnant (McClure-Griffiths et al.\\ 2001), and the Spitzer IRAC 8~$\\mu$m observation (Reach et al.\\ 2006) shows a large-scale environment with a density gradient increasing from west to east. However, our X-ray study of \\snr\\ does not show a clear trend of $\\NH$ (and spectral) variations along the density gradient (although E1 is the only region showing a significantly higher value; Table~\\ref{T:xspec}). This may indicate that the environmental hydrogen column density may take only a very low fraction in the intervening interstellar column ($\\sim2\\E{22}\\cm^{-2}$), either because the cloud complex in the vicinity is not very compact or bacause it is not very large.  The double X-ray rings are not expected in either the cloud evaporation or the thermal conduction model. In the latter (Cox et al.\\ 1999; Shelton et al.\\ 1999), thermal conduction in the remnant smooths out the temperature gradient from the hot interior to the cooler shell and increases the central density, resulting in luminous X-rays in the interior; this effect is dominant in the radiative stage. This model predicts a centrally peaked X-ray surface brightness for the SNR and, usually, a high ionization timescale for the hot gas, inconsistent with our observation. The cloud evaporation model (White \\& Long 1991) suggests that when an SNR expands in an inhomogeneous interstellar medium % whose mass is mostly contained in small clouds, the clouds engulfed by the blast wave can be evaporated to slowly increase the density of the interior hot gas; as a result, the SNR appears internally X-ray--brightened. This model can reproduce an X-ray--bright inner ring like that in the northeastern half of \\snr\\ and can also reproduce a fading-out radial surface brightness profile like that in the southwestern half (see Fig.`\\ref{f:rp}). A representative fit to the inner ring's brightness profile using the cloud evaporation model requires the ratio of the cloud evaporation timescale to the SNR's age to be $\\tau=10$ and the ratio of the mass of the cloudlets to the mass of the intercloud medium to be $C=50$. The fit is not unique, and combinations of $\\tau=8$--12 and $C=40$--60 produce similar profiles (relatively scaled to match the maximum brightness). A representative fit to the southwest radial brightness profile would need $\\tau=20$ (15--25) and $C=56$ (45--70). However, the generally higher inner gas temperature as compared with the outer region is not expected for such high $\\tau$- and $C$-values in the evaporation model. Moreover, this model overestimates the surface brightness inside the inner ring and cannot reproduce the outer X-ray ring in the northeast. It also cannot naturally explain why the physical conditions exemplified by the parameters $\\tau$ and $C$ should be different in the two halves and why there are not two rings in the southwestern half, in contrast to the northeastern.  \\begin{figure}[tbh!] \\centerline{ {\\hfil\\hfil \\psfig{figure=f9.eps,height=3.5in,angle=0, clip=} \\hfil\\hfil}} \\caption{ Diagram showing the structure and environment of the Kes~27 remnant. } \\label{f:diagram} \\end{figure}  Another scenario is required to account for the unusual X-ray morphology of \\snr, which is different from that of other thermal composite remnants. Actually, the explanation must be consistent with the following facts: (1) The large-scale environment around \\snr\\ has a density gradient increasing from west to east, as revealed in the infrared (Reach et al.\\ 2006). (2) As sketched in Figure~\\ref{f:diagram}, the remnant is surrounded by a ``crocodile's mouth''--like \\HI\\ shell (McClure-Griffiths et al.\\ 2001). (3) In X-rays, the remnant is brightened in the northeastern half and fades out to the southwestern low-density region. (4) The inner and outer X-ray arcs are both located in the northeastern half, where the ambient medium is denser than in the southwest. (5) The outer arc is in contact with the \\HI\\ shell. (6) The inner arc is hotter than the outer arc. (7) As noted by Enoguchi et al.\\ (2002), the radio polarization observed by Milne et al.\\ (1989) is strong in both the eastern rim and the center.  We suggest that both the outer and inner arcs represent shock waves, propagating outward and inward, respectively. The observed polarization in the eastern rim and the center is due to shocked and compressed matter in the two arcs. The present inward shock is not the initial reverse shock caused by expansion into low-density circumstellar material. The remnant is middle-aged, and this has long passed (McKee 1974). The present inward shock is a reverse shock due to the distant HI shell.   A comprehensive scenario is as follows: The massive progenitor exploded in a cavity excavated by the strong stellar wind and ionizing radiation. Because of the density gradient in the environmental medium, the wind cavity had a small radius in the east and a large radius in the west, or may even not have been confined in the west. After the supernova explosion, in the west and southwest, the blast wave propagated into a low-density medium, and the X-ray emission is faint. In the north and east, the blast wave has hit the cavity wall.  A transmitted shock is now propagating into the dense wall, and a reflected shock is traveling back toward the remnant's center, shocking the supernova ejecta to higher temperature. The outer arc contains gas heated by the transmitted shock, and the inner arc is the interior metal-enriched gas heated by the reflected shock (see again Fig.~\\ref{f:diagram}). Shock-heated S atoms in the compressed inner arc emit a strong He$\\alpha$ line.   Chen et al.\\ (2003) considered a thin-shell model for a blast shock hitting a wind-cavity wall but did not include shock reflections. Dwarkadas (2005, 2007) simulated shocks reflected back from a cavity wall, and a picture of shock reflection and transmission similar to what is seen in \\snr\\ was predicted. To capture the essential physics essence, we use the theory of reflected shocks described by Sgro (1975) to estimate some relevant parameters. For a shock reflected by a dense wall, the temperature ratio between the post--reflected-shock gas and post--transmitted-shock gas is given by \\begin{equation} \\frac{T_r}{T_t}=\\frac{A}{A_r}, \\label{eq:Tratio} \\end{equation} where $A_r$ is the density contrast between the post--reflected-shock gas and post--incident-shock gas and $A$ is the wall-to-cavity density contrast, which is related to $A_r$ as \\begin{equation} A=\\frac{3A_r(4A_r-1)}{\\{[3A_r(4-A_r)]^{1/2}-\\sqrt{5}(A_r-1)\\}^2}. \\end{equation} The pressure ratio of the post--transmitted-shock gas to the post--incident-shock gas is \\begin{equation} \\beta=\\frac{4A_r-1}{4-A_r}. \\end{equation} Using these equations, we plot the functional relations of $T_r/T_t$, $A_r$, and $\\beta$ with $A$ in Figure~\\ref{f:refl_shock}. Regions E2 versus C1 and N versus C2 represent two pairs of transmitted versus reflected shocks. The ranges of the temperature ratios obtained from Table~\\ref{T:xspec} are shown in Figure~\\ref{f:refl_shock}. Because there is a large scatter in the fitted temperatures, there are large uncertainties in the wall-to-cavity density $A$, which seems to be $\\sim$ 2. The density contrast between the post--reflected-shock gas and postincident shock gas $A_r$ is around 1.3. If regions W1, W2, and S can be taken as roughly indicative of a post--incident-shock gas density $\\sim(0.2-0.4)f^{1/2}\\du^{1/2}\\cm^{-3}$ (Table~\\ref{T:xspec}), then the post--reflected shock density $0.4-0.6f^{1/2}\\du^{1/2}\\cm^{-3}$ in C1 and C2 is consistent with the values of $A_r$ (although this estimate is crude, because the filling factor $f$ may vary from region to region).  \\begin{figure}[tbh!] \\centerline{ {\\hfil\\hfil \\psfig{figure=f10.ps,height=3.in,angle=0, clip=} \\hfil\\hfil}} \\caption{Parameters for the reflected shock. The large box represents the value ranges for region pair E2 and C1, and the small box is for region pair N and C2 (see text).} \\label{f:refl_shock} \\end{figure}  Such a reflected shock launched from the cavity wall is consistent with the observed properties of \\snr. As the remnant continues to age, the transmitted shock in the dense medium will become fainter in X-rays, while the reflected shock will converge in the center, producing a truly centrally X-ray--brightened morphology. The reflection of shocks is thus another mechanism that may produce thermal composite remnants."
    },
    "0711/0711.2510_arXiv.txt": {
        "abstract": "'Isolated planetary mass objects' (IPMOs) have masses close to or below the Deuterium-burning mass limit ($\\sim 15\\,M_\\mathrm{Jup}$) -- at {\\it the bottom of the stellar initial mass function}. We present an exploratory survey for disks in this mass regime, based on a dedicated observing campaign with the Spitzer Space Telescope. Our targets include the full sample of spectroscopically confirmed IPMOs in the $\\sigma$\\,Orionis cluster, a total of 18 sources. In the mass range $8\\ldots 20\\,M_\\mathrm{Jup}$, we identify 4 objects with $>3\\sigma$ colour excess at a wavelength of 8.0$\\,\\mu m$, interpreted as emission from dusty disks. We thus establish that a substantial fraction of IPMOs harbour disks with lifetimes of at least 2-4\\,Myr (the likely age of the cluster), indicating an origin from core collapse and fragmentation processes. The disk frequency in the IPMO sample is $29\\pm ^{16}_{13}$\\% at 8.0\\,$\\mu m$, very similar to what has been found for stars and brown dwarfs ($\\sim 30$\\%). The object SOri\\,70, a candidate  $3\\,M_\\mathrm{Jup}$ object in this cluster, shows IRAC colours in excess of the typical values for field T dwarfs (on a 2$\\sigma$ level), possibly due to disk emission or low gravity. This is a new indication for youth and thus an extremely low mass for SOri\\,70. ",
        "introduction": "\\label{intro}  Over the past five years, it has been firmly established that the evolution of brown dwarfs from 1-10\\,Myr follows the blueprint known from T Tauri stars: They show evidence for ongoing gas accretion and outflows \\citep[e.g.][]{2005ApJ...625..906M,2005ApJ...626..498M}, and harbour circum-sub-stellar disks with lifetimes of 5-10\\,Myr \\citep[e.g.][]{2003AJ....126.1515J,2007ApJ...660.1517S}, not vastly different from stars. This has often been interpreted as an indication for a common origin of stars and brown dwarfs -- sub-stellar objects as the natural extension of the stellar initial mass function (IMF) and the lowest-mass outcome of core collapse and fragmentation processes \\citep[see review by][and references therein]{2007prpl.conf..443L}. At the same time, ongoing deep surveys have pushed the sensitivity limits down to masses below the Deuterium burning limit at $\\sim 15\\,M_{\\mathrm{Jup}}$, finding evidence for a cluster population with masses comparable to giant planets \\citep{2000Sci...290..103Z,2000MNRAS.314..858L}. These objects {\\it at the bottom of the IMF} are inconsistently called planemos, sub-brown dwarfs, cluster planets, or isolated planetary mass objects in the literature; in the absence of a definite nomenclature, we will use the latter term (abbreviated IPMO) in the following.  The nature and origin of IPMOs have been debated intensely in recent years. Most of the problems raised in the discussions about the formation of brown dwarfs apply even more acutely to IPMOs \\citep[see review by][]{2007prpl.conf..459W}. For example, it is under debate whether IPMOs can form directly from the collapse of ultra-low mass cores; instead, they may be giant planets ejected from circumstellar disks or stellar embryos ejected from mini-clusters or decaying multiple systems.  As of today, the observational constraints on the nature of IPMOs are very limited. Evidence for accretion or disks has been reported for a few young sources with masses around the Deuterium burning limit \\citep{2002A&A...393..597N,2002ApJ...569L..99Z,2005ApJ...635L..93L,2006ApJ...644..364A,2006ApJ...647L.167J}. The next step is clearly to establish frequency and lifetimes of IPMO disks, and to push the mass limits of disk surveys towards the least massive objects in star forming regions. Here we present a Spitzer survey for disks around IPMOs in the $\\sigma$\\,Orionis cluster, which harbours the largest IPMO population identified thus far. With a well-explored stellar and brown dwarf population \\citep{2001ApJ...556..830B,2007A&A...470..903C} and a likely age of 2-4\\,Myr \\citep{2002A&A...384..937Z,2004AJ....128.2316S}, $\\sigma$\\,Ori is ideal for probing the disk frequency as a function of mass as well as the longevity of IPMO disks. ",
        "conclusions": "\\label{sum}  We have conducted a disk survey for isolated planetary mass objects, based on deep observations with the Spitzer Space Telescope. Our targets comprise the full sample of spectroscopically confirmed IPMOs in the $\\sigma$\\,Orionis cluster. As a pioneering study, the results presented here are not intended to be the last word on the subject. Instead, we intend this Letter to serve as an important step towards a more profound understanding of the nature and origin of the lowest mass objects found in isolation.  A substantial fraction of our targets exhibits mid-infrared excess emission indicative for the existence of a dusty disk. From our IRAC data, we derive a disk fraction of $29\\pm ^{16}_{13}$\\% (4/14) at 8.0\\,$\\mu m$. A similar value is found at 5.8\\,$\\mu m$. These results are fully consistent with the disk frequencies derived for T Tauri stars and brown dwarfs in the same cluster. The detection of dusty disks around IPMOs in $\\sigma$\\,Ori is another firm evidence for their nature as lowest mass members of this young cluster. The finding clearly establishes that objects at the bottom of the IMF can harbour dusty disks with lifetimes of at least 2-4\\,Myr, the most likely age of the $\\sigma$\\,Ori. Thus, our results fit into previous claims for a T Tauri-like phase in the planetary mass regime \\citep{2002A&A...393..597N,2005ApJ...635L..93L,2006ApJ...644..364A,2006ApJ...647L.167J}. Disk fractions and thus lifetimes are similar for objects spanning more two orders of magnitude in mass (0.008 to 2$\\,M_{\\odot}$), possibly indicating that stars, brown dwarfs, and IPMOs share a common origin. Star formation theory thus has to account for a number of objects with masses below the Deuterium burning limit.  Our sample includes SOri\\,70, the faintest candidate member so far in $\\sigma$\\,Ori, which is detected in our IRAC images from 3.6 to 5.8$\\,\\mu m$. Compared with field T dwarfs, the source stands out based on the IRAC colours, with excesses at 4.5 and and 5.8$\\,\\mu m$ at a $2\\sigma$ level. This may be an indication for youth and possibly disk occurence for an object with an estimated mass of 3$\\,M_\\mathrm{Jup}$ \\citep{2002ApJ...578..536Z}. Uncertainties are large, though, meaning that this should be seen as motivation for further follow-up rather than a definite confirmation of the 'cluster planet' nature of this object."
    },
    "0711/0711.2456_arXiv.txt": {
        "abstract": "We use a deep \\emph{XMM-Newton} Reflection Grating Spectrometer observation to examine the X-ray emission from the core of the Centaurus cluster of galaxies. We clearly detect Fe~\\textsc{xvii} emission at four separate wavelengths, indicating the presence of cool X-ray emitting gas in the core of the cluster. Fe ions from Fe~\\textsc{xvii} to \\textsc{xxiv} are observed. The ratio of the Fe~\\textsc{xvii} 17.1{\\AA} lines to 15.0{\\AA} line and limits on O~\\textsc{vii} emission indicate a lowest detected temperature in the emitting region of 0.3 to 0.45~keV (3.5 to $5.2 \\times 10^6$~K). The cluster also exhibits strong N~\\textsc{vii} emission, making it apparent that the N abundance is supersolar in its very central regions. Comparison of the strength of the Fe~\\textsc{xvii} lines with a Solar metallicity cooling flow model in the inner 17~kpc radius gives mass deposition rates in the absence of heating of $1.6-3 \\Msunpyr$. Spectral fitting implies an upper limit of $0.8\\Msunpyr$ below 0.4~keV, $4 \\Msunpyr$ below 0.8~keV and $8\\Msunpyr$ below 1.6~keV.  The cluster contains X-ray emitting gas over at least the range of 0.35 to 3.7~keV, a factor of more than 10 in temperature. We find that the best fitting metallicity of the cooler components is smaller than the hotter ones, confirming that the apparent metallicity does decline within the inner 1~arcmin radius. ",
        "introduction": "In the centres of many clusters of galaxies the mean radiative cooling time of the intracluster medium (ICM) is short, typically less than 1~Gyr (e.g. \\citealt{Voigt04}). This is also where the ICM is significantly cooler than the outskirts. It is therefore expected that a cooling flow \\citep{Fabian94} should be formed, where material cools out of the ICM.  High spectral resolution X-ray studies of nearby clusters of galaxies using the Reflection Grating Spectrometer (RGS) instruments on \\emph{XMM-Newton} has revealed a lack of cool X-ray emitting gas in these objects \\citep{Tamura01b, Tamura01a, Peterson01, Kaastra01, Sakelliou02, Peterson03}.  There are emission lines observed from gas down to a half or a third of the outer temperature, but with very little gas at cooler temperatures (see \\citealt{PetersonFabian06} for a review). The main spectral indicator of the lack of cool gas is the weakness of Fe~\\textsc{xvii} lines.  A cluster where there has been tentative evidence for the existence of cool gas is Abell~2597 \\citep{MorrisFabian05}, which shows possible Fe~\\textsc{xvii} emission in its RGS spectrum. NGC~4636 emits strong Fe~\\textsc{xvii} lines, but these are consistent with a small range of temperature as the group is cool \\citep{Xu02}. NGC~5044 is another cool system showing strong Fe~\\textsc{xvii} and hinting at O~\\textsc{vii} emission \\citep{Tamura03}.  The Centaurus cluster, Abell~3526, is a nearby bright cluster of galaxies. It lies at a redshift of 0.0104 \\citep{LuceyCurrieDickens86a}, and a 2-10 keV X-ray luminosity of $2.9 \\times 10^{43} \\ergps$ \\citep{Edge90}. It lies at a relatively low Galactic latitude of $21.6^\\circ$ with a Hydrogen column density towards the cluster of $\\sim 8.6 \\times 10^{20} \\psqcm$ (derived from the H~\\textsc{i} maps of \\citealt{Kalberla05}).  A cool X-ray emitting gas component has been seen in this object by low spectral resolution CCD quality observations. This component was first seen by \\emph{ASCA} \\citep{Fukazawa94} and \\emph{ROSAT} \\citep{AllenFabian94}, who found cool gas with a temperature of $\\sim 1$~keV in addition to a hot $\\sim 4$~keV component.  \\cite{Ikebe99} put forward a two-phase model fitting \\emph{ASCA} and \\emph{ROSAT} data, claiming 1.4 and 4.4~keV components over the inner 3~arcmin.  A colour analysis of \\emph{ROSAT} data \\citep{Sanders00} gave a mean temperature of 1.7~keV in the central 90~arcsec. \\cite{Allen01} found an improved fit to their \\emph{ASCA} spectra with two or more temperatures, with temperatures at 3.2 and 1.3~keV. There was also evidence for a hot or powerlaw component (seen previously by \\citealt{Allen00} and \\citealt{DiMatteo00}). \\emph{BeppoSAX} observations \\citep{Molendi02} found a peak temperature of 4~keV, dropping to below 2~keV in the core, and evidence for a hard component in the PDS instrument.  The excellent spatial resolution of \\emph{Chandra} provided a more complete picture of the temperature of the gas in the central regions \\citep{SandersCent02}. This observation showed a plume-like feature in the core of the cluster, emitting soft X-rays. Two-temperature component fits to projected spectra extracted from the plume were preferred to single-temperature fits, with temperatures of 0.7 and 1.5~keV. Outside this region the temperature increases quickly to the east to around 3.7~keV. To the west, there is a plateau of cooler gas at around 2.5~keV, before the temperature rises at a radius of around 190~arcsec. Deeper \\emph{Chandra} observations \\citep{Fabian05} confirmed this picture, highlighting the clear east-west asymmetry of the temperature distribution. \\emph{XMM-Newton} EPIC data confirmed the $\\sim 0.7\\keV$ component in the core \\citep{SandersEnrich06}, also finding the temperature of the ICM to drop to around 3.4~keV beyond 120~kpc radius. The \\emph{Chandra} observations also clearly showed that the metallicity of the ICM is inversely correlated with its temperature \\citep{SandersCent02,Fabian05,SandersEnrich06}.  An interesting feature of this cluster is that the metallicity of the gas towards the centre is significantly supersolar. The Fe metallicity peaks between 1.5 and 2\\Zsun \\citep{Fabian05,SandersEnrich06}. Si and S peak around 2\\Zsun, and Ni peaks around 4\\Zsun. In addition the metallicity of the gas appears to decline in the very central regions \\citep{SandersCent02}.  The multiple detections of cool X-ray emitting gas in the core of Centaurus provides an excellent opportunity to test the picture that there is only a range in 2--3 in X-ray gas temperature in clusters of galaxies. We therefore undertook a deep \\emph{XMM-Newton} RGS observation of Centaurus, which, when combined with the existing observation, gives a total exposure of around 160~ks.  In this paper we assume $H_0 = 70 \\kmpspMpc$, which gives an angular scale of 213~pc per arcsec for Centaurus. We assume the Solar metallicities of \\cite{AndersGrevesse89}. Error bars are quoted as 1$\\sigma$ and limits as 2$\\sigma$. ",
        "conclusions": "We clearly detect Fe~\\textsc{xvii} emission from the core of the Centaurus cluster. Fitting spectral models and measuring emission line ratios shows that the temperature of the gas declines to 0.3 to 0.45~keV. These results from this deep RGS observation show the widest range in ICM temperature detected in a cluster of galaxies, with a factor of 10 or more.  The data confirm that the metallicity of the ICM declines in the very central regions. Whether this is a real decline in metallicity, the result of an inhomogeneous metallicity distribution \\citep{MorrisFabian03} or excess continuum is unclear. We do, however, find strong nitrogen enhancement in the inner 6~kpc radius.  The very coolest gas is concentrated in the centre of the cluster, as previously found by \\emph{Chandra}, but there is much less cool material than would be expected from a simple cooling flow.  The problem of how to prevent this very cool gas, which has a mean radiative cooling time of $10^7$~yr, from cooling remains.  Any mechanism which heats the gas must preserve the sharp metallicity gradients and not switch off for timescales of greater than $10^7$~yr. Alternatively, the ICM in the cluster may be cooling in a non-radiative way below 0.4~keV. If this material formed stars, the rate of mass deposition is compatible with the metallicity of the ICM \\citep{SandersEnrich06}.  Using this very deep observation we are able to measure the amount of cool X-ray emitting gas very precisely in a cluster core. In order to get a more complete picture, it is necessary to make deep observations of several nearby clusters to see whether the large ranges in temperature are found in all objects."
    },
    "0711/0711.1872_arXiv.txt": {
        "abstract": "We discuss the relation between the highest energy cosmic rays (UHECR) and UHE neutrinos. The neutrinos produced in the sources of optically thin astrophysical sources have been linked to the UHECR emissivity of the Universe. The fluxes of cosmogenic neutrinos, generated in propagation by UHECR, also reflect the acceleration of these particles, the maximum acceleration energy, and the cosmological evolution of their sources. \\vspace{1pc} ",
        "introduction": "It is now a time when both subjects in the title are of high scientific interest in the field of particle astrophysics. The interest is related to the emergence of new set of experiments that are obviously of much better quality than what we know from the past.  The Southern Auger Observatory~\\cite{Auger} in Argentina is almost fully completed and is operating during construction. The Telescope Array~\\cite{TA} (TA) is under construction in Utah. Both these giant air shower arrays are designed and constructed to operate in hybrid mode, i.e. they employ both methods of air shower detection: surface array plus optical detectors that follow the fluorescent light generated by the giant air shower in the atmosphere. The optical detectors operate about 10\\% of the time but the energy assignment by integration over the shower longitudinal profile seems to be less model dependent than the classical method of relating the primary energy to the shower particle density on the ground. The latter depends much more on the hadronic interaction model used in Monte Carlo studies of the shower development. The fraction of showers detected in hybrid mode by both detectors can be used for a better calibration of the surface array. The sizes of the surface arrays, 3,000 km$^2$ for Auger and 1,000 km$^2$ for TA are very impressive, more than one order of magnitude higher than the previous largest array - Agasa.  There are also big news on the neutrino front. The IceCube neutrino observatory~\\cite{IceCube} on the South Pole is in the middle of  successful construction. Twenty two of the eighty underice strings are deployed and the deployment is now faster than initially expected. The Antares~\\cite{Antares} detector in the Mediterranean is also successfully close to completion and the km3net~\\cite{km3net} is working on the design of a cubic kilometer under water neutrino telescope in the Northern hemisphere. With these cubic kilometer neutrino detectors we approach for first time the dimensions needed for detection of astrophysical neutrinos - a gigaton of target matter or 6$\\times$10$^{38}$ target nucleons.  Experiments of this size and quality can certainly make the objects of their studies fashionable. The first high statistic results of the Auger observatory were posted on ArXive.org and reported at the 20th International Cosmic Ray Conference in Merida, Mexico in front of a huge audience. Auger has set~\\cite{Auger_gam} a strict limit of 2\\% on the gamma-ray contribution to the UHECR flux above 10$^{19}$ eV and in the following we shall assume that the highest energy cosmic rays originate in acceleration in powerful astrophysical objects.  IceCube has not yet detected any astrophysical neutrino signals but the limits already obtained have let to revisions of many models. ",
        "conclusions": "The new UHECR data sets agree that UHECR spectrum experiences a steepening energy spectrum at about 5-6$\\times$10$^{19}$ eV which looks consistent with the GZK effect which is due to the energy loss of these particles in photoproduction or photodisintegration interactions. It is not yet obvious what are the parameters of the acceleration process and what is the evolution of the cosmic ray sources. All these questions will probably have to wait until the total world statistics is increased by a large factor.  The composition of all particles above 10$^{19}$ eV is not yet established. Different experimental data are not far away from each other, but do not currently contribute to the understanding of the UHECR origin. The advance in that respect is the strict limit of 2\\% that the Auger collaboration has set on the fraction of $\\gamma$-rays in the cosmic ray flux above 10$^{19}$ eV.  Different UHECR models predict various fluxes of cosmogenic neutrinos that are generated by the cosmic rays at propagation from their sources to us. Detection of such neutrinos and a comparison of their fluxes to the direct observations of UHECR would contribute significantly to understanding of the UHECR origin. For this to happen, however, we will have to rely on new, bigger neutrino observatories."
    },
    "0711/0711.3046_arXiv.txt": {
        "abstract": "Determination of pulsar parallaxes and proper motions addresses fundamental astrophysical open issues. Here, after scrutinizing the ATNF Catalog searching for pulsar distances and proper motions,  we verify that for an ATNF sample of 212 Galactic run away pulsars (RAPs), which currently run across the Galaxy at very high speed and undergo large displacements, some gravitational-wave (GW) signals produced by such present accelerations appear to be detectable after calibration against the Advanced LIGO (LIGO II). Motivated by this insight, we address the issue of the pulsar kick at birth, or {\\sl short rise fling} from a supernova explosion, by adapting the theory for emission of GW by ultrarelativistic sources to this case in which Lorentz factor is $\\gamma \\sim 1$. We show that during the short rise fling each run away pulsar (RAP) generates a GW signal with characteristic amplitude and frequency that makes it detectable by current GW interferometers. For a realistic analysis, an efficiency parameter is introduced to quantify the expenditure of the rise fling kinetic energy, which is estimated from the linear momentum conservation law applied to the supernova explosion that kicks out the pulsar. The remaining energy is supposed to be used to make the star to spin. Thus, a comparison with the spin of ATNF pulsars having velocity in the interval $400-500$ km s$^{-1}$ is performed. The resulting difference suggests that other mechanisms (like differential rotation, magnetic breaking or magneto-rotational instability) should dissipate part of that energy to produce the observed pulsar spin periods. Meanwhile, the kick phenomenon may also occur in globular and open star clusters at the formation or disruption of very short period compact binary systems wherein abrupt velocity and acceleration similar to those given to RAPs during the short rise fling can be imparted to each orbital partner. To better analyzing these cases, pulsar astrometry from micro- to nano-arsec scales might be of much help. In case of a supernova, the RAP GW signal could be a benchmark for the GW signal from the core collapse. ",
        "introduction": "The detection of gravitational waves (GW) is one of the greatest goals of today's relativistic astrophysics. Because of the great success of Einstein's general relativity (GR) in explaining the dynamical evolution of the binary pulsar PSR 1913+16; continuously observed over almost 30 years by Taylor and Hulse \\cite{taylor2004}, most researchers in the field are confident that the very first GW signal to be detected must come from coalescing binary neutron star-neutron star or black hole-neutron star (NS-NS, BH-NS) systems. Ongoing LIGO science runs \\cite{Abbott2007} for the first time have set firm upper limits on the rate of binary coalescences in our Galaxy: at the level of a few events per year. Meanwhile the model-dependent estimate of such a rate at cosmic distance scales varies around $10^{-5}$ \\cite{postnov2006,kalogera2007,pacheco2006}. The expectation is enormous as well for an event like the supernova (SN) explosion.  { In this paper we focus on the radiation emitted when a pulsar is kicked out from the supernova remnant and is accelerated with very high speed, albeit non ultrarelativistic. Gravitational radiation is emitted whenever the nascent neutron star or pulsar changes its velocity, a condition that is fulfilled by the large thrust given to the nascent pulsar during the SN explosion, and also due to the pulsar dragging against the SN ejecta. (Recall that the gravitational field generated by a motion with constant velocity is nonradiative). We shall show that each individual event of launching a nascent pulsar to drift or run across the Galaxy, after a SN explosion, generates a gravitational wave signal that can be detected by current and planned GW interferometers such LIGO I, LIGO II, VIRGO, GEO-600, or TAMA-300. Nonetheless, it is stressed from the very beginning that the strain of the GW signal that we estimate below is {\\sl a lower bound}. The reason is that we are not taking into account the contribution to the GW amplitude, $h$, stemming from the effective acceleration that the pulsar receives at the kick at birth, which could be very large. Indeed, the distance over which the RAP is dragged against the supernova ejecta could also be astronomically large.}  At first glance, it appears that the detection of the GW signal from the short rise must have an event rate on the order of that for SN events in our galaxy, i. e. $\\sim$ 1 per 30-300 years. In principle such a rate has to be the same because the flinging of the pulsar could be just the aftermath of most supernova core collapses. However, in galactic globular clusters, where the frequency of formation and disruption of binary systems harboring compact stars is relatively high, the kick velocity of a pulsar partner at the binary disruption could be of the same magnitude as the one for galactic RAPs. Indeed, Vlemmings et al. \\cite{pulsars-vlemmings2005} are pursuing a VLBI astrometry program to perform measurements of pulsar proper motions and parallaxes with spatial resolution of microarsecond. As an illustration of the performance of the high precision pulsar astrometry system they simulated the detection of the disruption of a binary system in the stellar cluster Cygnus Superbubble. This is a quite estimulating perspective for pulsar astrometry in the near future.  Therefore, there is a potential to detect such a disruption or formation of a binary system through the GW signal emitted by the abrupt acceleration or sudden increase in velocity of a given pulsar. Observations in the range of micro- or nano-astrometry with planned satellites like GAIA and other instruments on Earth and space may help to identify such a phenomenon. An advanced technology useful for this purpose is currently operative or under planning (see Ref.\\cite{pulsars-gaia} and references therein).    \\begin{figure*} \\centering \\subfigure[\\label{k=a} Pulsar number (vertical axis) vs. distance (kpc) to Earth (horizontal axis) histogram ] { \\includegraphics[height=3.0in,width=3.0in]{atnf-histograma.eps}} \\hskip 0.5truecm \\subfigure[\\label{k=b} Pulsar number vs. Distance distribution of 212 pulsars ]{ \\includegraphics[height=3.0in,width=3.0in]{atnf-Dist-Vel.eps}} \\hskip 0.5truecm \\subfigure[\\label{PSR-VELOC-DIST} Velocity vs. Distance distribution of 212 ATNF pulsars having both parameters measured ] { \\includegraphics[height=2.750in,width=2.750in]{distrib_veloc.eps} } \\caption{Data taken from ATNF Pulsar Catalog} \\label{GW-HISTOGRAM-ATNF} \\end{figure*}   The physical reason for advocating in favor of this novel possibility for GW detection is based on a number of results on the dynamics and kinematics of galactic RAPs  and also on GR, as is shown next.  \\begin{itemize} \\item 1) In a seminal paper, Spruit  and Phinney \\cite{phinney-spruit98}  demonstrated that regardless of the actual physical engine driving the natal (instantaneous) pulsar kicks, that mechanism can only act upon the pulsar over a very restricted time scale, $\\Delta T^{\\rm kick}_{\\rm max}$, on the order of $\\simeq 0.32$~s. This is a fundamental piece of the dynamics of observed pulsars that seems to be overlooked by workers in the field.  \\item 2) As shown in the next Section, pulsar surveys have shown that the largest ($V^{\\rm max}_\\star$), smaller ($V^{\\rm min}_\\star$) and average ($V^{\\rm ave}_\\star$) spatial (3-D) velocity of cataloged pulsars are 5000 km s$^{-1}$, 84 km~s$^{-1}$, and 450-500 km~s$^{-1}$, respectively \\cite{lorimer93, arzoumian2002, Hobbs2005} (see Fig.\\ref{PSR-VELOC-DIST}). That is, $V_\\star$ changes by more than an order of magnitude.  \\item 3) Observations of NSs in binary systems have allowed to estimate their masses with very high accuracy. The  current values of this physical property gather around 1.4 M$_\\odot$ \\cite{chakra-thorsett}.  \\item 4) According to GR the acceleration of a mass ($M_\\star$) to a relativistic velocity ($V_\\star$); i. e. energy, should generate a GW pulse, the nature of which has a very distinctive imprint: the so-called gravitational wave memory \\cite{braginsky-thorne83,ori2001,sago04}, which makes the amplitude of this GW signal (the GW strain) not to fade away by the end of the traction phase (the kick), but rather leaves it variable in time as much as an energy source remains active in the source. \\end{itemize} ",
        "conclusions": "\\subsection{Computed rotation periods vs. Periods of ATNF sample of RAPs with $V=400-500$ km s$^{-1}$}  To the end, in order to make our analysis self-consistent we next discuss the relation between the translational speed and rotation frequency of ATNF pulsars. To do this, we select, as an example, the 3-D average pulsar velocity as determined by Hobbs et al. \\cite{Hobbs2005} (see Table-\\ref{periodos}). On this basis we determine the theoretical rotation periods that one could expect to find for that velocity sample if is factual that most of the explosion energy, as inferred from the momentum conservation law in Eq.(\\ref{momento}), goes into the pulsar rotation as discussed above. (Of course we compare with the actual periods measured by the ATNF surveys). Hence the torque applied to the just-born neutron star reads  \\begin{equation} \\tau \\sim \\frac {V}{2E_{k}} \\frac{\\frac{2}{5} M_{NS}R_{NS}^{2}}{d_{l}} \\end{equation} \\noindent  where $E_{k}$ is the kinetic energy of each velocity chosen for the present analysis, $V$ pulsar velocity, M$_{NS}=1.4$ M$_{\\odot}$, R$_{NS} = 20$ km and d$_{l} = 25$ km are the neutron star mass, radius and torque lever. Thus, for R$_{NS} = 20$ km and velocity $V=350$ km s$^{-1}$ we obtain p$ = 0.0182$ s, while for $V=433$ km s$^{-1}$ one gets p$=0.0182$ s. Meanwhile, for R$_{NS}=30$ km and $V=350$ km s$^{-1}$ we have p$=0.0411$ s, while for $V=433$ km s$^{-1}$ one gets p$=0.0332$ s.  One can then compare these results with those presented in Table-\\ref{periodos}. It appears that the theoretical prediction is about an order of magnitude smaller (much shorter periods) than the observed, and measured by the ATNF surveys. This discrepancy may suggest that \\begin{itemize} \\item 1) perhaps a large part of the initial rotational energy is dissipated by processes as differential rotation or magnetic braking, especially if processes as the magneto-rotational instability\\cite{ozernoy-somov-mri} indeed drives a large class of supernova explosions \\cite{MAGN-ROTAT-INSTAB}. \\item 2) it may also happen that convection and neutron fingers also becomes highly effective dissipative processes, thus reducing the pulsar initial spin to the level that we are currently finding in pulsar surveys. \\end{itemize}  \\begin{figure*}[tbh] \\centering \\subfigure[\\label{L1eff1=0.4a} Efficiency case $\\epsilon = 0.40$!! for parameters $a = 10^6$ km s$^{-2}$ and $V_{\\rm kick} \\sim$ 10$^4$ km s$^{-1}$, calibrated against LIGO I. ]{ \\includegraphics[height=3.0in,width=3.0in]{L1_a6_e4_nuevo.eps} } \\hskip 0.5truecm \\subfigure[\\label{L1eff1=0.4b} Efficiency case $\\epsilon = 0.10$!! for parameters $a = 10^6$ km s$^{-2}$ and $V_{\\rm kick} \\sim$ 10$^4$ km s$^{-1}$, calibrated against LIGO I. ]{ \\includegraphics[height=3.0in,width=3.0in]{L1_a6_e1_nuevo.eps} } {\\vskip 1.0truecm \\subfigure[\\label{L1eff1=0.4c} Efficiency case $\\epsilon = 0.05$!! for parameters $a = 10^6$ km s$^{-2}$ and $V_{\\rm kick} \\sim$ 10$^4$ km s$^{-1}$, calibrated against LIGO I. ]{ \\includegraphics[height=3.0in,width=3.0in]{L1_a6_e05_nuevo.eps} } \\hskip 0.5truecm \\subfigure[\\label{L1eff1=0.4d} Efficiency case $\\epsilon = 0.01$!! for parameters $a = 10^6$ km s$^{-2}$ and $V_{\\rm kick} \\sim$ 10$^4$ km s$^{-1}$, calibrated against LIGO I. ]{ \\includegraphics[height=3.0in,width=3.0in]{L1_a6_e01_nuevo.eps} } } \\caption{Pulsar GW signals for parameters: $a = 10^6$ km s$^{-2}$ and $V_{\\rm kick} \\sim$ 10$^4$ km s$^{-1}$} \\label{EFF=CALIB-LIGO-I} \\end{figure*}  In summary, from astronomical statistics the number of observed Galactic RAPs is definitely much larger than the corresponding to relativistic NS-NS binaries. The typical lifetime of canonical pulsars is $\\sim 10$ Myr, which is defined as the duration of the radio emitting phase for objects with $B \\sim 10^{12}$ G, or the time needed for the pulsar to emit less radiation or being entirely turned off. Hence, with a statistical Galactic rate of SNe, and the likely fling out rate of just-born pulsars, of about 1 per every 30-300 years, one expects that the GW signal of the kick at birth be observable as aftermath of the SN core collapse. In other words, the RAP GW signal may become a benchmark for looking for GW signals from the SN gravitational implosion. Besides, if the GW signal from the kick to a RAP were detected, for instance, from a source inside a globular, or from either an open star cluster \\cite{pulsars-vlemmings2005}, or a binary system harboring a compact star, such an observation might turn these RAPs the most compelling evidence for the existence of Einstein's gravitational waves.  We conclude by stating that the detection with advanced LIGO-type interferometers of individual GW signals from the short rise fling of Galactic RAPs is indeed possible.  \\begin{table} \\caption{Selected sample from the ATNF catalog of pulsars having speeds between 350-500 km s$^{-1}$, and their corresponding periods.} \\begin{center} \\begin{tabular}{llll} \\hline \\hline $\\#$ & JNAME & PERIOD & VTRANS  \\\\ { }&{ }& [s] & [km s$^{-1}$] \\\\ \\hline\\hline 1  & \\text{J0255-5304}  &   0.447708  &  381.648\\\\ 2  & \\text{J0538+2817}  &   0.143158  &  401.682\\\\ 3  & \\text{J0543+2329}  &   0.245975  &  377.151\\\\ 4  & \\text{J0837-4135}  &   0.751624  &  364.772\\\\ 5  & \\text{J1300+1240}  &   0.006219  &  350.637\\\\ 6  & \\text{J1645-0317}  &   0.387690  &  417.022\\\\ 7  & \\text{J1803-2137}  &   0.133617  &  351.252\\\\ 8  & \\text{J1902+0615}  &   0.673500  &  368.086\\\\ 9  & \\text{J1941-2602}  &   0.402858  &  351.027\\\\ 10 & \\text{J2048-1616}  &   1.961572  &  355.328\\\\ 11 & \\text{J2113+2754}  &   1.202852  &  386.791\\\\ 12 & \\text{J2219+4754}  &   0.538469  &  375.304\\\\ 13 & \\text{J2305+3100}  &   1.575886  &  373.546\\\\ 14 & \\text{J2326+6113}  &   0.233652  &  433.576\\\\ 15 & \\text{J2354+6155}  &   0.944784  &  357.846\\\\ \\hline \\hline \\label{periodos} \\end{tabular} \\end{center} \\end{table}"
    },
    "0711/0711.3793_arXiv.txt": {
        "abstract": "Large scale structure deflects cosmic microwave background (CMB) photons. Since large angular scales in the large scale structure contribute significantly to the gravitational lensing effect, a realistic simulation of CMB lensing requires a sufficiently large sky area. We describe simulations that include these effects, and present both effective and multiple plane ray-tracing versions of the algorithm, which employs spherical harmonic space and does not use the flat sky approximation. We simulate lensed CMB maps with an angular resolution of $\\sim 0.\\arcmin9$. The angular power spectrum of the simulated sky agrees well with analytical predictions.  Maps generated in this manner are a useful tool for the analysis and interpretation of upcoming  CMB experiments such as PLANCK and ACT. ",
        "introduction": "While the current generation of CMB experiments have  had a significant impact on cosmology by helping to establish a standard paradigm for cosmology \\citep{2003ApJS..148..175S, 2007ApJS..170..377S}, the upcoming generation of CMB experiments still has the potential to provide novel new insights into cosmology. PLANCK\\footnote{\\url{http://www.rssd.esa.int/index.php?project=Planck}} and ground based experiments, such as the Atacama Cosmology Telescope (ACT)\\footnote{\\url{http://www.physics.princeton.edu/act}}, will be mapping the CMB sky with significantly higher angular resolution than ever before. Secondary anisotropies on small angular scales encode important information about the late time interaction of CMB photons with structure in the Universe. One of the most basic of these interactions is the gravitational effect of the large scale structure potentials deflecting the paths of the photons, an effect justifiably referred to as the Gravitational Lensing of the CMB.  The effect of gravitational lensing can be thought of as a remapping of the unlensed CMB field by a line-of-sight averaged deflection field \\citep[for a recent review, see][]{2006PhR...429....1L}. Therefore, lensing does not change the one-point properties of the CMB. However, it does modify the two and higher-point statistics, and generates non-Gaussianity \\citep{1996ApJ...463....1S,1999PhRvD..59l3507Z,2000PhRvD..62f3510Z}. Although the typical deflection suffered by a CMB photon during its cosmic journey is about three arcminutes, the deflections themselves are coherent over several degrees, which is\tcomparable to the typical size of the acoustic features on the CMB. Thus lensing causes coherent distortions of the hot and cold spots on the CMB, and thereby broadens their size distribution. This leads to redistribution of power among the acoustic scales in the CMB, and shows up in the two-point statistics as a  smoothing of the acoustic peaks. At smaller scales, where the primordial CMB is well approximated by a {local} gradient, deflectors of small angular size produce small-scale distortions in the CMB, thereby transferring power from large scales in the CMB to the higher multipoles. Also, although the primordial CMB can be safely assumed to be a Gaussian random field \\citep{2003ApJS..148..119K}, and the large scale lensing potential can also be well approximated by  a Gaussian random field, the lensed CMB--- being a reprocessing of one Gaussian random field by another--- is itself not Gaussian. The effect of lensing on the power spectrum of the CMB is important enough that it should be taken into account while deriving parameter constraints with future higher resolution experiments. But what is even more interesting is that the non-Gaussianity in the lensed CMB field should enable us to extract information about the projected large scale structure potential, and thereby constrain the late time evolution of the Universe and Dark Energy properties. Therein lies the main motivation of studying this effect in utmost detail. Progress in this area has been slow.  Measurements of the CMB precise enough to enable a detection of weak lensing were not available in the pre-WMAP era. Also, picking out non-Gaussian signatures in the measured CMB sky by itself is extremely difficult, due to confusion from systematics, foregrounds, and limited angular resolution.  Rather than looking at signatures of lensing only in the CMB, one can also measure to what extent the deflection field estimated from the CMB correlates with tracers of the large scale structure which contributed to the lensing. It is easily realized that this approach is powerful \\citep{2000ApJ...540..605P} because many of the systematics disappear upon cross-correlating data sets. This approach was taken in recent years by \\cite{2004PhRvD..70j3501H} and \\cite{2007PhRvD..76d3510S}, using WMAP 1-year and 3-year data respectively. The former work looked at the cross correlation with SDSS luminous red galaxies (LRG), while the latter used the NRAO-VLA Sky Survey (NVSS) radio sources as their large scale structure tracers. As the lensing efficiency for the CMB is highest between redshifts of one and four, higher redshift tracers should show  greater cross correlation signal, which makes the NVSS radio sources better tracers for such study; \\cite{2007PhRvD..76d3510S} report a $3.4\\sigma$ detection. An independent analysis by \\citet{2008arXiv0801.0644H} looking for this effect in the WMAP 3-year data in cross correlation with SDSS LRG+QSO and NVSS sources find this signal at the $2.5\\sigma$ level.  With these pioneering efforts and with higher resolution CMB data from experiments such as ACT, PLANCK and the  South Pole Telescope (SPT)\\footnote{http://spt.uchicago.edu} on the horizon, we are entering an era where robust detection and characterization of this effect will become a reality. Also, with upcoming and proposed large scale structure projects (LSST\\footnote{\\url{http://www.lsst.org/lsst\\_home.shtml}}, SNAP\\footnote{\\url{http://snap.lbl.gov/}}, ADEPT\\footnote{\\url{http://universe.nasa.gov/program/probes/adept.html}}, DESTINY\\footnote{\\url{http://destiny.asu.edu/}}, etc.) there will in future be many more datasets to cross-correlate with the CMB.  One of the immediate results of such cross-correlation studies will be a measurement of the bias of the tracer population. Because such cross correlations tie together early universe physics from the CMB and late time evolution from large scale structure, they will also be sensitive to Dark Energy parameters \\citep{2006ApJ...650L..13H} and neutrino properties \\citep{2006PhRvD..74l3002S,2006PhRvD..73d5021L}, and can potentially break several parameter degeneracies in the primordial CMB (M.~Santolini,  S.~Das and D.~N.~Spergel, in preparation). Combination of galaxy or cluster lensing of the CMB with shear measurements from weak lensing of galaxies can also provide important constraints on the geometry of the Universe \\citep[S.~Das and D.~ N.~Spergel, in prep;][]{2007arXiv0708.4391H}. Again, with high enough precision of CMB data, it is possible to estimate, using quadratic \\citep{2003PhRvD..67h3002O}\tor maximum likelihood \\citep{2003PhRvD..67d3001H} estimators, the deflection field that caused the lensing. Such estimates can be turned into strong constraints of the power spectrum of the projected lensing potential \\citep{2002ApJ...574..566H}, which is also sensitive to the details of growth of structure. The estimated potential from the lensed CMB alone, or the potential estimated from weak lensing surveys \\citep{2007arXiv0710.2538M}, can be also used to significantly de-lens the CMB. This is particularly important in the detection of primordial tensor modes via measurements of CMB polarization. This is because (even though detection of the so-called B modes in CMB polarization is hailed as the definitive indicator of the presence of gravitational waves from the inflationary era) these mode can be potentially contaminated by the conversion of E-modes into B-modes via gravitational lensing. De-lensing provides a way of cleaning these contaminating B-modes produced by lensing and thereby probing the true gravitational wave signature.  In this paper, we describe a method for simulating the gravitational lensing of the CMB temperature field on a large area of the sky using a high resolution Tree-Particle-Mesh \\citep[TPM;][]{2000ApJS..128..561B,2003ApJS..145....1B} simulation of large scale structure to produce the lensing potential. The reason for considering a large area of the sky is twofold.\tFirst, the deflection field has most of its power on large scales (the power spectrum of the deflection field peaks at $\\ell\\sim 50$ in the best-fit cosmological model), and much of the power redistribution in the acoustic peaks of the CMB occurs via coupling of modes in the CMB with these large coherent modes in the deflection field. A large sky allows for several such modes to be realized. It is estimated that a small (flat) sky simulation that misses these modes would typically underestimate the lensing effect by about $10\\%$ in the acoustic regime, and more in the damping tail \\citep{2000PhRvD..62d3007H}. Second, one of the major goals of simulations such as this is to produce mock observations for upcoming CMB experiments.  PLANCK is an all-sky experiment, and many of the future CMB experiments (including ACT and SPT) will observe relatively large patches of the sky. Therefore, simulating CMB fields on a large area  of the sky is a necessity.  This method fully takes into account the curvature of the sky.  Although presented here for a polar cap like area, it can be trivially extended to the full sky.   The value of a simulation as described here is multifaceted, particularly in the development of algorithms for detection and characterization of the CMB lensing effect for a specific experiment. Since each experiment has a unique scanning mode, beam pattern, area coverage, and foregrounds, operations and optimizations performed on the data to extract the lensing information will have to be tailored to the specific experiment.  A large-sky lensed CMB map acting as an input for a telescope simulator provides the flexibility of exploring various observing strategies, and also allows for superposition of known foregrounds.  Another important aspect of this simulation is that the halos identified in the large scale structure simulation can be populated with different tracers of interest. Also, other signals, such as the Thermal and Kinetic Sunyaev Zel'dovich\teffects and {weak lensing of galaxies by large scale structure}, can be simulated using the same large scale structure. This opens up the possibility of studying the  cross-correlation of the CMB lensing signal with various indicators of mass, and thereby predicting the level of scientific impact that a specific combination of experiments can have.  As noted in \\cite{2005PhRvD..71h3008L}, the exact simulation of the lensed CMB sky, which requires the computation of spin spherical harmonics on an irregular grid defined by the original positions of the photons on the CMB surface, is  computationally expensive and requires robust parallelization. \\cite{2005PhRvD..71h3008L} suggested an alternative in which one would resample an unlensed CMB sky, generated with finer pixelation, at these unlensed positions. This method was implemented in the publicly available LensPix\\footnote{http://cosmologist.info/lenspix} code that was based on \\cite{2005PhRvD..71h3008L}. However, producing a high resolution lensed map requires a much higher resolution unlensed map, the generation of which becomes computationally more expensive as resolution increases. Here we put forward another alternative, in which we do the resampling with a combination of fast spherical harmonic transform on a regular grid followed by a high order polynomial interpolation. This interpolation scheme has been adapted from \\cite{2004PhRvD..70j3501H}, and is called the Non-Isolatitude Spherical Harmonic Transform (NISHT). This method is accurate as well as fast, and does not require parallelization or production of maps at a higher resolution. Another added advantage of this method is that the same algorithm can be used to generate the gradient of a scalar field on an irregular grid. Since the deflection field is a gradient of the lensing potential, this opens up the possibility of performing  a multiple plane ray tracing simulation. This is because the rays, as they propagate from one plane to another, end up on irregular grids, so the deflection fields on the subsequent planes have to be evaluated on irregular grids. At the time of the development of this project,  LensPix did not include an interpolation scheme, and used the methods as described originally in that paper. Concurrently with the completing of the current work, an interpolation scheme \\citep{232854}  different from the one described here has been added to that code. Another notable difference of our results with LensPix, is that while the latter uses a Gaussian Random realization of the deflection field, we have used a large scale structure simulation to produce the same, thereby including all higher order correlations due to non-linearities.  The paper is laid out as follows. In \\S\\ref{Algorithm} we explain the lensing algorithm,\tdescribing the governing equations in \\S\\ref{Equations} and their discretization in \\S\\ref{Discretization}. Then we discuss the effective lensing approach (\\S\\ref{Effective_Lensing}) as well as the multiple plane ray tracing approach (\\S\\ref{Multiple_Plane}). At the heart of the lensing algorithm lies the non-isolatitude spherical harmonic transform algorithm adapted from \\cite{2004PhRvD..70j3501H}, which is reproduced in some detail for completeness in \\S\\ref{Interpolation}. As discussed earlier, we have employed a light cone $N$-body  simulation and adopted a special polar cap like geometry for generating the lensing planes (\\S\\ref{Plane_Generation}). For comparison of the simulated fields with theoretical prediction, we compute the angular power spectra on the polar cap window; in \\S\\ref{Angular_Power_Spectra} we describe some of the subtleties involved in computing the power spectra. We present our results in \\S\\ref{Results} and describe the tests that we have performed in \\S\\ref{Tests}. Conclusions are presented in \\S\\ref{Conclusions}. ",
        "conclusions": "In this paper, we have put forward an algorithm for end-to-end simulation of the gravitational lensing of the cosmic microwave background, starting with an $N$-body simulation and fully taking into account the curvature of the sky. The method is applicable to maps of any geometry on the surface of the sphere, including the whole sky.  Our algorithm includes prescriptions for generating spherical convergence planes from an N-body light cone and subsequently ray-tracing through the planes to simulate lensing. The central feature of the algorithm is the use of a highly accurate interpolation method that enables sampling of both the deflection fields on intermediate lensing planes and the unlensed CMB map on an irregular grid. We have provided a detailed description of both a multiple plane ray tracing and an effective lensing version of the algorithm. The latter setting has been used to illustrate the algorithm, by generating an $\\sim$1$\\arcmin$ resolution lensed CMB map. We have compared the power spectra of the effective convergence map and the lensed CMB map with theoretical predictions, and have obtained good agreement. After this paper was completed, \\citet{2007arXiv0711.1540F} described a similar method of producing convergence maps in spherical geometry, and \\cite{2007arXiv0711.2655C} also described their techniques for simulating CMB lensing maps. The latter used broadly similar techniques to those described here, although they used a different method to obtain the deflection field.  Applications of the algorithm can be manifold.  The associated large scale structure planes can be populated with tracers of mass and foreground sources, in order to simulate cross-correlation studies and to investigate the effects of contamination. This lensing portion of the algorithm can be applied to generate lensed maps in large scale structure simulations that produce spherical shells \\citep{2007arXiv0711.1540F}. The multiple plane algorithm can be particularly useful, after trivial modifications, in simulating weak lensing of galaxies or the 21-cm background on a large sky. The lensed CMB maps can be used as inputs to telescope simulators for projects such as ACT and PLANCK, and will help in the analysis and interpretation of data. We intend to release all-sky high resolution lensed CMB maps made using this algorithm in near future."
    },
    "0711/0711.1105_arXiv.txt": {
        "abstract": "We present a new and completely general technique for calculating the fine-grained phase-space structure of dark matter throughout the Galactic halo.  Our goal is to understand this structure on the scales relevant for direct and indirect detection experiments.  Our method is based on evaluating the geodesic deviation equation along the trajectories of individual DM particles. It requires no assumptions about the symmetry or stationarity of the halo formation process. In this paper we study general static potentials which exhibit more complex behaviour than the separable potentials studied previously. For ellipsoidal logarithmic potentials with a core, phase mixing is sensitive to the resonance structure, as indicated by the number of independent orbital frequencies. Regions of chaotic mixing can be identified by the very rapid decrease in the real space density of the associated dark matter streams. We also study the evolution of stream density in ellipsoidal NFW halos with radially varying isopotential shape, showing that if such a model is applied to the Galactic halo, at least $10^5$ streams are expected near the Sun.  The most novel aspect of our approach is that general non-static systems can be studied through implementation in a cosmological N-body code. Such an implementation allows a robust and accurate evaluation of the enhancements in annihilation radiation due to fine-scale structure such as caustics.  We embed the scheme in the current state-of-the-art code GADGET-3 and present tests which demonstrate that N-body discreteness effects can be kept under control in realistic configurations. ",
        "introduction": "\\label{sect:intro}  Dark Matter (DM) is still to be directly detected. The first indirect indications of its existence came in the 1930s, with measurements of the velocities of galaxies in clusters. The cluster mass required to gravitationally bind the galaxies was found to be more than an order of magnitude larger than the sum of the luminous masses of the individual galaxies \\citep{1933AcHPh...6..110Z,1936ApJ....83...23S}. The early detection of large amounts of unseen matter associated with the Local Group \\citep{1959ApJ...130..705K} was followed in the 1970s by observations of the rotation curves of spiral galaxies which showed that these are flat, or even rising, at distances far beyond their stellar components \\citep{1970ApJ...159..379R,1979ARA&A..17..135F,1980ApJ...238..471R}. Studies of satellite systems suggested that the mass distributions of most galaxies might be an order of magnitude larger and more massive than their visible parts \\citep{1974ApJ...193L...1O,1974Natur.250..309E}.  All these discoveries led to the conclusion that a large fraction of mass in the Universe is dark. This has also been supported by recent gravitational lensing studies that demonstrate the existence of extended massive DM halos (e.g. Mandelbaum et al. 2006).  As the dominant mass component of galaxies and large-scale structures, DM has necessarily become a key ingredient in theories of cosmic structure formation. The most successful of these theories is the hierarchical paradigm.  In the current version of this model, the DM is composed of nonbaryonic elementary particles known as Cold Dark Matter (CDM) \\citep{1982ApJ...263L...1P}.  The term ``cold\" derives from the fact that the DM particles had negligible thermal motions at the time of matter-radiation equality. Their abundance was set when the interaction rate became too small for them to remain in thermal equilibrium with other species in the expanding Universe. The first objects to form in a CDM Universe are small galaxies. They then merge and accrete to give rise to the larger structures we observe today. Thus structure formation occurs in a ``bottom-up\" fashion \\citep{1984Natur.311..517B,1985ApJ...292..371D,2006Natur.440.1137S}.  The crucial test of this paradigm undoubtedly consists in the determination of the nature of DM through direct detection experiments. Among the most promising candidates from the particle physics perspective are axions and neutralinos.  Axions were introduced to explain the absence in Nature of strong-CP (Charge conjugation and Parity) violations \\citep{1977PhRvL..38.1440P}. The cosmological population formed out of equilibrium as a zero momentum Bose condensate.  They can be detected through their conversion to photons in the presence of a strong magnetic field, (e.g. Ogawa et al. 1996; Hagmann et al. 1998). Neutralinos are the lightest supersymmetric particles, and may be considered as a particular form of weakly interacting massive particles (WIMPs) \\citep{1985NuPhB.253..375S}. The most important direct detection process for neutralinos is through elastic scattering on nuclei.  Today many experiments are searching for these particles, \\citep{2004PhRvL..93u1301A, 2005PhRvD..71l2002S, 2006astro.ph.12565S, 2007JPhCS..60...58A, 2007arXiv0705.3345S}.  The main challenge lies in the large background they encounter.  Having an idea what the detector signal might look like can help substantially in fine-tuning the experiments in order to increase the chance of a detection. In addition, many experiments are attempting to detect WIMPs indirectly by searching for $\\gamma$-ray emission from their annihilation \\citep{2005NewAR..49..213D,2005PhRvL..95t9001D,2006PhRvD..73f3510B,2007JCAP...06...13H}. Predictions for this radiation are currently uncertain because very substantial enhancements are possible, at least in principle, from fine-scale structure such as caustics in the dark matter distribution \\citep{2001ApJ...561...35D,2007JCAP...05...15M}.  The differential detection rate for WIMPs in a given material (per unit detector mass) is sensitive to the local velocity distribution: \\begin{equation} \\frac{\\mathrm{d}R}{\\mathrm{d}Q} = \\frac{\\rho}{M_\\chi} \\int d^3 v f(\\mathbf{v}) v \\frac{\\mathrm{d}\\sigma}{\\mathrm{d}Q} \\label{eq:DIFF_WIMP_DET_RATE} \\end{equation} where $Q$ is the energy deposited in the detector and $\\mathrm{d}\\sigma/\\mathrm{d}Q$ is the differential cross section for WIMP elastic scattering with the target nucleus. $M_\\chi$ is the WIMP mass which lies in the GeV to TeV range for currently preferred models. $\\rho$ is the local mass density and $f(\\mathbf{v})$ the velocity distribution of WIMPs that reach the detector.  From Eq. (\\ref{eq:DIFF_WIMP_DET_RATE}) it is evident that the count rate in a direction-insensitive experiment depends on the velocity distribution of the incident particles and will be modulated by the orbital motion of the Earth around the Sun \\citep{1986PhRvD..33.3495D}. In most studies, an isotropic Maxwellian distribution relative to the galactic halo has been assumed, (e.g. Freese et al. 1988), although there are other models in the literature, discussing, for example, multivariate Gaussians \\citep{2000MNRAS.318.1131E}.  Some attempts at understanding the effect of fine-scale structure in the velocity distribution of DM particles have also been made \\citep{1998PhLB..432..139S,2001PhRvD..64h3516S,2001PhRvD..64f3515H}.  A significant signal could come from what are known as streams of DM \\citep{1995PhRvL..75.2911S,2002PhRvD..66f3502H,2005PhRvD..72h3513N}. Both axions and WIMPs are cold.  In the absence of clustering their present day velocity dispersion would be negligible ($\\delta v \\sim 10^{-10}c$ for WIMPs and $\\delta v \\sim 10^{-17}c$ for axions).  They are effectively restricted to a 3D hypersurface, a sheet in 6D phase-space.  The growth of structure results in continual stretching and folding in phase-space of this initially almost uniform sheet. This process is called mixing. The more strongly the system mixes, the more streams of DM particles are present at a given location in configuration-space. Mixing stretches each sheet and so its density decreases.  The existence of distinct streams is a direct consequence of the collisionless character and the coldness of CDM.  The situation is sketched in Fig. \\ref{fig:FLOW}. At the points where the number of streams changes by two, the local configuration-space density of dark matter becomes extremely high. These are caustics of the kind studied by catastrophe theory \\citep{0471050644,1999MNRAS.307..877T}.  Note that Liouville's Theorem prevents the CDM phase-space sheet from ever tearing, although it can be arbitrarily strongly stretched.   \\begin{figure} \\centerline{\\includegraphics[width=1.0\\linewidth]{fig1.eps}} \\caption{ \\label{fig:FLOW} Sketch of an idealised CDM phase-space sheet in the $x, \\dot{x}$ plane. The thickness of the line represents the local velocity dispersion within each stream. The small wiggles correspond to initial density perturbations and the multi-valued region reflects the multiple streams created by winding in non-linear regions.  Depending on $x$ position an observer sees one or three streams. At points where the number of streams changes by two, a caustic with a very high DM density is present. Such caustics may be significant for the total annihilation flux. The number of streams at each point is a measure of the local amount of mixing. The dashed line represents the Hubble Flow. The cross marks the phase-space coordinates of a particular CDM particle embedded in the flow.  } \\end{figure}  If the dark matter density in the solar neighbourhood is dominated by a small number of streams, its velocity distribution will effectively consist of a few discrete values, one for each stream.  If, on the other hand, the number of streams is very large, the velocity distribution will be smooth and individual streams will be undetectable. This issue has so far been addressed only under simplified conditions and divorced from its proper cosmological context \\citep{2002PhRvD..66f3502H,2005PhRvD..72h3513N}.  This is largely because the only tool capable of studying cosmological structure formation in full generality, namely N-body simulations, cannot resolve the relevant scales. For example, \\cite{2005PhRvD..72h3513N} estimate that of the order of $10^{12}$ particles would be required in a simulation of the Milky Way's halo to begin to resolve the streams in the solar neighbourhood.  Even the largest simulation so far published, the Via Lactea model \\citep{2007ApJ...657..262D}, is 4 orders of magnitude short of this minimum requirement. It is thus impossible to figure out the number of streams near the Sun and the properties of ``typical'' streams with current N-body capabilities.  A related issue that has been much discussed is whether caustics can affect direct or indirect detection experiments. Some authors claim that caustics play an important role as stable phenomena connected to any collapsing CDM halo and taking the form of relatively massive rings or shells \\citep{1999PhRvD..60f3501S, 2006PhRvD..73b3510N}.  Other authors have argued that caustics should be more abundant, weaker and dynamically negligible \\citep{2002PhRvD..66f3502H}. Even if caustics negligibly affect the gravitational potential, they may have very substantial effects on the annihilation rate of dark matter \\citep{2001PhRvD..64f3515H, 2001PhRvD..63h3515B, 2005JCAP...05..007P, 2006MNRAS.366.1217M, 2007JCAP...05...15M, 2007PhRvD..75l3514N}. All these papers were able to evaluate the enhancements due to caustics only under restrictive and unrealistic assumptions about symmetry, formation history, etc. While they demonstrate that large enhancement factors may be possible, they do not provide reliable estimates for the actual enhancement expected.  The method we present below is capable of providing such estimates for halos growing as predicted by the $\\Lambda$CDM model.  In this paper we will present a novel approach that directly analyses structure in the fine-grained phase-space distribution.  We circumvent the ``number of particle\" problem by solving the geodesic deviation equation (GDE) for every DM particle.  This allows us to calculate the local properties of the DM stream each simulation particle is embedded in, in particular, its configuration-space density and its local velocity distribution. The mass-weighted number of streams near any point is then the total local density divided by the mean density of the individual local streams. Caustic passages can be detected robustly from the properties of each particle's local stream, and the particle's instantaneous annihilation probability within this stream can be evaluated simply and integrated accurately through caustics.  This paper is the first of a series. Here we introduce our method and apply it to study the evolution of streams in static potentials that are too complex to be analysed using previous techniques.  We also implement our scheme in a state-of-the-art N-body simulation code, and use simple test problems to demonstrate that N-body discreteness effects can be kept under control in realistic applications. Later papers will address issues associated with mixing, caustics and annihilation radiation in the full cosmological context.  The outline of our paper is the following: In Section~2 we present a detailed derivation of the GDE and show how it can be used to quantify mixing, to locate caustics, and to calculate stream densities and annihilation rates. Section~3 describes our code, DaMaFlow, which is designed to solve the GDE for single orbits in a wide variety of potentials. In Section~4 we analyse static, separable potentials and compare results from our method to previous work. The fifth section applies our scheme to non-integrable, but still static potentials, revealing their complex phase-space structure.  In Section~6 we turn to more realistic non-spherical CDM halo potentials and discuss the influence of halo shape on stream density behaviour. Section~7 discusses the implementation of our method in an N-body code and presents results of simple tests of when discreteness effects compromise studies of stream densities and caustics. The final section summarises our results and gives some conclusions. ",
        "conclusions": "\\label{sect:discussion}  Direct DM detection experiments operate on length-scales far below the resolution of current structure-formation simulations.  The fine-grained phase-space structure on these scales will determine the signal they see. In addition, small-scale structure can substantially enhance the annihilation signal that is the target of current indirect detection experiments. A better understanding of such structure within the current concordance $\\Lambda$CDM cosmology is thus critical for analysing and interpreting all current DM searches.   We propose a new route to tackle these issues. Rather than improving simulations simply by increasing the number of particles, we attach additional information to each particle, namely a phase-space distortion tensor which allows us to follow the evolution of the fine-grained phase-space distribution in the immediate neighbourhood of the particle.  We introduce the Geodesic Deviation Equation (GDE) as a general tool for calculating the evolution of this distortion along any particle trajectory.  The projection from phase-space to configuration-space yields the density of the particular CDM stream that particle is embedded in and can also identify when the particle passes through a caustic.  This technique makes the fine-grained phase-space structure accessible. It enables studies of the phase-space structure of general non-integrable static potentials which reproduce all the results previously obtained using frequency analysis methods, identifying chaotic regions and finding substructure in regular regions in the form of resonances.  In addition, it can be used to quantify mixing rates and to locate caustics.  We demonstrate these capabilities for the complex phase-space structure of the ellipsoidal logarithmic potential with a core. All relevant phase-space regions could be identified by solving the GDE along the orbit. We have written a code, DaMaFlow, that allows us to carry out such stream density analyses for a wide variety of potentials in a very efficient way.  Stream density evolution is very sensitive to the shape of the underlying potential. We demonstrate this by comparing results for a realistic CDM halo with radially varying shape to those for a spherical halo with similar radial density profile. After 100 orbits the predicted stream densities in the inner regions differ by a factor of 100.  In general we expect the stream densities to decrease as $(t/t_{\\mathrm{orbital}})^{-3}$ for regular orbits and even faster for chaotic orbits, rather than as $(t/t_{\\mathrm{orbital}})^{-2}$, the result found for orbits in a spherical potential.  Scaling to the Milky Way leads us to estimate that there should be at least $10^5$ streams passing through the solar system.  The potentially revolutionary advantage of our approach, and our main reason for pursuing it, is that it applies equally well to non-symmetric, non-static situations of the kind that generically arise in CDM cosmologies. Indeed, it can be implemented in a relatively straightforward way in current state-of-the-art cosmological N-body codes.  We have carried out such an implementation in the {\\sm GADGET} code and have presented some tests based on equilibrium Hernquist models.  The N-body implementation is able to conserve 6-D phase-space density to high accuracy along individual particle orbits. In addition, it qualitatively reproduces the results found in the corresponding smooth potential for the evolution of stream density along individual orbits, and it reproduces the statistical results found for ensembles of orbits to impressive accuracy. The identification of caustic passages is particularly robust, showing very little dependence on N-body parameters such as particle number and softening. Thus discreteness effects appear to be well under control, at least for the large N systems studied here. The remarkably robust identification of caustic properties makes us optimistic that we will be able to calculate annihilation boost factors due to caustics in fully realistic $\\Lambda$CDM dark halos.  In future applications we will use these techniques to address mixing and DM detection issues within fully general simulations of the $\\Lambda$CDM structure formation model."
    },
    "0711/0711.2106_arXiv.txt": {
        "abstract": "We employ a two-dimensional grid-based hydrodynamic model to simulate upper atmospheric dynamics on extrasolar giant planets.  The hydrodynamic equations of motion are integrated on a rotating, irradiated sphere using a pseudospectral algorithm.  We use a two-frequency two-stream approximation of radiative transfer to model the temperature forcing.  This model is well-suited to simulate the dynamics of the atmospheres of planets with high orbital eccentricity that are subject to widely-varying irradiation conditions.  We identify six such planets, with eccentricities between $e=0.28$ and $e=0.93$ and semimajor axes ranging from $a=0.0508$ A.U. to $a=0.432$ A.U., as particularly interesting objects for study.   For each of these planets, we determine the temperature profile and resulting infrared light curves in the 8-$\\mu$m Spitzer bands.   Especially notable are the results for HD\\, 80606b, which has the largest eccentricity ($e=0.9321$) of any known planet, and HAT-P-2b, which transits its parent star, so that its physical properties are well-constrained.  Despite the variety of orbital parameters, the atmospheric dynamics of these eccentric planets display a number of interesting common properties.  In all cases, the atmospheric response is primarily driven by the intense irradiation at periastron.  The resulting expansion of heated air produces high-velocity turbulent flow, including long-lived circumpolar vortices.  Additionally, a superrotating acoustic front develops on some planets; the strength of this disturbance depends on both the eccentricity and the temperature gradient resulting from uneven heating.  The specifics of the resulting infrared light curves depend strongly on the orbital geometry.  We show, however, that the variations on HD 80606 b and HAT-P-2b should be readily detectable at  4.5 and 8 $\\mu$m using the Spitzer Space Telescope.  Indeed, these two objects present the most attractive observational targets of all known high-$e$ exoplanets. ",
        "introduction": "In July of 1994, the astronomical community was captivated by the collision of comet Shoemaker-Levy 9 and Jupiter.  The sudden deposition of energy into the Jovian atmosphere caused a dramatic response on a planetary scale, providing remarkable opportunities for observation and simulation \\citep{har01, dem01} and fostering the development of a deeper understanding of the structure and behavior of gas-giant atmospheres.   Within our own solar system, such an occurrence is a one-time event.  However, the discovery of extrasolar giant planets on orbits with substantial eccentricity affords the opportunity to study atmospheric dynamics on worlds where the sudden influx of a massive dose of energy occurs on a regular basis.  With the advent of Spitzer, direct observation of such planets has become feasible.  The Spitzer telescope has been an unexpectedly useful tool in the study of hot Jupiters, and has already been used to make a number of landmark observations of extrasolar planets. These began with the detections of the secondary eclipses of HD 209458 b \\citep{dem05} and TrES-1 \\citep{char05}, which confirmed that hot Jupiters have low albedos and day-side effective temperatures of order $T_{\\rm eff}\\sim1100\\,{\\rm K}$. The TrES-1 observations were carried out at 4.5$\\mu{\\rm m}$ and 8$\\mu{\\rm m}$, and thus provided an estimate of that planet's mid-infrared spectral slope. High S/N observations at 16 $\\mu{\\rm m}$ of the secondary transit of HD 189733 \\citep{dem06} produce a similar ($1117\\pm42\\,{\\rm K}$) day-side temperature.  More recently, Harrington et al. (2006) have detected strong orbital phase-dependent brightness variations for $\\upsilon$ And b at 24 $\\mu{\\rm m}$, which indicate a strong day-night temperature contrast in the far IR, and suggest that heat transfer between the hemispheres is inefficient at the level in the atmosphere that generates the observed $24\\mu{\\rm m}$ emission. Conversely, Cowan et al. (2007) found no significant orbital phase variations at 8 $\\mu{\\rm m}$ for 51\\,Peg b and HD\\,179949 b, indicating that for those two planets, surface flows {\\it are} effective at redistributing heat at the level probed by the 8 $\\mu{\\rm m}$ photosphere. Even more recently, Knutson et al. 2007 have presented results from a continuous 30-hour Spitzer observational campaign on HD 189733. The resulting 8$\\mu$m photometric time-series contains extraordinary (part in $10^{-4}$) signal-to-noise and clearly shows phase-dependant flux variations that indicate a day-night temperature contrast of order 200 K. Of particular interest in the Kuntson et al. (2007) time series is a $F/<F> \\sim 0.0005$ increase in planetary flux of $\\sim 5$hr duration just past the primary transit. This feature can be interpreted as arising from a hotspot near the dawn terminator of the planet.  In addition to providing infrared photometric time-series data, the Spitzer Space Telescope can also be used to obtain mid-infrared spectra of extrasolar planets. This has been done for the transiting planets HD\\,209458b (Richardson et al. 2007) and HD\\,189733b (Grillmair et al. 2007). In both cases, spectra in the 7-14$\\mu$m range were obtained by subtracting the stellar spectrum measured during secondary eclipse from the combined star+planet spectrum obtained just outside eclipse. Model atmospheres of extrasolar planets under strong insolation conditions (e.g. Sudarsky, Burrows \\& Hubeny 2003, Marley et al. 2006) predict that ${\\rm H_2O}$ absorption features would be prominent within the observed spectral range. Remarkably, however, neither planet displayed the predicted signature. This null result can be interpreted in several ways. First, it is possible that the atmospheres are dry, perhaps because available oxygen is sequestered in CO. This interpretation, however would require both planets to have unexpectedly large C to O ratios. Alternately, the effectively featureless spectrum could arise from a very clear atmosphere with an isothermal P-T profile down to large optical depth \\citep{for06}. This interpretation is favored by the fact that the hot Jupiters have extremely low albedos (Rowe et al. 2006). As a third possibility, \\citet{har07} suggest that HD\\,209458b may be displaying line emission at 10$\\mu$m. They interpret these features as possibly arising from silicate clouds. Any such clouds, however, would need to be highly non-reflective in order to be consistent with the low optical albedo.  At present, the existing observational results present an incomplete and somewhat contradictory overall picture. It is not understood how the wind vectors and temperature distributions on the observed planets behave as a function of pressure depth, and planetary longitude and latitude. Most importantly, the effective radiative time constant in the atmospheres of short-period planets remains unmeasured, and as a result, dynamical calculations of the expected planet-wide flow patterns (Cho et al. 2003, Cooper \\& Showman 2005, Burkert et al. 2005, Fortney et al. 2006, Langton \\& Laughlin 2007, Dobbs-Dixon \\& Lin 2007) have come to no consensus regarding how the surface flow should appear. The lack of agreement between the models stems in large part from the paucity of unambiguous measurements of the basic thermal structure of the atmospheres of short-period planets. Essentially, the problem comes down to the fact that the planets that have been observed so far have had their spin periods tidally synchronized to their orbits, so that each spot on the substellar surface sees a constant illumination from the parent star. The resulting flow pattern will therefore likely be in steady state, and the temperature distribution across the planetary surface will depend on both the radiative timescale as well as the windspeeds at optical depth unity.  In this paper, we discuss our modeling program for determining the hydrodynamic flows on the surfaces of planets traveling on significantly eccentric orbits.  Such planets are not spin-synchronous. Tidal effects cause the rotation rate to approximate the orbital angular velocity at periastron; the rotation is therefore pseudo-synchronous and the spin period is a sensitive function of the eccentricity \\citep{hut81}.  In contrast to the typical spin-synchronous hot Jupiter, on which one hemisphere is perpetually illuminated while the other receives no insolation, the entire surface of an eccentric planet receives radiation.  Pseudo-synchronization does, however ensure that a single hemisphere receives the preponderance of the illumination during the periastron passage.  Additionally, the intensity of the radiation received varies strongly over the course of an orbit.  The insolation of planets on eccentric orbits therefore exhibits significant time variation.  To our knowledge, other modeling efforts to date have focused exclusively on exoplanets that are in tidally synchronized orbits with a permanent day-night side asymmetry and essentially steady state irradiation conditions. With such bodies, it is difficult to disentangle the contributing effects of the effective radiative time constant in the planetary atmosphere, and the hydrodynamic flow velocities on the surface. In this paper, we propose that these ambiguities can be resolved by observing planets that have eccentric orbits, and which hence show different hemispheres (or partial hemispheres) to the star, and which experience a highly variable total insolation. We identify the highly eccentric planet HD 80606 b and the transiting planet HAT-P-2 b \\citep{bak07} as being a prime targets for such observations.  We propose that by monitoring these planets during their periastron passages, we can improve our overall understanding of the atmospheric properties of hot Jupiters. The planet-wide response to the rapid increase in stellar radiation will be readily observable in the 4.5$\\mu{\\rm m}$ and 8$\\mu{\\rm m}$ bands, and will provide a direct simultaneous measurement of the radiative time constant at two separate levels in the atmosphere.  The plan for this paper is as follows. In \\S 2, we describe our numerical model as applied to extrasolar planets. In \\S 3, we discuss the characteristics of the six planets considered in this paper, as well as relevant stellar properties of their parent stars.  In \\S 4, we present the results of this model as applied to these planets.   In \\S 5, we discuss the observational consequences of these results, in particular concentrating on the infrared light curves in the 8 micron Spitzer band.  We show that in general, the periastron passage produces a distinct flux peak, and that the shape of the overall curve can be used to constrain planetary characteristics such as rotation rate and atmospheric radiative time constant.  We conclude in \\S 6. ",
        "conclusions": "The advent of the Spitzer Space Telescope and the discovery of short-period planets on significantly eccentric orbits has presented astronomers with the opportunity to observe time-dependent weather patterns on extrasolar planets. Giant planets on eccentric orbits that experience strong variations in insolation have the potential to tell us a great deal about the structure, dynamics and chemistry of Jovian-like atmospheres. Eccentric planets can be profitably compared with short period planets on spin-synchronous circular orbits that are likely to have largely steady-state patterns of surface flow, and which are currently being observed with Spitzer.  Furthermore, when eccentricities are extreme -- as in the case of HD 80606 b  --  we are afforded the ability to directly observe planets that spend the bulk of their orbit in a region where the radiation environment is similar to that experienced on Earth. With current instruments, such planets would be completely beyond the possibility of direct observation.  Eccentric planets also have the best potential to give us an understanding of the dissipative mechanisms at work in planetary interiors. A planet such as HAT-P-2 b should be generating a very significant luminosity as a result of the strongly time-dependent tidal force on the planet. By observing the depth of the primary transit in the infrared, one can get a direct measurement of the planetary effective temperature, which can then be used to infer the amount of dissipation within the planet, and to put interesting constraints on the tidal quality factor, $Q$ associated with dynamical (as opposed to equilibrium) tides.  To take a first step toward understanding what to expect from the infrared light curves of short-period eccentric planets, we have created a two-dimensional hydrodynamic model with radiative forcing appropriate to the time-varying illumination received for a specified orbit. Our model handles radiation transfer in an approximate, but hopefully qualitatively correct manner by separating the transport of optical and IR radiation. We apply this model to six known giant planets on eccentric orbits with $e\\gtrsim0.3$. The irradiation intensity on these planets varies by a factor of three for the least eccentric planets up to a factor of 800 for HD 80606b. The effect of this tremendous variability on atmospheric flows is profound; the dominant feature in the evolution of each atmosphere studied is the subsolar heating through the perastron passage, with the subsequent flow being primarily driven by this dramatic forcing. For each planet, we find significant spatial and temporal variability in the state of the atmosphere, with turbulent flow being a common feature. Additionally, rapid increase in the thermal forcing leads to superrotating acoustic shocks, as seen in the atmospheres of HD 80606 b and HAT-P-2b.  Among the planets that we have modeled, HAT-P-2b and HD 80606b stand as the most promising candidates for time-series infrared observations in Spitzer's 8-$\\mu {\\rm m}$ band, and among these two planets, HAT-P-2 b is the best candidate. HAT-P-2b's secondary transit provides an opportunity for an overall flux calibration, and its orbital geometry (with $\\varpi=185^{\\circ}$) places the most interesting features of the light curve on a relatively short 35-hour segment within the full 5.6-day orbit. During this time span, we predict that the relative infrared flux will vary from $3.2\\times10^{-4}$ to $9.0\\times10^{-4}$.  HD 80606b, on the other hand, is predicted to produce a relative infrared flux ranging from $3.8 \\times 10^{-4}$ to $8.0 \\times 10^{-4}$. This variation easily satisfies Spitzer's sensitivity limit, and for this planet, the observational situation is greatly aided by HD 80607, the wide binary companion of near-identical spectral type to the parent star HD 80606. Billions of years ago, HD 80607 enabled Kozai migration. Today, it provides a stable photometric reference to gauge the rapid infrared flux rise from the planet during periastron passage. A measurement of the slope of this rise will allow a direct estimate of the atmospheric radiative time constant at the 4.5-$\\mu$ and 8-$\\mu$ photospheres.  A short-period Jovian planet on an eccentric orbit likely presents one of the Galaxy's most thrilling sights. One can imagine, for example, how HD 80606 b appears during the interval surrounding its hair-rising encounter with its parent star. The blast of periastron heating drives global shock waves that reverberate several times around the globe. From Earth's line of sight, the hours and days following periastron are characterized by a gradually dimming crescent of reflected starlight, accompanied by planet-wide vortical storms that fade like swirling embers as the planet recedes from the star. It's remarkable that we now have the ability to watch this scene (albeit at one-pixel and two-frequency resolution) from a vantage several hundred light years away.  Indeed, the observational and theoretical allure of short-period non-synchronous planets is worthy of much more detailed study than we have presented here. Our next steps on this problem, therefore, will be to (1) upgrade our radiative transfer scheme to allow many more frequency intervals to be modeled, and (2) incorporate the higher-resolution radiative transfer into a full three-dimensional hydrodynamic treatment."
    },
    "0711/0711.0759_arXiv.txt": {
        "abstract": "We consider the effects of eccentricity on the fragmentation of gravitationally unstable accretion disks, using numerical hydrodynamics.  We find that eccentricity does not affect the overall stability of the disk against fragmentation, but significantly alters the manner in which such fragments accrete gas.  Variable tidal forces around an eccentric orbit slow the accretion process, and suppress the formation of weakly-bound clumps.  The ``stellar'' mass function  resulting from the fragmentation of an eccentric disk is found to have a significantly higher characteristic mass than that from a corresponding circular disk.  We discuss our results in terms of the disk(s) of massive stars at $\\simeq0.1$pc from the Galactic Center, and find that the fragmentation of an eccentric accretion disk, due to gravitational instability, is a viable mechanism for the formation of these systems. ",
        "introduction": "\\label{sec:intro} The relative proximity of the Galactic Center (henceforth GC) provides a unique opportunity for ``close-up'' study of the processes that influence the formation of galaxies and black holes.  Recent advances in telescope technology have enabled the resolution of individual stars in the crowded GC environment, and these new data have produced several puzzles \\citep[e.g.,][]{ghez98,genzel03,ghez05,paumard06}.  In particular, the detection of large numbers of young, massive stars is strongly suggestive of recent star formation activity at or close to the GC, despite the fact that the extreme physical conditions at the GC (temperature, tidal shear, etc.) prohibit the formation of stars by the same mechanisms that form stars in the solar neighborhood.  The data show that a small number (few tens) of B-type stars exist very close to the GC (within $\\simeq0.01$pc; these are the so-called ``S-stars''), and a larger population ($>100$) of massive O, B and Wolf-Rayet (WR) stars exist at slightly larger radii ($\\simeq0.1$pc).  Many of this second population of stars are known to form a coherent ring or disk \\citep[e.g.,][]{genzel03} and, while the existence of a second such disk of stars remains controversial \\citep[e.g.,][]{paumard06,lu06}, the existence of at least one stellar disk seems secure.  A popular theory posits that the stellar disk(s) formed via fragmentation of an accretion disk due to gravitational instability \\citep[e.g.,][]{ps77,lb03,sb87,sb89,gt04,nayak06,levin07}.  However, we have recently demonstrated that the observed eccentricities of the stellar orbits \\citep[typically $\\gtrsim 0.3$,][]{bel06} are inconsistent with dynamical relaxation from a circular initial configuration \\citep{aba07}, unless the stellar mass function (MF) is extremely top-heavy, much more so than is inferred from observations \\citep{ns05,paumard06}.  We suggested that fragmentation of an eccentric accretion disk could explain this discrepancy, and here seek to investigate the effects of eccentricity on disk fragmentation.  In recent years gravitational instabilities in accretion disks have been the subject of much theoretical and numerical research, primarily in the context of protoplanetary disks and planet formation \\citep[see, e.g., the review by][]{durisen_ppv}.  Some details remain contested, but the basic physical processes are now well understood.  In order for a gravitationally unstable disk to fragment, two conditions must be satisfied.  Firstly, the disk must be sufficiently massive and/or cold that the \\citet{toomre64} criterion is satisfied: \\begin{equation} Q = \\frac{c_{\\mathrm s} \\Omega}{\\pi G \\Sigma} \\lesssim 1 \\end{equation} Here $c_{\\mathrm s}$ is the local sound speed, $\\Omega$ the local orbital frequency, and $\\Sigma$ the disk surface density.  This criterion ensures that the disk is gravitationally unstable, but in order for the instability to lead to fragmentation the disk must also cool rapidly, or else compressional heating will stabilize the disk against fragmentation.  \\citet{gammie01} used a local analysis to show that the cooling timescale of the disk must satisfy \\begin{equation} t_{\\mathrm {cool}} \\lesssim 3 \\Omega^{-1} \\, , \\end{equation} and subsequent work using global simulations has verified the validity of this criterion \\citep{rice03,lr04,lr05}.  The criterion also depends weakly on the adopted equation of state \\citep{gammie01,rice05}, and studies of protoplanetary disks have shown that this dependence can be non-trivial in real disks \\citep{boley07}.  In the case of an eccentric disk, the circularization energy provides an additional, and potentially important, source of heating, and this remains largely unstudied to date.  Theoretical models of disk fragmentation around black holes typically consider circular disks with near-equilibrium initial conditions, implicitly assuming that such a configuration results from some sort of ``accretion event'' \\citep[such as the capture of a molecular cloud, e.g.,][]{nayak06,nayak07,levin07}.  In the case of the GC, the so-called circumnuclear ring at a few pc from Sgr A$^*$ is thought to be the likely gas reservoir for such accretion events \\citep[][]{morris93,sanders98}, but the dynamics of the infall process is not well understood.  Moreover, some models of cloud infall predict the formation of coherent, eccentric disks around the central black hole \\citep[e.g.,][]{sanders98}.  We expect such disks to be gravitationally unstable, and star formation via fragmentation of an eccentric disk would provide an elegant solution to the problem of the stellar orbits discussed in \\citet{aba07}, although it may leave some issues, such as the $z$-distribution of the orbits, unresolved.  In this paper, therefore, we seek to understand the dynamics of gravitational instability and fragmentation in eccentric disks.  In Section \\ref{sec:sims} we describe our numerical simulations, and present the results in Section \\ref{sec:res}.  We discuss the implications of our results for models of star formation at the GC, as well as the limitations of our analysis, in Section \\ref{sec:dis}, and summarize our conclusions in Section \\ref{sec:summary}. ",
        "conclusions": "\\label{sec:dis}  \\begin{figure} \\includegraphics[angle=270,width=\\hsize]{f7.ps} \\caption{Comparison of clump MFs in the circular (black) and eccentric (red) disks.  As in Fig.\\ref{fig:MF_circ}, the dotted line shows the resolution limit, and the dashed line the Salpeter slope.  The MF for the eccentric disk is measured at the end of the simulation; in the circular case it is measured where the same fraction of the disk mass has been accreted into bound clumps.  The eccentric disk clearly forms fewer low-mass objects than the circular disk, and a KS test rejects the possibility that the two MFs are drawn from the same underlying distribution.}\\label{fig:MF_comp} \\end{figure}   While our simulations are scale-free, they have obvious applications to the formation of the stellar disk(s) observed at the GC.  Observations of the GC stellar disks find that the stellar MF must be much more ``top-heavy'' than the standard Salpeter MF.  The relatively low X-ray flux from the GC region suggests a significant deficit of low-mass ($\\lesssim 5$M$_{\\odot}$) stars compared to local star-forming regions \\citep{ns05}, and estimates of the stellar MF from the $K$-band luminosity function suggests a MF slope that is 1.0--1.5 dex flatter than the Salpeter slope \\citep{paumard06}.  Also, given that the total stellar mass in the disk(s) is $\\sim10^4$--$10^5$M$_{\\odot}$ \\citep[e.g.,][]{ndcg06,paumard06}, the large number ($\\gtrsim 100$) of observed massive O- \\& WR-stars suggests a characteristic stellar mass of $\\simeq 10$--50M$_{\\odot}$, much larger than that observed in the solar neighbourhood.  This is consistent with our results \\citep[and those of previous studies such as][]{nayak07}, which predict a characteristic stellar mass of $\\simeq 30$M$_{\\odot}$ (when scaled to the GC), and a MF that is significantly deficient in low-mass objects.  Previous studies have not been able to reproduce all of the observed properties of the disks (notably the eccentric stellar orbits), but our results suggest that fragmentation of moderately eccentric accretion disk can form disks of stars with masses and orbits consistent with those observed.  Our simulations suggest that eccentric disks fragment to produce MFs that are slightly more top-heavy than expected in the circular case (due to the tidal destruction of the lowest-mass clumps), but not so top-heavy as to be inconsistent with the observed stellar populations.  Moreover, previous simulations have shown that changes in the thermodynamic properties of the disk, such as the cooling rate, can alter the MF at a level at least as significant as that seen in our simulations \\citep{nayak07}.  Consequently, we are satisfied that our results are robust, within the range of parameter space explored by our simulations.  However, there are three major areas in which our simulations differ from real systems, and we discuss each of these in turn below.  In addition, we note that angular momentum transport (by spiral waves) is negligible in the simulations that exhibit fragmentation, as the fragmentation and subsequent accretion of gas occurs on the dynamical timescale (which is much shorter than the timescale for angular momentum transport).  It has long been thought that angular momentum transport due to gravitational instabilities can be responsible for significant accretion in disks around black holes \\citep[e.g.,][]{sfb89}.  This is indeed likely, but our results confirm those of \\citet{nayak07}, that star formation due to gravitational instability in accretion disks occurs on a timescale that is significantly shorter than the timescale for angular momentum transport.  Consequently, if star formation occurs at all radii in the disk, it seems likely that little or no gas will be accreted on to the central black hole.  We note, however, that real disks around black holes are likely only unstable to fragmentation in their outer regions \\citep[e.g.,][]{levin07,kp07}, and that whether or not gravitational instability in such disks results in fragmentation or transport depends critically on the issues of cooling discussed in Section \\ref{sec:dis_cool}.  Our simulations are most applicable to the GC stellar disk(s), but their scale-free nature allows us to consider the results in terms of other systems also.  It has been suggested that the young stellar population at the GC may be characteristic of more generic episodes of black hole (AGN) activity in other galaxies \\citep[e.g.,][]{lb03,tqm05}, and this has implications for models of AGN accretion and the growth of supermassive black holes.  The consequences of our models for AGN feeding depend on the disk thermodynamics (see Section \\ref{sec:dis_cool}), but we note in passing that our models fit naturally into, for example, the black hole growth scenario presented by \\citet{kp06,kp07}.  The applicability of our results to protoplanetary discs, however, is less clear.  It seems likely that protoplanetary disks are gravitationally unstable at early times in their evolution \\citep[e.g.,][]{durisen_ppv}, but it is less clear that such disks possess significant eccentricities.  If eccentric, self-gravitating disks do exist around young stars then our results could easily be applied to such systems, but at present there is no strong evidence for the existence of such systems.  \\subsection{Initial configuration}\\label{sec:dis_ics} The first area of concern is in our initial conditions.  We choose an initial configuration that is marginally stable, and allow the disks to cool into instability.  This approach has been followed by previous studies \\citep[e.g.,][]{rice03,rice05}, and recent work suggests that use of these artificial initial conditions does not affect the fragmentation boundary in circular disks \\citep{clarke07}.  What is less clear, however, is whether or not our choice of a uniform angles of pericenter is similarly valid.  As discussed in Section \\ref{sec:intro}, in systems like the GC we expect the disk to form via some type of accretion event, and while previous simulations \\citep[e.g.,][]{sanders98} have found that an eccentric configuration can result, there is no particular reason to expect a uniform eccentricity.  If the disk does not have uniform eccentricity then neighboring orbital streamlines will intersect, and the resulting shocks have the potential to provide significant additional heating.  Such heating may act to stabilize the disk against fragmentation, but a realistic treatment of this problem requires full (magneto-)hydrodynamic simulations of the infall and capture of a gas cloud by the black hole, and such modeling is clearly beyond the scope of this work.  What we have demonstrated, however, is that if an eccentric disk is able to evolve towards a configuration where it satisfies the (circular) conditions for fragmentation, then it will fragment before any such shock heating is able to stabilize the disk.  \\subsection{Cooling and thermodynamics}\\label{sec:dis_cool} The second major area of uncertainty in our models is our simplified cooling prescription.  Our cooling prescription (Section \\ref{sec:cooling}) parametrizes the cooling rate by semi-major axis only, and while this is ideal for the sort of numerical experiments conducted here, it is far from realistic.  In the case of protoplanetary disks, the subject of heating and cooling rates in gravitationally unstable disks is an area of on-going debate \\citep[e.g.,~the review by][]{durisen_ppv}, but in the case of a disk around a black hole we can justify the approximate form of our cooling prescription by a few simple scaling arguments.  Similar arguments have been discussed previously by other authors \\citep[e.g.,][]{lb03,jg03,gt04,levin07,rafikov07}, but we re-state them here for completeness.  A self-gravitating disk has $Q\\simeq 1$, and therefore if we know the disk surface density we can estimate the sound speed, and therefore the temperature, in the disk.  If we scale our model to the GC system (by adopting a characteristic radius of 0.1pc and $M_{\\mathrm {bh}}=3\\times10^6$M$_{\\odot}$), we find a typical disk surface density of $\\sim 100$g cm$^{-2}$ and a corresponding disk temperature of $\\sim 100$K (estimated at $a=2$, but these quantities do not vary significantly over the limited radial range of our simulations).  At such temperatures the dominant source of opacity is dust grains, and typical values for the Rosseland mean opacity are $\\kappa \\simeq 1$--10cm$^2$g$^{-1}$ \\citep[e.g.,][]{semenov03}.  The opacity does not change dramatically at temperatures below that at which the dust grains sublime (typically $\\simeq 1500$K).  As seen above, scaling our models to the GC system (assuming that the initial gas mass in the disk is comparable to the observed stellar mass) results in a much lower temperature than this, so significant variations in the disk opacity (such as the opacity gap discussed by, e.g., \\citealt{jg03,tqm05}) seem unlikely.  Only if star formation in such disks is very inefficient, contrary to what is seen in our (and other) simulations, would we expect the disk temperatures to be high enough for dust sublimation to be significant.  Therefore, we expect such a disk to be marginally optically thick\\footnote{Note that this also verifies our earlier assumption that the cooling rate is unlikely to change significantly around the orbit.}, and the resulting cooling timescale is typically comparable to the orbital period \\citep[e.g.,][]{rafikov07}.  However, we also note that our cooling model takes no account of changes in the cooling function as the local density increases (due to gravitational collapse).  A more realistic model would make use of an opacity-based cooling model that accounts for the density structure of the disc, and would also adopt a somewhat stiffer equation of state in the the collapsing clumps.  The application of such models to protoplanetary disks is still mired in controversy \\citep[e.g.,][]{boley06,mayer07}, but the physical conditions in a GC disk are known to be much more conducive to fragmentation than those in protoplanetary disks \\citep[see e.g.,~the discussion in][]{rafikov07}.  Consequently we expect that our disks will cool rapidly enough to be unstable to fragmentation.  Moreover, the uncertainties discussed here are not specific to our models of eccentric disks, so we do not expect our results regarding the differences between eccentric and circular disks to be affected by these simplifications.  A further simplification is that our cooling prescription depends only on the local conditions in the disk, and is not influenced at all by the global disk structure.  Previous work using this scale-free, local cooling prescription has found that it gives rise to instabilities that are well-approximated by a local description \\citep{lr04,lr05}, and therefore the global disk structure is unlikely to have a significant effect on the spatial distribution of clumps in our simulations.  However, in a real disk the cooling time is not scale-free and, in general, gravitational instabilities cannot be well-described locally \\citep{bp99}.  Numerical simulations of protoplanetary disks have found that transport can be dominated by low-order global modes in the non-scale-free case \\citep{boley06,cai07,durisen_ppv}.  It seems likely, however, that the local approximation will be more accurate for the very thin disks considered here than in the case of  protoplanetary disks (where $M_{\\mathrm d}/M_* \\sim 0.1$).  What is less certain is whether or not a realistic system is able to evolve into the configuration discussed here.  The virial temperature of a gas cloud that falls towards the GC from, for example, the circumnuclear ring is orders of magnitude larger than the disk temperatures considered here \\citep[e.g][]{morris93,sanders98}.  An infalling cloud may not virialize before it forms a coherent disk structure, but there is clearly sufficient energy available to heat the gas significantly.  At temperatures from $\\simeq 1000$--$10^4$K the typical Rosseland opacities are much smaller than those at $\\sim 100$K, as the primary coolants (dust, and also molecular species) are destroyed at such temperatures \\citep{semenov03}.  Thus, we expect the cooling rates at higher temperatures to be much lower than those considered here, and it may be difficult for such a disk to cool into instability in a more realistic model \\citep[see also the discussion in][]{jg03}.  Our simple thermodynamic model is viable once the disk reaches the state considered here but, as discussed above, we do not address the issue of whether or not a real system will be able to reach such a configuration.  \\subsection{Accretion physics at stellar scales} The last major area neglected by our simulations is the small-scale physics of star-formation.  Our simulations have sufficient numerical resolution to treat the fragmentation process correctly, but when scaled to the GC our minimum length resolution is a few tens of AU.  This is 2--3 orders or magnitude larger than the typical stellar radius, and we make no attempt to model the physics of star formation on small scales.  Essentially our MFs are MFs of pre-stellar cores, and as such they may not be representative of the stellar MFs that result from disk fragmentation.  However, we note that the fragmentation into bound clumps occurs on orbital timescales, and the results presented in Fig.\\ref{fig:number_high} suggest that all of the disk gas will be bound to clumps within $\\simeq10$ orbital periods.  The orbital period of the GC stellar disks is $\\sim 1000$yr, but the typical timescales for such cores to collapse to form stars are $\\sim 10^4$--$10^5$yr \\citep[e.g.,][]{mt03}.  Consequently we do not expect stellar feedback processes to have a significant effect on the process of disk fragmentation.  However, individual bound clumps may well form more than one star each, and we also expect to see some mergers between clumps \\citep{levin07}.  More detailed simulations, using higher numerical resolution and more realistic thermal physics, will be required to address these issues further.    We have constructed hydrodynamic models of gravitationally unstable disks, and investigated the effects of eccentricity on disk fragmentation.  We find that the fragmentation boundary is unaffected by eccentricity, and that conditions which lead to fragmentation in circular disks do likewise in eccentric disks.  However, we find that the fragmentation process is altered by eccentricity, as the variable tidal forces around an eccentric orbit have a strong effect on the accretion of gas on to bound fragments.  We find that the formation of low-mass fragments is suppressed, and that the growth (via accretion) of more massive fragments proceeds much more slowly in an eccentric disk than in a circular one.  We consider our results in the context of the GC stellar disks, and find good agreement between the MFs in our simulations and those observed.  We find that fragmentation of eccentric accretion disks is a viable mechanism for the formation of the GC stellar disks, with the resulting stellar MF and orbits consistent with observations."
    },
    "0711/0711.2920.txt": {
        "abstract": "We give a concise formula for the non-Gaussianity of the primordial curvature perturbation generated on super-horizon scales in multi-scalar inflation model without assuming slow-roll conditions. This is an extension of our previous work. Using this formula, we study the generation of non-Gaussianity for the double inflation models in which the slow-roll conditions are temporarily violated after horizon exit, and we show that the non-linear parameter $f_{NL}$ for such models is suppressed by the slow-roll parameters evaluated at the time of horizon exit.  %We also briefly discuss the possibility to get large $f_{NL}$ %by violating the slow-roll conditions. ",
        "introduction": "Non-Gaussianity of the primordial curvature perturbation is a potentially useful discriminator of the many existing inflation models \\cite{Komatsu:2001rj,Bartolo:2004if}. Planck \\cite{:2006uk} is expected to detect the primordial non-Gaussianity if the so-called non-linear parameter, $f_{NL}$, which parameterizes the magnitude of the bispectrum, is larger than $3\\sim 5$ \\cite{Komatsu:2001rj, Babich:2004yc}. Higher order correlation functions such as trispectrum would also be a useful probe of the primordial non-Gaussianity \\cite{Okamoto:2002ik, Bartolo:2005fp, Kogo:2006kh}. Hence it is important to theoretically understand the generation of non-Gaussianity.  Standard single slow-roll inflation model predicts rather small level of the non-linear parameter, $f_{NL}$, suppressed by the  slow-roll parameters \\cite{Maldacena:2002vr}. In multi-scalar field inflation models with separable potential, explicit calculations in the slow-roll approximation show that $f_{NL}$ is also of the order of the slow-roll parameters~\\cite{Kim:2006te,Vernizzi:2006ve,Battefeld:2006sz,Battefeld:2007en,Choi:2007su,Rigopoulos:2005xx,Rigopoulos:2005us,Rigopoulos:2005ae}. This feature is likely to hold for more general potential that satisfies the slow-roll conditions \\cite{Seery:2005wm,Seery:2005gb,Yokoyama:2007uu}. On the other hand, the generation of non-Gaussianity due to the violation of the slow-roll conditions remains an open problem.  In this paper, based on $\\delta N$ formalism~\\cite{Sasaki:1995aw,Sasaki:1998ug,Lyth:2005fi,Lyth:2004gb,Starobinsky:1986fx}, we give a useful formula for calculating the non-linear parameter in the multi-scalar inflation models including the models in which the slow-roll approximation is (temporarily) violated after the cosmological scales exit the horizon scale during inflation. Current observations do not exclude such models. We also study the possibility of the large non-Gaussianity for double inflation model~\\cite{Silk:1986vc,Turner:1986te}, as an example of such a model.  In section \\ref{formulation}, we derive the formula for the non-linear parameter $f_{NL}$, which is an extension of our previous paper~\\cite{Yokoyama:2007uu}. In section \\ref{example}, we analyze the primordial non-Gaussianity in the double inflation model using our formula. %In section \\ref{discussion}, %we study the generation of $f_{NL}$ in the hybrid inflation model %and also discuss the possibility to get large $f_{NL}$ %by violating the slow-roll conditions. We give a summary in section \\ref{summary}. ",
        "conclusions": "\\label{summary} Based on the $\\delta N$ formalism, we have derived a useful formula for calculating the primordial non-Gaussianity due to the super-horizon evolution of the curvature perturbation in multi-scalar inflation without imposing slow-roll conditions. This formula can apply for the inflation models with general field space metric, $h_{IJ}$, as far as super-horizon contributions are concerned. Generally, when one calculates the non-Gaussianity of the curvature perturbations, one has to solve the second order perturbation equations. In doing so for a multi-scalar inflation, there appear tensorial quantities with respect to the indices of the field components. Our formula reduces the problem of calculating the non-linear parameter $f_{NL}$ to solving only first order perturbation equations for two vector quantities. This reduces ${\\cal O}({\\cal N}^2)$ calculations to ${\\cal O}({\\cal N})$ ones where ${\\cal N}$ is the number of the scalar field components. Hence our formalism has a great advantage for the numerical evaluation of $f_{NL}$ in the inflation model composed of a large number of fields.  We have also studied the primordial non-Gaussianity in double inflation model as an example that violates slow-roll conditions by using our formalism. We found that, although $f_{NL}$ defined for the curvature perturbation on a constant Hubble hypersurface exceeds $1$ for a moment around the time when the slow-roll conditions are violated, the final value of $f_{NL}$ is suppressed by the slow-roll parameters evaluated at the time of horizon exit. We have shown that this can be understood even analytically in the $\\delta N$ formalism. This result is straightforwardly extended to more general double inflation model and ${\\cal N}$-flation model."
    },
    "0711/0711.2089_arXiv.txt": {
        "abstract": "{BL Lacertae objects have been the targets for numerous recent multiwavelength campaigns, continuum spectral variability studies, and theoretical spectral and variability modeling.  A meaningful interpretation of the results of such studies requires a reliable knowledge of the objects' redshifts; however, the redshifts for many are still unknown or uncertain.  } {Therefore, we hope to determine or constrain the redshifts of six BL Lac objects with unknown or poorly known redshifts.  } {Observations were made of these objects with the MDM 2.4 m Hiltner telescope.  Although no spectral features were detected, and thus no redshifts could be measured, lower redshift limits were assigned to the objects based on the expected equivalent widths of absorption features in their host galaxies.  Redshifts were also estimated for some objects by assuming the host galaxies are standard candles and using host galaxy apparent magnitudes taken from the literature. } {The commonly used redshift of $z=0.102$ for \\object{1219+285} is almost certainly wrong, while the redshifts of the other objects studied remain undetermined.  } {} ",
        "introduction": "Blazars are among the most violent, highly variable astrophysical high-energy phenomena, with rapidly-varying emission from radio through $\\gamma$-ray energies.  They are thought to consist of relativistic jets from supermassive black holes, closely aligned with our line of sight.  BL Lacertae (BL Lac) objects, a subclass of blazars, are defined by their quasar-like continuum and the weakness or absence of broad emission lines in their optical spectra.  These spectra are thought to be dominated by nonthermal emission from the highly relativistic jet, masking the contributions of the stars in the host galaxy, and line emission from gas clouds near the supermassive black hole.  There are various models for the source of the nonthermal emission from blazars.  Of critical importance in distinguishing between the models is setting the energy scale, on which many parameters (e.g., the jet's speed, density, magnetic field, etc.) depend.  This can only be done if the objects' redshifts, and hence distances, are known.  Many well-observed BL Lac objects have unknown or poorly known redshifts based on $\\sim$1--3 emission or absorption lines in low signal-to-noise ($S/N$) spectra.  In several cases, a poorly determined redshift of a BL Lac is repeatedly cited throughout the literature until its reliability is no longer questioned.  Thus, a project was undertaken to obtain the optical spectra of several of these objects at the MDM Observatory with the goal of obtaining redshift estimates or improved constraints.  Unfortunately, no definitive spectral features were revealed, and thus no definitive redshifts could be assigned. However, using the method of \\citet{sbarufatti06}, we were able to estimate a minimum redshift based on the expected equivalent widths of absorption features in the host galaxy.  Also, assuming the host galaxies are standard candles \\citep[see, e.g.,][]{nilsson03}, we estimated the redshifts of several objects based on the observed magnitude of the host galaxies, found from the literature, based on the method of \\citet{sbarufatti05}. ",
        "conclusions": "\\label{summary}  The spectra of six BL Lac objects with poorly known or unknown redshifts have been obtained. For several objects, these spectra have higher $S/N$ than any previously published.  Based on this papers' results, the commonly used redshift of $z=0.102$ for \\object{1219+285} is almost certainly wrong. The redshifts of the other objects studied remain undetermined."
    },
    "0711/0711.2685_arXiv.txt": {
        "abstract": "We present multi-wavelength observations of stellar features in the H\\,{\\sc i} tidal bridge connecting M81 and M82 in the region called Arp's Loop. We identify eight young star-forming regions from Galaxy Evolution Explorer ultraviolet observations. Four of these objects are also detected at H$\\alpha$.  We determine the basic star formation history of Arp's Loop using F475W and F814W images obtained with the Advanced Camera for Surveys on board the Hubble Space Telescope. We find both a young ($<$ 10 Myr) and an old ($>1$ Gyr) stellar population with a similar spatial distribution and a metallicity $Z\\sim0.004$. We suggest that the old stellar population was formed in the stellar disk of M82 and/or M81 and ejected into the intergalactic medium during a tidal passage ($\\sim$ 200--300 Myr ago), whereas the young UV-bright stars have formed in the tidal debris. The UV luminosities of the eight objects are modest and typical of small clusters or OB associations. The tidal bridge between M81--M82 therefore appears to be intermediate between the very low levels of star formation seen in the Magellanic bridge and actively star-forming tidal tails associated with major galaxy mergers. ",
        "introduction": "The M81 group contains three major closely interacting galaxies: M81, M82 and NGC 3077. Studies of this triple system at radio wavelengths have shown an extremely disturbed H{\\sc i} distribution with tidal bridges connecting the three galaxies (e.g. Gottesman \\& Weliachew 1975; van der Hulst 1979; Appleton, Davies \\& Stephenson 1981; Yun, Ho \\& Lo 1994). Within the tidal bridges, H{\\sc i} knots with optical counterparts are seen. The most prominent of these to the east/north-east of M81 are Holmberg IX and Arp's Loop or A0952$+$69 (Arp 1965). It has been suggested (e.g. Boyce et al. 2001) that these are forming tidal dwarf galaxies. Numerical simulations of the system by Yun (1999) can successfully reproduce the the H{\\sc i} tidal debris if the closest approaches of M82 and NGC 3077 to M81 took place 220 and 280 Myr ago respectively. Arp's loop, the topic of this paper, occurs at the intersection of the three tidal streamers in the simulation of Yun (1999) who gives an H{\\sc i} mass of $2.6 \\times 10^8$ M$_\\odot$. Sun et al. (2005) present very deep wide-field optical continuum images of the M81/M82 galaxy group, and show that groups of young stars are distributed across the entire region between the two galaxies.  Although tidal tails usually have blue optical colors (e.g. Schombert et al. 1990), it is not clear whether these colors are due to tidally-disrupted stellar disk material or active star formation in the tails. Recently, ultraviolet (UV) imaging obtained with the Galaxy Evolution Explorer (GALEX) satellite has shown that regions of bright UV emission, containing recently formed stars, correlate well with regions of high H{\\sc i} column density in tidal tails (e.g. Hibbard et al. 2005; Neff et al. 2005). These findings suggest that the combination of UV and H{\\sc i} data provides a powerful technique for identifying and studying star formation in the intergalactic medium (IGM) between interacting galaxies. This topic is of importance for earlier epochs in the universe because galaxy interactions, and thus tidal tails were more common, leading to a higher star formation rate in the IGM. This may well have implications for the enrichment of the IGM; Ryan-Weber et al. (2004) have shown that stars forming in the IGM offer a viable alternative to galactic wind enrichment on large scales.  We have thus started a study of star formation in the IGM of the M81 galaxy group with the overall aim of elucidating the importance of this mode of star formation. We present results for Arp's loop in this paper and for Holmberg IX in Sabbi et al. (2007). Arp's loop is located 17 arcmin north-east of M81 and its composition was first discussed by Efremov et al. (1986). They resolved the ring-like structure into blue diffuse and star-like objects which they interpreted as young star clusters and associations. Makarova et al. (2002) investigated the stellar population of Arp's loop using photometry obtained with the Wide Field and Planetary Camera 2 (WFPC2) on board the Hubble Space Telescope (HST). They were unable to clarify the evolutionary nature of Arp's loop because of photometric limitations and WFPC2's small field of view. They find evidence for: enhanced star formation between 20 and 200 Myr ago, coincident with the likely epoch of tidal interactions in the M81 group; intermediate age stars which were either formed prior to the interaction in the Arp's loop or were torn out of the massive galaxies during the interaction; and also faint stars with ages $\\ge$ 8 Gyr but which were too close to the detection limit for detailed analysis. In this paper, we use multi-wavelength observations, including deep photometry obtained with the Advanced Camera for Surveys (ACS), to resolve the stellar population, and thus determine the evolutionary status of Arp's loop.  \\section {The Data}  \\subsection{GALEX, ACS, H{\\sc i} and WIYN}  The M81 group was observed with the Galaxy Evolution Explorer (GALEX) mission in the far (FUV: $\\lambda$$_{\\rm eff}$=1528\\AA; $\\Delta$$\\lambda$$_{FUV}$=269\\AA) and near ultraviolet (NUV: $\\lambda$$_{\\rm eff}$=2271\\AA; $\\Delta$$\\lambda$$_{NUV}$=616\\AA). We obtained the data from the Multimission Archive at STScI (MAST) as part of the team General Releases. The GALEX field of view is 1$^{\\circ}$.28 and 1$^{\\circ}$.24 in the FUV and NUV bands, and the pixel scale is 1.5$''$/pixel. The exposure times were 3093s for each band.  Fluxes were calculated using the equations of Morrissey et al. (2005): m$_{\\lambda}=-2.5$ log[F$_{\\lambda}$/a$_{\\lambda}$] + b$_{\\lambda}$, where a$_{FUV}$ = 1.4 $\\times$ 10$^{-15}$ erg s$^{-1}$ cm$^{-2}$ \\AA$^{-1}$, a$_{NUV}$=2.06$\\times$ 10$^{-16}$ erg s$^{-1}$ cm$^{-2}$ \\AA$^{-1}$, b$_{FUV}$=18.82 and b$_{NUV}$=20.08 for FUV and NUV, respectively. Fluxes were multiplied by the effective filter bandpass ($\\Delta$$\\lambda$$_{FUV}$=269\\AA\\, and $\\Delta$$\\lambda$$_{NUV}$=616\\AA) to give units of erg s$^{-1}$ cm$^{-2}$ and luminosities were calculated for a distance of 3.6~Mpc (Freedman et al. 1994).   Figure \\ref{m81m82hi} shows the distribution of the atomic hydrogen in the M81 group of galaxies (Yun, Ho \\& Lo 1994) on top of the GALEX FUV image. We have identified 8 FUV objects (Fig.~\\ref{nuv_feet}, Table 1) inside the H{\\sc i} cloud located on the northeast side of the bridge between M81 and M82 and north of Holmberg~IX. This region, called concentration~II by Yun et al., is the brightest H{\\sc i} complex in the tidal bridge. It includes Arp's loop, and may not be a single cohesive feature, containing tidally stripped gas from M81 and M82. Makarova et al. (2002) found 250 stars in the HST/WFPC2 F606W/F814W images of Arp's loop. However, the WFPC2 observations were centered on the H{\\sc i} concentration and cover only objects \\#7 and \\#8, missing the other FUV objects. We searched the HST archive at MAST and found new pointings of the same area but taken with the  Wide Field Channel (WFC) of ACS using the F475W and F814W filters and covering a larger field of view (GO 10915: PI Dalcanton). Fig.~\\ref{acs_feet} shows the ACS F475W image where we have plotted the WFPC2 footprint, the GALEX contours, and marked the 8 FUV objects. The high resolution of ACS (0.05$''$) resolves the FUV objects into hundreds of starlike objects as can be seen in Fig.~\\ref{blobszoom}. However, the ACS pointing misses two of the FUV objects (\\#3 and \\#4) and part of another object (\\#1).  We have also obtained an H$\\alpha$ image of Arp's loop with the OPTIC imager at the WIYN{\\footnote {The WIYN Observatory is a joint facility of the University of Wisconsin-Madison, Indiana University, Yale University, and the National Optical Astronomy Observatories.}} 3.5m telescope. The observations were taken with WIYN filter \\# 13 ($\\Delta$$\\lambda$ =  47 \\AA) and a standard Johnson $R$ filter for continuum subtraction. OPTIC is a mosaic of two 2kx4k orthorgonal transfer CCDs, resulting in a field of view of 9.5$'$$\\times$9.5$'$. In Fig.~\\ref{wiyn} we present only an image of chip \\#2 which was bias-corrected and flat-fielded using IRAF{\\footnote{IRAF is distributed by the National Optical Astronomy Observatory, which is operated by AURA, Inc., under cooperative agreement with the National Science Foundation.}} The two images taken in each filter were combined with cosmic-ray rejection, and the resulting scaled $R$-band image was subtracted from the combined H$\\alpha$ image to remove the continuum signal. Our data confirm Karachentsev \\& Kaisin's (2007) detection of objects \\#1, \\#8, \\#7, and \\#3, with \\#1 being the brightest and \\#3 the faintest object (see fig.1 in Karachentsev \\& Kaisin). As shown in Fig.~\\ref{wiyn}, although we have detected the faintest object, we have not detected two of their H$\\alpha$ sources which are marked as squares in Fig.~\\ref{wiyn}. This could be due to the different bandwidths, since Karachentsev \\& Kaisin used a broader filter ($\\Delta$$\\lambda$ = 75 \\AA) than ours. H$\\alpha$ detections of the FUV sources suggest that these are small H{\\sc ii} regions where hot, blue OB stars are emitting ultraviolet radiation. ",
        "conclusions": "We have analyzed the UV and optical light within the H{\\sc i} bridge between M81 and M82 system and identified 8 FUV regions in an area known as Arp's loop. Four of these objects were also detected in H$\\alpha$. High resolution ACS images show six of the regions in great detail and allow us to trace their star formation history. The synthetic CMD modeling shows the presence of a young population ($<$ 10 Myr) together with an older component with ages $>$ 1 Gyr. It is possible that these two populations are not related. In any case the young population must have formed ``in situ\" whereas the old population probably was formed in the M82 and/or M81 galaxies, and removed from their disks during a tidal passage $\\sim$ 200--300 Myr ago which also triggered star formation in the H{\\sc i} debris.  The FUV luminosities of the eight objects we detect are modest and typical of small clusters or associations of O and B stars. It is likely that these young stars have condensed recently from the tidally stripped gas from M81 and M82. The tidal bridge between M81--M82 appears to be intermediate between the very low levels of star formation seen in an object like the Magellanic bridge and very active tidal tails like those found in interacting galaxies such as the Mice and the Antennae. Ongoing star formation in relatively diffuse structures, such as Arp's Loop, is consistent with expectations for gas collisions and subsequent turbulence to lead to inefficient conversion of gas into stars in low density systems such as tidal tails.  The fate of these star-forming clumps is not clear. They can grow to become TDGs, or make star clusters that survive to live in the distant halos of their hosts, or even become intergalactic objects. Finally, they may dissolve and not remain gravitationally bound, yielding only very sparse star streams. In any case, they will have a composite population of old and young stars formed in very different environments such as the M81/M82 disks and the tidal bridge. More observations of interacting systems similar to the M81 group will help establish the importance of these systems in galaxy evolution and in the enrichment of the intergalactic medium."
    },
    "0711/0711.2016_arXiv.txt": {
        "abstract": "{ I review the latest results on the calibration of the Tip of the Red Giant Branch as a standard candle, in the optical and in the near infrared. The agreement among different and independent empirical calibrations is rather good, if all the uncertainties are taken into account. The possible extension of the calibration to SDSS photometric bands ({\\tt i,z}) is also discussed. ",
        "introduction": "The use of the Tip of the Red Giant Branch (TRGB) as a standard candle is a mature technique that is currently used to reliably estimate distances to galaxies of all morphological types, from the Local Group \\cite[see, f.i.,][and references therein]{mcc} up to the Virgo cluster and beyond \\cite[see, f.i.,][and references therein]{ferra,durre}. The underlying physics is well understood \\citep[][hereafter MF98]{scw,mf98} and the observational procedure is operationally well defined and robust to contamination by fore/background stars and/or stars brighter than the TRGB \\citep[MF98,][]{sakai}. The key observable is the sharp cut-off occurring at the bright end of the Red Giant Branch (RGB) Luminosity Function (LF) that can be easily detected with the application of either non-parametric \\citep{mf95,sakai} or parametric \\cite[see, for example][]{mendez,mcc} methods. The necessary condition for a safe application of the technique is that the observed RGB LF be well populated, with more than $\\sim$100 stars within 1 mag from the TRGB \\citep{mf95,b02}. Under these conditions, the typical uncertainty on the estimate of the apparent magnitude of the tip $I^{TRGB}$  is of order of a few hundreds of mag ($\\le 0.05$), and the uncertainties on $M_I^{TRGB}$ dominates the error budget of distance estimates obtained with the TRGB technique. In general, the magnitude of the TRGB depends quite weakly on the age of the considered population \\citep{bark,scw} and, in the Cousins' I band, it depends weakly also on the metallicity, at least for relatively metal poor systems \\citep{DA90,lfm93}. In Near Infra Red (NIR) passbands the dependence on metallicity is strong but it can be accounted for with very simple models \\cite[see below, and also][and references therein]{f00,iva,ele,b04}.  Here, I will briefly review the status of the calibration of this standard candle and the perspectives for future applications. For obvious reasons of space this review cannot be and is not intended to be exhaustive or complete; for example, I will not report on the hundreds of applications of the method, I will focus on the most recent results, I will deal more with calibrations in the optical than in  NIR passbands, and, also, I will not discuss the differences between the prediction of the various theoretical models \\citep{scw}. More detailed discussions and references can be found in MF98, \\citet{sc98}, \\citet[][hereafter B01]{b01}, \\citet[][hereafter B04]{b04}, and \\citet[][hereafter R07]{r07}. The TRGB method in the broader context of  distance scales is discussed by \\citet{W03} and \\citet{alves}, among the others. The behaviour of the TRGB in complex stellar populations is investigated by \\citet{bark} and  \\citet{sg05}. A detailed and up-to-date report of stellar evolution on the Red Giant Branch is given by \\citet{scw}.  Nature provides serious limitations in our possibility of studying the behaviour of the TRGB as a function of metallicity, age, elemental abundance pattern, etc., since we lack populous star clusters covering the whole extension of the space of parameters. For this reason, to have an idea of - at least - the differential changes of the TRGB magnitude in a given passband in response to variations of the above quoted parameters we must recur to grids of theoretical models \\cite[which may suffer from several problems and unadequacies, see][]{scw}. Also in this case a complete overview of all the available sets of tracks/isochrones would be impossible: here I use mainly BASTI models \\citep{basti,cordier}, but also Padova models \\citep{g00,g04}, Yale-Yonsei (Y$^2$) models \\citep{YY}, the old-ages isochrones by \\citet{scl97}, and, finally, the recently published Darthmouth models \\citep{dotter}.     \\begin{figure}[] \\resizebox{\\hsize}{!}{\\includegraphics[clip=true]{compz.eps}} \\caption{\\footnotesize Comparison between different empirical calibrations of $M_I^{TRGB}$ as a function of metallicity (upper panel). The large filled circle is the observed tip level of $\\omega$ Cen, the smaller circle corresponds to the tip of 47 Tuc, according to B04. The R07 calibration, as a function of $(V-I)_0$ color has been expressed as a function of [Fe/H] using the $(V-I)_0^{TRGB}$ vs. [Fe/H] relation by B01. The lower panel shows the difference between the B04 and R07 calibrations as a function of [Fe/H].} \\label{compz} \\end{figure} ",
        "conclusions": "The future of the TRGB as a distance indicator looks very promising, as new telescopes and surveys will give access to more distant systems, larger areas of the sky and more favourable wavelenght windows. In general applications to distant galaxies, the lack of knowledge of many parameters that may influence the TRGB magnitudes will probably prevent the method to reach exquisite levels of accuracy (i.e., better than 10\\% uncertainties); on the other hand the {\\em weakness} of all these dependencies ensures a good level of robustness and reliability of the method. The Zero Point of current empirical calibrations is uncertain at the $\\pm 0.1$ level, but the results the GAIA astrometric mission should reduce such uncertainty virtually to zero. An extension of the empirical calibration to passbands like SDSS {\\tt i} and {\\tt z} and a detailed comparison between the prediction of the different sets of theoretical models are suggested as the most interesting problems to afford, at the present stage."
    },
    "0711/0711.1223_arXiv.txt": {
        "abstract": "Spectral Spatial Fluctuations (SSF) of the Cosmic  Microwave Background Radiation (CMBR) temperature are considered as a result of an interaction of primordial atoms and molecules with CMBR in proto-objects moving with peculiar velocities relative to the CMBR. Expected optimistic values of $\\Delta T/T$ are $2\\times 10^{-5}-2\\times 10^{ - 6}$ for SSF caused by HeH$^{+}$ at z =20-30 which are possible redshifts of early reionization scenario. The bandwidth of the lines is 0.1-2{\\%} depending on the scale of proto-objects and redshifts. For the SSF search CMBR maps in different spectral channels are to be observed and then processed by the Difference method. Simulation of the experiment is made for MSRT (Tuorla Observatory, Finland) equipped with a $7\\times4$ beam cryo-microbolometer array with a chopping flat and frequency multiplexer providing up to 7 spectral channels in each beam (88-100 GHz). Expected $\\Delta $T/T limit in the experiment is 2$\\times $10 $^{-5}$ with 6$'$-7$'$ angular and 2{\\%} frequency resolution. Simulation shows that SSF may be recognized in the angular power spectrum when S/N in single frequency CMBR maps is as small as 1.17 or even something less for white noise. Such an experiment gives us a possibility to set upper limit of SSF in MM band and prepare future SSF observations. ",
        "introduction": "Observational effects caused by primordial molecules seem to be most promising in investigating Dark Ages epoch of the Universe. The basic properties of the molecules are discrete narrow energy levels and high efficiency of their interaction with CMBR. This leads to forming SSF if proto-objects, containing these molecules, move with peculiar velocities V$_{p}$ relative to CMBR (Dubrovich,1977,1982, Puy,1993). We  may consider SSF as manifestation  of the proto-objects at high redshifts 300 $> z >$ 10 as scattering of CMB photons by primordial molecules must leave imprints on the CMBR temperature distribution.   The most abundant chemical elements predicted  by the pure Big Bang model are H, He, D, $^{3}$He, Li and their ions. They give some molecules in the primordial matter at z=100-200 such as H$_{2}$,H$_{2}^{+}$,HD, HD$^{+}$, HeH$^{+}$, LiH, LiH$^{+}$, H$_{2}$D$^{+}$. There are other molecules which should be considered if non-standard nuclear synthesis at the early times (z=10$^9$) or star formation at z = 100-200 took place.  For the high efficiency of interaction between  molecules and CMBR photons two values are important - the cross-section for scattering and the concentration (or relative abundance) of this molecule. The first parameter depends on the specific quantum structure of the molecule - its symmetry and dipole moment. The second one depends on the abundance of the chemical elements of which it is composed and on the rate of the appropriate chemical reactions. According to these constraints, we should take into consideration only those molecules which have a large enough dipole moment and relatively high abundance. These are HD$^{+}$, HeH$^{+}$, LiH, H$_{2}$ D$^{+}$. We consider them as the basic molecules interacting with CMBR at high z (Dubrovich,1977,1982,Puy et al.,1993,Stancil et al.,1996).   \\section {Theoretical predictions of SSF}  Let us remind physical conditions in Dark Ages  epoch when the Universe was younger than 300 mil. years. The Universe was in $(z+1)^{3/2}$ times smaller and in $(z+1)^3$ times more dense than now, temperature of CMBR and molecular gas followed $T_{z=0}(z+1)$ law till z $\\approx $150. Then, temperature of matter should be lower than CMBR temperature due to different adiabatic indexes of radiation and matter (Varshalovich, Hhersonskii,  Sunyaev, 1981) and it must follow $T_{z=0}(z+1)^2$ law till first star formation. Calculated temperatures of CMBR and molecular gas in Dark Ages epoch are given in Table 1, we used $T_{z=0}$ = 2.726 K (Mather et al., 1994). In epoch before secondary ionization molecules are formed in accordance with kinetic rate of specific chemical reactions when matter and radiation are cooling. After reionization molecular formation follows a more complicated law due to growth of ionization degree of hydrogen and existence of heavy chemical elements produced by first stars (Basu, Hernandez-Monteagudo, Sunyaev, 2004).  \\begin{table} \\caption{Temperatures of CMBR and molecular gas in Dark Ages epoch z=300-10}\\label{tab1}   \\begin{tabular}{lllll} \\hline z & Temperature of & Temperature of & Commentary & \\\\ & CMBR, & molecular gas, & & \\\\ & K & K & & \\\\ \\hline 1100 & 3000 & 3000 & Recombination & \\\\ 300 & 815 & 815 & Beginning of Dark Ages & \\\\ 150 & 415 & 415 & Middle of Dark Ages & \\\\ 30 & 85 & 17 & Dark Ages + reionization & \\\\ 20 & 57 & 8 & Dark Ages + reionization & \\\\ 10 & 30 & 100-10000 & End of Dark Ages, first stars and quazars & \\\\ \\hline \\end{tabular} \\end{table}    There are two mechanisms responsible for SSF  formation: simple scattering (Dubrovich,1977) and luminescence (Dubrovich,1997). In the simple mechanism opacity in narrow molecular lines and peculiar motion of the proto-clouds results in CMBR disturbances. This can be more visible in the rest frame of the proto-object. The CMBR in this frame becomes non-isotropic and out of thermal equilibrium. On the side towards which the proto-cloud moves, the temperature of the CMBR will be higher than average and on the opposite side it will be lower. After reflection, the photons are distributed isotropically in this frame that leads to non-isotropic distribution in laboratory frame. This explains the principal role of opacity and peculiar velocity. The amplitude of the effect depends on spectral index of the reflecting radiation (Dubrovich, 1977, Maoli et al., 1994). This effect corresponds to the elastic scattering between the molecules and photons, i.e. the total number of the photons does not change.   In fact all molecules have quite a complicated  energy level structure. This allows for the possibility of a non-elastic process. It is the well known luminescence process which plays an important role in the formation of radiation from a reflection nebula. Luminescence may enhance the intensity of rotation lines due to breaking of vibrational line quants (Dubrovich, 1997).  The magnitude of the CMBR temperature fluctuations  due to the pure reflection of photons by a moving object depends on peculiar velocity $V_p$ of a proto-object and optical depth $\\tau$(Sunyaev and Zel'dovich, 1970, Dubrovich, 1977, Maoli et al., 1994). The value of the distortions produced by molecules in CMBR is: \\begin{equation} \\Delta T/T = (V_{p}/c) \\cdot \\tau \\cdot K, \\end{equation} where $c$ is the speed of light and $\\tau $ is optical depth of the proto-object ( $\\tau << 1$, $V_{p}/c < 10^{ - 3})$, K - gain coefficient due to  luminescence effect, K=1-1000(K=1 for simple scattering). For the linear stage of perturbation evolution $V_p$ could be estimated as (Zel'dovich, Novikov, 1975): \\begin{equation} V_{p}=V_{p0}/(z+1)^{1/2} \\end{equation} where $V_{p0}$=600 km/s at large scales of CMB anisotropy (CMBA) and factor 2-5 higher at small CMBA scales.  Luminescence effect is tempting for SSF  search with limited sensitivity but high K may be expected  in narrow z intervals (Dubrovich,1997) while the simple mechanism of SSF formation operates in the wide field of z.  The optical depth $\\tau$ of the proto-object  may be calculated as (Dubrovich, Lipovka, 1995): \\begin{equation} \\tau=8 \\pi^{5/2}d^2(1+z)^{3/2}n_{H}\\alpha \\cdot [1-exp^{-h \\nu /kT}]/3hH_0, \\end{equation} \\begin{equation} h\\nu /kT=h c/[k T_0 \\lambda_0(1+z)] \\end{equation} $H_0$ - Hubble constant, h-Planck constant,  k-Bolcman constant, $\\lambda_0$ - the rest wavelength, $T_0$ - CMBR temperature, $n_H=n_{0H}\\cdot \\beta $, where $n_H$ - density of $H$, $n_{0H}$ - critical density of atomic hydrogen HI($10^{-5}$), $\\beta$ - a real density of $H$ in proto-object which may be less or much more than unity ($\\beta \\ge4\\cdot 10^{-2}$), $d$ - dipole moment, $\\alpha$ - abundance, $\\alpha = 10^{-8}-10^{-17}$ for different molecules and conditions in Early Universe.  Let us consider LiH, H$_{2}$D$^{ + }$, HeH$^{ + }$ HD$^{ + }$ primordial molecules.  \\underline{LiH} is a very important molecule, because it consists of primordial Li. Its abundance is a good test for the epoch of nuclear synthesis in the Early Universe. Its large dipole moment and relatively low frequency of the rotational and rovibrational transitions lead to a high value of K. But, unfortunately, some difficulties with the chemical processes of forming this molecule lead to a not too optimistic prediction for its abundance $\\alpha=10^{-11}-10^{-17}$. The rest wavelengths of LiH are $\\lambda_r$=676 $\\mu$m(rotational transition), $\\lambda_v$=7.1 $\\mu$m (rovibrational transition).  \\underline{HD$^{+}$} is also an important molecule due to the presence of primordial deuterium D. The abundance of D is about 5 orders of  magnitude larger than that of Li. But HD$^{+}$ has a dipole moment about 10 times less than LiH and a cross-section which is 100 times smaller. Another small factor is the abundance of H$^{+}$ at redshift z = 200, which might be about relative to that of neutral hydrogen. Due to the relatively high frequency of the rotational and rovibration transitions, the resulting value of K is not very large. Yet, if high sensitivity were reached, this molecule might be seen. The rest wavelengths of HD$^{+}$ are $\\lambda_r$=227 $\\mu$m, $\\lambda_v$=3.0 $\\mu$m.  \\underline{HeH$^{+}$} does not have any low abundance species. There are only two small factors which lead to a low abundance: a high rate coefficient for destruction (by electron recombination and collisions with the neutral atoms of hydrogen) compared with the very small rate coefficient of formation, and small abundance of $H^+$ at high redshift. It might be the most likely molecule to be searched for due to more optimistic predicted abundance $\\alpha = 10^{-8}-10^{-15}$. The rest wavelengths of HeH$^{+}$ are $\\lambda_r$=149 $\\mu$m, $\\lambda_v$=3.3 $\\mu$m.  \\underline{H$_{2}$ D$^{+}$} is the simplest triatomic molecule with a high dipole moment. It contains primordial deuterium D. Due to the presence of very low frequency transitions in its spectrum the value of K can be very high. It is very important that the redshift of the H$_{2}$ D$^{+}$ recombination is relatively high. The rest wavelengths of H$_{2}$ D$^{+}$ are $\\lambda_r$=1920 $\\mu$m, $\\lambda_v$=4.5 $\\mu$m.  The expected abundances of these molecules in  the Early Universe are discussed by many authors (Lepp and Shull,1984, Puy et al, 1993, Palla et al, 1995, Maoli et al, 1996, Stancil et al, 1996).  One of the main observational effects caused by  expansion of the Universe is the redshift of photons, expressed as \\textbf{ $\\lambda $}$_{obs}$ = (z + 1) \\textbf{$\\lambda $}$_{em}$, where \\textbf{$\\lambda $}$_{em}$ is the emitted or rest wavelength of the photon. So, if an object contains a gas of any molecules and has a $z$ redshift , it can only be seen at different but discrete frequencies $\\nu _{i}=\\nu _{oi}$/(1 + z), where $\\nu_{oi}$ is a discrete set of the molecule's rest transition frequencies. And vice versa, if some object is seen at given frequency $\\nu _{i}$ this object can have the redshifts z$_{i}=\\nu_{oi}$/$\\nu_i$ -1.   The final spectrum depends on the angular size of the object $\\theta $. The angular size $\\theta $ and the frequency bandwidth $\\Delta \\nu/\\nu $ depend on linear size L (for a spherical shape object) and z of the object to be observed. For small z the angular size of the object(if the linear size is constant and equal to L) is decreasing with moving away, but it grows for z larger than some value. Exact relations between $\\theta$ and $z$ in this model contain the $\\Omega_m$ and $\\Omega_{\\Lambda}$ parameters, which are the ratios of the matter $n_m$ and \"vacuum\" densities to the critical density $n_c$ respectively ($\\Omega_m+\\Omega_{\\Lambda}=1$). The fact that the cloud occupies the interval of redshifts z means that if it radiates or reflects the radiation locally in sufficiently narrow lines, all the radiation occupies the frequency interval\\textbf{ $\\Delta $}$\\nu $, moreover \\textbf{$\\Delta $}$\\nu $/$\\nu $= \\textbf{$\\Delta $}z/(1 + z). Scattering by molecules in several lines leads to superposition of images of different proto-objects at different z. Taking into account {$\\Delta $}z dependence on L and $\\Omega$ (Sahni, Starobinsky, 2000) we have:  \\begin{equation} \\Delta \\nu/ \\nu\\approx \\theta \\phi(z)[\\Omega_m(1+z)^3+\\Omega_{\\Lambda}]^{1/2}](1+z)^{-1}, \\end{equation} \\begin{equation} \\phi(z)=\\int_0^z 1/\\sqrt{\\Omega_m(1+z)^3+\\Omega_{\\Lambda}}dz, \\end{equation}  \\begin{equation} \\theta =(H_0 L/c)(1+z)/\\phi(z), \\end{equation} \\begin{equation} \\Delta \\nu/\\nu\\approx (H_0 L/c)\\sqrt{\\Omega_m(1+z)^3+\\Omega_{\\Lambda}}. \\end{equation}  If we simultaneously observe CMBR in two receiver  channels with frequency difference $\\Delta \\nu $, the SSF with an angular size larger than some critical $\\theta_c $ are seen at both receiver channels. If the size of the fluctuations is smaller, they are only seen in one of the receiver channels, another one will detect background noise only (Dubrovich,1982). In this case the correlation function of these two observations will be within noise to zero at small scales and different from zero at angular scales greater than $\\theta_c $. Then the difference signal of the two channels is not zero at the scales less than $\\theta_c$($L_c$). Frequency separation $\\Delta \\nu/\\nu$ of receiver channels determines a critical scale of proto-objects $L_c$ and scales $L < L_c$ remain in difference of two channels. Frequency resolution or the channel bandwidth $\\Delta \\nu/\\nu$ determines the optimal scale $L_{opt}$ which has the best S/N at given z. The channel bandwidth should be a few times less than frequency separation and $L_{opt} < L_c$.  The above consideration is correct for objects  freely expanding over the Universe in accordance with the Hubble law. But at the sufficiently late stages the deceleration of their expansion takes place and contraction begins due to self-gravitation. In this case the width of the line can be considerably smaller, and the amplitude larger (Zel'dovich, 1978).  Fig.1. shows expected $\\tau $($\\alpha$, $n_H$,z),  $\\Delta T/T$($\\alpha,n_H,z$), $\\phi(\\Omega_m, \\Omega_L,z)$, $\\theta(L,z)$, $\\Delta \\nu /\\nu (L,z)$ for HeH$^{+}$ ($\\lambda_0$=1.49 10$^{-2}$ cm, d=1.66 deb ) and K=1. According to Fig.1 an optimistic level of $\\Delta T/T=2\\times10^{-5}-2\\times10^{-6}$ may be achieved with $\\alpha =10^{-9}- 10^{-10}$.  Expected observational frequencies f(z) are given in Fig.1. for rotational transitions(I=0) of LiH, H$_{2}$D$^{ + }$, HeH$^{ + }$ HD$^{ + }$. For $\\phi$($\\Omega_m$, $\\Omega_L,z)$ calculation we used $\\Omega_m$ and $\\Omega_L$ which are discussed in several Cosmological models and recent CMBR experiments (Ruhl et al, 2004, Miller, 2004).  \\begin{figure} \\caption{$\\phi(\\Omega_m, \\Omega_L,z)$, $\\tau (\\alpha, n_H,z)\\theta(L,z)$, $\\Delta \\nu /\\nu (L,z)$ $\\Delta T/T$, $f(z)$ for HeH$^{ + }$} \\includegraphics[scale=1]{fsim.eps} % \\end{figure}  Let us limit possible physical proto-object scales L  in the Early Universe: 2 Mpc $>L>$0.002 Mpc. The most of Cosmological molecules considered may exist at z$<$200 and some at z$<$300. Then possible angular scales $\\theta$ for different $\\Delta \\nu/\\nu$, L and z are calculated with (7),(8) and presented in Table 2. Table 2 shows that modern angular scales $\\theta=5'-10'$ are still reasonable to search for large scale SSF at z=20-30 with $\\Delta \\nu/\\nu$=1-2$\\%$. Possible observational frequencies $f(z)$ for HeH$^{ + }$(I=0) are given in Table 3. Most of primordial molecules can radiate at several rotational transitions I that leads to SSF formation at set of frequencies $f_{i}(z)$. LiH has radiating transitions at I=0,1,2,3,4,5,6,7. The optical depth $\\tau$ in molecular lines depends on I and in the case of thermodynamic equilibrium for optically thin layer is (Dubrovich,1982):  \\begin{equation} \\tau_{I,I+1}= \\tau_0 \\cdot (I+1)\\cdot exp(-T_0I(I+1)/T) \\end{equation} where optical depth $\\tau_0$=$\\tau$(I=0) is given by (3), $T_{0}$ -  temperature correspondent to energy of transition I=0-1, T- temperature of radiation at given z. For LiH extremum of $\\tau$ lies at I=6 ($\\lambda_0 =0.95 \\cdot 10^{-2}$ cm) and $\\tau (I=6)/ \\tau_0 $=2-4. Possible observational frequencies $f_i(z)$ for LiH(I=6) are given in Table 4. Expected $\\Delta T/T$ level for SSF caused by LiH  with existing estimates of LiH abundance is less than $10^{ - 5}$ that can not be achieved in this experiment but we can set an upper limit of $\\Delta T/T$ for LiH at z=30 as well.  \\begin{table} \\caption{Angular and spatial scales for HeH$ ^{ + }$(I=0)}\\label{tab1}  \\begin{tabular}{lllllllll} \\hline $\\Delta \\nu/\\nu,$\\% & $z_{min}$ & $z_{max}$ & $L_{max}$,Mpc & $L_{min}$,Mpc &$\\theta_{max}(')$ & $\\theta_{min}(')$ \\\\ \\hline 0.1 & 5 & 200 & 0.5& 0.0025& 1.4 &0.14 & &\\\\ 0.5 & 10 & 200& 1&0.015&4.3 &0.85 & & \\\\ 1 & 10 & 200& 2&0.025 & 8.6 & 1.4 & & \\\\ 2 & 15 & 200& 2&0.05 & 11.5 & 0.3 & & \\\\ 4 & 25 & 200& 2& 0.1& 17.3 & 5.7 & & \\\\ 5 & 30 & 200& 2& 0.15 & 20.1 & 8.5 & &\\\\ 6 & 35 & 200& 2 & 0.17 & 22.9 & 9.6 & & \\\\ 7 & 40 & 200& 2 & 0.2 & 25.8 & 11.3 & & \\\\ 10 & 50 & 200& 2 & 0.25 & 31.4 & 14.2 & & \\\\ \\hline \\end{tabular} \\end{table}  \\begin{table} \\caption{ Observational frequencies f(z)in GHz for HeH$ ^{ + }$(I=0) }\\label{tab1} \\begin{tabular}{llllllllllllllllllllllllllllll} \\hline \\ z&5&10&15 &19&20&22&52&60&65& \\\\ \\hline f &335.57&183.30&125.84&100.67&95.88&87.54&37.99&33.00&30.51& \\\\ \\hline \\end{tabular} \\end{table} \\begin{table}  \\caption{Possible observational frequencies f(z) in GHz for LiH(I=6) }\\label{tab1} \\begin{tabular}{llllllllllllll} \\hline \\ z&10&15&20&30&32&34&35&80&101& \\\\ \\hline f&287.08&197.37&150.38&101.87&98.68&90.23&87.71&38.9&30.96&\\\\ \\hline \\end{tabular} \\end{table}  \\begin{table} \\caption{Spatial scales L of peaks in CMBR power spectrum for different z} \\begin{tabular}{llllllll} \\hline $z$ & $L_1$, Mpc & $L_2$,Mpc & $L_3$,Mpc & $L_4$,Mpc \\\\ \\hline 1000& 0.22 & 0.11 & 0.073& 0.043& \\\\ 100& 1.99 & 0.99 &0.66 & 0.39 & \\\\ 50& 3.72 & 1.86& 1.24 & 0.74 & \\\\ 30 & 5.79 & 2.91 & 1.93 & 1.16 & \\\\ 20& 8.08 & 4.04 & 2.69 & 1.62 & \\\\ \\hline \\end{tabular} \\end{table}  If we realize a channel bandwidth of $\\Delta \\nu/\\nu$=2{\\%}  in 3 mm band preferable or optimal linear scales for SSF search are L=1.5 Mpc for HeH$ ^{ + }$(I=0) at z=20  and L=0.75 Mpc for LiH(I=6) at z=30, correspondent angular sizes according to (7) are 10.7$'$ and 7.5$'$.  Recent WMAP polarization measurements have shown strong probability of an early reionization scenario with $z_{r}=20(+10/-9)$ (Kogut et al, 2003, Page et al, 2006) due to Pop III stars (Wyithe, Loeb,2003) in Galaxy halos. SSF detection at z=20-30 may provide us with important information about ionization history of the Universe.  The question is which maximum proto-object scales  are reasonable, could exist at z=20-30 or are connected with some known scales in the Universe. Let us compare them with spatial scales of primary CMBA. The angular scale of the first peak in primary CMBR power spectrum $\\theta_1 \\approx$ 50$'$, the following CMBA peaks are $\\approx$ 1/2$\\theta_1$,1/3$\\theta_1$, and 1/5$\\theta_1$ (see Fig.2 down). Then possible spatial scales L of 1-4 peaks in CMBR power spectrum for different $z$ are given in the Table 5 with (7) and (6).  The above consideration shows that at $\\Delta $T/T=2$\\cdot $10 $^{-5}$  level  and with $\\Delta \\nu/\\nu$=2{\\%} we can search for SSF which are the result of large scale proto-object motion with peculiar velocity relative to CMBR at z=20-30. These possible proto-objects are still less than spatial scale of the fourth peak in CMBR power spectrum (Table 5) which might be related to a proto-galaxy scale. With less $\\Delta \\nu/\\nu$, higher sensitivity and multipole resolution we can approach to smaller spatial scales of proto-objects up to proto-star scales.  SSF firstly proposed as a new phenomena in (Dubrovich,  1977) is one of most promising secondary CMBA. SSF is  non-isotropic spectral distortions of CMBR which are to be present at z=300-10. Homogeneous and isotropic spectral distortions of CMBR may be present at z=6000-1000 due to primordial Helium and Hydrogen recombination (Dubrovich, 1975, Dubrovich, Stolyarov, 1995, Dubrovich, Shahvorostova, 2004, Dubrovich, Grachev, 2004). Non-isotropic distortions of CMBR similar to SSF may be present in this epoch at lower $\\Delta T/T$ level as well(Dubrovich, 1977/2). Calculation of this effect in more details at z=1000-1100 is done in (Rubino-Martin, Mernandez-Monteagudo and Sunyaev, 2005). ",
        "conclusions": "Observational effects caused by  interaction of primordial molecules with CMBR seem to be most promising in investigating of Dark Ages epoch of the Early Universe and ionization history of the Universe. Strategy and concept of the experiment are given. Simulation of the experiment in an atmospheric transparency window at 3 mm with MSRT equipped with a 7x4 beam cryo-microbolometer array gives us the expected sensitivity of 50-55 $\\mu$K in the differential CMBR maps near NCP with angular resolution 6$'$-7$'$ (l=1200-1500) and frequency resolution \\textbf{$\\Delta $}$\\nu $/$\\nu $=2{\\%}. The expected integral sensitivity in the experiment is $\\Delta $T/T=2$\\cdot $10 $^{\\mbox{--}5}$ which is near to optimistic estimates of SSF caused by molecules HeH$^{+}$ at z=20-30. Simulation shows that SSF may be recognized in the angular power spectrum when S/N in single frequency CMBR maps is as small as 1.17  for a white noise. Such an experiment will give us a possibility to set upper limits in SSF observations in MM band and prepare future SSF experiments."
    },
    "0711/0711.4630_arXiv.txt": {
        "abstract": "We consider the no-boundary proposal for homogeneous isotropic closed universes with a cosmological constant and a scalar field with a quadratic potential. In the semi-classical limit, it predicts classical behavior at late times if the initial scalar field is more than a certain minimum. If the classical late time histories are extended back, they may be singular or bounce at a finite radius. The no-boundary proposal provides a probability measure on the classical solutions which selects inflationary histories but is heavily biased towards small amounts of inflation. This would not be compatible with observations. However we argue that the probability for a homogeneous universe should be multiplied by $\\exp(3N)$ where $N$ is the number of e-foldings of slow roll inflation to obtain the probability for what we observe in our past light cone. This volume weighting is similar to that in eternal inflation.  In a landscape potential, it would predict that the universe would have a large amount of inflation and that it would start in an approximately de Sitter state near a saddle-point of the potential. The universe would then have always been in the semi-classical regime. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.4589_arXiv.txt": {
        "abstract": "The interaction between an accreting satellite and the Andromeda galaxy (M31) has been studied analytically and numerically, using a high-resolution $N$-body simulation with $4\\times10^7$ particles. For the first time, we show the self-gravitating response of the disk, the bulge, and the dark matter halo of M31 to an accreting satellite. We reproduce the Stream and the shells at the East and West side of M31, by following the evolution of the collision 4 Gyr into the future, and show that recently discovered diffuse arcs on the south side of the minor axis are the remnants of a similar collision that occurred 3-4 Gyr earlier than the Stream event. The present day integrity of the M31 disk constrains the satellite mass to be $M \\leq 5\\times10^9 M_\\odot$. The stars that were originally in the center of the satellite are now in the east shell. Therefore, observations in this region might reveal additional clues about the nature of satellite, such as the central core and any metallicity gradient. ",
        "introduction": "Large spiral galaxies like the Andromeda galaxy (M31) are believed to have formed in part from the merger of many less massive galaxies, and their halos are modeled as the debris of accreted dwarfs (e.g. Bullock \\& Johnston 2005). The giant stream of Ibata et al. (2001) and the starcount maps of Ferguson et al. (2002) establish the complex merger origin of the M31 halo population; evidence for coherent structures now extends to $\\approx 100 $ kpc (Ibata et al. 2007). Further support comes from the spectroscopic side e.g. the disk population of Ibata et al. (2005). Deep HST/ACS photometry shows that three deep fields in the M31 halo exhibit unequivocal evidence for intermediate age populations (Brown et al. 2006). Starcount maps and radial velocities of red giant stars near M31 exhibit a giant stellar stream to the south of this galaxy, as well as giant stellar shells to the east and the west of M31's center (Ibata et al. 2001; Ferguson et al. 2002; McConnachie et al. 2003; Ibata et al. 2004, 2005; Guhathakurta et al. 2006). Recently Ibata et al. (2007) have carried out a deep, wide field photometric survey of the minor axis toward M33. They find a series of arc-like structures parallel to the major axis of M31 toward M33, and spaced by $\\approx 10$ kpc.  $N$-body simulations of the interaction between the progenitor of the giant stream and M31 suggest that the stream and shells are the tidal debris formed in the last pericentric passage of a satellite on a radial orbit (Ibata et al. 2004; Font et al. 2006; Geehan et al. 2006; Fardal et al. 2006, 2007). However, these models have always assumed a fixed potential to represent the influence of M31, and no models consider the possible effect of a live ($N$-body) disk. The discovery of the widespread intermediate age halo population has raised the question of whether these stars belong to the infalling satellite or are instead disk stars that were ejected in the stream collision or an earlier interaction. Present constraints on the progenitor mass are weak; consequently we look to the present day thickness of the disk to limit the progenitor mass. In addition, the newly discovered large scale arcs (Ibata et al. 2007) have not yet been modeled.  These issues motivate us to explore the effect of the accreting satellite using the first self-consistent, $N$-body model of M31 that has a disk, bulge, and dark matter halo. ",
        "conclusions": "\\begin{figure} \\begin{center} \\includegraphics[scale=0.8]{fig4.eps} \\caption{ Phase space distribution of the satellite particles. The color of each particle corresponds to the initial total energy with the unit of $10^{51}$ erg, with more negative values having originally been deeper in the potential. Note that most of the particles that were originally deep in the potential of the satellite are found at present day in the East shell. } \\label{f4} \\end{center} \\end{figure}  In this work, we follow the approach of earlier studies (Ibata et al. 2004; Font et al. 2006; Geehan et al. 2006; Fardal et al. 2006, 2007) in simulating the infall of a massive satellite into M31. However, we add a live $N$-body disk and bulge, and this permits us to ask under what conditions this collision might eject stars from these populations into the halo and quantifies the effect of these populations on the satellite. We also follow the evolution of this encounter for 4 Gyr beyond the present. As is the case in Fardal et al. (2007), our simulations successfully reproduce the giant stream and (for a similar but much older collision) the apparent shell structures that are observed in the starcount maps. The large number of test particles in our satellite affords a good comparison with observations (Koch et al. 2007). We conclude that the event responsible for the giant stream most likely has not populated the halo with ejected disk stars; hence the intermediate age metal rich populations observed in deep HST fields by Brown et al. (2006) are not disk stars ejected {\\it by the stream progenitor.} An ancient radial collision involving a more massive progenitor ($\\approx 10^{10} M_\\odot$), or a collision with a different approach angle, might have been capable of ejecting disk stars into the halo. The present-day thickness of the disk following this collision constrains the progenitor mass to be $\\leq 5\\times10^9 M_\\odot$.  Our simulation gives additional predictions, such as the existence of high velocity stars in the central region of M31. Figure 4 shows that the tidal debris in the M31 central bulge has velocity $ > 300 {\\rm ~km s}^{-1}$. The particles that were originally deepest in the potential (see color-coding in $10^{51}$ erg units) are likely, at present, to lie in the East shell. Observations in this region might reveal additional clues about the nature of satellite, such as the central core and the metallicity gradients (cf. Mori et al. 1997, 1999). It is also noteworthy that $>3$ Gyr after the collision, the debris more or less uniformly fills the halo.  Simulations such as Font et al. (2007) may wish to consider the disk, bulge, halo, and satellite stellar populations and their effect on the long term evolution of streams; encounters with stars in the disk and bulge may disrupt the coherence of streams and scatter the stars in a manner analogous to disk and bulge shocking of globular clusters.  A series of arc-like structures discovered by Ibata et al. (2007) are similar to the large-scale shell structures shown in Fig. 1f-1h. We conclude that these arcs may be the fossils of previous radial mergers several Gyr in the past (cf. Hernquist \\& Quinn 1988).  Our model predicts that similar arcs should be found on the opposite side of the M31 disk, in the Northwest quadrant, if in fact they are the shells of ancient radial infall collisions. Our model also offers a natural explanation for stars with $R>100$ kpc that are observed in the M31 halo near the systemic velocity of M31 (Gilbert et al. 2006; Koch et al. 2007). Some of these stars at $R\\approx 165$ kpc are found to have [Fe/H]=$-1$ by Koch et al. (2007).  We propose that these stars are {\\it not} a pressure supported halo but are rather the remnants of ancient radial collisions (note that in Fig. 1, the final extent of the collision is of order 100 kpc). Future studies will examine collisions from different attack angles, stellar disks, and the fate of dust and gas."
    },
    "0711/0711.0775_arXiv.txt": {
        "abstract": "{Bright-rimmed clouds (BRCs) are potential examples of triggered star formation regions, in which photoionisation driven shocks caused by the expansion of HII regions induce protostellar collapse within the clouds.} {The main purpose of the paper is to establish the level of star formation occuring within a known set of BRCs. A secondary aim is to determine the extent, if any, to which this star formation has been promulgated by the process of photoionisation triggering.} {A primary set of observations is presented obtained with submillimeter SCUBA observations and archival data from near-IR and mid- to far-IR have been explored for relevant observations and incorporated where appropriate.} {SCUBA observations show a total of 47 dense cores within the heads of 44 observed BRCs drawn from a catalogue of IRAS sources embedded within HII regions, supportive of the scenario proposed by RDI models. The physical properties of these cores indicate star formation across the majority of our sample. This star formation appears to be predominately in the regime of intermediate to high mass and may indicate the formation of clusters. IR observations indicate the association of early star forming sources with our sample. A fundamental difference appears to exist between different morphological types of BRC, which may indicate a different evolutionary pathway toward star formation in the different types of BRC.} {Bright-rimmed clouds are found to harbour star formation in its early stages. Different evolutionary scenarios are found to exist for different morphological types of BRC. The morphology of a BRC is described as type `A', moderately curved rims, type `B', tightly curved rims, and `C', cometary rims. `B' and `C' morphological types show a clear link between their associated star formation and the strength of the ionisation field within which they are embedded.  An analysis of the mass function of potentially induced star-forming regions indicate that radiatively-driven implosion of molecular clouds may contribute significantly toward the intermediate to high-mass stellar mass function.} ",
        "introduction": "Bright-rimmed clouds (BRCs) are small molecular clouds found at the edges of large HII regions. These clouds are potential examples of triggered star-forming regions, whereby shocks driven into the BRC by the photoionisation of the cloud exterior results in the collapse (and perhaps even the formation) of sub-critical molecular cores within the cloud (e.g.~\\citealt{Deharveng2005}; \\citealt{Elmegreen1998}; \\citealt{Sugitani1991}). The process of photoionisation-induced collapse is usually known as radiative-driven implosion or RDI  \\citep{Bertoldi1989,Lefloch1995,Lefloch1994}.  The RDI of molecular cores at the periphery of HII regions may thus be responsible for subsequent generation of star formation, amounting to a possible cumulative total of several hundred new stars per HII region \\citep{Ogura2002} and perhaps 15\\% or more of the low-to-intermediate mass stellar mass function \\citep{Sugitani1991}. IRAS-selected protostars associated with HII regions are systematically more luminous than those that are not \\citep{Dobashi2001}. This highlights a trend toward higher mass star formation and greater star formation efficiency. Confirming bright-rimmed clouds as star-forming and quantifying the nature and location of their star formation thus provides important insights into the clustered mode of star formation and the overall star formation efficiencies of molecular clouds.  \\citet{Sugitani1991}, hereafter referred to as SFO91, searched the Sharpless HII region catalogue \\citep{Sharpless1959} for bright-rimmed clouds associated with IRAS point sources, their selection was based upon the FIR colours thought to identify YSOs/protostars. Later, \\citet{Sugitani1994} -- SO94 -- extended their search to  include bright-rimmed clouds from the ESO(R) Southern Hemisphere Atlas. A total of 89 optically identified bright-rimmed clouds have been found to be associated with IRAS point sources. For brevity (and consistency with SIMBAD), we will refer to the combined  SFO91 and SO94 catalogues as the SFO catalogue. Whilst a few individual clouds from the SFO catalogue have been studied in detail \\citep[e.g.~][]{Lefloch1997,Megeath1997,Codella2001,Thompson2004a,Thompson2004b,Urquhart2004,Urquhart2006,Urquhart2007} and shown to harbour protostellar cores, the question of whether star formation is a common occurence within BRCs (irrespective of the formation mechanism) is still unresolved.  We are in the process of carrying out a star formation census of SFO bright-rimmed clouds in order to determine the current level of star formation within these sources and to relate it to their physical properties and morphologies with respect to their ionising star(s). In two previous publications we reported observations of the ionised gas surrounding the bright-rimmed clouds \\citep{Thompson2004,Morgan2004}. Here we report a Submillimetre Common User Array (SCUBA) imaging survey of the submillimetre continuum emission from the clouds, in order to reveal the presence of  compact, potentially star-forming, dust cores within the clouds and to contrast the properties of these cores with those found in other star-forming regions. We use NIR archival data from the 2 millimetre all-sky survey (2MASS) database to construct colour-colour diagrams to investigate the current level of star formation within these clouds and to identify YSOs.  We examine the kinematics and temperature of the dust cores revealed by our SCUBA observations to show that a number of the SFO bright-rimmed clouds are associated with active star formation. ",
        "conclusions": "Our main findings are as follows: \\begin{enumerate} \\item The SCUBA observations presented here generally show centrally condensed cores interior to the rims of BRCs. The morphology of the cores themselves suggests the presence of bound condensations, indicative of star formation. The morphology of the larger area, incorporating the entire bright rim, is, in general, supportive of the scenario seen in RDI models; a dense core at the head of an elongated column. Support for the interaction of multiple layers is given by observations at multiple wavelengths i.e. radio observations tracing the hot ionised layer bordering the PDR observed in the near infrared which borders the molecular material, the dense regions of which are seen at submillimetre wavelengths.  \\item The physical properties of the submillimetre cores were determined using modified blackbody analyses. The dust temperatures found for the sources are high compared to starless cores (average T$_{\\rm{d}}$ = 24 K), indicating the presence of internal stellar sources in these clouds. Based on their L$_{\\rm{submm}}$/L$_{\\rm{bol}}$ ratio these sources are consistent with being Class 0 protostars. In addition, the indicated stellar sources embedded within the clouds may be classified via the hypothesis that all illumination is due to a single illuminating source. This leads to spectral classifications of intermediate to high mass stellar sources within the cores, whether these sources are isolated singular sources or weaker sources, distributed in clusters.  \\item Observations of mid-infrared sources taken from the 2MASS further support the proposed scenario of early star formation across our sample by revealing the presence of young stars in the vicinity of the majority (89\\%) of our sample. The spectral type of the associated sources is not determinable from the current data, however the observed stars are known to be in the relatively early stages of their formation.  \\item Both the mass functions and correlations of source luminosity with ionisation field strength suggest some fundamental difference between the subsets of `A' type rims and `B' and `C' type rims. It is suggested that `B' and `C' type rims represent true `triggered' star formation as enabled by the RDI process. There is no reason to rule out the possibility that `A' type rims are also the result of triggering, though perhaps via a different mechanism i.e. the `collect and collapse' model \\citep{Elmegreen1977}. The correlation between ionisation field strength and source luminosity is highly indicative of `B' and `C' type rims forming with at least some degree of interaction with their UV surroundings. By increasing the size of our sample and incorporating other BRCs that have yet to begin their star formation this hypothesis may be investigated further.  \\end{enumerate}  $Acknowledgements$ The authors would like to thank both an anonymous referee and Malcolm Walmsley for their useful and illuminating comments. They have improved the quality of this work.  The authors would like to thank the JCMT support staff for a pleasant and productive observing run; Helen Kirk for assistance with data reduction issues and D.J. Pisano for his insight and helpful discussions.  The Digitized Sky Surveys were produced at the Space Telescope Science Institute under U.S. Government grant NAG W-2166. The images of these surveys are based on photographic data obtained using the Oschin Schmidt Telescope on Palomar Mountain and the UK Schmidt Telescope. The plates were processed into the present compressed digital form with the permission of these institutions.  This publication makes use of data products from the Two Micron All Sky Survey, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology, funded by the National Aeronautics and Space Administration and the National Science Foundation.  This research made use of data products from the Midcourse Space Experiment. Processing of the data was funded by the Ballistic Missile Defense Organization with additional support from NASA Office of Space Science.  This research has made use of the NASA/ IPAC Infrared Science Archive, which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration.  This research has made use of the SIMBAD database, operated at CDS, Strasbourg, France  This research has made use of NASA's Astrophysics Data System."
    },
    "0711/0711.4823.txt": {
        "abstract": "We use the UV-optical color magnitude diagram in combination with spectroscopic and photometric measurements derived from the SDSS spectroscopic sample to measure the distribution of galaxies in the local universe (z$<$0.25) and their physical properties as a function of specific star formation rate (\\ssfr) and stellar mass ($M_\\star$).  Throughout this study our emphasis is on the properties of galaxies on and off of a local ``star-forming sequence.''  We discuss how the physical characteristics of galaxies along this sequence are related to scaling relations typically derived for galaxies of different morphological types.   We find, among other trends that our measure of the star formation rate surface density, \\sfrkpc ~is nearly constant along this sequence.  We discuss this result and implications for galaxies at higher redshift.   For the first time, we report on measurements of the local UV luminosity function versus galaxy structural parameters as well as inclination.   We also split our sample into disk-dominated and bulge-dominated subsamples using the $i$-band Sersic index and find that disk-dominated galaxies occupy a very tight locus in \\ssfr~ vs. \\Mstar~ space while bulge-dominated galaxies display a much larger spread of \\ssfr~ at fixed stellar mass.  A significant fraction of galaxies with \\ssfr~ and \\sfrkpc~ above those on the ``star-forming sequence'' are bulge-dominated.  We can use our derived distribution functions to ask whether a significant fraction of these galaxies may be experiencing a final episode of star formation (possibly induced by a merger or other burst), soon to be quenched, by determining whether this population can explain the growth rate of the non-star-forming galaxies on the ``red sequence.''  We find that this is a plausible scenario for bulge-dominated galaxies near the characteristic transition mass under reasonable assumptions regarding quenching timescales.  Similarly we use this technique to estimate the rate of mergers/starbursts that take galaxies off of the star-forming sequence and show that the implied merger rates are consistent with local measurements. ",
        "introduction": "What determines the star formation rate (SFR) of a galaxy?   Ample evidence suggests that it is the quantity and distribution of cold gas \\citep{Schmidt1959, Kennicutt1998a} and the gas-dynamical processes responsible for triggering, regulating or quenching new star formation.  In the context of a hierarchical clustering scenario for galaxy formation \\citep[e.g.][]{Kauffmann1993} these mechanisms are necessarily linked to the flow of dark and baryonic matter over a wide range of scales, densities and temperatures \\citep{Keres2005}.   Given this complexity it is intriguing that the integrated light from many galaxies can be explained using simple star formation histories (SFH)  \\citep[][although see also Kauffmann et al. 2006b]{Tinsley1968,Searle1973,Bruzual2003}.  Such work has led to an apparent understanding of star formation on cosmological scales \\citep{HopkinsA2006}, although accurate physical models embedded within a realistic framework for galaxy assembly \\citep[e.g. ][]{Stringer2007} are required to understand star formation in individual galaxies.   Measurements of the colors and structure of a galaxy should guide these models by providing insight into the connection between star formation and assembly.   The fact that galaxies in the local universe appear to show a remarkable correlation between their star formation history and their structure---disk-dominated galaxies show higher present to past-averaged star formation rates than bulge-dominated galaxies \\citep{Kennicutt1998b}---would appear to suggest a straightforward link, but we now know that the explanation must be non-trivial \\citep{De Lucia2006}.  A crucial component of these analyses is a quantitative and representative description of the galaxy population.  In this regard color-magnitude distributions and their derivatives have emerged as useful tools because they can be easily interpreted in terms of the star formation history and the stellar mass content and therefore are easily connected to models of galaxy assembly  and the build up of massive galaxies along the ``red sequence'' \\citep{Faber2007}.   In this paper, the second in a series, we explore how the UV-optical properties of a large sample of galaxies in the local universe can be used to understand the distribution of SFR and the connection with assembly history across the population.  We accomplish this by expanding on the analysis of the UV-color magnitude diagram \\citep[][hereafter Paper I]{Wyder2007} using observations of 26241 galaxies from the {\\it Galaxy Evolution Explorer (GALEX)} Medium Imaging Survey (MIS), combined with the Sloan Digital Sky Survey (SDSS) primary spectroscopic survey and  incorporating additional ``value-added'' data related to the morphology/structure, star formation history and dust attenuation in each galaxy.  Our analysis has many similarities to recent studies conducted using SDSS \\citep{Blanton2003c,Kauffmann2003a, Kauffmann2003b, Shen2003}, local surveys \\citep{Driver2006, Jansen2001} and high-z investigations \\citep{Bell2005, Noeske2007a, Noeske2007b}, although it extends those studies in several ways, described below.  As discussed in \\citetalias{Wyder2007}, a notable feature of the UV-optical color diagram is the very wide separation between the peaks of the blue and the red galaxy populations.  Of central importance is the strongly peaked locus of star-forming blue galaxies that has been variously called the ``blue sequence'' \\citet{Blanton2006} and the ``main sequence'' \\citet{Noeske2007a}.  \\citetalias{Wyder2007}, \\citet{Noeske2007a} and \\citet[][hereafter S07]{Salim2007} show that the majority of star-forming galaxies of a given stellar mass possess a narrow range of SFR, a result already noted by \\citet{Brinchmann2004}, \\citet{Feulner2006} and \\citet{Cattaneo2007}.  This stands in marked contrast to the optical color-based view, which emphasizes a tight ``red sequence'' and a scattered ``blue cloud'' \\citep{Bell2005}.    It suggests strong similarities among star-forming galaxies and a greater diversity of SFR for those galaxies that optically appear  ``red-and-dead.''  In many respects, this alternate view is reminiscent of the progression from the Hubble classification scheme, with a rich description of spirals and only a few elliptical classes, towards work in recent years that revealed that ellipticals possess a greater range of structure (at low and high spatial frequency) than originally thought \\citep{de Zeeuw1991}.   \\begin{figure*}[p] \\plotone{fig1.eps} %\\plotone{figv3/morph_magnitude.ps} \\caption{Main galaxy sample bivariate number density distribution: {\\it Top}: $NUV$ magnitude and $FUV-NUV$ color vs. log $n_i$. {\\it Bottom}:  $NUV$ and $FUV-NUV$ vs. log $\\mu_\\star$.  Contours enclose 50\\% and 90\\% of the distribution, with outliers plotted individually.} \\label{fig:morph_magnitude} \\end{figure*}   \\begin{figure*} \\plotone{fig2.eps} %\\plotone{figv3/morph_params.ps} \\caption{Main galaxy sample bivariate number density distribution. {\\it Top}: $i$-band concentration and $NUV-r$ color vs. log $n_i$. {\\it Bottom left}: $g-r$ color vs. $NUV-r$. {\\it Bottom right}: log $\\mu_\\star$ vs. log $n_i$.  Contours enclose 50\\% and 90\\% of the distribution, with outliers plotted individually.} \\label{fig:morph_params} \\end{figure*}     Here we quantify the structural and physical properties of a local ``star-forming sequence'' (SF sequence) defined by a relationship between stellar mass and star formation rate and use it to understand the characteristics of the dominant galaxy population including the slope of the sequence itself.  Some of this analysis is quite complementary to the work of \\citetalias{Salim2007}.  We also focus on the distribution disk-dominated and bulge-dominated galaxies across the full range of \\ssfr~and \\Mstar~and investigate the structure of outliers to the SF sequence.  While in \\citet[][hereafter Paper III]{Martin2007}, specific attention has been given to galaxies in the intermediate region between the blue and red galaxy population (the ``green valley''), our ultimate focus will be on galaxies with specific star formation rates higher than those on the SF sequence.  This is the population in the \\ssfr~vs. \\Mstar~diagram that has evolved most dramatically since z$\\sim$1.  We briefly describe the outline of this paper.  After presenting the data in section 2 we investigate in section 3 the physical properties of galaxies on and off of the SF sequence. In section 4 we investigate the relationship between star formation history and structure and its connection with the evolution of galaxies on and off of the SF sequence.   Throughout this paper, we make use of the flat $\\Lambda$CDM cosmology with $H_0$ = 70 km s$^{-1}$ Mpc$^{-1}$ and $\\Omega_\\Lambda =0.7$. ",
        "conclusions": "We have generated a catalog of galaxies in the local (z$<$0.25) universe with a combination of UV-optical photometry, spectroscopic measures, structural parameters, and value-added and physical quantities, and have used it for an investigation into the distribution of star formation across galaxies of different morphologies and stellar masses.  Our chief results are as follows,  1.  We have derived a new set of physical properties of the galaxies in our samples, including star formation and stellar mass rates and surface densities, dust attenuations, and gas fractions.  Our measurements incorporate a slightly modified prescription for dust attenuation, designed to use the best available data to derive star formation rates across the whole galaxy sample.  In general, this follows most closely the approaches described in \\citet{Johnson2007} and \\citet{Kauffmann2003a} but ultimately we hope to develop it as a refinement over current methodology.  2.  For the first time we have measured the local UV luminosity function against galaxy structural parameters as well as inclination.  Among our key results is that we have shown that the fraction of intermediate and early-type galaxies is highest for the most UV luminous galaxies, dropping off to low fractions for the least luminous galaxies.  3.  Throughout this study our emphasis has been on the properties of galaxies on and off of a local ``star-forming sequence'' defined by $\\log \\ssfr = -0.36 \\log M_\\star - 6.4$.   We find, among other trends, that our measure of the star formation rate surface density, \\sfrkpc~(measured within $r_{u,50}$) is nearly constant along this sequence.  4.   We have split our sample into disk and bulge-dominated galaxies using the $i$-band Sersic index, and find that disk galaxies occupy a very tight locus in SFR vs. $M_\\star$ space, while bulge-dominated galaxies display a much larger spread of SFRs at fixed stellar mass.  In particular, a significant fraction of galaxies with SFR and \\sfrkpc~ above those on the ``star-forming sequence'' are bulge-dominated.  5. We have used our derived distribution functions to ask whether a significant fraction of these galaxies may be experiencing a final episode of star formation (possibly induced by merger of other bursts), soon to be quenched, by determining whether this population can explain the growth rate of the non-star-forming population We find that this is a plausible scenario for bulge-dominated galaxies near the characteristic transition mass under reasonable assumptions regarding quenching timescales.  We use this technique to estimate the rate of mergers/starbursts that take galaxies off of the star-forming sequence and show that the implied merger rates are consistent with local measurements.   %% Included in this acknowledgments section are examples of the %% AASTeX hypertext markup commands. Use \\url without the optional [HREF] %% argument when you want to print the url directly in the text. Otherwise, %% use either \\url or \\anchor, with the HREF as the first argument and the %% text to be printed in the second."
    },
    "0711/0711.2293_arXiv.txt": {
        "abstract": "\\noindent{\\bf Supernova 2006gy in the galaxy NGC 1260 is the most luminous one recorded \\cite{2006CBET..644....1Q,2006CBET..647....1H,2006CBET..648....1P,2006CBET..695....1F}. Its progenitor might have been a very massive ($>100$\\,\\msun) star \\cite{2006astro.ph.12617S}, but that is incompatible with hydrogen in the spectrum of the supernova, because stars $>40$\\,\\msun\\, are believed to have shed their hydrogen envelopes several hundred thousand years before the explosion \\cite{2005A&A...429..581M}. Alternatively, the progenitor might have arisen from the merger of two massive stars \\cite{2007ApJ...659L..13O}. Here we show that the collision frequency of massive stars in a dense and young cluster (of the kind to be expected near the center of a galaxy) is sufficient to provide a reasonable chance that SN 2006gy resulted from such a bombardment. If this is the correct explanation, then we predict that when the supernova fades (in a year or so) a dense cluster of massive stars becomes visible at the site of the explosion. } ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.0069_arXiv.txt": {
        "abstract": "The fragmentation process of primordial-gas cores during  prestellar collapse is studied using three-dimensional nested-grid hydrodynamics. Starting from the initial central number density of $\\nc\\sim10^3\\cm$, we follow the evolution of rotating spherical cores up to the stellar density $\\nc \\simeq 10^{22}\\cm$. An initial condition of the cores is specified by three parameters: the ratios of the rotation and thermal energies to the gravitational energy ($\\beta_0$, and $\\alpha_0$, respectively), and the amplitude of the bar-mode density perturbation ($\\ap$). Cores with rotation $\\beta_0 > 10^{-6}$ are found to fragment during the collapse. The fragmentation condition hardly depends on either the initial thermal energy $\\alpha_0$ or amplitude of bar-mode perturbation $\\ap$. Since the critical rotation parameter for fragmentation is lower than that expected in the first star formation, binaries or multiples are also common for the first stars. ",
        "introduction": "More stars are in binaries or multiples than in singles in the solar vicinity (Duquennoy \\& Mayor 1992; Fischer \\& Marcy 1992). Is this also true of the first-generation stars? Being very massive ($\\ga 100M_{\\sun}$; see e.g., Bromm \\& Larson 2004), first-star binaries, if they are formed, can contribute greatly as progenitors of gamma-ray bursts (Bromm \\& Loeb 2006; but see Belczynski et al. 2006 for the opposite view) or gravitational wave sources (Belczynski \\etal 2004). However, binary formation in primordial gas has sometimes been considered unlikely, based on the fact that no cooling mechanism induces fragmentation of high-density ($\\ga 10^{4}{\\rm cm^{-3}}$) cores (e.g., Ripamonti \\& Abel 2004).  Rotation can change this picture. Unlike the case of present-day star formation, where the angular momentum is effectively transported by magnetic braking \\citep{basu94} or outflow \\citep{tomisaka02}, in  first-star formation, the magnetic field is extremely weak \\citep[e.g.,][]{takahashi05,ichiki06} and thus expected to play little role in the angular momentum transport. With conserved angular momentum, the effect of rotation is enhanced in the course of the collapse of the core, and a disk is expected to form. Do such disks fragment to binaries or multiples? To answer this question, by three-dimensional hydrodynamics, Saigo \\etal (2004) demonstrated that primordial-gas cores with rapid rotation indeed fragment into binaries. Nonetheless, their analysis is not satisfactory in the following respects: First, they started calculation at rather high density $\\sim 10^{10} {\\rm cm^{-3}}$, and thus were forced to assume somewhat ad hoc initial conditions of high angular-momentum cores. In practice, the angular momentum might have been transferred by the non-axisymmetric structure before reaching this density. In addition, they presented only one model of a core experiencing fragmentation. Recently, Clark, Glover, \\& Klessen (2007) performed three-dimensional hydrodynamical simulation of star formation in a low-metallicity turbulent medium and found that rotation causes fragmentation even in metal-free gas, as suggested by Saigo et al. (2004). However, they do not discuss fragmentation conditions of the cores in their analysis.    In this paper, we calculated the evolution of primordial-gas rotating cores for twenty orders of magnitude in the density contrast, starting from the formation of dense cores ($n\\simeq 10^3\\cm$) up to stellar-core formation ($n\\simeq 10^{22}\\cm$). We adopt a realistic equation of state derived by Omukai et al. (2005), rather than a simple polytropic relation $P\\propto \\rho^{1.1}$ used by Saigo et al. (2004). To determine the fragmentation conditions for the primordial-gas cores, we calculated 38 models with different values of rotation, thermal energy, and amplitude of non-axisymmetric perturbation. Using the initial ratio of rotation to gravitational energies $\\beta_0$, the conditions for a primordial-gas core to fragment are $\\beta_0 > 10^{-6}$, while dependence of these conditions on either the thermal energy or the amplitude of non-axisymmetric perturbation is very weak. These conditions are easily satisfied in the first-star forming cores, and thus binary formation is also important for first stars.  The following is the plan of this paper. In this paper, after first summarizing the basics of our model in \\S 2, we describe the numerical method  in more detail in \\S 3. We present the results for the initially quasi-stable cores in \\S 4, and those for highly gravitationally unstable cores in \\S5. The fragmentation criterion is examined in \\S 6. Finally in \\S 7, summary and discussions are presented. ",
        "conclusions": "We have studied the gravitational collapse of zero-metallicity star-forming cores for different initial conditions, which are specified by the following three parameters, i.e., thermal energy, rotation energy, and amplitude of non-axisymmetric perturbations. Since the thermal energy of star-forming cores is expected to be similar to the gravitational energy, particular attention is paid to such cores. Initial conditions of the cores where fragmentation occurs are surveyed. We found that the cores with initial rotation energy higher than $10^{-6}-10^{-5}$ times the gravitational energy fragment into a few pieces. Dependence on the other parameters is weak and is not observed clearly in our calculation. The physical explanation for this fragmentation condition is also presented.  How much angular momentum do initial cores have in reality? According to cosmological simulations by \\citet{yoshida06}, first-star forming cores have the rotation parameter $\\beta\\simeq 0.1$. In discussing the fragmentation condition, the rotation parameter near the center should be used rather than that averaged over the core. Matter in the outer part tends to have larger rotation energy than gravitational energy, the value of the rotation parameter $\\beta$ becomes larger if outer material is included in calculations. Even though the central value of the rotation parameter might be significantly smaller than 0.1 owing to this effect, it is probably higher than our threshold value for fragmentation ($10^{-5}-10^{-6}$) with a wide margin. Therefore, most of the cores are expected to fragment and proto-binaries are formed even in primordial-gas clouds.    In Figure~\\ref{fig:12}, we plot the separations between fragments $R_{\\rm sep}$ at the end of the calculation against the fragmentation epoch $n_{\\rm f}$ for all the models where fragmentation is observed. The large crosses show ranges of fragmentation epochs (horizontal) and fragmentation scales (vertical) for models with $\\beta_0 >0.02$ (dotted line) ($\\beta_0<0.02$, solid line, respectively). The Jeans length (left axis) and gas temperature (right axis) are also shown. We can see that the separation between fragments is approximately equal to the Jeans length at the fragmentation. Another striking feature is that fragmentation densities are clustered in two epochs of $n \\approx 10^{12}\\cm$ and $10^{21}\\cm$. Cores with rapid rotation ($\\beta_0 > 0.02$; symbols '$+$') fragment in the earlier epoch ($\\nc \\simeq 10^{10}-10^{14}\\cm$) with larger separation $\\simeq 10-1000$\\,AU, while those with slower rotation ($\\beta_0 < 0.02$; '$\\times$') in the later epoch ($\\nc \\simeq 10^{18}-10^{21}\\cm$) with smaller separation $\\simeq 0.01-1$\\,AU. These two fragmentation epochs have their origin in thermal evolution of primordial gas. The lower-density fragmentation is a result of a temperature drop owing to the three-body H$_2$ formation at $\\nc \\simeq 10^{10}\\cm$, while the higher-density one results from another temperature drop at $\\nc \\simeq 10^{17}\\cm$, owing to the H$_2$ collision-induced cooling. Recall that a quasi-rotation supported disk can continue to collapse for $\\gamma <1$ while conserving angular momentum. Thus even if collapse of the disk is about to halt, the disk can collapse further owing to the temperature drops. Such disks stop their collapse and fragment when the temperature starts increasing again. It is known that these two cooling mechanisms do not cause fragmentation of the core by simple thermal instability (Omukai \\& Yoshii 2003; Ripamonti \\& Abel 2004). In our case, where a disk is present, these cooling mechanisms result in the fragmentation of the disk.  There is a caveat in our results on fragmentation epochs. We adopted the equation of state derived from one-zone calculation. For example, the temperature variation in one-zone models tends to be exaggerated: the two temperature dips are deeper than in hydrodynamical calculations (e.g., Omukai \\& Nishi 1998). In addition, in hydrodynamical models, temperature in the envelope tends to be higher than that in the center for the same density \\citep{whitehouse06,stamatellos07}. Despite these possible differences, we expect that the fragmentation condition itself (eq. \\ref{eq:beta0}) remains valid. In deriving this condition (Sec. 6), we only assumed that once a quasi-rotation supported disk forms, it eventually fragments. This is correct if the EOS has a sufficiently large density range with $\\gamma >1$, which is indeed the case in the primordial gas. In fact, in the context of present-day star formation, the same fragmentation criterion has been reported \\citep{matsu03,machida05b}. On the other hand, the fragmentation epoch and scale can be subject to change if EOS is altered.  Although the majority of cores are predicted to fragment before the stellar core formation, whether they finally evolve into bona fide binaries/multiples remains obscure. The separation between fragments changes significantly during the mass accretion phase as a result of angular momentum influx to the fragments. An accreted gas in a later stage tends to have a larger specific angular momentum, while spiral arms in a density pattern, which transfer angular momentum outward, might be invoked in such a phase. To answer whether fragments evolve into binaries/multiples or merge into a single star, we need to extend calculations over the entire accretion phase.  When more than two fragments appear (e.g., models 1 and 2 in Fig.~\\ref{fig:final}), some might be ejected from the parental core as a result of three-body interactions \\citep{bate02,bate03}. Ejected objects stop gaining mass by accretion thereafter. By this mechanism, a spectrum in mass of metal-free stars might be formed. For example, if an ejection event takes place in the very early accretion phase, the ejected object can eventually evolve into a metal-free brown dwarf. Or if a stellar core with its envelope of $\\sim 1\\msun$ is ejected, a low-mass metal-free star is the outcome. Considering the origin of the two hyper metal-poor ([Fe/H]$<-5$) stars (Christlieb et al. 2001; Frebel et al. 2005), \\citet{suda04} showed that their abundance patterns can be explained by nucleosynthesis and mass transfer in first-generation low-mass binary stars. In our scenario, such a system can be formed naturally. In some models (e.g., see Fig.~\\ref{fig:final}{\\it a} and {\\it b}), multiple pairs of binaries appear. If one such pair is ejected from the parental core, it would result in a low-mass metal-free binary system.  With neither ejection nor merging to a single star, our binary systems of fragments are expected to become very massive star binaries. Such systems are considered to be important sources of gravitational waves (Belczynski et al. 2004). Also they can be gamma-ray burst progenitors if formed in close binaries \\citep{bromm06,yoon06}.    In this study, we did not include magnetic effects. According to recent studies, however, a low level of magnetic field could be present even in the early universe \\citep{langer03,ichiki06}. If an initial primordial-gas core has $B>10^{-13.5}$\\,G of the magnetic field, a protostellar jet is launched \\citep{machida06c}, and this jet significantly affects the mass accretion process. Even with a lower level, a magnetic field can suppress fragmentation of the core, by transferring angular momentum via magnetic braking \\citep[e.g.,][]{machida07, price07}. So far those effects on primordial star formation have been largely overlooked. Further studies along this line are clearly needed in future."
    },
    "0711/0711.0891_arXiv.txt": {
        "abstract": "{ In the  context  of dwarf spheroidal galaxies   it is hard   to firmly disentangle a  genuine Blue Stragglers  (BSS) population from a normal young main (MS) sequence. This difficulty is persistent. For a sample   of 9 non-star  forming  Local Group  dwarf  galaxies we compute the ``BSS frequency'' ($F^{\\rm BSS}_{\\rm HB}$) and compare it with that found in the Milky Way globular/open clusters and halo. The comparison shows that $F^{\\rm BSS}_{\\rm HB}$ in dwarf galaxies, at any given ${ M_{V}}$, is always  higher than that in globular clusters of similar  luminosities.  Moreover, the  estimated $F^{\\rm  BSS}_{\\rm HB}$ for  the   lowest  luminosity  dwarf   galaxies  is  in excellent agreement with that observed in the Milky Way halo and open clusters. We conclude that  the low density, almost  collision-less environment, of our dwarf galaxy sample point to their very low dynamical evolution and consequent negligible production of collisional BSS. ",
        "introduction": "First identified by Sandage   \\cite{sandage53}, Blue Stragglers  are usually defined  as  a  hotter and  bluer   extension  of  normal main sequence  stars) The origin of BSS is sought as either {\\it primordial binaries} coeval with  the globular/open  cluster formation epoch,   or to a continuous production (in successive epochs) of {\\it collisional binaries} due to dynamical collisions/encounters  experienced    by single/binary stars throughout the life of the cluster.  Piotto et  al.  \\cite{piotto04} presented a  homogeneous compilation of $\\sim3000$ BSS (based on HST  observations of 56 globular cluster), and derived a significant and rather puzzling anti-correlation between the  BSS specific frequency and  the cluster total absolute luminosity (mass). That  is  to say   that more  massive  clusters are  surprisingly  BSS deficient, as if their higher collision rate had no correlation with the production of collisional BSS. Another puzzling observable  is  that the BSS  frequency in  Milky Way field  (Preston \\&  Sneden  2000) is  at  least an order of  magnitude larger than that of globular clusters.  In the context of  dwarf galaxies one  {\\em cannot} exclude  that blue plume stars may include genuinely young main sequence (MS) stars, i.e. a residual star forming    activity (e.g.  Held 2005,  and  references therein). A useful indicator of  the presence of a   recent star formation in  a galaxy  is the detection of a  vertical extension in correspondence of the red HB region.   Stars  forming this  sequence are usually  called vertical clump stars (VC, Gallart et al. 2005). These are helium-burning stars  of few hundred  Myr to $\\sim1$ Gyr old population whose progenitors are to be searched in the blue plume. Nevertheless, we note that Ferraro et al.  \\cite{ferraro99} identify a similar sequence in  the M80 globular cluster   and ascribe it to  the evolved-BSS population.  We  therefore  cautionly conclude  that   the detetction of VC star in a dwarf galaxy is not a clear-cut evidence of the presence of a young MS population.  To shed  light on the BSS-MS ambiguity  in  dwarf galaxies, we measure the BSS frequency in 9 non-star  forming dwarf galaxies and compare it with that derived in other stellar systems. Dwarf spheroidals/irregulars in which there is {\\em current} or recent ($\\le500$ Myr) star formation are not considered. Dwarf spheroidals/irregulars showing  evidence of  recent $\\le500$ Myr star formation (e.g.  Fornax  dwarf, Saviane et al. 2000), or young  MS  stars reaching the  horizontal  branch [e.g. Sagittarius irregular, Momany et al. \\cite{momany05}] level have been excluded. The  Canis Major dwarf  galaxy (Martin et   al. 2004 and Bellazzini et al.  2004)  was   not  included [see    arguments   in Momany et   al. \\cite{momany04w} and \\cite{momany06w}]. The  Carina dwarf (Monelli  et al.    2003),  is one of those  showing evidence  of star  formation in  recent epochs  ($\\sim1$ Gyr),  and we include in our  sample so as to compare  with other  dwarf galaxies of similar luminosities. ",
        "conclusions": "For a sample of 8 non star-forming dwarf galaxies,  we have tested the hypothesis that the blue plume  populations are made  of a genuine BSS population (as   that observed in   open and   globular clusters)  and estimated their frequency with respect to HB stars. Should this assumption be incorrect (and  the blue plume population is made of young MS stars) then one  would not expect an anti-correlation between the   galaxies  total luminosity   (mass) and  the  blue plume frequency, but rather a correlation between the two. Instead, and within the limits of this and similar analysis, we detect a statistically significant anti-correlation between $F^{\\rm BSS}_{\\rm HB}$ and $M_{V}$. Thus, should  a dwarf galaxy  ``obey''  the $F^{\\rm BSS}_{\\rm  HB}-{ M_{V}}$ anti-correlation   displayed by our     sample then  its blue    plume population is probably made of blue stragglers.   We  also    estimated  stellar  specific   collision parameter  (${\\rm log}~\\Gamma_{\\star}$: the number of collisions per star per year). The   mean collisional   parameter  of   the  9  studied   galaxies is $\\simeq-19$.  The  lowest value  is  that  for Sagittarius with  ${\\rm log}~\\Gamma_{\\star}\\simeq-20.2$, and this is due  to its very extended galaxy core. Compared with  the  mean value  of $-14.8$  for the  globular clusters sample   [see  lower   right  panel   of     Fig.~3  in  Momany     et al. \\cite{momany07}] the estimated number of  collisions per star per year in a dwarf spheroidal is $10^{-5}$ times lower.   To summarize, from Fig.~\\ref{fig1}  one finds that  $F^{\\rm BSS}_{\\rm HB}$ in  dwarf galaxies  is (i)  always higher than  that in  globular clusters,  (ii) very close, for the  lowest luminosity dwarfs, to that observed in the MW field and open  clusters, (iii) the Carina specific $F^{\\rm   BSS}_{\\rm HB}$ frequency     probably sets a threshold   for star-forming galaxies, and most interestingly, (iv)  shows a hint of a $F^{\\rm BSS}_{\\rm HB}-{ M_{V}}$ anti-correlation. This almost precludes the occurrence of  collisional binaries in dwarf galaxies, and one may  conclude that  genuine  BSS sequences  in dwarf galaxies are mainly made of primordial binaries.  Lastly, it  is  interesting  to note   how  the BSS frequency  in  the low-luminosity  dwarfs      and open  clusters    (${\\rm  \\log}(N_{\\rm BSS}/N_{\\rm  HB})\\sim0.3-0.4$) is very close  to  that derived for the Galactic halo (${\\rm \\log}(N_{\\rm BSS}/N_{\\rm HB})\\sim0.6$) by Preston \\& Sneden. The latter   value however  has  been derived  relying on  a composite sample of only  62 blue metal-poor  stars that are  (i) distributed at different line  of  sights;  (ii)  at  different distances; and   most importantly,  (iii)  for which  no observational   BSS-HB star-by-star correspondence can be established. Thus, allowing for all these uncertainties in the field BSS frequency, it is safe to conclude that the  observed open clusters-dwarf galaxies BSS frequency sets  a realistic, and observational  upper limit to the primordial BSS frequency in stellar systems."
    },
    "0711/0711.4010_arXiv.txt": {
        "abstract": "We present long-term optical and {\\it RXTE} data of two X-ray binary pulsars in the Small Magellanic Cloud, SXP46.6 and SXP6.85.  The optical light curves of both sources show substantial ($\\sim0.5-0.8$ mag) changes over the time span of the observations.  While the optical data for SXP6.85 do not reveal any periodic behaviour, by detrending the optical measurements for SXP46.6 we find an orbital period of $\\sim 137$ days, consistent with results from the X-ray data.  The detection of Type I X-ray outbursts from SXP46.6, combined with the fact that we also see optical outbursts at these times, implies that SXP46.6 is a high orbital eccentricity system.  Using contemporaneous optical spectra of SXP46.6 we find that the equivalent width of the H$\\alpha$ emission line changes over time indicating that the size of the circumstellar disc varies.  By studying the history of the colour variations for SXP6.85 we find that the source gets redder as it brightens which can also be attributed to changes in the circumstellar disc.  We do not find any correlation between the X-ray and optical data for SXP6.85.  The results for SXP6.85 suggest that it is a low eccentricity binary and that the optical modulations are due to the Be phenomenon. ",
        "introduction": "Observations of the Small Magellanic Cloud (SMC) at multi-wavelengths have uncovered numerous high-mass X-ray binaries, the majority of which are Be/X-ray transients \\citep[see e.g.][]{hab04,coe05,mcg07}.  These systems comprise of an OB star and a neutron star.  As the compact object orbits the companion star in a wide and eccentric orbit it passes close to the circumstellar disc of the Be star and X-ray outbursts are seen \\citep{oka01}. In some cases optical outbursts coincident with the X-ray outbursts are also detected \\citep[see e.g.][]{alc01,coe04}.  \\begin{figure} \\includegraphics[width=84mm]{fig1.eps} \\caption{The OGLE III light curve of SXP46.6.  The times when optical spectra were taken are marked with solid lines.} \\label{fig:sxp46_lc} \\end{figure}  \\xte\\ monitoring is providing a wealth of information on the X-ray pulsars in the SMC \\citep[see e.g.][]{lay05,gal07}.  However, to investigate the sources in detail, in particular at multi-wavelengths, more precise locations must be determined, via e.g. \\chan\\ and \\xmm.  In this paper we present observations and analysis of optical and X-ray data for two Be/X-ray binaries in the SMC, SXP46.6 and SXP6.85, whose positions have been identified recently using \\chan\\ and \\xmm.  \\begin{figure*} \\includegraphics[width=100mm]{fig2.eps} \\caption{The detrended OGLE III light curve (top) and \\xte\\ light curve (bottom) for SXP46.6.  Only X-ray data with identical collimator responses (see text) and covering the time of overlap with the optical data have been plotted.  The times corresponding to the orbital period determined from the \\xte\\ data, $P=137.74$ d, are marked (dashed lines).} \\label{fig:sxp46_ox} \\end{figure*} ",
        "conclusions": "\\label{sect:disc}  \\subsection{SXP46.6}  The optical light curve of SXP46.6 displays long-term variations that are most likely due to the formation and depletion of the Be star's circumstellar disc.  It is reasonable to assume, if the equatorial plane of the Be star coincides with the orbital plane of the compact object, that as the neutron star orbits close to this disc at periastron X-ray outbursts should be observed.  Our results show that not only do we detect X-ray outbursts, but optical outbursts are also seen.  SXP46.6 is therefore one of only a handful of systems that displays optical bursts at the binary period \\citep[see e.g.][]{alc01,sch03,cow03,coe04,edg05}.  We find that the amplitude of the optical outbursts ($\\sim 0.04$ mag) is similar to those of other sources \\citep[see e.g.][]{sch03,cow03}.  In Be/X-ray transients in general, the depletion of the circumstellar disc leads to a decrease in the optical brightness of the system.  For SXP46.6, even when the optical light curve goes into a sharp decline there must be a residual disc present as outbursts, both at optical and X-ray wavelengths, are seen.  In contrast, as the optical brightness recovers, which would normally be attributed to the disc reforming, an X-ray outburst is evident but there does not appear to be a corresponding burst in the optical.  The reason for this is unclear, but could possibly be related to the amount of material in the circumstellar disc available for accretion and the quantity that is transferred to the neutron star.  If less matter is accreted the optical outbursts could have a lower amplitude and may not be detectable above the intrinsic variability in the optical light curve.  Another possibility for the lack of an optical outburst at this time is that there remains a wind that is sufficient to generate X-ray outbursts only.  Using the H$\\alpha$ equivalent width as an indicator for the circumstellar disc size \\citep{dac86}, we find that for the first three years of OGLE III monitoring the disc size is increasing, anti-correlated with the optical brightness.  Unfortunately we do not have spectra during the optical minimum.  After the optical minimum the equivalent width of the H$\\alpha$ emission line seems to reflect the recovery of the circumstellar disc, with the width of the line increasing as the optical brightness increases.  The anti-correlated behaviour initially may be due to the disc covering part of the Be star leading to a reduction in brightness.  This has been seen in another source, A 0538-66 \\citep{alc01}.  It is unfortunate that we do not have multi-band data that would allow us to investigate the colours of the source at these times.  The shape of the optical and X-ray outbursts for SXP46.6 are both asymmetric, with a faster rise and slower decline, reminiscent of other Be/X-ray binaries \\citep[see e.g.][]{alc01,coe04}.  We propose that the optical bursts could be due to the circumstellar disc being perturbed at each passage causing the disc to grow slightly, leading to an enhancement in the optical brightness.  We can calculate a rough estimate for this increase in brightness.  If we assume that the 0.5 magnitude global change in the light curve represents the difference between no disc and a full disc, then this corresponds to an increase in light of 58\\%.  The optical outbursts we detect have an amplitude of $\\sim 0.04$ mag, which equals a total increase from no disc of 64\\%.  Assuming that the increase in brightness is equivalent to the disc size, then during outburst the disc grows by $\\sim 10$\\%.  We note that the increase in optical brightness at outburst may not be due to geometric effects alone.  We find that the luminosity of the optical outbursts is $\\sim3\\times10^{34}$ erg s$^{-1}$ and the luminosity of the X-ray outbursts is $\\sim2\\times10^{36}$ erg s$^{-1}$.  This implies that reprocessing could be contributing to the change in the optical light.  \\citet{stel86} distinguished between the various kinds of X-ray behaviour observed from Be/X-ray binaries, placing them into three groups: (i) persistent, low-level (quiescent) emission, (ii) relatively bright, short-lived outbursts recurring on the orbital period at the time of periastron passage, known as Type I, and (iii) large, long-lived (weeks to months) outbursts that occur irregularly, known as Type II.  We are most likely detecting quiescent emission from SXP46.6 in the \\chan\\ observations. \\citet{oka01} have shown that the class of X-ray outbursts detected from a Be/X-ray transient, Type I and/or Type II, is dependent on the orbital eccentricity of the system, and the viscosity of the disc.  The disc in binaries with low eccentricity is truncated at the 3:1 resonance radius.  This means that too little material is captured (if any) by the compact object as it orbits the Be star as the disc is too far away.  Therefore, regular Type I outbursts cannot be produced by these systems, but they may show infrequent Type II outbursts.  In contrast, periodic Type I outbursts are generated in high eccentricity binaries as the disc is truncated at a larger resonance radius, which permits material to be accreted by the neutron star every orbit.  The detection of Type I X-ray outbursts from SXP46.6, and the corresponding optical outbursts, and its similarity to other sources that have shown optical outbursts on the orbital period, implies that SXP46.6 has a high orbital eccentricity.  \\subsection{SXP6.85}  The long-term optical light curve of SXP6.85 is similar to that of SXP46.6 in the fact that the large changes observed are not periodic, leading to the conclusion that they are related to the Be phenomenon.  The variations in the source's colour support this.  Even after detrending, we find little evidence for any coherent modulations in the optical light curve of SXP6.85.  Analysis of SXP6.85's X-ray light curve does not firmly demonstrate the presence of periodic behaviour either \\citep{gal07}.  More data are needed to confirm the tentative orbital period proposed.  There is a lack of correlation between the optical and X-ray data.  By comparing the two datasets we find that one X-ray outburst occurs just prior to optical maximum, while another two seem to occur as the optical source is fading.  This implies that the neutron star is not accreting directly from the stellar wind, and could indicate that the X-ray emission is triggered by discrete episodes of disc material being lost radially \\citep[see e.g.][]{cla99}.  The characteristics of SXP6.85 are very similar to those of EXO 0531-66 and H 0544-665 \\citep{mcg02}.  All three sources are highly variable in the optical over long timescales but do not show variations that are orbital in origin, they become redder as they brighten in the optical and no strong correlation is found between the X-ray and optical measurements.  Due to the similarities of these sources, and with comparison to SXP46.6, we therefore propose that SXP6.85 is a low eccentricity system."
    },
    "0711/0711.3540_arXiv.txt": {
        "abstract": "I review the long-term survival chances of young massive star clusters (YMCs), hallmarks of intense starburst episodes often associated with violent galaxy interactions. In particular, I address the key question as to whether at least some of these YMCs can be considered proto-globular clusters (GCs). In the absence of significant external perturbations, the key factor determining a cluster's long-term survival chances is the shape of its stellar initial mass function. I conclude that there is an increasing body of evidence that GC formation appears to be continuing until today; their long-term evolution crucially depends on their environmental conditions, however. ",
        "introduction": "\\label{intro.sec}  Young, massive star clusters (YMCs) are the hallmarks of violent star-forming episodes triggered by galaxy collisions and close encounters. Their contribution to the total luminosity induced by such extreme conditions completely dominates the overall energy output due to the interaction-induced star formation (e.g., de Grijs \\& Parmentier 2007; and references therein).  The question remains, however, whether or not at least a fraction of the compact YMCs seen in abundance in extragalactic starbursts, are potentially the progenitors of ($\\ga 10$ Gyr) old globular cluster (GC)-type objects -- although of higher metallicity than the present-day GCs. If we could settle this issue convincingly, one way or the other, such a result would have far-reaching implications for a wide range of astrophysical questions, including our understanding of the process of galaxy formation and assembly, and the process and conditions required for star (cluster) formation. Because of the lack of a statistically significant sample of YMCs in the Local Group, however, we need to resort to either statistical arguments or to the painstaking approach of case-by-case studies of individual objects in more distant galaxies.  \\subsection{The stellar initial mass function}  The evolution to old age of young clusters depends crucially on their stellar initial mass function (IMF). If the IMF slope is too shallow, i.e., if the clusters are significantly deficient in low-mass stars compared to, e.g., the solar neighbourhood, they will likely disperse within about a Gyr of their formation (e.g., Chernoff \\& Shapiro 1987; Chernoff \\& Weinberg 1990; Goodwin 1997b; Smith \\& Gallagher 2001; Mengel et al. 2002). As a case in point, Goodwin (1997b) simulated the evolution of $\\sim 10^4 - 10^5$ M$_\\odot$ YMCs similar to those observed in the LMC, with IMF slopes $\\alpha = 2.35$ (Salpeter 1955; where the IMF is characterised as $\\phi(m_\\ast) \\propto m_\\ast^{-\\alpha}$, as a function of stellar mass, $m_\\ast$) and $\\alpha = 1.50$, i.e., roughly covering the range of (present-day) mass function slopes observed in LMC clusters at the time he performed his {\\it N}-body simulations (see also de Grijs et al. 2002a,b). The stellar mass range covered ranged from 0.15 to 15 M$_\\odot$; his {\\it N}-body runs spanned at most a few 100 Myr. Following Chernoff \\& Weinberg (1990), and based on a detailed comparison between the initial conditions for the LMC YMCs derived in Goodwin (1997b) and the survival chances of massive star clusters in a Milky Way-type gravitational potential (Goodwin 1997a), Goodwin (1997b; see also Takahashi \\& Portegies Zwart 2000, their fig. 8) concluded that -- for Galactocentric distances $\\ga 12$ kpc -- some of his simulated LMC YMCs should be capable of surviving for a Hubble time if $\\alpha \\ge 2$ (or even $\\ga 3$; Mengel et al. 2002), but not for shallower IMF slopes for any reasonable initial conditions (cf. Chernoff \\& Shapiro 1987; Chernoff \\& Weinberg 1990). More specifically, Chernoff \\& Weinberg (1990) and Takahashi \\& Portegies Zwart (2000), based on numerical cluster simulations employing the Fokker-Planck approximation, suggest that the most likely survivors to old age are, additionally, characterised by King model concentrations, $c \\ga 1.0-1.5$. Mengel et al. (2002; their fig. 9) use these considerations to argue that their sample of YMCs observed in the Antennae interacting system might survive for at least a few Gyr, but see de Grijs et al. (2005), and Bastian \\& Goodwin (2006) and Goodwin \\& Bastian (2006), for counterarguments related to environmental effects and to variations in the clusters' (effective) star-formation efficiencies (SFEs), respectively.  In addition, YMCs are subject to a variety of additional internal and external drivers of cluster disruption. These include internal two-body relaxation effects, the nature of the stellar velocity distribution function, the effects of stellar mass segregation, disk and bulge shocking, and tidal truncation (e.g., Chernoff \\& Shapiro 1987; Gnedin \\& Ostriker 1997). All of these act in tandem to accelerate cluster expansion, thus leading to cluster dissolution -- since expansion will lead to greater vulnerability to tidally-induced mass loss.  \\subsection{Survival Diagnostics: the Mass-to-Light Ratio versus Age Diagram}  \\begin{figure}[h!] \\plotone{mlratio_pavel.ps} \\caption{Updated version of the YMC $M/L$ ratio versus age diagnostic diagram. The numbered data points were taken from Bastian et al. (2006; and references therein); overplotted are the SSP predictions for a Salpeter (1955) and a Kroupa (2001) stellar IMF. We have included four new YMCs, NGC 1140-1 (Moll et al. 2007) the Galactic Centre Arches cluster, R136 in 30 Doradus (Goodwin \\& Bastian 2006) and Westerlund 1 (denoted `Wld 1'; see de Grijs \\& Parmentier 2007). The evolution expected for SSPs governed by IMFs as defined by both Salpeter (1955) and Kroupa (2001) is shown as the dolid and short-dashed lines, respectively. The long-dashed lines represent the evolution expected for SSPs with a Kroupa (2001)-type IMF, but a range of effective star-formation efficiencies (Goodwin \\& Bastian 2006) } \\label{diagnostic.fig} \\end{figure}  With the ever increasing number of large-aperture ground-based telescopes equipped with state-of-the-art high-resolution spectrographs and the wealth of observational data provided by the {\\sl Hubble Space Telescope}, we may now finally be getting close to resolving the issue of potential YMC longevity conclusively. To do so, one needs to obtain (i) high-resolution spectroscopy, in order to obtain dynamical mass estimates, and (ii) high-resolution imaging to measure their sizes (and luminosities). As a simple first approach, one could then construct diagnostic diagrams of YMC mass-to-light ($M/L$) ratio versus age, and compare the YMC loci in this diagram with simple stellar population (SSP) models using a variety of IMF descriptions (cf. Smith \\& Gallagher 2001; Mengel et al. 2002; Bastian et al. 2006; Goodwin \\& Bastian 2006). In Fig. \\ref{diagnostic.fig} I present an updated version of the $M/L$ ratio versus age diagram, including all of the YMCs for which the required observables are presently available (see also de Grijs \\& Parmentier 2007). However, such an approach, while instructive, has serious shortcomings. The viability of this approach depends, in essence, on the validity of the virial equation to convert line-of-sight velocity dispersions, $\\sigma_{\\rm los}$, to dynamical mass estimates, $M_{\\rm dyn}$, via (Spitzer 1987):  \\begin{equation} \\label{virial.eq} M_{\\rm dyn} = \\frac{\\eta \\sigma^2_{\\rm los} r_{\\rm h}}{G} \\quad , \\end{equation} where $r_{\\rm h} = 1.3 \\; R_{\\rm eff}$ are the half-mass and effective (or half-light) radii of the cluster, respectively, and $\\eta = 3 a$; $a \\approx 2.5$ is the factor required to convert the half-mass to the gravitational radius, $r_{\\rm g}$. More specifically, following Fleck et al. (2006), we write \\begin{equation} \\label{gravrad.eq} r_{\\rm g} = \\frac{5}{2} \\times \\frac{4}{3} r_{\\rm h} \\quad , \\end{equation} where the factor 5/2 provides an approximate conversion for a large range of clusters characterised by King (1966) mass profiles; the second numerical factor in Eq. (\\ref{gravrad.eq}) results from projection on the sky, assuming that light traces mass throughout the cluster. The use of both Eq. (\\ref{virial.eq}) and the $M/L$ ratio versus age diagram rely on a number of assumptions and degeneracies, however, which I will discuss in some detail below.  \\subsubsection{IMF degeneracies}  In the simplest approach, in which one compares the YMC loci in the $M/L$ ratio versus age diagram with SSP models, the data can be described by {\\it both} variations in the IMF slope {\\it and} variations in a possible low-mass cut-off (e.g., Sternberg 1998; Smith \\& Gallagher 2001; Mengel et al. 2002); the models are fundamentally degenerate for these parameters. For instance, Sternberg (1998) derived for the YMC NGC 1705-I that it must either have a flat mass function ($\\alpha < 2$) or a low-mass truncation between 1 and 3 M$\\odot$ (see also Smith \\& Gallagher 2001); in both cases, it is unlikely that this cluster may be capable of surviving for a Hubble time.  However, the conclusion that the IMFs of such starburst clusters may be unusual must be regarded with caution. As Smith \\& Gallagher (2001) point out, previous claims for highly abnormal (initial) mass functions have often proven incorrect. If anything, the shape of the mass function may vary on the size scales of the individual clusters, but once one considers their birth environments on larger scales the present-day mass function appears to be remarkably robust (e.g., Scalo 1998; Kroupa 2001), with the possible exception of the resolved starburst clusters in the Milky Way (e.g., Stolte et al. 2005, 2006), NGC 3603 and -- in particular -- the Galactic Centre Arches cluster.  Despite this controversy (particularly for some of the youngest clusters), it appears that most of the YMCs for which high-resolution spectroscopy is available are characterised by `standard' Salpeter (1955) or Kroupa (2001) IMFs (e.g., Larsen et al. 2001; McCrady, Gilbert \\& Graham 2003; Maraston et al. 2004; Larsen, Brodie \\& Hunter 2004; Larsen \\& Richtler 2004; Bastian et al. 2006; see also de Grijs et al. [2005] for a comparison of dynamical and photometric masses, the latter based on `standard' IMF representations).  \\subsubsection{Mass segregation}  While the assumption that these objects are {\\it approximately} in virial equilibrium is probably justified at ages greater than a few $\\times 10^7$ yr and for realistic SFEs $\\ga 30$ per cent (at least for the stars dominating the light; see, e.g., Goodwin \\& Bastian 2006), the {\\it central} velocity dispersion (as derived from luminosity-weighted high-resolution spectroscopy) does not necessarily represent a YMC's total mass. It is now well-established that almost every YMC exhibits significant mass segregation from very young ages onwards, so that the effects of mass segregation must be taken into account when converting central velocity dispersions into dynamical mass estimates (see also Fleck et al. 2006; Moll et al. 2007).  By ignoring the effects of mass segregation, as is in essence done if one simply applies Eq. (\\ref{virial.eq}), the underlying assumption is then that of an isotropic stellar velocity distribution, i.e., $\\sigma^2_{\\rm total} = 3 \\sigma^2_{\\rm los}$, where $\\sigma^2_{\\rm total}$ is the cluster's mean three-dimensional velocity dispersion. In the presence of (significant) mass segregation in a cluster, the central velocity dispersion will be dominated by the higher-mass stars populating the cluster core. If we focus on dynamical evolution as the dominant cause of mass segregation in clusters (as opposed to the possible preferential formation of the higher-mass stars close to the cluster core, also known as `primordial' mass segregation; e.g., Bonnell \\& Davies 1998; de Grijs et al. 2002a), it follows that for the high-mass stars to migrate to the cluster core, i.e., to the bottom of the gravitational potential well, they must have exchanged some of their kinetic energy with their lower-mass counterparts on more extended orbits. As a consequence, the velocity dispersion dominating the observed high-resolution spectra will be {\\it lower} than expected for a non-mass-segregated cluster of the same mass. In addition, measurements of $r_{\\rm h}$ will also be biased to smaller values, and not to the values associated with the cluster as a whole. Mass segregation will thus lead to an {\\it under}\\,estimate of the true cluster mass.  I also note that the assumption of virial equilibrium only holds to a limited extent, even in old GCs, because cluster-wide relaxation time-scales of massive GC-type objects are of order $10^9$ yr or longer (Djorgovski 1993). In fact, full global, or even local, energy equipartition among stars covering a range of masses is never reached in a realistic star cluster, not even among the most massive species (e.g., Inagaki \\& Saslaw 1985; Hunter et al. 1995). As the dynamical evolution of a cluster progresses, low-mass stars will, on average, attain larger orbits than the cluster's higher-mass stars, and the low-mass stars will thus spend most of their time in the cluster's outer regions, at the extremes of their orbits. For this reason alone, we would not expect to achieve global energy equipartition in a cluster.  The time-scale for the onset of significant dynamical mass segregation is comparable to the cluster's dynamical relaxation time (Spitzer \\& Shull 1975; Inagaki \\& Saslaw 1985; Bonnell \\& Davies 1998; Elson et al.  1998). A cluster's characteristic time-scale may be taken as its half-mass (or median) relaxation time, i.e., the relaxation time at the mean density for the inner half of the cluster mass for cluster stars with stellar velocity dispersions characteristic for the cluster as a whole (Spitzer \\& Hart 1971; Lightman \\& Shapiro 1978; Meylan 1987; Malumuth \\& Heap 1994).  Although the half-mass relaxation time characterises the dynamical evolution of a cluster as a whole, significant differences are expected locally within the cluster. The relaxation time-scale will be shorter for higher-mass stars than for their lower-mass companions; numerical simulations of realistic clusters confirm this picture (e.g., Aarseth \\& Heggie 1998; Kim et al. 2000; Portegies Zwart et al. 2002). From this argument it follows that dynamical mass segregation will also be most rapid where the local relaxation time is shortest, i.e., near the cluster centre (cf. Fischer et al. 1998; Hillenbrand \\& Hartmann 1998). Thus, significant mass segregation among the most massive stars in the cluster core occurs on the local, central relaxation time-scale (comparable to just a few crossing times; cf. Bonnell \\& Davies 1998).  The combination of these effects will lead to an increase of the dimensionless parameter $\\eta$ in Eq. (\\ref{virial.eq}) with time, if the characteristic two-body relaxation time of a given (massive) stellar species is short (Boily et al. 2005; Fleck et al. 2006), and thus to an {\\it under}\\,estimate of the true cluster mass. However, we note that Goodwin \\& Bastian (2006) point out that a large fraction of the youngest clusters in the $M/L$ ratio versus age diagram appear to have dynamical masses well in excess of their photometric masses, and that, therefore, the result of Boily et al. (2005) and Fleck et al. (2006) does not seem applicable to these YMCs.  \\subsubsection{Stellar masses}  Estimating dynamical masses, $M_{\\rm cl}$, via Eq. (\\ref{virial.eq}) assumes, in essence, that all stars in the cluster are of equal mass. This is clearly a gross oversimplification, which has serious consequences for the resulting mass estimates. The straightforward application of the virial theorem tends to {\\it under}\\,estimate a system's dynamical mass by a factor of $\\sim 2$ compared to more realistic multi-mass models (e.g., Mandushev, Spassova \\& Staneva 1991; based on an analysis of the observational uncertainties). Specifically, Mandushev et al. (1991) find that the mass-luminosity relation for GCs with mass determinations based on multi-component King-Michie models (obtained from the literature) lies parallel to that for single-mass King models, but offset by $\\Delta \\log M_{\\rm cl} ({\\rm M}_\\odot) \\simeq 0.3$ towards higher masses. Farouki \\& Salpeter (1982) already pointed out that cluster relaxation and its tendency towards stellar energy equipartition is accelerated as the stellar mass spectrum is widened; mass segregation will then take place on shorter time-scales than for single-component (equal-mass) clusters, and thus this will once again lead to an {\\it under}\\,estimate of the true cluster mass (see also Goodwin 1997a; Boily et al. 2005; Fleck et al. 2006; Kouwenhoven \\& de Grijs 2008, for multi-mass {\\it N}-body approaches).  We also point out that if the cluster contains a significant fraction of primordial binary and multiple systems, these will act to effectively broaden the mass range and thus also speed up the dynamical evolution of the cluster (e.g., Fleck et al. 2006; Kouwenhoven \\& de Grijs 2008). ",
        "conclusions": ""
    },
    "0711/0711.4374_arXiv.txt": {
        "abstract": "We show that the rate for upward showers from an isotropic cosmic neutrino flux at neutrino telescopes like IceCube is independent of the neutrino-nucleon cross section. For bins that span a relatively narrow range in energy, neither scaling the cross section, nor changing its power-law energy behavior affects the upward shower rate, which depends only on the flux. The neutrino flux can be completely known since its spectral shape can be determined by comparing the rates in neighboring bins. We also show that the downward shower rate varies linearly with cross section with a proportionality constant determined by the energy-dependence of the cross section, independent of the power-law behavior of the flux.  The normalization and energy dependence of the cross section can be found by comparing the downward rates in neighboring bins. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.3281_arXiv.txt": {
        "abstract": "A system of equations governing the structure of a steady, relativistic radiation dominated shock is derived, starting from the general form of the transfer equation obeyed by the photon distribution function.  Closure is obtained by truncating the system of moment equations at some order.  The anisotropy of the photon distribution function inside the shock is shown to increase with increasing shock velocity, approaching nearly perfect beaming at upstream Lorentz factors $\\Gamma_{-}>>1$.  Solutions of the shock equations are presented for some range of upstream conditions.  These solutions are shown to converge as the truncation order is increased. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.1883_arXiv.txt": {
        "abstract": "PSR~J0514$-$4002A is a 5-ms pulsar is located in the globular cluster NGC~1851; it belongs to a highly eccentric ($e\\,=\\,0.888$) binary system. It is one of the earliest known examples of a numerous and fast-growing class of eccentric binary MSPs recently discovered in globular clusters. Using the GBT, we have obtained a phase-coherent timing solution for the pulsar, which includes a measurement of the rate of advance of periastron $\\dot{\\omega} = 0.01289(4)^\\circ \\rm yr^{-1}$, which if due completely to general relativity, implies a total system mass of $2.453(14)\\,\\rm M_{\\odot}$; we also derive $m_p\\,<\\,1.5\\,M_{\\odot}$ and $m_c\\,>\\,0.96\\,M_{\\odot}$. The companion is likely to be a massive white dwarf star. ",
        "introduction": "PSR~J0514$-$4002A was discovered in a 327-MHz search for steep-spectrum pulsars in globular clusters (GCs) carried out with the Giant Metrewave Radio Telescope (GMRT) at Khodad, near Pune, India \\cite{fgri04}. The pulsar is located in the globular cluster NGC~1851 and we will hereafter refer to it as NGC~1851A. It has a spin frequency of 200 Hz and is a member of a binary system, with a massive $\\sim\\,1\\,\\rm\\,M_{\\odot}$ companion and a highly eccentric orbit ($e = 0.888$).  With the notable exception of PSR~J1903+0327 (see Champion et al., these proceedings), all binary MSPs ($P\\,<\\,20$ms) in the Galaxy are low-eccentricity systems where the companions are white dwarf (WDs) stars. Some Galactic pulsars have eccentric orbits and massive companions, but their spin periods are tens, not hundreds, of Hz (see Fig. \\ref{fig:eccentric}). For these systems, the short evolution times of their massive companions is not long enough for accretion to spin up the pulsar to higher spin frequencies \\cite{lor05}.  The fast rotation of NGC~1851A indicates that the pulsar itself probably formed in the same manner as a normal MSP, ending up with a light WD companion in a nearly circular orbit. A subsequent stellar interaction then disrupted the original binary system, replacing the lowest mass component (the original WD) with the present companion star in a highly eccentric orbit. Such exchange interactions are only likely to happen in environments with very high stellar densities, like the central regions of dense globular clusters. The same process could also form even more exotic systems, such as MSP/MSP or MSP/Black Hole binaries. An alternative scenario for the formation of such binaries is the nearly head-on collision of a neutron star (possibly even an MSP) with a giant star \\cite[e.g.][]{rs91}; this can form an $e \\sim 0.9$ binary system consisting of the MSP and the stripped core of the giant star. We discuss this possibility at length in \\cite{frg07}, since it might provide an explanation for the unusual flux density variations observed in NGC~1851A. ",
        "conclusions": ""
    },
    "0711/0711.2622_arXiv.txt": {
        "abstract": "We show that the cosmic microwave background (CMB) data of WMAP can give subelectronvolt limit on the neutrino mass: $m_\\nu < 0.63$\\,eV (95\\% CL). We also investigate its degeneracy with other cosmological parameters. In particular, we show the Hubble constant derived from the WMAP data decreases considerably when the neutrino mass is a few times 0.1\\,eV. ",
        "introduction": "\\label{sec:introduction}  It is known that cosmological considerations give more stringent constraints on the neutrino mass than the present tritium $\\beta$-decay experiments. For example, \\cite{Spergel:2003cb} obtained $m_\\nu < 0.21$\\,eV and \\cite{Tegmark:2003ud} obtained $m_\\nu < 0.58$\\,eV, using the WMAP 1st year data (WMAP1) combined with the galaxy power spectrum (the former used the data from the 2dFGRS and the latter from the SDSS main sample. The difference can be ascribed to the use of the bias information by \\cite{Spergel:2003cb}). Meanwhile, we found $m_\\nu < 0.66$\\,eV from WMAP1 alone as reported in \\cite{Ichikawa:2004zi}. We note that this is the first to point out CMB data (precisely speaking, WMAP1) alone can give a sub-eV upper bound on the neutrino mass, which is comparable to the limits obtained from the CMB and galaxy clustering data combined.  Before \\cite{Ichikawa:2004zi} (and some time after that too), there seems to be a lack of consensus about whether the CMB experiment with the WMAP-level precision can derive a sub-eV neutrino mass limit and, in fact, the WMAP1 alone limit reported in \\cite{Tegmark:2003ud}, $m_\\nu < 3.8$\\,eV, which allows 100\\% HDM was seemingly accepted (incidentally, the WMAP group did not report the WMAP1 alone limit on the neutrino mass). We, on the contrary, have derived the upper limit 0.66\\,eV as quoted above from the same data by the $\\chi^2$ minimization method which is independent from the MCMC method adopted by \\cite{Tegmark:2003ud}. Our conclusion is later confirmed by \\cite{MacTavish:2005yk,Hannestad:2006zg,Lesgourgues:2006nd} (\\cite{MacTavish:2005yk} does not report the WMAP1 alone limit in a number but judging from their likelihood figure, it looks less than 1\\,eV. \\cite{Hannestad:2006zg} and \\cite{Lesgourgues:2006nd} gives $m_\\nu < 0.70$\\,eV and $m_\\nu < 0.63$\\,eV respectively).  Below, we first discuss the WMAP alone limit on the neutrino mass comparing results from WMAP1 and the WMAP 3rd year data (WMAP3). Then, we will investigate its degeneracy with other cosmological parameters, especially the Hubble constant.  We here summarize our notations. We derive neutrino mass constraint in the flat $\\Lambda$CDM model with the power-law adiabatic perturbations. Namely, cosmological parameters we consider are: baryon density $\\omega_b$, matter density $\\omega_m$, hubble parameter $h$, reionization optical depth $\\tau$, spectral index of primordial spectrum $n_s$, its amplitude $A$ and massive neutrino density $\\omega_\\nu$. Here, $\\omega \\equiv \\Omega h^2$ where $\\Omega$ is the energy density normalized to the critical density and $\\omega_m \\equiv \\omega_b + \\omega_c$ where $\\omega_c$ is CDM density (Caution that many literatures define $\\omega_m \\equiv \\omega_b + \\omega_c +\\omega_\\nu$ to include the massive neutrino in the matter density). $\\omega_\\nu$ is related to neutrino masses by $\\omega_\\nu = \\sum m_\\nu/(94.1$\\,eV) and we assume three mass degenerate generations of neutrinos so that $\\omega_\\nu =m_\\nu/(31.4$\\,eV). ",
        "conclusions": ""
    },
    "0711/0711.1692_arXiv.txt": {
        "abstract": "{}% {The realism of current models of the penumbra is assessed by comparing their predictions with the plasma properties of penumbral filaments as retrieved from spectropolarimetric observations.} {The spectropolarimeter onboard Hinode allows us to distinguish for the first time the fine structure of the penumbra. Therefore, we can use one-component inversions to obtain the stratifications of plasma parameters in each pixel. The correlations between the plasma parameters and the continuum intensity are studied.} {We find that, in the outer penumbra, the stronger flows and higher values of magnetic field inclination tend to be located in dark filaments. This finding does not seem to be compatible with the scenario of a field-free gappy penumbra.} {} ",
        "introduction": "There are two models of the penumbral fine structure under discussion nowadays. The first is the uncombed model proposed by \\citet{Solanki:1993} and further developed by \\citet{vmp:2000} to explain the broad-band circular polarisation observed in sunspots.  It envisages the penumbra as a collection of horizontal flux tubes embedded in a more vertical and stronger background field. \\citet{Schlichenmaier:1998} made simulations of the temporal evolution of such flux tubes. Initially placed at the sunspot magnetopause, they rise by magnetic buoyancy and quickly become horizontal in the photosphere. At the same time, a strong Evershed flow develops due to the gas pressure gradient that builds up during the rise. The tubes are hotter than the surrounding plasma and appear as bright penumbral filaments which gradually cool down due to radiative losses, i.e. they become dark further from the bright penumbral grains where the tubes cross the $\\tau = 1$ layer. Hereafter, this model will be called the rising flux tube model (RFT).  The second model was suggested by \\citet{Spruit:2006}. These authors argued that the RFT model cannot produce sufficient heating of the penumbra. As an alternative, they proposed a field-free gap model (FFG) which consists of field-free material protruding into the penumbral magnetic field from below. With such a configuration, the bright filaments would be heated all along their lengths and not only at the point of emergence as in the case of the RFT model. While the FFG model does not offer any explanation for the Evershed flow, there is only one place where the flow can reside, namely on top of the field-free gaps, since this is the sole region in the model where horizontal fields exist.  Both models have problems from a theoretical point of view which need to be clarified \\citep[see][and references therein]{Bellot:2007}. However, as illustrated in Fig.~\\ref{models}, RFTs and FFGs should produce similar stratifications of the plasma parameters if they are placed near $\\tau = 1$ (marked by the gray line). The two models predict weaker and more inclined fields low in the atmosphere, along with increased line-of-sight (LOS) velocities. Using Hinode spectropolarimetric data, \\citet{Jurcak:2007} found this configuration in the bright filaments of the inner penumbra.  \\begin{figure}[!b] \\centering \\includegraphics[width=\\linewidth]{jurcak_fig1.eps} \\caption{Sketch of a RFT (left) and a FFG (right) surrounded by the background magnetic field of the penumbra. The gray line represents the possible position of the optical depth unity layer.} \\label{models} \\end{figure}  \\begin{figure}[!t] \\centering \\includegraphics[width=\\linewidth]{jurcak_fig2.eps} \\caption{Maps of continuum intensity at 630.3 nm reconstructed from the SP raster scans. North is up and West to the right. The maps cover the limb side part of the penumbra in evolving AR 10930. The arrows point to the disc centre. The black and white bands indicate the areas selected for analysis.} \\label{imaps} \\end{figure}  In this Letter we concentrate on the mid and outer penumbra, where differences between the models could be more pronounced. According to the FFG model, horizontal fields and increased LOS velocities should be associated with bright areas everywhere in the penumbra. By contrast, the RFT model expects such characteristics in the dark filaments of the outer penumbra. ",
        "conclusions": "The most important result of our analysis is that we do not find positive correlation coefficients between $I_{\\rm c}$ and $\\gamma$ or $v_{\\rm LOS}$. This indicates that bright penumbral filaments show the more vertical fields and weaker flows, in contradiction with the predictions of the FFG model. The horizontal fields and stronger flows tend to be associated with dark filaments, as has been reported previously from lower resolution data. The fact that the correlation coefficients between $I_{\\rm c}$ and $\\gamma$ or between $I_{\\rm c}$ and $v_{\\rm LOS}$ are not $-1$ can be easily explained by the mixing of flux tubes and background magnetic fields. In the mid and outer penumbra, dark areas may represent either the darkened tails of RFTs or background fields.  The latter are characterised by more vertical fields and weaker flows. Such an opposite behaviour decreases the correlation coefficients, which nevertheless remain negative because of the overall dominance of dark RFTs in the outer parts of the penumbra.  This argument does not influence the correlation coefficients between $\\gamma$ and $v_{\\rm LOS}$, which are indeed much higher than those involving $I_{\\rm c}$.  \\citet{Schlichenmaier:2005} and \\citet{Ichimoto:2007} came to a similar conclusion (regarding the velocities) from simpler interpretations of filtergrams and spectropolarimetric data, respectively. They also reported that the horizontal fields and the flows are associated with bright filaments in the innermost penumbra \\citep[see also][]{Hirzberger:2005,Jurcak:2007}, where the dark areas most likely represent background magnetic fields. Thus, the correlation coefficients between continuum intensity and LOS velocity/inclination change sign from the inner to the outer penumbra, due to the fact that RFTs predominantly show up as bright structures in the inner penumbra but darken toward the outer edge of the spot. This sign reversal cannot be explained by the FFG model. The existence of magnetic field lines pointing to the solar interior in the outer penumbra, also confirmed by our observations, represents another problem for the FFG model \\citep{Bellot:2007}.  The RFT simulations of \\citet{Schlichenmaier:1998} seem to be able to explain all the results of our analysis, including the sign reversal of the correlation coefficients and the larger mean value of the LOS velocity in the outer penumbra as compared with the mid penumbra. Overall, this suggests that the RFT model provides a better description of the penumbra than the FFG model."
    },
    "0711/0711.3697_arXiv.txt": {
        "abstract": "{} {In this paper we investigate theoretical pulsation models for the $\\delta$ Scuti star 44~Tau. The star was monitored during several multisite campaigns which confirmed the presence of radial and nonradial oscillations. Moreover, its exceptionally low rotational velocity makes 44~Tau particulary interesting for an asteroseismic study. Due to the measured $\\log g$ value of $3.6 \\pm 0.1$, main sequence and post-main sequence models have to be considered.} {We perform mode identification based on photometric and spectroscopic data. A nonadiabatic pulsation code is used to compute models that fit the identified modes. The influence of different opacity tables and element mixtures on the results is tested. } {The observed frequencies of 44 Tau can be fitted in both the main sequence and the post-main sequence evolutionary stage. Post-main sequence models are preferable as they fulfill almost all observational constraints (fit of observed frequencies, position in the HRD and instability range). These models can be obtained with normal chemical composition which is in agreement with recent spectroscopic measurements. The efficiency of envelope convection (in the framework of the mixing-length theory) is predicted to be very low in 44 Tau. We show that the results are sensitive to the choice between the OPAL and OP opacities. While the pulsation models of 44 Tau computed with OP opacities are considerably too cool and too faint, the use of OPAL opacities results in models within the expected temperature and luminosity range. } {} ",
        "introduction": "\\label{sec:intro}  $\\delta$ Scuti stars are pulsating stars situated at the lower part of the classical instability strip in the HR diagram. Their spectral types range from A0 to F2 and they are in the hydrogen-core or hydrogen-shell burning stage. With masses between 1.5 and 2.5~M$_{\\odot}$ their periods of pulsation are between 30 min and 8~hours. In these stars pulsation is driven by the opacity mechanism ($\\kappa$ mechanism) acting in the HeII ionization zone.  Among the $\\delta$ Scuti stars two subgroups are distinguished: high-amplitude $\\delta$ Scuti stars (HADS) with peak-to-peak amplitudes $\\geq$ 0.3 mag in $V$ with mainly radial pulsation and low-amplitude $\\delta$ Scuti stars which oscillate predominantly non-radially. HADS stars show rather slow rotation in comparison with the average $v \\sin i$ of $96~\\rm km~s^{-1}$ of low amplitude $\\delta$ Scuti stars (Solano \\& Fernley \\cite{solano}). 44~Tau (HD 26322, spectral type: F2 IV) exhibits characteristics of both groups: radial as well as nonradial modes are excited to observable amplitudes and its projected rotational velocity is very small ($v \\sin i = 2 \\pm 1~\\rm km~s^{-1}$).  As discussed by Antoci et al. (\\cite{antoci}) (hereafter ABR07) this indicates either pole-on view and/or intrinsic slow rotation. Zima et al. (\\cite{zima}) recently confirmed that 44~Tau is an intrinsic slow rotator. This makes this star particulary interesting for asteroseismology.  Previous theoretical studies on this star have been carried out by Civelek et al. (\\cite{civelek}) and K{\\i}rb{\\i}y{\\i}k et al. (\\cite{kirbiyik}). By means of fitting adiabatic oscillation frequencies to the observed data they conclude that 44~Tau is a 1.89 M$_{\\odot}$ star at the hydrogen-shell burning stage. Recently, Garrido et al. (\\cite{garrido}) concluded that the pulsations of 44~Tau could be explained by a main sequence model with a mass of 1.94 M$_{\\odot}$.  This paper extends the work of Lenz et al. (\\cite{lenz}) and presents a comprehensive seismic study based on the detailed frequency analysis of ABR07 and on additional photometric data obtained in 2004/05.  The paper is structured into the following parts: Section~\\ref{sec:obs} summarizes the results of available photometric and spectroscopic data. Section~\\ref{sec:models} is devoted to the theoretical modelling of the frequencies of 44~Tau. After the determination of the spherical degrees of the observed frequencies we compute models to fit the detected frequencies and the observed frequency range of excited modes. Finally, the effect of using different opacity tables and element mixtures is studied. In Section~\\ref{sec:conclusions} we present our conclusions. ",
        "conclusions": "\\label{sec:conclusions}  We presented a comprehensive asteroseismic study for the slowly rotating $\\delta$ Scuti star 44~Tau. The spectroscopically determined $\\log g$ value of $3.6 \\pm 0.1$ allows no definitive determination of the evolutionary stage, so the main sequence and post-main sequence case were considered.  Our mode degree identification based on photometric and spectroscopic data determines f$_1$ and f$_5$ as radial fundamental and first overtone modes. This confirms the results of previous studies. Four frequencies, f$_2$, f$_3$, f$_4$ and f$_6$, were identied as $\\ell=1$ modes, and three frequencies, f$_7$, f$_8$ and f$_{10}$, as $\\ell=2$ modes. We used the simplest model of convection (mixing-length theory) and the simplest convection-pulsation treatment (frozen convective flux assumption) which may not be adequate for cold $\\delta$~Scuti stars, as discussed by Montalb\\'{a}n \\& Dupret (\\cite{montalban}). However, our results are in good agreement with those obtained by Garrido et al (\\cite{garrido}) using time-dependent convection.  The identification of the radial fundamental and first overtone frequency poses an important constraint to the models as it reduces the number of possible models significantly. Moreover, it is likely that an avoided crossing of nonradial modes in 44~Tau is observed. This fact puts additional strong constraints on the parameter space for appropriate models. For both considered evolutionary stages it is possible to fit the observed radial period ratio.  If standard chemical composition and no overshooting from the convective core are applied, only post-main sequence models of 44~Tau are located within the photometric error box in the HR diagram. Main sequence models can only be obtained by assuming convective overshooting. Examining such hydrogen-core burning models within a wide parameter range we find that all these models are significantly too cool. This fact is a major problem for main-sequence models. Nevertheless, it is remarkable that all four identified $\\ell=1$ modes can be fitted in both evolutionary stages.  The predicted range of unstable modes is in good agreement with the observed frequencies in both cases. For main sequence models the instability range extends to lower frequencies than observed while for post-main sequence models higher frequencies are found to be unstable. The lowest frequency f$_{10}$ is predicted to be stable in the post-main sequence models, but it is close to the theoretical instability region. Our results on the instability ranges can be used to put a further constraint on the determination of the evolutionary stage if new frequencies are detected that extend the frequency range. Additional frequencies beyond the highest observed frequency may also help to constrain the efficiency of convection more precisely.  Considering the good fit of the observed frequencies and physical parameters such as effective temperature, we find that the post-main sequence stage is preferable to explain the pulsation spectrum of 44~Tau.  In this investigation, we also tested the effect of opacities and element mixtures on the radial modes. The choice between OPAL and OP tables causes a significant difference in the Petersen diagram. In the case of 44~Tau, the OPAL opacities are preferable for the asteroseismic analysis, because with the OP data it is not possible to fit the radial modes of 44~Tau within or close to the observational error box in the HR diagram. The use of the new solar element mixture results in a slightly cooler post-main sequence model. We do not find indications in favor of the old or the new element mixture. However, it is important to note that the influence of different opacity data on the period ratio should also be considered when using the Petersen diagram as a diagnostic tool."
    },
    "0711/0711.4886_arXiv.txt": {
        "abstract": "{% The control of systematic effects when measuring background galaxy shapes is one of the main challenges for cosmic shear analyses. } {% Study the fundamental limitations on shear accuracy due to the measurement of the point spread function (PSF) from the finite number of stars that are available. We translate the accuracy required for cosmological parameter estimation to the minimum number of stars over which the PSF must be calibrated. } {% We characterise the error made in the shear arising from errors on the PSF. We consider different PSF models, from a simple elliptical gaussian to various shapelet parametrisations. First we derive our results analytically in the case of infinitely small pixels (i.e. infinitely high resolution), then image simulations are used to validate these results and investigate the effect of finite pixel size in the case of the elliptical gaussian PSF. } {% Our results are expressed in terms of the minimum number of stars required to calibrate the PSF in order to ensure that systematic errors are smaller than statistical errors when estimating the cosmological parameters. On scales smaller than the area containing this minimum number of stars, there is not enough information to model the PSF. This means that these small scales should not be used to constrain cosmology unless the instrument and the observing strategy are optimised to make this variability extremely small. The minimum number of stars varies with the square of the star Signal-to-Noise Ratio, with the complexity of the PSF and with the pixel size. In the case of an elliptical gaussian PSF and in the absence of dithering, 2 pixels per PSF full width at half maximum (FWHM) implies a 20\\% increase of the minimum number of stars compared to the ideal case of infinitely small pixels; 0.9 pixels per PSF FWHM implies a factor 100 increase. } {% In the case of a good resolution and a typical Signal-to-Noise Ratio distribution of stars, we find that current surveys need the PSF to be calibrated over a few stars, which may explain residual systematics on scales smaller than a few arcmins. Future all-sky cosmic shear surveys require the PSF to be calibrated over a region containing about 50 stars. Due to the simplicity of our models these results should be interpreted as optimistic and therefore provide a measure of a systematic `floor' intrinsic to shape measurements. } ",
        "introduction": "\\label{sec:intro}  Gravitational lensing by large scale structure (or `cosmic shear') has grown rapidly as a research field over the last decade  \\citep{refregier03,hoekstra03,munshi06}. In the coming years, we can expect that the field will continue to make important contributions to cosmology. For instance, \\cite{2006astro.ph..9591A} and \\cite{2006Msngr.125...48P} have singled out cosmic shear as potentially the most powerful probe for constraining dark energy.  A great deal of work is currently under way to develop techniques that will allow the maximum possible potential to be reached and not limited by systematic measurement errors.  This can be done by designing and building instruments with weak lensing as the central science driver. This approach leads to a top-down mission design process which begins with science requirements, which are translated into technical requirements and then into an instrument design.  This method for tackling the problem of high accuracy weak lensing measurements is fundamentally different to the bottom-up approach that is currently used. In the latter approach an existing telescope is used and the science team is assigned the task of extracting the maximum possible information from their surveys (until they hit the systematic limit of the instrument).  The field of gravitational lensing lends itself most naturally to top-down design, since most (but not all) of the systematic errors are not astronomical but are associated with the instrument and the atmosphere. These can therefore be controlled through instrument design and an optimised survey strategy. For this reason, many of the ambitious future imaging surveys that are currently under development have placed weak lensing as primary science driver, including the Dark UNiverse Explorer\\footnote{http://www.dune-mission.net} (DUNE), the SuperNovae Acceleration Probe\\footnote{http://snap.lbl.gov/} (SNAP), the Panoramic Survey Telescope \\& Rapid Response System\\footnote{http://pan-starrs.ifa.hawaii.edu} (Pan-STARRS), the Dark Energy Survey\\footnote{http://www.darkenergysurvey.org} (DES) and the Large Synoptic Survey Telescope\\footnote{http://www.lsst.org} (LSST). The issue is then to establish the instrumental requirements needed to reach the full statistical potential of the survey.  In \\cite{2007MNRAS.381.1018A}, a wide range of survey parameters, such as area and depth, were considered.  Their impact on the statistical potential of a lensing survey was calculated and summarised in a scaling relation that can be used for survey designs to trade-off one property of a survey against another.  In brief, this work finds that, once the median redshift of a survey is sufficiently high ($z \\gtrsim 0.7$), the optimal survey strategy is to make the lensing survey as wide as possible. Similar results have been found by \\cite{2006MNRAS.373..105H}. In the same spirit, \\cite{2007arXiv0710.5171A} looked at the requirements an ultra-wide field survey places on the control of systematics and concluded with a set of scaling relations that show the tolerance on residual systematic errors as a function of survey parameters. In particular, two types of shape measurement systematics were considered: multiplicative and additive systematics which are, respectively, correlated and uncorrelated with the lensing signal. \\cite{2006MNRAS.366..101H} have also studied the impact of multiplicative and additive errors.  In the present study, we link the above systematic requirements to errors associated with measurement of the point spread function (PSF) of the instrument. Indeed, since galaxy images need to be PSF-corrected before shapes can be measured, errors in the estimation of the PSF are propagated into errors on galaxy shapes and can mimic the shear signal. The PSF calibration is driven by 3 factors: \\begin{enumerate} \\item \\textbf{The PSF model:} The model that is chosen to describe the PSF will never perfectly describe the true PSF hence the choice of PSF model introduces an `a priori' systematic;  \\item \\textbf{The interpolation scheme:} The PSF will vary over the field which means that the PSF at a galaxy position needs to be interpolated from the PSFs of the nearby stars;  \\item \\textbf{The finite information available from each star:} A star provides an image of the PSF that is noisy and pixelated. Thus, the PSF information that we are able to extract from each star is finite. As a result, to reach high precision, it is necessary to combine the information from several stars. \\end{enumerate} This paper is devoted to the last factor. We quantify the accuracy of the PSF calibration with analytical predictions in the case of infinitely small pixels and use image simulations to quantify the pixelation effects. Previous work has also studied the impact that PSF errors are likely to induce in cosmic shear measurements. For instance,  \\cite{2007arXiv0710.3355S} looked specifically at the PSF from the current  SNAP design and translated this into systematic errors.  The study that we present here sets out to be more general  and uses simplified PSFs to try and quantify the systematic floor of an instrument. We investigate the extent to which the pixel scale will degrade the amount of information relative to our analytic predictions but we do not explore the optimal methods for combining multiple exposures, such as the one proposed by \\cite{2006JCAP...02..001J}. \\cite{2007astro.ph..3471H} have studied the impact of the pixel scale on the number of useful galaxies (assuming the PSF is known perfectly). While they find a weak dependence, we expect that the driving factor for pixel scale is shape measurement systematics rather than statistical errors.  For a survey with a given statistical potential we use the results of \\cite{2007arXiv0710.5171A} to find the upper limit on the systematic errors in the shears such that the induced bias in cosmological parameters remain subdominant compared to the statistical (marginalised Fisher matrix) errors when estimating cosmological parameters. We then convert the requirements on shear systematics into the minimum number of stars ($N_*$) needed to measure the PSF to the level of accuracy required. On scales smaller than the area containing $N_*$, there is not enough information from the stars to calibrate the PSF. Therefore, the systematic errors in the cosmological parameters can be dominant over the statistical ones. This means that these small scales should not be used to constrain cosmology unless the PSF variability is known to be extremely small. $N_*$ clearly depends on the Signal-to-Noise Ratio distribution of stars, since bright ones contain more information about the PSF than faint noisy ones. It also depends on the stability and the complexity of the PSF, characterised by the number of degrees of freedom that must be estimated or interpolated from the stars. Each degree of freedom increases $N_*$.  Figure \\ref{fig:cartoonns} shows a graphical example to illustrate the calibration requirement we are considering. We see a central galaxy surrounded by several stars. The grey shaded region shows the area over which the PSF must be calibrated and in this example contains 11 stars. Stars do not provide enough information to model the PSF variations on scales smaller than this area.  \\begin{figure} \\resizebox{\\hsize}{!}{ \\includegraphics{9150fig1.ps} } \\caption{ \\label{fig:cartoonns} Illustration of the required number of stars $N_*$ characterised in this paper. When measuring the shape of a galaxy (marked by a red spiral), the PSF needs to be calibrated with at least $N_*$ nearby stars (black asterisks) contained in the shaded region. In this example $N_*=11$. On scales smaller than this there is not enough information coming from the stars to measure the PSF variations. } \\end{figure}  This paper is organised as follows. Section \\ref{sec:systematics} explains the weak lensing and cosmological context and summarises the issue of accurate PSF calibration. Section \\ref{sec:analyticalmodel} shows the analytical predictions of the PSF calibration accuracy, expected in the case of infinitely small pixels. Section \\ref{sec:simulations} describes the simulations we use to validate these predictions and extend them to finite pixels. In section \\ref{sec:discussion} we give the final accuracy of the PSF calibration and derive the number of stars it requires. ",
        "conclusions": "\\label{eq:conclusion}  We have studied the PSF calibration requirements for cosmic shear to measure cosmological parameters. We connect the finite information that we are able to extract from stars to the statistical error of the PSF calibration, then to the error on the galaxy ellipticity estimation. We express our results in terms of the minimum number of stars ($N_*$) that are needed to calibrate the PSF in order to keep the systematic errors below statistical errors for cosmological parameter estimation. On scales smaller than the area containing $N_*$ stars there is not enough information coming from the stars to calibrate the PSF. Therefore, the systematic errors in the cosmological parameters may dominate over the statistical errors. This means that these small scales should not be used to constrain cosmology unless the variability is known to be extremely small. Our results show that for current cosmic shear surveys this scale is about an arcminute, which may explain residual systematics found on these scales in current analyses. In future all sky cosmic shear surveys, the data set will be increased by several orders of magnitude and a tight control of the PSF behaviour will be required to prevent this scale from being increased and to reach smaller scales. This places strong requirements on the hardware and observing strategy. For instance, for ground observations where the atmosphere prevents the PSF from being constant over several stars in a single exposure, \\cite{2004astro.ph.12234J,2006JCAP...02..001J} suggest the shear correlation functions should be calculated by cross-correlating galaxies in different exposures. The PSF calibration error contribution would then be uncorrelated and averaged down.  We also demonstrate how $N_*$ can depend on the complexity of the PSF. We define a `complexity factor' for two different shapelet parametrisations and an elliptical gaussian parametrisation. This complexity factor is a number lower than the number of parameters used to describe the PSF and increases with the PSF complexity. There is a different complexity factor associated with each PSF shape parameter. For accurate galaxy shear measurements, we find that the complexity factor for the PSF ellipticity is more important than that for the size. We summarise the dependence of $N_*$ as a function of complexity, star Signal-to-Noise Ratio, galaxy size and cosmological requirements in a convenient scaling relation (see equation \\ref{eq:reqnstar2}). Note that we work under the setup that the PSF can be perfectly described by the model, thus we do not include in our error budget the systematics due to a poor choice of the PSF model.  We find that pixelation degrades the PSF calibration accuracy. Our calculation of $N_*$ holds for the optimistic case of infinitely small pixels (i.e. infinitely high resolution). A finite pixel size increases $N_*$.  We have performed simulations to predict the increase in the case of an elliptical gaussian PSF. We show that for 2 pixels (or more) per FWHM, the rms scatter agrees with the standard deviation predicted for infinitely small pixels to within 10\\%, which in turn can be translated into a 20\\% increase in the required number of calibration stars.  For large pixels we see that the pixelation effect can be dramatic, for instance having 0.9 pixels per PSF FWHM would lead to a factor of 100 increase in the number of stars needed. This effect may be mitigated to some extent by using dithering, which is beyond the scope of this paper. The main results of this paper are for the high resolution case that provides a strict lower limit on $N_*$.   \\clearpage  \\begin{figure*} \\includegraphics[bb=10 10 400 400,width=3.5cm]{9150fig7.eps} \\hspace*{-0.8cm} \\includegraphics[bb=10 10 400 400,width=3.5cm]{9150fig8.eps} \\hspace*{-0.8cm} \\includegraphics[bb=10 10 400 400,width=3.5cm]{9150fig9.eps} \\hspace*{-0.8cm} \\includegraphics[bb=10 10 400 400,width=3.5cm]{9150fig10.eps} \\hspace*{-0.8cm} \\includegraphics[bb=10 10 400 400,width=3.5cm]{9150fig11.eps} \\hspace*{-0.8cm} \\includegraphics[bb=10 10 400 400,width=3.5cm]{9150fig12.eps} \\vspace*{-0.8cm}\\\\ \\vspace*{-0.8cm}\\\\ \\vspace*{-0.8cm} \\includegraphics[bb=10 10 400 400,width=3.5cm]{9150fig13.eps} \\hspace*{-0.8cm} \\includegraphics[bb=10 10 400 400,width=3.5cm]{9150fig14.eps} \\hspace*{-0.8cm} \\includegraphics[bb=10 10 400 400,width=3.5cm]{9150fig15.eps} \\hspace*{-0.8cm} \\includegraphics[bb=10 10 400 400,width=3.5cm]{9150fig16.eps} \\hspace*{-0.8cm} \\includegraphics[bb=10 10 400 400,width=3.5cm]{9150fig17.eps} \\hspace*{-0.8cm} \\includegraphics[bb=10 10 400 400,width=3.5cm]{9150fig18.eps} \\vspace*{-0.8cm} \\includegraphics[bb=10 10 400 400,width=3.5cm]{9150fig19.eps} \\hspace*{-0.8cm} \\includegraphics[bb=10 10 400 400,width=3.5cm]{9150fig20.eps} \\hspace*{-0.8cm} \\includegraphics[bb=10 10 400 400,width=3.5cm]{9150fig21.eps} \\hspace*{-0.8cm} \\includegraphics[bb=10 10 400 400,width=3.5cm]{9150fig22.eps}  \\begin{center} \\includegraphics[width=15cm]{9150fig23.eps} \\end{center}  \\vspace*{-9.5cm} \\begin{tabular}{p{2.5cm}p{2.5cm}p{2.5cm}p{2.5cm}p{2.5cm}p{2.5cm}} \\hspace*{0.7cm}(a)\\vspace*{2.2cm} & \\hspace*{0.7cm}(b) & \\hspace*{0.7cm}(c) & \\hspace*{0.7cm}(d) & \\hspace*{0.7cm}(e) & \\hspace*{0.7cm}(f)\\\\ \\hspace*{0.7cm}(g)\\vspace*{2.2cm} & \\hspace*{0.7cm}(h) & \\hspace*{0.7cm}(i) & \\hspace*{0.7cm}(j) & \\hspace*{0.7cm}(k) & \\hspace*{0.7cm}(l)\\\\ \\hspace*{0.7cm}(m) & \\hspace*{0.7cm}(n) & \\hspace*{0.7cm}(o) & \\hspace*{0.7cm}(p) & & \\\\ \\end{tabular} \\vspace*{3.3cm} \\caption{ \\label{fig:psf} Non-gaussian PSFs used to study the effect of complexity in section \\ref{sec:spacesteppsf} (10 pixels per PSF FWHM, logarithmic color scale, the normalisation is such as the total flux is 10$^3$). Each of them is described with a different shapelet model, described in table \\ref{tab:psf}. } \\end{figure*}  \\begin{table*} \\begin{center} \\begin{tabular}{|c|cc|ccc|} \\hline label & $n_{{\\rm max}}$ & diamond & $N_{\\rm coef}$ & $N_{\\rm coef}^{m=2(0)}$ & complexity factor\\\\ & & (yes or no) & & & $\\Psi_\\ellipticity$($\\Psi_{\\size^2}$)\\\\ \\hline \\hline gauss & - & - & - & - & $\\sqrt{2}(\\sqrt{2})$\\\\ \\hline a & 4 & no & 15 & 2(3) & 3.27(2.58)\\\\ b & 5 & no & 21 & 2(3) & 3.27(2.58)\\\\ c & 6 & no & 28 & 3(4) & 4.47(3.74)\\\\ d & 7 & no & 36 & 3(4) & 4.47(3.74)\\\\ e & 8 & no & 45 & 4(5) & 5.66(4.90)\\\\ f & 9 & no & 55 & 4(5) & 5.66(4.90)\\\\ g & 10 & no & 66 & 5(6) & 6.83(6.06)\\\\ h & 11 & no & 78 & 5(6) & 6.83(6.06)\\\\ i & 12 & no & 91 & 6(7) & 8.00(7.21)\\\\ \\hline j & 4 & yes & 9 & 1(3) & 1.63(2.58)\\\\ k & 6 & yes & 16 & 2(4) & 2.83(3.74)\\\\ l & 8 & yes & 25 & 3(5) & 4.00(4.90)\\\\ m & 10 & yes & 36 & 4(6) & 5.16(6.06)\\\\ n & 12 & yes & 49 & 5(7) & 6.32(7.21)\\\\ o & 14 & yes & 64 & 6(8) & 7.48(8.37)\\\\ p & 16 & yes & 81 & 7(9) & 8.64(9.52)\\\\ \\hline \\end{tabular} \\caption{ \\label{tab:psf} Properties of the PSF models investigated in section \\ref{sec:spacesteppsf}. The first one is the elliptical gaussian (labelled `gauss'). For each of the 16 other basis sets, we study one PSF (labelled a to p) shown in figure \\ref{fig:psf}. From left to right, the columns are: $n_{\\textrm{max}}$ and diamond defined in section \\ref{sec:shapelets} and by \\cite{2005MNRAS.363..197M}; the number $N_{\\rm coef}$ of coefficients in the basis; the numbers $N_{\\rm coef}^{m=2(0)}$ of coefficients with $m=2$ (and $m=0$) in the basis and the complexity factors defined by equations \\ref{eq:psidef1} to \\ref{eq:psidef3}. } \\end{center} \\end{table*}"
    },
    "0711/0711.3232_arXiv.txt": {
        "abstract": "A common feature of hierarchical galaxy formation models is the process of ``inverse'' morphological transformation: a bulge dominated galaxy accretes a gas disk, dramatically reducing the system's bulge-to-disk mass ratio. During their formation, present-day galaxies may execute many such cycles across the Hubble diagram.  A good candidate for such a ``hermaphrodite'' galaxy is NGC 3108: a dust-lane early-type galaxy which has a large amount of HI gas distributed in a large scale disk. We present narrow band H$\\alpha$ and R-band imaging, and compare the results with the HI distribution from the literature. The emission is in two components: a nuclear bar and an extended disk component which coincides with the HI distribution. This suggests that  a stellar disk is currently being formed out of the HI gas. The spatial distributions of the H$\\alpha$ and  HI emission and the HII regions are consistent with a barred spiral structure, extending some $20$ kpc in radius.  We measure an extinction-corrected star formation rate of 0.42 \\msunyr. The luminosity function of the HII regions is similar to other spiral galaxies, with a power law index of -2.1, suggesting that the star formation mechanism is similar to other spiral galaxies.  We measured the current disk mass and find that it is too massive to have been formed by the current star formation rate over the last few Gyr. It is likely that the star formation rate (SFR) in NGC 3108 was higher in the past. With the current SFR, the  disk in  NGC 3108 will grow to be $\\sim 6.2\\times 10^9$ \\msun\\ in stellar mass within the next 5.5 Gyr. While this is substantial, the disk will be insignificant compared with the large bulge mass: the final stellar mass disk-to-bulge ratio will be $\\sim 0.02$. NGC 3108 will fail to transform into anything resembling a spiral without a boost in the SFR and additional supply of gas. ",
        "introduction": "It is widely accepted that galaxies may change their morphology as they evolve. Usually, such morphological transformation drives galaxies from late to early types, often as the result of a galaxy's interaction with its environment (Dressler 1980). For example, galaxy mergers are believed to transform spiral galaxies into ellipticals (Tommre \\& Tommre 1972, Barnes 1990) and the removal of gas from galaxies is thought to drive the transformation of spiral galaxies into the S0 types (Gunn \\& Gott, 1972, Larson et al.\\ 1980).  The ``inverse'' of such transformations is less widely considered. Yet, the regeneration of spiral disks in elliptical galaxies is a key component of current galaxy formation models (White \\& Frenk 1991, Cole et al.\\ 1994, Kauffmann et al.\\ 1993, Sommerville \\& Primack, 1999, De Lucia et al.\\ 2006, Bower et al.\\ 2006). In such models, galaxies vary in morphology as they grow potentially passing through several cycles back and forth across the Hubble diagram. This paper is the first in a series examining the observational evidence of such morphological cycles. Here, we focus on the gas-rich dust-lane elliptical NGC~3108, asking whether this could be an example of a ``hermaphrodite''\\footnote{In biology, a hermaphrodite is an organism that possesses both male and female genitalia during its life.  Some are able to change sex multiple times.  {\\it Sequential hermaphrodites} are organisms born as one sex and then later change into the other sex, an example being the Clownfish (source: www.wikipedia.org).  We find hermaphrodite a useful analogy to galaxies undergoing multiple transformations in a hierarchical Universe. We decided to use this term because our criteria go beyond the observationally motivated labels; for examaple, ``disky'' and ``S0'' do not describe the direction of the transformation.} galaxy under-going such an ``inverse'' morphological transformation.  Of course, on-going star formation in elliptical galaxies is not a new discovery. Many ellipticals have small amounts of ionized gas and star formation (eg., Caldwell 1984, Knapp et al.\\ 1985, Philips et al.\\ 1986, Goudfrooji et al.\\ 1994, Sadler et al.\\ 2000, Buyle et al.\\ 2006). However, the star formation is usually low-level and centrally concentrated, or associated with features indicating a recent merger event: of the galaxies in the SAURON HI survey with significant ionized gas content, only the polar ring galaxy NGC~2685 has an HI gas mass greater than $4\\times 10^7$ \\msun (Morganti et al., 2006). Low levels of on-going star formation are very likely to be driven by the reprocessing of stellar mass-loss during the late phases of stellar evolution, or the accretion of small gas rich companions (Emsellem et al.\\ 2007). Recent observations from GALEX seem to confirm this view. The large scatter in the near-ultraviolet colour-magnitude relation suggests that roughly 15\\% of $z < 0.13$ bright ($M(r) < -22$) early-type galaxies show signs of recent ($\\ale$ 1 Gyr) star formation at 1--2\\% level in mass compared to the total stellar mass, which implies that recent low-level residual star formation may be common even in bright early-type galaxies (Yi et al.\\  2005, Salim et al.\\ 2005). Such ``frosting'' is also supported by other studies (e.g. Trager et al 2000; Terlevich \\& Forbes 2002; Koo et al 2005) and the existence of young stellar disks at the centres of many elliptical galaxies is also becoming well established (Scorza et al 1998, Hau, Carter \\& Balcells 1999, Emsellem et al.\\ 2007).  However, while such residual activity is relatively common, these small amounts of star formation generate far fewer stars than the existing bulge and will never result in a significant change in the galaxy's morphology or bulge-to-disk ratio. In order to result in a genuine ``inverse'' transformation, the new episode of star formation must have an extent greater than the existing spheroid, and must eventually generate sufficient new stars to significantly lower the relative importance of the bulge light. Such galaxies are likely to be much more common in the distant universe, where the rates of merging and gas accretion are much higher than today. However, the concept of such cycles of morphology will receive considerable support if even one {\\it bona fide} example of an ``inverse'' morphological transformation could be established in the local universe.  Dust-lane ellipticals may be a good starting point for this search. Osterloo has observed such galaxies, discovering that several contain substantial HI gas mass distributed in a macroscopic disk (Oosterloo et al.\\ 2002, 2007, J\\'ozsa et al.\\ 2004). One prominent example is NGC~3108. In this paper we report on deep optical observations of this galaxy, probing the star formation associated with the HI gas disk.  In another paper, we will report the results of parallel search in a sample of isolated elliptical galaxies.  The structure of this paper is as follows. We summarise the available data on NGC~3108 in \\S2 and present our H$\\alpha$ imaging in \\S3. We derive star formation rates in \\S4, and consider the HII region luminosity function and the age of the stellar disk in \\S5 and~6. In the discussion (\\S7) we examine whether NGC~3108 is a plausible candidate for a hermaphrodite galaxy and compare it with two other nearby candidates NGC~5128 and the Sombrero galaxy (M104). We also briefly consider the implications in the context of current galaxy formation models. The conclusions are summarised in \\S8. ",
        "conclusions": "We investigated whether NGC 3108, one of the more promising candidates for a ``hermaphrodite'', is an early type galaxy which has been slowly transforming into a spiral galaxy. We have argued that such ``inverse morphological transformations\" are a fundamental component of galaxy formation models but are yet to be proven observationally.  Using H$\\alpha$ and R band imaging, we derived an extinction-corrected SFR of 0.42 \\msunyr. The emission is in two components: a nuclear bar which coincides with the hole (and offset from it) in the HI distribution, and an extended disk component which coincides with the HI distribution and mainly confined to HI densities of $31 \\times 10^{19}$ atoms cm$^{-2}$. This suggests that a stellar disk is currently being formed out of the HI gas. The HII regions are located close to the dust lanes and define a spiral structure; comparison with the HI distribution suggests that the HI is also distributed in a spiral arm structure. The mode of star formation is similar to that in other spirals, with the HII regions luminosity function most resembling that of M33 in shape.  We measured the current disk mass and find that it is probably too massive to be formed by the current star formation rate over the last few Gyr. While we cannot rule out that the current star formation may be more obscured than estimated, the amount of extinction required is too large. More likely that the SFR in NGC 3108 was higher in the past. With the current SFR, the  disk in  NGC 3108 will grow to be $6.2\\times 10^9$ \\msun\\ in stellar mass within a Hubble time. While this is substantial, the disk will be insignificant compared with the large bulge mass: the final stellar mass disk-to-bulge ratio will be $\\sim 0.02$. NGC 3108 will fail to transform into anything resembling a spiral without a boost in SFR and additional supply of gas.  NGC 3108 is compared with two other ``hermaphrodite'' candidates. Although there is a good morphological resemblance with NGC 5128, a ``symbiotic'' galaxy resembling a barred spiral residing inside an elliptical, the latter is more typical of a merger remnant with a short period of intense activities. The Sombrero galaxy with its dissolved outer bar and more massive stellar disk may represent the end point of a ``failed hermaphrodite''. All 3 galaxies will be unable to grow a sufficiently large disk to lower their bulge-to-total ratio unless there is a fresh supply of HI gas and the SFR can be boosted. In contrast, the observations suggest that the current episode of star formation will be short lived and ineffective."
    },
    "0711/0711.1237_arXiv.txt": {
        "abstract": "With a non local shell model of magnetohydrodynamic turbulence we investigate numerically the turbulent dynamo action for low and high magnetic Prandtl numbers ($Pm$). The results obtained in the kinematic regime and along the way to dynamo saturation are understood in terms of a phenomenological approach based on the local ($Pm\\ll 1$) or non local ($Pm\\gg 1$) nature of the energy transfers. In both cases the magnetic energy grows at small scale and saturates as an inverse `` cascade ''. ",
        "introduction": "Dynamo action is generally believed to be at the origin of the magnetic field in most astrophysical objects. The conducting fluid which produces the magnetic field, is generally strongly turbulent ($Re \\gg 1$) with an extended inertial range. In addition as it is electrically conducting, the fluid is characterized by the magnetic Prandtl number $Pm=\\nu / \\eta$, where $\\nu$ and $\\eta$ are the fluid viscosity and magnetic diffusivity. In a liquid metal as in planetary cores or stellar convective zones $Pm \\ll 1$, while in the warm interstellar medium, coronal and cluster plasmas $Pm\\gg 1$ \\cite{Schekochihin02a}. In direct numerical simulations both conditions,  $Re \\gg 1$ and $Pm$ much different from unity, are still out of reach of the present day computers. An alternative is to use shell models.  Hydrodynamical shell models are a rough approximation of the Navier-Stokes equations formulated on a discrete set of real wave numbers $k_n=\\lambda^n$, corresponding to a Fourier space divided into shells of logarithmic width $\\lambda$. They involve (complex) scalar quantities reminiscent of the velocity Fourier components of an isotropic flow. The model satisfies the conservation of both kinetic energy and kinetic helicity when the viscosity is set to zero. It leads to the resolution of a system of ordinary differential equations, requiring much less computing power than the direct numerical simulations of the original equations. It is then possible to run a shell model for realistic values of the viscosity such as $\\nu=10^{-6} m^2\\cdot s^{-1}$ (for a review on shell models, see Biferale (2003) and references therein).  The results obtained with shell models are in general agreement with the Kolmogorov phenomenology including intermittency \\cite{Leveque97}. Firstly at any scale $k^{-1}$ lying in the inertial range, the flux rate of kinetic energy is equal to the injection rate of kinetic energy at the forcing scale. This writes $ku^3(k)\\sim \\varepsilon$, leading to $u(k)\\sim \\varepsilon^{1/3}k^{-1/3}$ and to the Kolmogorov spectrum $E(k)\\sim \\varepsilon^{2/3}k^{-5/3}$. Secondly, the inertial range is found to extend to the viscous scale $k_{\\nu}^{-1}$ which can be evaluated saying that at this scale the turn over time $k^{-1}u^{-1}$ compares with the viscous time $\\nu^{-1}k^{-2}$, leading to $k_{\\nu}\\sim \\varepsilon^{1/4}\\nu^{-3/4}$. Thirdly at low viscosity some deviation of the Kolmogorov power scaling is found  due to intermittency. The power scaling of the statistical moments are given by $<|u(k)|>^p\\sim k^{-\\zeta_p}$ with scaling exponents $\\zeta_p$ deviating from Kolmogrov's mean field theory $\\zeta_p=p/3$. The value of these anomalous scaling exponents given by the shell models compare well with those found in experiments \\cite{Leveque97}. Therefore shell models appear to be a useful tool to study fully developed turbulence at high Reynolds numbers. They are also used in their MHD version to tackle astrophysical issues \\cite{Frick06,Buchlin07,Galtier07}.  In a previous paper \\cite{Plunian07}, denoted after by PS07, we introduced a non local shell model which is not only in agreement with the previous features of turbulence but in addition gives an appropriate slope of the infrared spectrum in freely decaying turbulence (which is not the case of the local shell models described in Biferale 2003). The magnetohydrodynamic (MHD) version of this model, also introduced in PS07, permits to study turbulent dynamo action for arbitrary low or high values of $Pm$. In PS07 we calculated the energy transfer functions of the MHD system in a (statistically stationary) saturated state. We found that for $Pm \\le 1$ the energy transfers are mainly local, eventually strengthening our previous results obtained with a local shell model of MHD turbulence \\cite{Stepanov06}. For $Pm \\gg 1$ the dominant transfers are also mainly local except the ones from the flow scales lying in the inertial range to the magnetic scales smaller than the viscous scale. In that case the use of a non local model is definitely necessary.  Here our goal is to describe the transient properties of the turbulent dynamo, namely the kinematic regime during which the magnetic energy grows exponentially, and the route to saturation, starting when the kinematic regime stops and ending when a (statistically stationary) saturated state is reached. For that it is not sufficient to make time averages as we did in PS07 when studying the saturated dynamo states. Indeed the transient properties we are interested in are, by definition, not statistically stationary. Then instead of time averages we have made ensemble averages over $M=12\\cdot10^3$ independent realisations. The initial conditions have the same energy but random phases. In addition the forcing has also a random phase in time (therefore different from one realisation to the other). Then the results between two realisations are different. Every $t_0=10^{-2}$ unit of time we freeze the calculation. For each shell we average the magnetic and kinetic energy over the $M$ realisations without modifying the phases of the magnetic, kinetic and forcing quantities in order to keep the randomness between every realisations. All the calculation were done for $N=60$ shells, corresponding to more than 12 decades. Before presenting the results (sections \\ref{lowpm} and \\ref{highpm} for low and high $Pm$) we briefly describe in the next section the shell model (previously introduced in PS07). ",
        "conclusions": "In the kinematic regime, we found that the magnetic growthrate satisfies $\\gamma \\sim (\\varepsilon / \\eta)^{1/2}$ for $Pm\\ll 1$ and $\\gamma \\sim (\\varepsilon / \\nu)^{1/2}$ for $Pm\\gg 1$ and is therefore always fast \\cite{Childress95}. In both cases this corresponds to a small-scale dynamo (or fluctuation dynamo in the terminology of Schekochihin \\textit{etal.} 2007) in opposition to a large-scale (or mean-field dynamo). At low $Pm$ the magnetic spectrum is peaked at $k=k_{kin} \\approx 0.4 k_{\\eta}$ and the energy transfer from kinetic to magnetic is mainly local (as shown in PS07 the non local transfers from small kinetic to large magnetic scales is less than 20 $\\%$ than the local transfers). At high $Pm$ the magnetic spectrum is almost flat between the viscous and resistive scale and the energy transfer is mainly non local from $u(k_{\\nu})$ to $b(k)$ with $k_{\\nu}<k<k_{\\eta}$. In the dynamic regime we found an inverse ``cascade'' mechanism which explains the route to saturation. The relevance of this mechanism does not depend on $Pm$. When all magnetic scales are saturated the resulting equilibrium between kinetic and magnetic energy is not equipartition with a slight excess of magnetic energy \\cite{Haugen03}. We find that the ratio $|b(k)/u(k)|$ is constant for $k_F<k<\\min\\{{k_{\\nu},k_{\\eta}}\\}$ corresponding to a residual energy with a Kolmogorov scaling \\cite{Muller05}.  Dynamo action is usually found in natural objects with strong rotation. A transition from a Kolmogorov spectrum $E(k)\\sim k^{-5/3}$ for weak rotation to $E(k)\\sim k^{-2}$ for strong rotation is predicted by weak turbulence theory, and observed in experiments \\cite{Baroud02,Baroud03} and numerical simulations \\cite{Smith99} including shell models \\cite{Hattori04}. It is then of interest to see how the previous phenomenology is changed. At low $Pm$ for $u(k)\\sim k^{-\\zeta}$ with $1/3\\le \\zeta \\le 1/2$, we derive the quantities given in table \\ref{summary} and give the numerical values expected for the three scalings, Kolmogorov, intermittent and strongly rotating. The change of the power scalings from weak to strong rotation should be sufficient to be identified in measurements or numerical calculations. On the other hand $k_{kin}/k_{\\eta}$ and $Rm(k_{kin})$ do not change much. \\begin{table} \\begin{tabular}{@{\\hspace{0.8cm}}r@{\\hspace{0.8cm}}c@{\\hspace{0.8cm}}c@{\\hspace{0.8cm}}c@{\\hspace{0.8cm}}c@{}} &Phenomenology&Kolmogorov&Intermittency&Strong rotation\\\\*[0cm] $\\zeta=$&             &$1/3$&$0.369$&$0.5$\\\\ $k_{kin}/k_{\\eta}=$ & $\\left(\\frac{1-\\zeta}{2}\\right)^{\\frac{1}{1+\\zeta}}$&0.439&0.431&0.397\\\\ $Rm(k_{kin})=$ & $\\frac{2}{1-\\zeta}$ & 3 & 3.17 & 4 \\\\ $k_{kin},k_{\\eta} \\sim$ & $\\eta^{-\\frac{1}{1+\\zeta}}$ & $\\eta^{-0.75}$ & $\\eta^{-0.73}$ & $\\eta^{-0.66}$ \\\\ $k_{\\nu} \\sim$ & $ \\nu^{-\\frac{1}{1+\\zeta}}$ & $\\nu^{-0.75}$ & $\\nu^{-0.73}$ & $\\nu^{-0.66}$\\\\ $\\gamma_{kin} \\sim$ & $ \\eta^{-\\frac{1-\\zeta}{1+\\zeta}}$ & $\\eta^{-0.5}$ & $\\eta^{-0.46}$ & $\\eta^{-0.33}$ \\\\ $\\frac{d(2\\gamma)^{-1}}{dt} \\sim$ & $  \\frac{1-\\zeta}{a+\\zeta}$&0.143&0.133&0.1 \\end{tabular} \\caption{Summary of the numerical values expected from phenomenological predictions at low $Pm$ for different velocity spectra $u(k)\\sim k^{-\\zeta}$.} \\label{summary} \\end{table}"
    },
    "0711/0711.4138_arXiv.txt": {
        "abstract": "The majority of pulsar population synthesis studies performed to date have focused on isolated pulsar evolution. Those that have incorporated pulsar evolution within binary systems have tended to either treat binary evolution poorly of evolve the pulsar population in an ad-hoc manner. Here we present the first model of the Galactic field pulsar population that includes a comprehensive treatment of both binary and pulsar evolution. Synthetic observational surveys mimicking a variety of radio telescopes are then performed on this population. As such, a complete and direct comparison of model data with observations of the pulsar population within the Galactic disk is now possible. The tool used for completing this work is a code comprised of three components: stellar/binary evolution, Galactic kinematics and survey selection effects. Here we give a brief overview of the method and assumptions involved with each component. Some preliminary results are also presented as well as plans for future applications of the code. ",
        "introduction": "We evolve n systems by first selecting the initial parameters of each single star (mass) or binary (masses, period) from appropriate distributions, e.g.- stellar masses drawn from the Kroupa, Tout \\& Gilmore (1993) initial mass function. Limits must also be chosen where we typically take $5-80 M_\\odot$ and $0.1-40 M_\\odot$ for the binary primary and secondary masses. The algorithm loops over the n systems in \\textbf{binpop} which follows the stellar/binary evolution. Any systems of interest that are produced (in this case pulsars) are passed to \\textbf{binkin} which follows the kinematic evolution of each system. The resulting population is then `surveyed' in \\textbf{binsfx} which produces synthetic data for analysis. ",
        "conclusions": ""
    },
    "0711/0711.2528_arXiv.txt": {
        "abstract": "We study the individual contribution to secondary lepton production in hadronic interactions of cosmic rays (CRs) including resonances and heavier secondaries. For this purpose we use the same methodology discussed earlier \\cite{Huang07}, namely the Monte Carlo particle collision code DPMJET3.04 to determine the multiplicity spectra of various secondary particles with leptons as the final decay states, that result from inelastic collisions of cosmic-ray protons and Helium nuclei with the interstellar medium of standard composition. By combining the simulation results with parametric models for secondary particle (with resonances included) for incident cosmic-ray energies below a few GeV, where DPMJET appears unreliable, we thus derive production matrices for all stable secondary particles in cosmic-ray interactions with energies up to about 10 PeV.  We apply the production matrices to calculate the radio synchrotron radiation of secondary electrons in a young shell-type SNR, RX J1713.7-3946, which is a measure of the age, the spectral index of hadronic cosmic rays, and most importantly the magnetic field strength. We find that the multi-mG fields recently invoked to explain the X-ray flux variations are unlikely to extend over a large fraction of the radio-emitting region, otherwise the spectrum of hadronic cosmic rays in the energy window 0.1-100~GeV must be unusually hard.  We also use the production matrices to calculate the muon event rate in an IceCube-like detector that are induced by muon neutrinos from high-energy $\\gamma$-ray sources such as RX J1713.7-3946, Vela Jr. and MGRO J2019+37. At muon energies of a few TeV, or in other word, about 10~TeV neutrino energy, an accumulation of data over about five to ten years would allow testing the hadronic origin of TeV $\\gamma$-rays. ",
        "introduction": "\\label{Section:Introduction} It is known that hadronic cosmic rays interacting with the interstellar medium (ISM) can produce $\\gamma$-rays and leptons. The principal production mechanisms for $\\gamma$-rays in high-energy astrophysics are inelastic CR+ISM interactions with subsequent decays of the secondaries (mostly the neutral $\\pi^0$ and $\\eta$) into $\\gamma$-rays. The decays of charged $\\pi^\\pm$ mesons and other secondaries will result in the production of high-energy neutrinos and secondary leptons, electrons and positrons. In GeV-scale Galactic cosmic rays, the secondary electrons and positrons comprise around 20\\% of the total electron flux and significantly contribute to the non-thermal electromagnetic radiation of $\\gamma$-ray sources. The flux of cosmic-ray neutrinos is determined by the cosmic-ray abundance and the production cross sections of the parent particles ($\\pi^\\pm,~K^\\pm,~K^0_S,~K^0_L,~\\Sigma^\\pm,~\\Sigma^0,~\\Lambda,~\\bar{\\Lambda}$) in the collisions of cosmic-rays with the interstellar gas nuclei.  The purpose of this work is to carefully calculate the lepton production in cosmic-ray interactions. For this purpose we use the high-energy physics event generator DPMJET-III \\cite{Roesler00} to simulate all secondary productions in both p-generated and He-generated interactions, similar to earlier work on hadronic $\\gamma$-ray production \\cite{Huang07}. Our study includes direct lepton production and the decays of all relevant secondary particles with leptons as the final decay products. For the cosmic rays, protons and helium nuclei are taken into account. We assume the composition of the ISM as 90\\% protons, 10\\% helium nuclei, 0.02\\% carbon, and 0.04\\% oxygen. At energies below 20 GeV, where DPMJET appears unreliable, we combine the simulation results with parametric models for $\\pi^\\pm$ and $\\pi^0$ production \\cite{Blattnig00,Kamae06} that include the production of the resonances $\\Delta$(1232) and $\\Delta$(1600) and their subsequent decays, thus deriving a lepton production matrix for cosmic rays with energies up to about 10 PeV that can be easily used to interpret the spectra of cosmic lepton sources.  We use our production matrices to calculate the CR-induced electron/positron production in young shell-type supernova remnants (SNR), for which the neutrino fluxes are also calculated. Observations of non-thermal X-ray synchrotron radiation from the SNRs SN~1006 \\cite{Koyama95}, RX J1713.7-3946 \\cite{Koyama97}, IC 443 \\cite{Keohane97,Slane99}, Cas A \\cite{Allen97}, and RCW 86 \\cite{Borkowski01} support the hypothesis that the Galactic cosmic-ray electrons may be accelerated predominantly in SNR, although other acceleration sites of high-energy electrons might exist in the Galaxy. The existence of electrons in SNRs with an energy up to about 100 TeV implies TeV-scale $\\gamma$-ray emission \\cite{Pohl96} besides that possibly produced by hadronic cosmic rays, and TeV-scale gamma-rays have indeed been detected from the SNRs Vela Junior \\cite{Aharonian05VelaJr} and RX J1713.7-3946 \\cite{Aharonian06RXJ1713.7-3946}. Although the hadronic interpretation is often favored to explain the $\\gamma$-ray emission from SNRs such as RX J1713.7-3946 \\cite{Aharonian06RXJ1713.7-3946}, there is still no direct observational evidence for nucleon acceleration in SNRs \\cite{Reimer02, Huang07}. Currently, both hadronic and leptonic models can describe the observed TeV $\\gamma$-rays in different astrophysical objects, but have problems of similar magnitude \\cite{Katz07}.  RX J1713.7-3946 is a unique SNR in the sense that its X-ray emission is strongly dominated by a non-thermal component, which is presumed to be synchrotron radiation of ultrarelativistic electrons (see references \\cite{Hiraga05, Drury01,Uchiyama07} and references therein for details). The recent broadband X-ray spectroscopy performed with Chandra \\cite{Uchiyama07} and the Suzaku experiment \\cite{Takahashi07} provides evidence of very effective acceleration of particles in the shell of RX J1713.7-3946 and also reveals that the X-ray emission from two compact regions is variable in flux, which the authors interpret as indication of multi-mG magnetic fields. If this conclusion is correct and applies to the entire remnant, then the observed TeV-band $\\gamma$-ray emission must be of hadronic origin. Here we show that the synchrotron flux from the secondary electrons in RX~J1713.7-3946 would in fact exceed the observed radio flux from this object, if multi-mG magnetic fields would permeate the remnant, so a magnetic field of such magnitude will likely exist at most in a small fraction of emission region.  In a hadron accelerator, TeV neutrinos should be produced in roughly the same number as TeV $\\gamma$-rays. The Universe is generally more transparent for high-energy neutrinos than for TeV and PeV $\\gamma$-rays, that can be absorbed by pair production with either the microwave background (CMB) or the infrared/optical photon background (IRB). Therefore the detection of astrophysical neutrinos can provide invaluable complementary information on the existence of hadronic processes at astrophysical objects. Here we calculate the observable TeV-scale neutrino flux from RX~J1713.7-3946 and other sources. This is not the first attempt to do such work. In fact, several calculations of the expected $\\gamma$-ray and neutrino spectra were made for parametrized hadronic particles in the energy range of neutrino detectors \\cite{AlvarezMuniz02,Costantini05,Vissani06,Lipari06,Kelner06,Kappes07,Beacom07}. In contrast to the earlier works, we include all relevant channels of neutrino production and also we do not rely on assuming a simple power-law proton spectrum or a power-law with an exponential cutoff, because our neutrino production matrix can be applied to energetic particles with arbitrary spectrum. ",
        "conclusions": "In addition to Galactic $\\gamma$-ray emission, the Galactic neutrino emission could provide complementary information on the existence of hadronic processes in astrophysical objects and also on the origin of Galactic cosmic-rays. Here we have calculated the $e^\\pm$, $\\nu_e$ and $\\nu_\\mu$ production in cosmic-ray interactions and derived matrices that allow to estimate the spectra of secondary leptons for arbitrary spectra of the parent hadronic cosmic rays.  With the production matrices, we calculate the synchrotron radiation of secondary electrons in the shell-type SNR RX~J1713.7-3649 under the assumption that the observed TeV-scale gamma-ray emission is hadronic in origin. We find that the radio synchrotron flux of the secondary electrons exceeds the observed flux level if the magnetic field strength is too high. The multi-mG fields recently invoked to explain the X-ray flux variations \\cite{Uchiyama07} are therefore unlikely to extend over a large fraction of the radio-emitting region or the spectrum of hadronic cosmic rays in the 0.1-100~GeV energy window must be unusually hard.  Our calculations for neutrinos show that the neutrino detection rates from the TeV $\\gamma$-ray sources RX J1713.7-3946 and Vela Jr. are promising for muon energies of a few TeV, or in other words, about 10~TeV neutrino energy. An accumulation of data over about five to ten years would allow testing the hadronic origin of TeV $\\gamma$-rays.  The production matrices for $e^-$, $e^+$, $\\nu_e$, $\\bar{\\nu}_e$, $\\nu_\\mu$ and $\\bar{\\nu}_\\mu$ will be made available for download at {\\em http://cherenkov.physics.iastate.edu/lepton-prod}."
    },
    "0711/0711.2002_arXiv.txt": {
        "abstract": "The amplitude of the Compton-Getting (CG) anisotropy contains the power-law index of the cosmic-ray energy spectrum. Based on this relation and using the Tibet air-shower array data, we measure the cosmic-ray spectral index to be $-3.03 \\pm 0.55_{stat} \\pm < 0.62_{syst}$ between 6 TeV and 40 TeV, consistent with $-$2.7 from direct energy spectrum measurements. Potentially, this CG anisotropy analysis can be utilized to confirm the astrophysical origin of the ``knee'' against models for non-standard hadronic interactions in the atmosphere. ",
        "introduction": "\\label{intro} The cosmic-ray (CR) energy spectrum has been observed using ground-based air shower detectors and various instruments on board satellites and balloons. It features a power law d$N$/d$E$ $\\propto$ $E^{- \\gamma}$ in the energy range from 30 GeV to 100 EeV, where $N$ denotes the number of cosmic rays at the top of the atmosphere, $E$ the cosmic-ray particle energy. The spectral index $\\gamma$ changes from 2.7 to 3.1 around the ``knee'' at $E \\simeq$ 4 PeV. Various models have been proposed for the origin of this knee; most models suppose that it is of astrophysical origin \\cite{Hoerandel2}, while there are some models that attribute it to non-standard high-energy nucleon interactions in the Earth's atmosphere.  For a power-law CR energy spectrum, the Compton-Getting (CG) anisotropy involves the spectral index $\\gamma$. When a detector observing cosmic rays moves relative to the rest frame of the CR plasma in the Earth orbit around the Sun, a fractional CR intensity variation $\\Delta I$ can be observed in the solar time frame. For a presumed power-law CR energy spectrum, this CG anisotropy \\cite{CG-1,CG-2} is expressed as \\begin{equation} \\frac{\\Delta I}{<I>} = (\\gamma + 2)\\frac{v}{c}\\cos\\theta, \\label{eq-1} \\end{equation} where $I$ denotes the CR intensity, $\\gamma$ the power-law index of the cosmic-ray energy spectrum, $v$ the orbital velocity, $c$ the speed of light and $\\theta$ the angle between the arrival direction of CRs and the moving direction of the observer.  At multi-TeV energies, a clear evidence for the CG anisotropy in the solar diurnal variation was first reported in \\cite{C-G}. The observed variation was in reasonable agreement with a sinusoidal curve expected from the CG anisotropy with the maximum phase of the curve deviating from 6:00 h by $+$2 hours at 2$\\sigma$ significance. This deviation was ascribed to meteorological effects on the underground muon intensity. The Tibet air-shower experiment performed the first study on energy dependence of the solar diurnal variation and showed that it was consistent with the CG anisotropy above multi-TeV energies \\cite{UDO,COROTATE}.  Conventionally, the spectral index above $\\sim$100 TeV is obtained by indirect measurements with air-shower arrays, in which each CR particle energy is determined from its shower size. The relation between particle energy and shower size is deduced from air-shower simulations assuming hadronic interaction models as a reasonable extrapolation of the accelerator results to the higher energy range. The spectral index thus obtained is valid only when no violation of the standard hadronic interaction is expected in the relevant energy range. This fact leaves a room for an argument whether the knee is due to non-standard high-energy nucleon interactions in the atmosphere or not.  On the other hand, the CG anisotropy does not depend on CR particle energy by Eq.(\\ref{eq-1}). We can then measure the power-law index of the all-particle energy spectrum in the way it is virtually uninfluenced by hadronic interaction models, because air-shower simulations are used only for the purpose of estimating the rough energy range where we are observing the CG anisotropy. Potentially, this method makes it possible to confirm that the bend of the energy spectrum at the knee is of astrophysical origin, not due to non-standard hadronic interactions in the atmosphere. ",
        "conclusions": "The CR events were divided into eight $\\sum \\rho_{\\mbox{\\scriptsize FT}}$ bins in order to find out the energy region where the observed solar daily variation is free from effects other than the CG anisotropy. The fitting curve for the observed solar daily variation is expressed as \\begin{equation} f(\\lambda) = \\alpha \\cos\\left(\\frac{\\pi}{12}(\\lambda - \\phi) \\right), \\label{eq-2} \\end{equation} where $\\alpha$ denotes the amplitude of the CG anisotropy, $\\lambda$ [hr] the local solar time and $\\phi$ [hr] the phase where the sinusoidal curve reaches its maximum. Note that in a differential form, the phase is shifted earlier by 6 hours (1/4 cycle) from the corresponding actual daily variation and the amplitude is $\\pi$/12 times that of the physical CG anisotropy.  Figure \\ref{fig1}(a) shows $\\sum \\rho_{\\mbox{\\scriptsize FT}}$ dependence of the phase $\\phi$ in Eq.(\\ref{eq-2}). Although there exists deviation due to some contamination in the region of $\\sum \\rho_{\\mbox{\\scriptsize FT}} < 50$, the measured phase is consistent with the expected value of the CG anisotropy in $\\sum \\rho_{\\mbox{\\scriptsize FT}} \\ge 50$. Therefore, we regard the observed variation in $\\sum \\rho_{\\mbox{\\scriptsize FT}} \\ge 50$ as caused by the CG anisotropy. Figure \\ref{fig1}(b) shows $\\sum \\rho_{\\mbox{\\scriptsize FT}}$ dependence of the observed amplitude $\\alpha$. Since the CR energy spectrum in the multi-TeV region is expressed by a single power law, the amplitude should be consistent with a constant line if there is no contamination. This figure confirms that we do observe the CG anisotropy in the region of $\\sum \\rho_{\\mbox{\\scriptsize FT}} \\ge 50$. Thus, we report here on the spectral index using $9.7 \\times 10^{9}$ events with $\\sum \\rho_{\\mbox{\\scriptsize FT}} \\ge 50$. This region is roughly between 6 TeV and 40 TeV, corresponding to 20\\% and 80\\% of the median energy distribution, respectively.  Figure \\ref{fig2}(a) shows the observed solar daily variation data fitted with a sinusoidal curve. The $\\chi^2$-fitting results are summarized in Table I. From the amplitude of the solar daily variation, we obtained $\\gamma$ of CRs using the following relation, \\begin{equation} \\alpha = (\\gamma + 2) \\frac{v}{c} \\frac{\\pi}{12} F, \\label{eq-3} \\end{equation} where $\\alpha$ denotes the amplitude, $v = 2.978 \\times 10^4$ [m/s] the average orbital velocity of the Earth and $F = 0.827$ an effective geometrical factor by which $\\alpha$ decreases, calculated according to the latitude at Yangbajing site (see Appendix). The factor $\\pi$/12 emerges, because we measure the anisotropy in the differential form Eq.(\\ref{EWeq}). As a result, we found $\\alpha = (10.9 \\pm 1.2_{stat}) \\times 10^{-5}$. In Eq.(\\ref{eq-3}), we omitted a minor correction. Strictly speaking, the orbital velocity of the Earth is time dependent. Although it varies $\\pm$2\\%, we used the average $v$ instead. It may become necessary to take this correction into account, when more data are accumulated.  A seasonal change of the sidereal daily variation due to the galactic anisotropy could produce a spurious variation in the solar time frame. The differential variation in the local extended-sidereal time (367.2422 cycles/yr) is shown in Figure \\ref{fig2}(b). The insignificant variation in the extended-sidereal time $(0.7 \\pm 1.2_{stat}) \\times 10^{-5}$ ensures that contamination in the solar daily variation due to the seasonal change of the sidereal daily variation is less than 12\\% of the CG anisotropy. Adding this as a systematic error, we found $\\gamma = 3.08 \\pm 0.55_{stat} \\pm < 0.62_{syst}$. This systematic error would become smaller in the future when more data are accumulated.  As the detection efficiency of the Tibet air-shower array depends on CR nuclei in this energy range (6 $-$ 40 TeV), the correction for it should be taken into account. An air-shower simulation employing direct observational data for primary cosmic rays revealed that $-$0.05 should be added to the $\\gamma$ above \\cite{HD4}. This correction is not needed, however, above $\\sim$60 TeV where the detection efficiency becomes independent of nuclei, reaching $\\sim$100\\%.  Finally, the spectral index $\\gamma$ between 6 and 40 TeV turned out to be $3.03 \\pm 0.55_{stat} \\pm < 0.62_{syst}$, consistent with 2.7 from direct energy spectrum measurements.  Using this CG anisotropy method, a future high-statistics experiment with a huge effective area will be able to measure the power-law index of the all-particle energy spectrum at higher energies. The spectral index $\\gamma$ measured by this method can be compared not only with those $\\gamma$ values by direct measurements below $\\sim$100 TeV but also with those by indirect air-shower analyses above $\\sim$100 TeV. This would check the mutual consistency between different ways to measure the CR energy spectrum. Following the same strategy, this CG anisotropy analysis can be applied to much higher energies to confirm the astrophysical origin of the knee."
    },
    "0711/0711.0007_arXiv.txt": {
        "abstract": "Our understanding of fundamental processes in the solar corona has been greatly progressed based on the space observations of SMM, Yohkoh, Compton GRO, SOHO, TRACE, RHESSI, and STEREO. We observe now acoustic waves, MHD oscillations, turbulence-related line broadening, magnetic configurations related to reconnection processes, and radiation from high-energy particles on a routine basis. We review a number of key observations in EUV, soft X-rays, and hard X-rays that innovated our physical understanding of the solar corona, in terms of hydrodynamics, MHD, plasma heating, and particle acceleration processes. ",
        "introduction": "Our grasp to understand the fundamental physical processes in the solar corona in more depth mostly benefits from the new high-resolution imaging capabilities that became available in extreme ultraviolet (EUV), soft X-rays, hard X-rays, and gamma rays. These wavelengths are particularly revealing in the following aspects: The corona is filled with high-temperature plasma at temperatures of $T\\approx 1-2$ MK, which radiates free-free emission with its emission measure peak in EUV wavelengths. Hot active regions and postflare loops have plasma temperatures of $T\\approx 2-40$ MK with their free-free emission peak in soft X-ray wavelengths. The nonthermal emission of flare-accelerated high-energy particles radiate bremsstrahlung in hard X-ray wavelengths and have nuclear de-excitation lines in gamma rays. Because of the absorbing properties of the Earth atmosphere, the Sun can be observed in these wavelengths only from space.  We discovered waves in the solar corona from oscillatory plasma motions and from propagating density disturbances. Turbulence is somewhat harder to detect, but we mostly infer if from line broadening in EUV or from scintillation experiments in radio wavelengths. Magnetic reconnection is most conspicuously inferred from the post-reconnection topology of post-flare loop systems, but we detect also plasma inflows and outflows in reconnection regions. Particle acceleration is still the biggest black box because it occurs on microscopic scales, but we can localize the acceleration regions from direct bremsstrahlung in the acceleration region itself, or from reconstructed particle trajectories using time-of-flight measurements. Here we review the fundamental physics of these four processes, applied to solar and stellar coronae, and touch on some new challenging observations. ",
        "conclusions": "Our insights into the fundamental physical processes in the solar corona is exponentially growing thanks to all the new high-resolution imaging and spectral observations. We have identified almost all principal MHD wave modes in the solar corona, and theoretical studies are exploring now second-order effects, which provide previously unknown physical parameters. Turbulence is a process that is established in the lower solar atmosphere (driven by the subsurface magneto-convection), as well as in the solar wind (caused by interacting outgoing and reflected Alfv\\'en waves), but seems to be less important in the closed-field corona (due to the large dissipation length of Alfv\\'en waves). Magnetic reconnection is an ubiquitous process in the solar corona, ranging from small-scale phenomena in the transition region (explosive events, nanoflares, etc) to catastrophic large-scale reconfigurations during flares and CMEs, but modeling problems deal mostly with the relative importance of the drivers (magnetic flux emergence, cancellation, loss-of-equilibrium, magnetic break-out, etc). The process of particle acceleration is, due to its microscopic physics, the most difficult to study, and most conclusions are drawn from secondary processes, such as the collisional thin- and thick-target interactions of the accelerated particles, or in-situ particle detections in interplanetary space. It is imperative to get a grip on the 3D geometry and topology of magnetic reconnection regions and shock waves in order to get a glimpse on the electromagnetic fields that accelerate high-energy particles. It is clear that only multi-wavelength observations with comprehensive modeling can lead to a deeper physical understanding and to more reliable conclusions about these fundamental physical processes in the solar corona. Conclusions about the operation of the same physical processes in stellar coronae can only be drawn by inference, but the sensitivity threshold of stellar observations generally imply a bias towards more energetic flare processes."
    },
    "0711/0711.2234_arXiv.txt": {
        "abstract": "We study the consequences of limited recoil sense reconstruction on the number of events required to reject isotropy and detect a WIMP signal using a directional detector. For a constant probability of determining the sense correctly, 3-d read-out and zero background, we find that as the probability is decreased from 1.0 to 0.75 the number of events required to reject isotropy using the mean angle statistic is increased by a factor of a few. As the probability is decreased further the number of events required using this statistic increases sharply, and in fact isotropy can be rejected more easily by discarding the sense information and using axial statistics. This however requires an order of magnitude more events than vectorial data with perfect sense determination. We also consider energy dependent probabilities of correctly measuring the sense, 2-d read-out and non-zero background. Our main conclusion regarding the sense determination is that correctly determining the sense of the abundant, but less anisotropic, low energy recoils is most important for minimising the number of events required. ",
        "introduction": "Weakly Interacting Massive Particle (WIMP) direct detection experiments aim to directly detect non-baryonic dark matter via the elastic scattering of WIMPs on detector nuclei~\\cite{DD}, and are presently reaching the sensitivity required to detect the lightest neutralino (which in most supersymmetry models is the lightest supersymmetric particle and an excellent WIMP candidate). Since the expected event rates are very small ( ${\\cal O} (10^{-5} - 1)$ counts ${\\rm kg^{-1} day^{-1}}$) distinguishing a putative WIMP signal from backgrounds due to, for instance, neutrons from cosmic-ray induced muons or natural radioactivity, is crucial. The direction dependence of the WIMP scattering rate (due to the Earth's motion with respect to the Galactic rest frame)~\\cite{dirndep} provides a potential WIMP `smoking gun' and low pressure gas time projection chambers (TPCs), such as DRIFT ({\\bf D}irectional {\\bf R}ecoil {\\bf I}dentification {\\bf F}rom {\\bf T}racks)~\\cite{drift,sean:drift} and NEWAGE~\\cite{newage}, see also Ref.~\\cite{martoff}, seem to offer the best prospects for a detector capable of measuring the directions of sub-100 keV nuclear recoils.  Early studies found that in principle as few as 30 events would be required to distinguish a WIMP induced signal from isotropic backgrounds~\\cite{copi:krauss,lehner:dir}. In reality the number of events, and hence the exposure required, depends on the detector properties including the energy threshold, background event rate, whether the read-out measures the recoil momentum in 2 or 3 dimensions (and if 2-d in which plane)~\\cite{pap1,pap2,copi2d,pap3} and whether the sense (i.e. the absolute sign $+ {\\vec{p}}$ or $- { \\vec{p}}$) of the recoil momentum vectors can be measured~\\cite{pap1,pap2,copi2d,pap3,vf}.  The factor with the biggest effect on the number of events required is whether or not the sense of the recoils can be measured~\\cite{pap1,pap2,pap3}. If the data is axial (i.e. the sense can not be measured) the exposure is increased, compared to the vectorial case, by one order of magnitude for 3-d read-out and at least two orders of magnitude for 2-d read-out in the optimal plane~\\footnote{This can be improved to a factor $\\sim 30$ if the reduced angles (with the direction of solar motion subtracted) are calculated and analyzed.}~\\cite{pap3}. At low energies, the energy deposition dE/dx of nuclear recoils, and hence the density of ionisation created, is predicted to slowly decrease with decreasing recoil energy~\\cite{hitachi-dedx}. Thus the sense of a recoil track is, in principle, measurable by determining the direction in which the ionisation density decreases along the track.  Previous work on simulating recoil tracks using SRIM2003~\\cite{srim} indicated that the ionization density distributions reconstructed from recoil tracks were close to uniform due to both fluctuations in the production of ionisation and diffusion during the drift of this ionisation to the read-out plane~\\cite{pap1,pap3}.  It was therefore not possible to determine the absolute signs of the reconstructed recoil vectors in these simulations. It should be noted however that while SRIM2003 predicts sulfur recoil ranges and quenching factors (fraction of recoil energy going into ionisation) in agreement with experimental data~\\cite{drift:neutron}, it was not designed to model recoils in gaseous targets.  Whether or not the sense of low energy nuclear recoils can be determined is therefore an important question which needs to be resolved experimentally. Ongoing studies~\\cite{cygnus,senseexpt} indicate that the expected decrease in dE/dx is observable, but fluctuations in the ionization distribution caused by the ionization process and diffusion may wash out the information in the ionization distribution leading to errors in sense reconstruction. In other words, if the sense of a recoil of energy $E$ is determined via a hypothesis test on some parameter $M$ derived from its reconstructed ionization distribution, then the PDFs $g(M| {\\rm forward \\, sense},E)$ and $g(M| {\\rm backward \\, sense},E)$ will not be cleanly separated, i.e. the significance level, $\\mathcal{S}(E)$, is greater than zero and the power, $\\mathcal{P}(E)$, is less than one. Here, the significance level is the probability of rejecting the forward sense hypothesis when it is true, and the power is the probability of rejecting the backward sense hypothesis when it is false~\\cite{cowan:stats}. Provided $\\mathcal{S}(E)<0.5$ and $\\mathcal{P}(E)>0.5$, the correct recoil sense will be reconstructed more often than not with probability $P({\\rm correct \\, sense} |E)=1-\\mathcal{S}(E)$. In this Brief Report we therefore study the effect of this probabilistic sense reconstruction on the number of events required to distinguish a WIMP induced recoil signal from isotropic backgrounds. ",
        "conclusions": "In this Brief Report we have investigated the effect of probabilistic recoil sense determination on the number of events required by a directional detection experiment to reject isotropy, and detect a WIMP signal. We have considered a constant sense determination probability and also increasing probability with increasing recoil energy (higher energy recoils have longer tracks and hence it is likely that it will be easier to determine their sense). For a constant probability of determining the sense correctly, 3-d read-out and zero background, we find that as the probability is decreased from 1.0 to 0.75 the number of events required to reject isotropy using the mean angle statistic is increased by a factor of a few. As the probability is decreased further the number of events required using this statistic increases sharply, and in fact for $P_{100 \\,{\\rm keV}}<0.7$ isotropy can be rejected more easily by discarding the sense information and using axial statistics (this does however require an order of magnitude more events than vectorial data with perfect sense determination).  The linearly increasing sense determination probability with $P_{cs}(E_{\\rm th})=0.75 \\, (0.50)$ and $P_{100 \\, {\\rm keV}}= 1.0 $ requires $\\sim 3$ times (an order of magnitude) more events than the constant probability with the same value of $P_{100 \\, {\\rm keV}}$. This suggests that maximising the probability of correctly determining the sense of the low energy, more common, but less anisotropic events is most important for minimising the number of events required. We have confirmed this conclusion by considering step function sense determination probabilities.  We also considered non-zero background and 2-d read-out. As the signal to noise ratio is decreased the numbers of WIMP events required in each case are increased, by factors of $ \\sim 1.1-1.2 , \\, 2-4$ and $10$ for $S/N= 10, \\, 1 $ and $0.1$ respectively.  For 2-d read-out with perfect angular resolution, the numbers of events required are roughly a factor of $\\sim 2$ larger than for 3-d read-out with the recoil direction reconstruction uncertainty taken into account (with the same sense determination probability and signal to noise ratio).        \\begin{figure} \\includegraphics[width=8.0cm]{step.ps} \\caption{\\label{step} $N_{95}$ for a step function probability of correctly determining the sense, $P_{cs}(E<E_{\\rm step})=0.5 \\, (1.0)$, $P_{cs}(E>E_{\\rm step}) =1.0 \\, (0.5)$ plotted using circles (triangles), for 3d read-out and zero background.  The dotted line shows the number of events required using the value of $\\langle |\\cos{\\theta}| \\rangle$, which does not require sense information.  }  \\end{figure}"
    },
    "0711/0711.0231_arXiv.txt": {
        "abstract": "We present ammonia observations of 193 dense cores and core candidates in the Perseus molecular cloud made using the Robert F. Byrd Green Bank Telescope.  We simultaneously observed the NH$_3$(1,1), NH$_3$(2,2), C$_2$S ($2_1\\to 1_0$) and C$_2 ^{34}$S($2_1\\to 1_0$) transitions near $\\nu = 23$ GHz for each of the targets with a spectral resolution of $\\delta v \\approx 0.024 \\mbox{ km s}^{-1}$.  We find ammonia emission associated with nearly all of the (sub)millimeter sources as well as at several positions with no associated continuum emission.  For each detection, we have measured physical properties by fitting a simple model to every spectral line simultaneously.  Where appropriate, we have refined the model by accounting for low optical depths, multiple components along the line of sight and imperfect coupling to the GBT beam.  For the cores in Perseus, we find a typical kinetic temperature of $T_k=11$~K, a typical column density of $N_{\\mathrm{NH3}}\\approx 10^{14.5}\\mbox{ cm}^{-2}$ and velocity dispersions ranging from $\\sigma_v = 0.07\\mbox{ km s}^{-1}$ to 0.7$\\mbox{ km s}^{-1}$.  However, many cores with $\\sigma_v >0.2\\mbox{ km s}^{-1}$ show evidence for multiple velocity components along the line of sight. ",
        "introduction": "Ammonia remains one of the best molecules for studying the cool, dense molecular cores where most stars form.  The utility of ammonia was recognized early in the pursuit of molecular line astronomy \\citep[and references therein]{nh3-araa} and it remains the standard for identifying and studying the internal conditions of dense molecular cores. The unique quantum structure of the molecule coupled with its relative abundance allows for a host of measurements to be made from the hyperfine transitions among the multiple metastable states, which emit near $\\nu=23$ GHz.  A single spectrum can be used to determine the line-of-sight velocity, velocity dispersion and gas kinetic temperature for a dense core.  This set of properties form an excellent complement to surveys of submillimeter emission \\citep[e.g.][]{motte-andre,testi-sargent} which readily study the size and distribution of the dust emission in cores, while yielding no information about the kinematics and temperatures.  Recently, submillimeter surveys of dense cores have been extended to cover large fractions of molecular clouds \\citep{hatchell05,bolocam-perseus}, making complete surveys of dense cores possible.  In addition, large scale mapping projects have surveyed several nearby molecular clouds at high resolution in emission from the isotopomers of CO \\citep[e.g., the COMPLETE Surveys of Serpens, Ophiuchus, Perseus;][]{complete-data}.  The CO surveys establish the cores in the larger context of the molecular cloud. However, these surveys have raised many questions about the properties of cores and their relationship to the larger molecular environment.  To measure the kinetic temperature and kinematics of an unbiased sample of dense cores embedded in the same molecular complex, we have conducted a survey of dense cores in ammonia across the Perseus molecular cloud using the Green Bank Telescope.  There are several advantages to adopting Perseus as a target.  It has been extensively studied in several observational campaigns: the COMPLETE survey of star forming regions which surveyed the molecular gas using the FCRAO 14-m \\citep{complete-data}; the SCUBA survey of submillimeter emission \\citep{hatchell05,perseus-scuba}; the BOLOCAM survey of the region in the 1.1 mm continuum \\citep{bolocam-perseus}; and in all the Spitzer bands by the c2d project \\citep{c2d-pers-irac,c2d-pers-mips}.  The locations of dense cores have been identified in the (sub)millimeter maps, and their protostellar content has been explored \\citep{jorgensen-cores,hatchell-cores}.  In addition, Perseus shows a wide range of star forming environments, ranging from the newly formed clusters IC 348 and NGC 1333 to more isolated star forming regions such as B5 and L1448.  A substantial portion of the molecular mass in the cloud is not currently forming stars.  Several previous observational studies provide context for the observations of Perseus.  The observational results are homogenized in \\citet{jijina}.  The typical cores in Perseus have $\\log(\\mathrm{NH}_3/\\mbox{cm}^{-2})=14.5$, $R_{pc}=0.09~\\mbox{ pc}$ (after rescaling to our preferred distance of 260 pc), velocity dispersion $\\sigma_v=0.17\\mbox{ km s}^{-1}$, $T_{k}=11~\\mbox{K}$. However, these studies have primarily observed the well known star forming regions with less concern for objects in the sterile portions of the molecular cloud.  This survey presents observations of NH$_3$ and C$_2$S emission from a variety of sources in Perseus including millimeter-bright dense cores as well as otherwise unremarkable high column density features selected from far infrared emission.  The two tracers present complimentary views of the chemical evolution of the cloud.  C$_2$S is regarded as an ``early-time'' tracer formed in the initial conversion of atomic to molecular gas and excited at high densities \\citep[$n_{cr}\\sim 10^{4.5}\\mbox{ cm}^{-3}$][]{langer95,pp5-difran}. As carbon species are depleted, C$_2$S disappears.  In contrast, NH$_3$ is regarded as a late-time tracer like N$_2\\mbox{H}^+$, with both species requiring the relatively slow formation of N$_2$ as a precursor.  The molecules are excited at similar densities as C$_2$S but should not appear in significant amounts until later in the protostellar collapse \\citep[$t\\sim 10^{5.5}\\mbox{ yr}$][]{flower06,pp5-difran}.  \\citet{tafalla04} note that the abundance of NH$_3$ varies by a factor of several across starless cores while N$_2\\mbox{H}^+$ remains constant.  Such variations can complicate using ammonia as a structural tracer, but the utility of having a direct temperature measurement from the NH$_3$ is immense.  Our survey of dense cores in Perseus adopted the form of a spectral survey to maximize the number of cores we could sample in a limited amount of time.  Even with a single spectrum, we are able to determine core kinematics, velocity dispersion, kinetic temperatures, and chemical abundances of ammonia and C$_2$S.  We also spent a significant amount of time surveying core candidates derived through a variety of methods to find ammonia emission from objects not bright in the submillimeter.  In this paper, we present the results of our survey and derive physical parameters from the ammonia spectra.  A detailed comparison of the core properties to other tracers will be presented elsewhere. ",
        "conclusions": "Summary of Source Origins} \\tablewidth{0pt} \\tablehead{ \\colhead{Origin} & \\colhead{Abbrev.} & \\colhead{Number} & \\colhead{NH$_3$(1,1) Det. Frac.} & \\colhead{NH$_3$(2,2) Det. Frac.} & \\colhead{C$_2$S Det. Frac.} } \\startdata Bolocam & B & 115 & 98\\% & 85\\% & 67\\% \\\\ SCUBA & S & 16 & 100\\% & 94\\% & 56\\% \\\\ Literature & L & 5 & 100\\% & 80\\% & 100\\% \\\\ Dust & D & 38 & 61\\% & 8\\% & 13\\% \\\\ Weak Submm & W & 19 & 26\\% & 11\\% & 16\\% \\\\ \\hline Total & \\nodata & 193 & 84\\% & 63\\% & 51\\% \\enddata \\end{deluxetable*}  \\begin{figure*} \\plotone{f1.eps} \\caption{\\label{findchart1} Locations of GBT pointings (red points) overlaid upon a map of extinction for the Perseus region.  Seven subregions have been defined and have their sources labeled in Figure \\ref{findchart2}.  For sources outside the defined subregions, the sources are labeled on the map with the number of the source given in Table \\ref{obstable}.  The extinction map is derived from applying the NICER algorithm to 2MASS data \\citep{complete-data} and the grayscale has a square-root transfer function running from $A_V=0$ to 7 mag.} \\end{figure*}  \\begin{figure*} \\plotone{f2.eps} \\caption{\\label{findchart2} Locations of GBT pointings (red points) overlaid upon a maps of millimeter emission from the BOLOCAM survey of the region \\citep{bolocam-perseus}. The seven maps correspond to the subregions defined in Figure \\ref{findchart1}.  Individual pointings are labeled with the number of the source given in Table \\ref{obstable}.  The grayscale has a square-root transfer function spanning $f_\\nu=0$ to 400 mJy beam$^{-1}$.} \\end{figure*}    The C$_2 ^{34}$S line was not detected along any of the lines of sight.  Given the ISM isotopic ratio $^{32}$S/$^{34}$S~$\\sim 22$ \\citep{ism-abund}, the lack of any detections, particularly when the C$_2$S line is strong suggests that the main line is usually optically thin.  For NH3SRC 42, we establish a lower limit on the line ratio of 11.7 (the maximum in our population) by setting the C$_2 ^{34}$S amplitude to the 3$\\sigma$ limit.  Assuming excitation conditions are the same for both isotopomers, this implies $\\tau_{\\mathrm{CCS}}<1.4$ for one of the brightest C$_2$S lines in our sample.  In Figure \\ref{sampleplot} we show three spectra from our sample. Source 47 was observed every night as a consistency check on our flux calibration and is thus the best observed ammonia source in our sample (integrated S/N of 530).  Source 89 shows a typical narrow line spectrum ($\\sigma_v=0.12\\mbox{ km s}^{-1}$) illustrating the resolution of the data. Source 31 is a multi-component spectrum which must be analyzed in more detail (see \\S\\ref{multicomp}). In Source 31, the multiple components are also visible in the NH$_3$(2,2) and C$_2$S lines.  \\begin{figure*} \\plotone{f3.eps} \\caption{\\label{sampleplot} Three sample spectra from the ammonia sample including the source used for a flux check (Source 47), a narrow-line spectrum (Source 89) and a multicomponent spectrum (Source 31).} \\end{figure*}   We present a summary of our observations in Table \\ref{obstable}.  The properties of the cross-referenced names of the submillimeter cores are given in \\citet{bolocam-perseus} and \\citet{perseus-scuba} respectively.  Since ammonia has several hyperfine components that are well-separated in velocity, the integrated intensities reported for the (1,1) and the (2,2) lines are the sum of the integrated intensities over all channels that are within 3~km~s$^{-1}$ of any hyperfine component.  We subtract the mean intensity in the off-line channels between the hyperfine components from the intensity of each channel in the on-line windows to offset any low-lying baseline residuals.  When there is no ammonia emission or C$_2$S from which the line velocity can be determined, the main component of the ammonia line is assumed to lie at the the mean $^{13}$CO velocity along the line of sight, as derived from the COMPLETE $^{13}$CO data.      We have searched for NH$_3$(1,1), NH$_3$(1,1), C$_2$S($2_1\\to 1_0$) emission along 193 lines of sight towards the Perseus molecular cloud.  The lines of sight were selected based on positions that were detected in (sub)millimeter emission or had large dust column densities implied by far infrared (FIR) emission.  We detect ammonia emission along 162 (84\\%) of the lines of sight and C$_2$S along 96 (51\\%) of the lines of sight.  We estimate the physical properties of the gas by fitting a model emission profile to all spectral lines simultaneously. The emission is modeled as a uniform slab of gas that completely fills the beam, has a Gaussian intrinsic line width, and a single excitation temperature for all lines.  Where appropriate, we refined the model to account for low optical depths, incomplete coupling to the GBT beam and multiple velocity components along the line of sight.  Nearly all (98\\%) bright, (sub)millimeter cores have strong ammonia emission associated with them and the exceptions appear to be artifacts in the submillimeter map based on examining the original BOLOCAM data.  In addition, we detected emission towards 23 sources selected based on FIR emission that implies large dust column densities and low temperatures.  Twenty-one objects are not seen in the (sub)millimeter, suggesting that the submillimeter emission is not a perfect tracer of the dense gas (the remaining two sources are outside the bounds of the continuum surveys).  However, the FIR-based ammonia detections have lower line intensities than (sub)millimeter-bright source, as well as lower optical depths and larger line widths.  It remains to be shown whether this could be an evolutionary effect or whether the (sub)millimeter-weak sources simply trace isolated pockets of gas not associated with the dense cores traced by the dust continuum.  We find that the ammonia implies dense gas temperatures in Perseus are predominantly cold ($T_k\\sim 11~$K).  Ammonia column densities are typical for cores presented in the literature \\citep[$N_{\\mathrm{NH3}}\\sim 10^{14.5}\\mbox{ cm}^{-2}$,][]{jijina} though we also find several lines-of-sight with very low ammonia column densities ($N_{\\mathrm{NH3}}\\lesssim 10^{13.5}\\mbox{ cm}^{-2}$) associated with the IC 348 region.  Forthcoming work will examine the properties of these objects in more detail including comparison with the (sub)millimeter emission, protostellar content, and the velocity structure of the dense core population."
    },
    "0711/0711.0188_arXiv.txt": {
        "abstract": "LISA is a planned space-based gravitational-wave (GW) detector that would be sensitive to waves from low-frequency sources, in the band of roughly $(0.03 - 0.1)\\,{\\rm mHz} \\lesssim f \\lesssim 0.1\\,{\\rm Hz}$. This is expected to be an extremely rich chunk of the GW spectrum --- observing these waves will provide a unique view of dynamical processes in astrophysics.  Here we give a quick survey of some key LISA sources and what GWs can uniquely teach us about these sources. Particularly noteworthy science which is highlighted here is the potential for LISA to track the moderate to high redshift evolution of black hole masses and spins through the measurement of GWs generated from massive black hole binaries (which in turn form by the merger of galaxies and protogalaxies).  Measurement of these binary black hole waves has the potential to determine the masses and spins of the constituent black holes with percent-level accuracy or better, providing a unique high-precision probe of an aspect of early structure growth.  This article is based on the ``Astrophysics and Relativity using LISA'' talk given by the author at the Seventh Edoardo Amaldi Conference on Gravitational Waves; it is largely an updating of the author's writeup of a talk given at the Sixth International LISA Symposium {\\cite{me_lisa6}}. ",
        "introduction": "Stochastic GWs are ``random'' waves, arising from the superposition of many discrete, uncorrelated sources.  They typically combine to produce a broadband spectrum that is nearly flat, or weakly peaked near some fiducial frequency.  Some of the most anticipated stochastic sources are cosmological in origin.  By far the most eagerly sought waves are GWs arising from primordial fluctuations in the universe's spacetime metric which have been parametrically amplified by inflation.  In energy density units (commonly used to parameterize these waves; see {\\cite{allen96}} for a very readable discussion), these waves are flat across an extremely wide band, ranging from frequencies $f \\sim 10^{-16}\\,{\\rm Hz}$ to roughly $10^{10}\\,{\\rm Hz}$.  The lower end of this scale corresponds to the inverse Hubble scale when the universe became matter dominated; below this frequency, the spectrum grows as $f^{-2}$.  At the high end, the spectrum is cut off by the (short but finite) time scale over which inflation ends and the universe enters a hot, radiation-dominated phase; see Fig.\\ 12 of {\\cite{allen96}} and associated discussion.  What's particularly exciting about these waves is that the level of this spectrum is set by --- and thus encodes --- the potential which drives inflation.  Thus, measuring these waves would {\\it directly probe inflationary physics}.  Unfortunately, these waves have such small amplitudes that direct detection with LISA (or any other near term GW experiment) is completely out of reach.  LISA should be able to detect a background at a level $\\sim 10^{-11}$ of the universe's closure density; current estimates suggest that this is at least 4 or 5 orders of magnitude from the sensitivity needed to measure these waves.  Direct detection at the relatively high frequencies of man-made facilities requires finding a relatively quiet band (without too many ``foreground'' sources) and a much more sensitive instrument than LISA.  At the lowest frequencies, detection is plausible in the next decade or so through cosmic microwave background (CMB) studies.  The tensor nature of GWs has a unique impact on the polarization of CMB photons {\\cite{zs97,kks97}}.  The signal is weak compared to other signals in the CMB, and potentially contaminated by foreground effects. Searching for it is nonetheless one of the most important directions in CMB physics today.  Although inflationary waves are out of reach, other cosmological backgrounds may not be.  Backgrounds are associated with phase transitions that occur as the universe expanded and cooled.  For example, when the mean temperature was roughly 1 TeV, the electroweak interaction separated into electromagnetic and weak interactions. This separation was not spatially homogeneous --- some regions underwent the transition before others.  Different regions were at different densities; boundaries between regions collided and coalesced, producing a background {\\cite{allen96,kmk02}}.  More recently, there has been a flurry of work in superstring theory suggesting that cosmic strings may have been produced in the early universe and then expanded to cosmic sizes (see, for example {\\cite{p05}} for an overview).  A network of such objects could produce a background accessible to LISA; Hogan summarizes such backgrounds in a previous brief review {\\cite{craig}}.  Though a relatively speculative source, cosmological stochastic backgrounds are easily searched for, and there is a tremendous amount of discovery space.  The probability of payoff may be low, but the potential reward for discovery is extraordinarily high. ",
        "conclusions": "LISA is an astrophysics mission that will make possible unprecented {\\it precision} measurements of dynamic, strong-gravity phenomena. Perhaps most exciting will be the ability to track the mergers of massive black holes, precisely weighing their masses and spins.  Such measurements will directly probe the early growth of black holes, opening a new window onto structure growth."
    },
    "0711/0711.0970_arXiv.txt": {
        "abstract": "I review the subjects of non-solar cosmic rays (CRs) and long-duration gamma-ray bursts (GRBs). Of the various interpretations of these phenomena, the one best supported by the data is the following. Accreting compact objects, such as black holes, are seen to emit relativistic puffs of plasma: `cannonballs' (CBs).  The inner domain of a rotating star whose core has collapsed resembles such an accreting system. This suggests that core-collapse supernovae (SNe) emit CBs, as SN1987A did.  The fate of a CB as it exits a SN and travels in space can be studied as a function of the CB's mass and  energy, and of `ambient' properties: the encountered matter- and light- distributions, the composition of the former, and the location of intelligent observers. The latter may conclude that the interactions of CBs with ambient matter and light generate CRs and GRBs, all of whose properties can be described by this  `CB model' with few parameters and simple physics. GRB data are still being taken in unscrutinized  domains of energy and timing. They agree accurately with the model's predictions. CR data are centenary. Their precision will improve, but new striking predictions are unlikely. Yet, a one-free-parameter description of all CR data works very well. This is a bit as if one discovered QED today and only needed to fit $\\alpha$. ",
        "introduction": "This is a version of an introductory talk to high-energy physicists. Cosmic rays (CRs) were the first item in their field, and will remain the energy record-breakers for the foreseeable future. I shall argue that nothing `besides the standard model' is required to understand CRs of any energy, subtracting from their interest. `Long' gamma-ray bursts (GRBs) are flashes of mainly sub-MeV photons, originating in supernova (SN) explosions. The $\\gamma$-rays are highly collimated. Hence, GRBs are not the publicized `highest-energy explosions after the big bang', but more modest torches occasionally pointing to the observer. GRBs are of interest because their understanding is intimately related to that of CRs. It might have been more precise to say `my understanding' of GRBs and CRs, for the work of my coauthors and I is viewed as unorthodox.  What a start! I have already admitted that our stance is not trendy and that the subject is of no post-standard  interest. But our claims are based on clear hypothesis, which may be proven wrong, and very basic physics, which is precise enough, very pretty, understandable to undergraduates, and successful.  The information about GRBs and CRs is overwhelming. GRBs are known since the late 60's and CRs since 1912. Surprisingly, no theories have arisen that are both accepted (`standard') and acceptable (transparent, predictive and successful). I cannot refer to a representative subset of the $\\sim\\!70\\!+\\!70$ kilo-papers on CRs and GRBs. For reviews of the standard views on CRs, see e.g.~Hillas\\cite{Hillas} or Hoerandel\\cite{Hoer2}. For the accepted `fireball' model of GRBs see e.g.~Meszaros\\cite{Mes} or Piran\\cite{Pir}. Fewer self-citations and many more references, particularly to data, appear in DDD02\\cite{AGoptical}, DDD03\\cite{AGradio}, DD04\\cite{DDGRB}and DD06\\cite{DDCR}. ",
        "conclusions": "We do not have a solid understanding of accretion onto black holes or neutron stars. But such processes are observed to result in the ejection of relativistic and highly collimated jets. We assumed that a similar process takes place as a stellar core collapses, leading to a supernova event. We posited that the SN's relativistic ejecta --two jets of cannonballs-- are the sources of GRBs and CRs. The association between SNe and (long) GRBs is now established. We argued that the electrons in a CB, by inverse Compton scattering on the illuminated surroundings of the exploding star, generate the  $\\gamma$ rays, X-rays, UV and optical light of the `prompt' phase of a GRB. The ensuing results are the simplest and most predictive, they are a firm `theory'. In this paper I have, however, followed the historical development, in which the CB parameters were first extracted from the observations of the afterglow of GRBs. This involves a `model', a set of arguable but simple hypothesis leading to the prediction of the properties of the AG --dominated by synchrotron radiation by the ISM electrons that a CB intercepts-- as a function of frequency and time. In the historical order the `prompt' results for GRBs are predictions of the theory. Some results for Cosmic Rays --the ISM particles that CBs scatter in their journey-- are also `theory', others can be viewed as further tests of the `model'.  The results for GRB afterglows may be based on a simplified model, but they work with no exception all the way from radio to X-rays (DDD02,03). In particular, they describe correctly GRB 980425, located at a redshift two orders of magnitude closer than average. Its associated SN is the one we `transported' to conclude  --thanks to the reliability of our AG model-- that core-collapse SNe generate long GRBs (DDD02). The X-ray light curve of GRB 980425 and a few others, with extremely scarce data, was described with the `canonical' properties later observed in detail in many SWIFT-era GRBs. It is not recognized that the two CBs of GRB 030329 were seen, or that their separation in the sky was the predicted `hyper-luminal' one. In view of the overall success of the CB model, this is a durable hurdle: GRBs so close and luminous are very rare.  The accuracy of the predictions for the prompt phase of GRBs amazes even the CB-model's proponents. The typical values and the correlations between the $\\gamma$-ray prompt observables leave little doubt that the production mechanism is inverse Compton scattering on `ambient' light of $\\sim\\!1$ eV energy. The approximate scaling law $E\\,dN_\\gamma/dE dt\\!\\propto\\!F(E\\,t^2)$ --spectacularly  confirmed in the case of XRF 060218-- demonstrates that the light is that of a `glory': the early SN light scattered by the `windy' pre-SN ejecta. A GRB spectrum that works even better\\cite{HessLady} than the phenomenological `Band' expression is also predicted. The flux and its spectral evolution during the prompt and rapid-decline phases are the expected ones, as we tested in minute detail with SWIFT data.  In the internal-external fireball model of GRBs, highly relativistic thin conical shells of $e^+e^-$ pairs, sprinkled with a finely tuned baryon `load', collide with each other generating a shock that accelerates their constituents and creates magnetic fields. Each collision of two shells produces a GRB pulse by synchrotron radiation. The ensemble of shells collides with the ISM to produce the AG by the same mechanism. The energy available to produce the GRB pulse --as two shells moving in the {\\it same} direction collide-- is more than one order of magnitude smaller than that of the merged shells as they collide with the ISM at rest. The ratio of observed GRB and AG energies is more than one order of magnitude, but in the opposite direction. This `energy crisis' in the comparison of bolometric prompt and AG fluences\\cite{Pir} is not resolved. Moreover, the GRB spectrum cannot be accommodated on grounds of synchrotron radiation\\cite{GCL}, the `standard' prompt mechanism. The SWIFT-era observations also pose decisive problems to the standard model, whose microphysics\\cite{Petal}, reliance on shocks\\cite{Ketal} and correlations based on the jet-opening angle\\cite{Setal} have to be abandoned, according to the cited authors.  In spite of the above, the defenders of the fireball model are not discouraged. Their attitude towards the CB model, whose observational support is so remarkable, is not equally supportive. This may be due to cultural differences. Particle physicists believe that complex phenomena may have particularly simple explanations. They thrive on challenging their standard views. Doubting or abandoning a previous consensus in astrophysics is less easy.  The CB-model description of Cosmic Rays is also simplistic: there is only one source of (non-solar) CRs at all energies, and only one parameter to be fit. The model has a certain inevitability: if CBs with the properties deduced from GRB physics are a reality, what do they do as they scatter the particles of the interstellar medium? We have argued that they transmogrify them into CRs with all of their observed properties. The mechanism is entirely analogous to the ICS responsible for a GRB's prompt radiation. Suffice it to substitute the CB's electron, plus the ambient photon, by a moving-CB's inner magnetic field, plus an ambient nucleus or electron.  After a century of CR measurements, the CB-model results lack the glamour of predictions. Yet, the expectations for the knee energies, and for the relative abundances of CRs, are `kinematical', simple, and  verified. They constitute evidence, in my opinion, that the underlying model is basically correct. The prediction for the shape of the spectra: the low energy hips, the large energy stretch very well described by a power-law of (source) index $\\beta_s\\!=\\!13/6$, and the steepening at the knees, are also verified. The index $\\beta_s$ is measured well enough for the adequacy of the prediction  to be sensitive to the details of the underlying model, such as the form of the function $R(\\gamma)$ in Eq.~(\\ref{Radius}). I cannot claim that the fact that the prediction is right on the mark is much more than a consistency test, for the physics underlying this aspect of the problem may be terrifyingly complex. The same CR source --cannonballs from supernovae, this time extragalactic-- satisfactorily describes the CR data above the ankle. Finally, the properties of CR electrons, and of the high-latitude Gamma `Background' Radiation, are also correctly reproduced.   Most CR scholars agree with the `standard' paradigm that the flux well below the knee is produced by the acceleration of the ISM in the frontal shocks of the nonrelativistic ejecta of SNe. In spite of recent observations of  large magnetic fields\\cite{Yasu} in collisions of SN shells and molecular clouds, nobody has been able to argue convincingly that this process can accelerate particles up to the (modest) energy of the knee, {\\it and} to show that the number and efficiency of the putative sources suffices to generate the observed CR luminosity (to my satisfaction, I  add, to make these statements indisputable). From this point on, there is no `standard' consensus on the origin of CRs, e.g.,~of the highest-energy ones. In this sense, the CB model is regarded as yet another model, which it is. After all, we are only saying that CRs are accelerated by the jetted relativistic ejecta of SNe, as opposed to the quasi-spherical, non-relativistic ones. Yet, the CB model is also rejected by the CR experts, sometimes even in print\\cite{Hillas}, though it survives the critique\\cite{HResponse}. But, concerning CRs, the model does not trigger the same indignant wrath as in the GRB realm.  I have shown that the problem of GRBs is convincingly --i.e.~predictively-- solved and that, on the same simple basis, all properties of CRs can be easily derived. Only an overwhelmed minority recognizes these facts, in contradiction with Popper's and Ockham's teachings. I would conclude with a dictum attributed to Lev Landau: {\\it `In astrophysics, theories never die, only people do.'}  \\vskip .4cm \\noindent {\\bf Acknowledgement}  \\noindent I thank Shlomo Dado, Arnon Dar and Rainer Plaga for our collaboration."
    },
    "0711/0711.1093_arXiv.txt": {
        "abstract": "{We report on the results of an IRAM-30m search for CO emission lines in three galaxies at intermediate redshifts. The idea was to investigate the molecular content of galaxies bright in the infrared at $z=0.4-1.5$, a redshift desert for molecular line studies, poorly investigated as of yet. We integrated 8-10h per source and did not succeed in detecting any of the sources. {From} our upper limits, we are able to constrain the molecular gas content in these systems to less than $4$ to $8\\times 10^9 M_\\odot$, assuming a CO-to-H$_2$ conversion factor ($\\alpha=0.8$~M$_\\odot$/(K\\,km\\,s$^{-1}$\\,pc$^{2}$)). We stress the current difficulty of selecting sources with a detectable molecular content, a problem that will be faced by the ALMA First Science projects.} ",
        "introduction": "Our current knowledge of the molecular content of galaxies at $z>0.4$ is currently limited to a fraction of the submillimetre bright objects \\citep{Solomon:2005,Greve:2005}, usually selected on the basis of their strong infrared (IR) luminosity. The extent to which active galactic nuclei (AGN) contribute to these extreme infrared luminosity is currently {a} matter of debate. While the interpretation of the IR-radio correlation in terms of on-going star formation is commonly accepted \\citep{Condon:1991,Condon:1992,Yun:2001}, this explanation remains uncertain \\citep[e.g.][]{Vlahakis:2007}.  The evidence that, in the strongest infrared sources, a significant fraction of the infrared luminosity is due to AGN \\citep{Genzel:2000,Alexander:2005} further complicates the interpretation of this correlation. \\cite{Farrah:2003} have shown a correlation of the AGN and starburst luminosities over a wide range of IR luminosities.  The recent detection of a molecular torus in Arp 220 \\citep{Downes:2007} demonstrates that the true source of at least part of its infrared luminosity is due to a black hole accretion disc, while this galaxy was considered as a prototypical starburst. Hence, there is the possibility that the infrared luminosity is not a good tracer of star formation activity, which is quite troublesome because this is nevertheless the most reliable and easy tracer used so far (at least unbiased by dust extinction).  In this complex context, the detection of molecular gas emission is essential to get information about the mass and dynamics of these galaxies, and to confirm the huge star formation rates (750-1000\\,M$_\\odot$\\,yr$^{-1}$) usually derived for (sub)millimetre galaxies. Ultimately, this will contribute to further constrain the scenario of hierarchical galaxy formation and evolution. However, the current sensitivities are relatively low and require configurations with huge amount of gas.  Different types of samples have been investigated so far, with various success. On the one hand, \\citet{Evans:2006}, in their study of molecular gas in quasi-stellar objects $z<0.15$, found a low CO to infrared luminosity ratio, suggesting that the infrared luminosities of these QSO might be dominated by the AGN component. \\citet{Saripalli:2007} failed to detect any molecular gas in restarting radio galaxies at $z<0.15$. On the other hand, for the high-z galaxies with a CO detection, starbursts and AGN do not exhibit significant differences in their molecular gas content \\citep{Solomon:2005,Greve:2005}. However, there is a clear difference in the CO line widths, a factor of 2.3 narrower in velocity in QSO host galaxies with respect to submillimetre galaxies, while powerful radio galaxies fall in between \\citep{Greve:2005,Carilli:2006}. The origin of this effect is not yet clear: it could be due to a systematic in the orientation of the QSO, but the possibility of merger signatures in submillimetre galaxies is not excluded.  The redshift range $z=0.4-1.5$ corresponds to a key period of the star formation history of the Universe: the end of the Star Formation plateau \\citep[e.g.][]{Madau:1998,Dahlen:2007} converging to the nearby galaxy population. This range has been poorly investigated in CO so far due to the lack of appropriate detectors and is known as a redshift desert for molecular studies. This desert is due to the presence of atmospheric lines (O$_2$), which prevent the detection of CO(1-0) below 81\\,GHz. (A similar situation was observed in optical spectroscopy for galaxies with ${z=1.5-2.5}$ \\citep{Steidel:2004}.) However, while most $z>1.5$ sources detected in CO are magnified, {the} lower redshift range {is} relatively more favourable to CO detection because there is no negative K-correction for CO \\citep{Combes:1999}.  In this paper, we discuss the search for CO emission lines in three galaxies at intermediate redshifts, in order to investigate the molecular gas content of galaxies bright in the infrared at $z=0.4-1.5$.  Throughout this paper, we adopt a flat cosmology, with $\\Omega_\\mathrm{m}=0.24$, $\\Omega_\\Lambda=0.76$ and $H_0=73$~km\\,s$^{-1}$\\,Mpc$^{-1}$ \\citep{Spergel:2007}.   \\begin{table*} \\centering \\caption{\\protect Characteristics of the galaxies studied in this paper. Listed from left to right are the names of the sources, their J2000.0 positions, infrared luminosities, rest-frame B-band luminosities, galaxy types, redshifts, and luminosity distances. The appropriate references for the listed infrared luminosities and redshifts are given in the rightmost column. The values of L$^\\mathrm{rest}_B$ are from \\citet{Weiner:2005} for the CFRS sources, and computed with the 2MASS J-band flux for NGP9\\,F268-0341339. } \\begin{tabular}{lllllllll} \\hline Source & RA (J2000) & DEC (J2000) & L$_{IR}$  ($10^{11}$\\,L$_{\\sun}$)& L$^\\mathrm{rest}_{B}$ ($10^{11}$\\,L$_{\\sun}$) $[$ L$_*^B$ $]^a$ & Type & Redshift & $D_L$ (Gpc)& Ref. \\\\ \\hline \\object{CFRS\\,14.1329} & 14:17:34.8 & +52:27:52.0     & $1.3\\pm 0.2$$^b$ & $0.04$ $[2]$& LIRG  & 0.375 &1.96  & 1, 2 \\\\ \\object{CFRS\\,14.1157} & 14:17:41.9 & +52:28:24.0    & $67\\pm 4$ &$0.53$  $[25]$ & ULIRG & 1.149 & 7.84 &3, 2 \\\\ \\object{NGP9\\,F268-0341339} & 12:28:47.4 & +37:06:12.3 & $ -- $ & $6.9$ $[328]$ & QSO, synch. & 1.515 & 11.12 & 4\\\\ \\hline \\end{tabular} \\begin{minipage}{160mm} 1: \\cite{Zheng:2004}; 2: \\cite{Hammer:1995}; 3: LeFloch, private communication; 4:  \\cite{Snellen:2002}. \\\\ $^a$ $ L_*^B = 2.1\\times 10^9 L_{\\sun}$ \\citep{Marzke:1998}. $^b$ \\citet{Flores:1999b} initially estimated L$_{IR} = $ 6.3$\\pm$0.3$\\times 10^{11} L_{\\sun}$. \\end{minipage} \\label{tab:data} \\end{table*} ",
        "conclusions": "\\label{sect:dis} The non-CO detections discussed here reveal the difficulties of selecting favourable candidates. The scatter of the correlation between the IR and CO luminosities (see Fig. \\ref{fig:plotgao04}) is large and probably linked to the fact that the IR luminosities does not trace the star formation activity only. In addition, it is difficult to {obtain} IR luminosities, given the lack of all-sky survey since IRAS. {The only solution has been here to rely on IR luminosities derived from shorter wavelength measurements \\citep[namely 15$\\mu$m][]{Flores:1999b,Zheng:2004}.}  {The lack of CO detections towards \\object{CFRS\\,14.1329} and \\object{NGP9~F268-0341339} is probably due to the fact that they are not as good candidates as we initially thought.} As explained in Appendix \\ref{sec:cfrs1}, \\object{CFRS\\,14.1329} had been selected on the basis of a favourable $L_{IR}$ luminosity \\citep{Flores:1999b}, which was subsequently revised \\citep{Zheng:2004}. We also discussed in Appendix \\ref{app:ngp9} that \\object{NGP9~F268-0341339} is behind a foreground galaxy at $z=0.138$ with SFR$\\sim 95\\,M_{\\sun}$\\,yr$^{-1}$.  {From our upper limits on the CO(1-0) and CO(3-2) luminosities, we estimate that the molecular gas content\\footnote{Following \\citet{Solomon:2005}, our molecular gas content estimates account for the He mass.} of \\object{CFRS\\,14.1329} is smaller than $5.5\\times 10^9 M_{\\sun}$ and $4.2\\times 10^9 M_{\\sun}$, assuming a molecular gas mass to CO luminosity ratio $\\alpha=0.8\\,M_{\\sun}\\,($K\\,km\\,s$^{-1}$pc$^2)^{-1}$, typical of {ultraluminous infrared galaxies (ULIRG)} \\citep[e.g.][]{Solomon:2005}. These total molecular gas masses are consistent with the gas content of local normal spiral galaxies \\citep[]{Gao:2004}, while the infrared luminosity (derived from 15\\,$\\mu$m data) was suggesting a luminous infrared galaxy (LIRG). Hence, the value of $\\alpha$ is uncertain and probably underestimated. For the more distant galaxy \\object{NGP9~F268-0341339}, the limits are less stringent and it is still compatible with an (U)LIRG. With the same assumptions, the molecular mass is smaller than $33\\times 10^9 M_{\\sun}$ and $8.6\\times 10^9 M_{\\sun}$ according to our upper limits on the CO(2-1) and CO(5-4) luminosities.  The luminosities derived from higher CO line transitions are more stringent than those derived from the lower ones as we assumed that the lines were optically thick (see also Sect.  \\ref{sec:upper}).}  In contrast, \\object{CFRS\\,14.1157} was a very favourable candidate that has remained as such. \\citet{Zheng:2004} estimated $L_{IR}=178.6 \\times 10^{11}\\, L_{\\sun}$ (SFR$\\sim 3054$\\,M$_{\\sun}$\\,yr$^{-1}$), relying on the IR-15$\\mu$m luminosities correlation measured by \\citet{Elbaz:2002}. \\citet{LeFloch:2007} published a panchromatic spectral energy distribution of this galaxy with a very good wavelength coverage. They derived a direct estimate\\footnote{Please note that \\citet{LeFloch:2007} worked with the 1-1000 $\\mu$m range. In this paper, we adapted their figures to the 8-1000$\\mu$m range used here, thanks to the corresponding value kindly provided by E. Le Floc'h.} of the infrared luminosity computed over the range 8-1000$\\mu$m : $L_{IR}=67\\pm 4\\,\\times 10^{11} L_{\\sun}$, which would correspond to SFR$\\sim 1150$\\,M$_{\\sun}$\\,yr$^{-1}$. However, \\citet{LeFloch:2007} have estimated that 67.2$\\%$ percent of the infrared flux is due to the AGN component. Accordingly, the SFR is probably of order 375\\,M$_{\\sun}$\\,yr$^{-1}$.  {From our upper limits on the CO(2-1) and CO(4-3) luminosities of \\object{CFRS\\,14.1157}, we find that the molecular gas content of this galaxy is smaller than $15\\times 10^9 M_{\\sun}$ and $7.5\\times 10^9 M_{\\sun}$} (with $\\alpha=0.8\\,M_{\\sun}\\,($K\\,km\\,s$^{-1}$pc$^2)^{-1}$).  This supports the idea that the observed infrared luminosity is probably dominated by the AGN component, as estimated by \\citet{LeFloch:2007}. In addition, one can note that the SFR derived on the basis of the $[$\\ion{O}{II}$]$ line is significantly lower: $2M_{\\sun}$\\,yr$^{-1}$ relying on the equivalent width and rest-frame absolute magnitude measured by \\citet{Weiner:2005} and the formula of \\citet{Guzman:1997}, while we estimate $3.5M_{\\sun}$\\,yr$^{-1}$ with an integration of $[$\\ion{O}{II}$]$ the spectra published by \\citet{LeFloch:2007} (see their figure 2a) and normalised to SDSS fluxes. This could be compatible with the IR-derived SFR only with a factor of extinction of $105-190$, while $L_{IR}/L_B=41$ if one assumes that only one third of the infrared luminosity contributes to the SFR. {While a large scatter is known to affect $[$\\ion{O}{II}$]$ luminosities as discussed by \\citet{Weiner:2007}, the previous comparison stresses the importance of the actual fraction of the infrared luminosity due to the starburst: this fraction is often overestimated.}"
    },
    "0711/0711.3575_arXiv.txt": {
        "abstract": "We present our numerical comparisons between the BSSN formulation widely used in numerical relativity today and its adjusted versions using constraints. We performed three testbeds: gauge-wave, linear wave, and Gowdy-wave tests, proposed by the Mexico workshop on the formulation problem of the Einstein equations. We tried three kinds of adjustments, which were previously proposed from the analysis of the constraint propagation equations, and investigated how they improve the accuracy and stability of evolutions. We observed that the signature of the proposed Lagrange multipliers are always right and the adjustments improve the convergence and stability of the simulations.  When the original BSSN system already shows satisfactory good evolutions (e.g., linear wave test), the adjusted versions also coincide with those evolutions; while in some cases (e.g., gauge-wave or Gowdy-wave tests) the simulations using the adjusted systems last 10 times as long as those using the original BSSN equations. Our demonstrations imply a potential to construct a robust evolution system against constraint violations even in highly dynamical situations. ",
        "introduction": "Numerical integration of the Einstein equations is the only way to investigate highly dynamical and nonlinear gravitational space-time. The detection of gravitational wave requires templates of waveform, among them mergers of compact objects are the most plausible astrophysical sources.  Numerical relativity has been developed with this purpose over decades.  For neutron star (NS) binaries, a number of scientific numerical simulations have been done so far, and we are now at the level of discussing the actual physics of the phenomena, including the effects of the equations of state, hydrodynamics, and general relativity by evolving various initial data~ \\cite{Shibata:2003ga,Shibata:2006nm,Marronetti:2003gk,Marronetti:2003hx,Faber:2006qc}. Mergers of black holes (BHs) are also available after the breakthrough by Pretorius~\\cite{Pretorius:2005gq} in 2004. Pretorius's implementation had many novel features in his code; among them he discretizes the four-dimensional Einstein equations directly, which is not a conventional approach so far. However, after the announcements of successful binary BH mergers by Campanelli et al.~\\cite{Campanelli:2005dd} and Baker et al.~\\cite{Baker2006} based on the standard 3+1 decomposition of the Einstein equations, many groups began producing interesting results~\\cite{Diener2006,Herrmann2006,Baker:2006yw,Gonzalez:2006md,Etienne:2007hr,Tichy:2007hk,Campanelli:2007cg, Gonzalez:2007hi,Campanelli:2007ew,Thornburg:2007hu}. The merger of NS-BH binary simulations has also been reported recently, e.g. ~\\cite{Shibata:2006bs}.   Almost all the groups which apply the above conventional approach use the so-called BSSN variables together with ``$1+\\log$''-type slicing conditions for the lapse function and ``$\\Gamma$-driver'' type slicing conditions for the shift function.  BSSN stands for Baumgarte-Shapiro~\\cite{BS} and Shibata-Nakamura~\\cite{SN}, the modified Arnowitt-Deser-Misner formulation initially proposed by Nakamura~\\cite{Nakamura}. (The details are described in \\S \\ref{subsec:BSSN}.) There have already been several efforts to explain why the combination of this recipe works from the point of view of the well-posedness of the partial differential equations (e.g.~\\cite{Beyer2004,Gundlach2006}). However, the question remains whether there exists an alternative evolution system that enables more long-term stable and accurate simulations. The search for a better set of equations for numerical integrations is called the formulation problem for numerical relativity, of which earlier stages are reviewed by one of the authors~\\cite{novabook}.  In this article, we report our numerical tests of modified versions of the BSSN system, the {\\it adjusted BSSN systems}, proposed by Yoneda and Shinkai~\\cite{Yoneda:2002kg}. The idea of their modifications is to add constraints to the evolution equations like Lagrange multipliers and to construct a robust evolution system which evolves to the constraint surface as the attractor. Their proposals are based on the eigenvalue analysis of the constraint propagation equations (the evolution equations of the constraints) on the perturbed metric.  For the ADM formulation,  they explain why the standard ADM does not work for long-term simulations by showing the existence of the constraint violating mode in perturbed Schwarzschild space-time \\cite{Shinkai2002}. For the BSSN formulation, they analyzed the eigenvalues of the constraint propagation equations only on flat space-time ~\\cite{Yoneda:2002kg}, but one of their proposed adjustments was immediately tested by Yo et al.~\\cite{YoBaumgarteShapiro} for the numerical evolution of Kerr-Schild space-time and confirmed to work as expected. (The details are described in \\S \\ref{subsec:adjBSSN}.)  Our numerical examples are taken from the proposed problems for testing the formulations of the Mexico Numerical Relativity Workshop 2001 participants~\\cite{Alcubierre:2003pc}, which are sometimes called the Apples-with-Apples test. To concentrate the comparisons on the formulation problem, the templated problems are settled so as not to require technical complications; e.g., periodic boundary conditions are used and the slicing conditions do not require solving elliptical equations. Several groups already reported their code tests using these Apples tests (e.g. \\cite{Jansen:2003uh,Zlochower:2005bj,Boyl2007}), and we are also able to compare our results with theirs.  This article is organized as follows. We describe the BSSN equations and the {\\it adjusted} BSSN equations in Sec.~\\ref{subsec:BSSN} and \\ref{subsec:adjBSSN}. We give our three numerical test problems in Sec.~\\ref{sec:setup}. Comments on our coding stuff are in Sec.~\\ref{sec:code}. Sec.~\\ref{sec:num} is devoted to showing numerical results for each testbeds, and we summarize the results in Sec.~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary} In this article, we presented our numerical comparisons of the BSSN formulation and its adjusted versions using constraints. We performed three testbeds: gauge-wave, linear wave, and collapsing polarized Gowdy-wave tests with their evolutions by three kinds of adjustments, which were previously proposed by Yoneda and Shinkai~\\cite{Yoneda:2002kg} based on their constraint propagation analysis.  The idea of the adjusted systems is to construct a system robust against constraint violations by modifying the evolution equations using the constraint equations.   We can summarize our tests as follows: \\begin{itemize} \\item When the plain (original) BSSN evolutions already show satisfactory good evolutions (e.g., the linear wave test), the constraint violations (i.e., adjusted terms) are also small or ignorable. \\end{itemize} Therefore the {adjusted} BSSN equations become quite similar to the plain BSSN equations, and their results coincide with the plain BSSN results. \\begin{itemize} \\item Among the adjustments we tried, we observed that the adjusted BSSN system with the $\\tilde{A}$-eq. (\\ref{B1-adj}) is the most robust for all the testbeds examined in this study. It gives us an accurate and stable evolution compared to the plain BSSN system. Quantitatively, the life-time of the simulation becomes 10 times longer for the gauge-wave testbed and 5 times longer for the Gowdy-wave testbed than the life-time of the plain BSSN system. However, it should be noted that for the gauge-wave testbed, the convergence feature is lost at a comparatively early time, the 200 crossing-time in the Hamiltonian constraint and the 50 crossing-time in the momentum constraint. \\end{itemize} Recently, it has been claimed that the set up of the gauge wave problem in Apples-with-Apples has a problematic point~\\cite{Babiuc:2007vr}, which arises from the harmonic gauge condition. In \\cite{Babiuc:2004pi}, it is argued that this gauge has a residual freedom in the form $H \\to e^{{\\lambda}t}H$, where $\\lambda$ is an arbitrary and $H$ is a function in Eq.~(\\ref{eq:gwave-metric}). Of course, our set up corresponds to the $\\lambda=0$ case, but numerical error easily excites modes that result in either exponentially increasing or decaying metric amplitude. Actually, we find the amplitude of the error decays with time in this testbed. So, we conclude that due to the adjustment, the growing rate of the gauge mode is suppressed and the life-time of the simulation is extended as a result. \\begin{itemize} \\item The other type of adjustments (\\ref{B2a-adj} and \\ref{B2b-adj}) show their apparent effects while depending on a problem. The $\\tilde{\\Gamma}$-adjustment for the gauge-wave testbed makes the life-time longer slightly. The $\\tilde{\\gamma}$-adjustment for the Gowdy-wave testbed makes possible a simulation twice as long as the plain BSSN system. \\end{itemize} We can understand the effect of the adjustments in terms of adding dissipative terms. By virtue of the definition of the constraints, we can recognize that the adjusted equation corresponds to the diffusion equation (see, for example, Eq.~(\\ref{B1-adj})) and the signature of $\\kappa$ determines whether the diffusion is positive or negative. In the adjusted $\\tilde{A}$-eq. system, (\\ref{B1-adj}), the adjustment term corresponds to the positive diffusive term, due to the definition of ${\\cal M}_i$ and the positiveness of $\\kappa_{A}$ (see Eq.~(\\ref{eq:cal Mi}) and (\\ref{B1-adj})). This fact might explain why the adjusted $\\tilde{A}$-eq. system works effectively for all the testbeds.  In contrast, why are not all the adjustments effective in all testbeds? As we mentioned in Sec. IIB, the eigenvalue analysis was made on the linearly perturbed violation of constraints on the Minkowski space-time. Since the constraint violation grows non-linearly as seen in the Appendix of \\cite{Yoneda:2002kg}, the candidates may not be the best in their later evolution phase.   We remark upon two more interesting aspects arising from our study. The first is the mechanism of the constraint violations. As was shown in the appendix of~\\cite{Yoneda:2002kg}, each constraint propagation (behavior of their growth or decrease) depends on the other constraint terms together with itself. That is, we can guess ${\\cal A}$ and ${\\cal S}$ constraints (\\ref{eq:cons-calA} and \\ref{eq:cons-calS}) in this article, propagate independently of the other constraints, while the violation of the ${\\cal G}$-constraint, (\\ref{eq:cons-calG}) is triggered by the violation of the momentum constraint, and both the Hamiltonian and the momentum constraints are affected by all the other constraints. Such an order of the constraint violation can be guessed in Fig.~\\ref{fig:C-vio-Gauge} (earlier time), where we plot the rate of constraint violation normalized with its initial value, $||\\delta{\\cal C}||_2(t)/||\\delta{\\cal C}||_2(0)$, as a function of time, for the gauge-wave testbeds with the plain BSSN evolution. (Note that the constraints at the initial time, $\\delta{\\cal C}(0)$, are not zero due to the numerical truncation error. ) The parameters  are the same as those shown in Sec.~\\ref{sec:Gaugewave}, and the lowest resolution run is used. From this investigation, we might conclude that to monitor the momentum constraint violation is the key to checking the stability of the evolution. \\begin{figure*} \\begin{center} \\includegraphics[width=90mm]{C-violation.eps} \\caption{\\label{fig:C-vio-Gauge} The violation of all constraints normalized with their initial values, $||\\delta{\\cal C}||_2(t)/||\\delta{\\cal C}||_2(0)$, are plotted with a function of time.  The evolutions of the gauge-wave testbeds with the plain BSSN system are shown. The parameters of the test are the same as those shown in Sec.~\\ref{sec:Gaugewave}, and the lowest resolution run, $\\rho=2$, is applied. By observing which constraint triggers the other constraint's violation from the constraint propagation equations, we may guess the mechanism by which the entire system is violating accuracy and stability.  See the text for details. } \\end{center} \\end{figure*}  The second remark is on the Lagrange multipliers, $\\kappa$, used in the adjusted systems. As discussed in Sec.~\\ref{subsec:adjBSSN}, the signatures of the $\\kappa$s are determined {\\it a priori}, and we confirmed that all the predicted signatures of $\\kappa$s in~\\cite{Yoneda:2002kg} are right to produce positive effects for controlling constraint violations. However, we have to search for a suitable magnitude of $\\kappa$s for each problem. Therefore we are now trying to develop a more sophisticated version, such as an auto-controlling $\\kappa$ system, which will be reported upon in the future elsewhere.  Although the testbeds used in this work are simple, it might be rather surprising to observe the expected effects of adjustments with such a slight change in the evolution equations. We therefore think that our demonstrations imply a potential to construct a robust system against constraint violations even in highly dynamical situations, such as black hole formation via gravitational collapse, or binary merger problems. We are now preparing our strong-field tests of the adjusted BSSN systems using large amplitude gravitational waves, black hole space-time, or non-vacuum space-time, which will be reported on in the near future."
    },
    "0711/0711.1746_arXiv.txt": {
        "abstract": "We present a strong case for a transiting Hot Jupiter planet identified during a single-field transit survey towards the Lupus Galactic plane. The object, Lupus-TR-3b, transits a V=17.4  K1V host star every 3.91405d.  Spectroscopy and stellar colors indicate a host star with effective temperature 5000 $\\pm$ 150K, with a stellar mass and radius of 0.87 $\\pm$ 0.04M$_{\\odot}$ and 0.82 $\\pm$ 0.05R$_{\\odot}$, respectively. Limb-darkened transit fitting yields a companion radius of $0.89 \\pm 0.07$ R$_{\\rm{J}}$ and an orbital inclination of $88.3 ^{+1.3}_{-0.8}$ deg.   Magellan 6.5m MIKE radial velocity measurements reveal a 2.4$\\sigma$ K$=$114 $\\pm$ 25m/s sinusoidal variation in phase with the transit ephemeris. The resulting mass is 0.81 $\\pm$ 0.18M$_{\\rm{J}}$ and density 1.4 $\\pm$ 0.4g/cm$^{3}$. Y-band PANIC image deconvolution reveal a V$\\geqslant$21 red neighbor 0.4$''$ away which, although highly unlikely, we cannot conclusively rule out as a blended binary with current data. However, blend simulations show that only the most unusual binary system can reproduce our observations. This object is very likely a planet, detected from a highly efficient observational strategy. Lupus-TR-3b constitutes the faintest ground-based detection to date, and one of the lowest mass Hot Jupiters known. ",
        "introduction": "Transiting extrasolar planets provide unique opportunities to study the structure of exoplanets. Transit photometry, coupled with radial velocities, allows direct measurement of the exoplanetary radius, total mass and density.  Currently $\\sim$30 transiting planets are known\\footnote{http://exoplanet.eu/}, the majority identified in wide-field transit surveys and confirmed by radial velocities. Due to their easy detection and transit-related bias, nearly all the currently known eclipsing exoplanets are large ($\\sim R_{\\rm J}$), the exception being GJ 436b \\citep{G2007b}, and in close orbits. Transits identified against bright stars (V$\\leqslant$12) permit investigation into planetary atmosphere \\citep{VM2004,TIN2007} and infrared emission \\citep{C2005,D2005}. Deeper surveys using more numerous fainter stars can detect large numbers of planets for statistical studies of exoplanet populations, and additional system planets through large aperture photometry.  Transit surveys uncover many candidates, most of which are astrophysical false-positives which must undergo strict follow-up procedures \\citep{OD2006,OD2007}. Distinguishing hierarchical triples and blended binaries \\citep{T2004,Mand2005} requires large time investment and careful analysis with high-resolution, large-aperture spectroscopy and photometry.  We present the detection of Lupus-TR-3b, a strong case for a transiting 0.81 $\\pm$ 0.18M$_{\\rm{J}}$ Hot Jupiter planet identified from a deep, 0.66deg$^{2}$ transit survey towards Lupus. This result required 12.8 hours of Magellan 6.5m time spread over 5 nights (MIKE spectroscopy shared with another project for increased efficiency, and PANIC imagery including overheads), and 1 hour of ANU 2.3m time (DBS spectroscopy) for a target at V$=$17.4.  We detail the likely nature of this object and the overall system parameters. ",
        "conclusions": "We announce the detection of Lupus-TR-3b, a strong candidate for a transiting Hot Jupiter orbiting a V$=$17.4 K1V star with a period of 3.91405d. Our observations are consistent with a planet of radius R$_{\\rm{p}}$ $=$ 0.89 $\\pm$ 0.07 R$_{\\rm{J}}$ (assuming a Jovian equatorial radius of 71,492 kms). MIKE radial velocities reveal a 2.4$\\sigma$ low-amplitude variation in phase with the transit ephemeris, yielding a planetary mass of 0.81 $\\pm$ 0.18M$_{\\rm{J}}$ and thus density of 1.4 $\\pm$ 0.4g/cm$^{3}$, equivalent to that of Jupiter.  Deep imaging, image deconvolution techniques and extensive binary modeling argue strongly against a blended or triple configuration for the system that can reproduce the photometric observations. With current data, we cannot rule out the unlikely possibility that the nearby faint, red star E is a blended binary responsible for the transit, but this is testable with minimal amounts of planned high resolution imaging. The low-mass, Hot Jupiter Lupus-TR-3b is an example of the extraordinary discoveries that await if innovative computational methods and time-efficient follow-up observations are brought to bear on candidates selected from wide and deep transit surveys."
    },
    "0711/0711.4355_arXiv.txt": {
        "abstract": "This letter presents the first \\emph{ab initio}, fully electromagnetic, kinetic simulations of magnetized turbulence in a homogeneous, weakly collisional plasma at the scale of the ion Larmor radius (ion gyroscale). Magnetic and electric-field energy spectra show a break at the ion gyroscale; the spectral slopes are consistent with scaling predictions for critically balanced turbulence of \\Alfven waves above the ion gyroscale (spectral index $-5/3$) and of kinetic \\Alfven waves below the ion gyroscale (spectral indices of $-7/3$ for magnetic and $-1/3$ for electric fluctuations). This behavior is also qualitatively consistent with \\emph{in situ} measurements of turbulence in the solar wind.  Our findings support the hypothesis that the frequencies of turbulent fluctuations in the solar wind remain well below the ion cyclotron frequency both above and below the ion gyroscale. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.4814_arXiv.txt": {
        "abstract": "We investigate the spatial correlation properties of the solar wind using simultaneous observations by the ACE and WIND spacecraft. We use mutual information as a nonlinear measure of correlation and compare this to linear correlation. We find that the correlation lengthscales $\\lambda$ of fluctuations in density $\\rho$ and magnetic field magnitude $|B|$ vary strongly with the solar cycle, whereas correlation lengths of fluctuations in B field components do not. We find $\\lambda(|B|) \\sim 120 R_E$ and $\\lambda(|B|) \\sim 270 R_E$ at solar minimum and maximum respectively and $\\lambda(\\rho) \\sim 75 R_E$ and $\\lambda(\\rho) \\sim 170 R_E$ at minimum and maximum. The components of the B field have $\\lambda \\sim \\lambda(|B|)$ at minimum. ",
        "introduction": "\\textit{In situ} solar wind plasma observations over the last 30 years show that its local properties at 1 a.u. are modulated by the solar cycle \\cite{HapgoodSWV} and that this has an impact on the magnetosphere and ionosphere \\cite{Johnson}. There is evidence for more mixing of fast and slow solar wind plasma at solar maximum \\cite{Bame} and changing magnetic for complexity over the cycle \\cite{Kiyani} \\cite{HnatGRL}. One method for quantifying the effect of solar activity on the solar wind is to determine the correlation length scale of fluctuations measured \\textit{in-situ}. In the past this was done using the Taylor hypothesis \\cite{Taylor} and long timeseries from single spacecraft (e.g. Bruno \\textit{et al} \\cite{Bruno}). More recently, multiple spacecraft studies \\cite{Matthaeus} have used linear correlation to calculate a typical correlation length for magnetic field in the solar wind. The ACE and WIND spacecraft give an excellent range of separations and enable long baseline correlations to be calculated for magnetic field vector \\textbf{B} and magnitude $|B|$ as well as solar wind bulk velocity and density $\\rho$. These observations have been used to show anisotropy in the solar wind \\cite{Milano} and measure the Taylor microscale \\cite{Matthaeus}. Quantitative knowledge of the spatial correlation of the fluctuations of the flow is relevant to attempts to understand the solar wind in terms of locally evolving turbulence \\cite{Goldstein} \\cite{Veltri}. It also has implications for cosmic ray propagation  \\cite{Zank} \\cite{Parhi} which shows solar cycle dependence \\cite{Minnie}. \\par Previous studies used linear correlation measures and did not investigate solar cycle effects on the measurements of correlation length. Here we use a nonlinear measure of correlation, mutual information \\cite{Johnson} \\cite{Shannon} \\cite{TKMarch} \\cite{Sello} \\cite{Wicks}, alongside linear correlation coefficient. We calculate correlation lengths for the components of $\\textbf{B}$, as well as $|B|$ and $\\rho$. We also investigate the behaviour of the correlation length of these quantities over the recent solar cycle, using ACE and WIND spacecraft data. ",
        "conclusions": "We have used simultaneous data from the ACE and WIND spacecraft separated by between $30$ and $220 R_E$, to calculate the spatial correlation of solar wind $\\rho$, $|B|$ and \\textbf{B} components. Cross correlation and normalised mutual information were used as alternative linear and nonlinear measures of correlation. At solar minimum we have $48$ days of contemporaneous WIND and ACE observations. We determine the correlation lengthscale $\\lambda(\\rho)$ for $\\rho$ to be $\\lambda(\\rho) = 75 R_E$. This is smaller than $\\lambda(|B|)$ by $\\sim$ 1.6. For the components of \\textbf{B}, we find $\\lambda(B_z)$ is within error of $\\lambda(|B|)$ and $\\lambda(B_y)$ is slightly above the range of errors. \\par At solar maximum we have $15$ days of contemporaneous WIND and ACE data. This allows us to investigate the effect of the solar cycle. We find that $\\lambda(|B|)$ and $\\lambda(\\rho)$ at solar maximum are larger than at minimum by $\\sim$ a factor of $2$. In contrast the components of the magnetic field show weak variation with solar cycle. \\par We have used two window lengths; one within the inertial range ($200$ min) and the other on longer timescales ($960$ min).The window on inertial range timescales gives values of $\\lambda$, that are systematically shorter, but within the errors, of the values from the larger window. The solar cycle dependence of the correlation length is in all cases independent of our chosen window sizes. \\par Our result that $\\rho$ and $|B|$ show variation with solar cycle, whereas the components of \\textbf{B} do not, is consistent with the idea that the correlation in $|B|$ and $\\rho$ is, at least in part, of solar origin. Intriguingly this behaviour persists when we restrict our analysis to timescales within the inertial range. This is consistent with recent single-spacecraft results \\cite{Kiyani} \\cite{HnatGRL} which show solar cycle dependence within the inertial range of magnetic energy density fluctuations."
    },
    "0711/0711.0926.txt": {
        "abstract": "We continue our studies of alignment of atoms by radiation in diffuse media and their realignment by the ambient magnetic field. We understand atomic alignment as the alignment of atoms or ions in their ground or metastable states that have more than two fine or hyperfine sublevels. In particular we consider the alignment of in interstellar  and circumstellar media, with the goal  of developing new diagnostics of magnetic fields in these environments.  We provide predictions of the polarization that arises from astrophysically important aligned atoms (ions) with fine structure of the ground level, namely, OI and SII and Ti II. Unlike our earlier papers which dealt with weak fields only, a part of our current paper is devoted to the studies of atomic alignment when magnetic fields get strong enough to affect the emission from the excited level.  This is a regime of Hanle effect, but modified by the atomic alignment of the atomic ground state. We discuss also ground Hanle effect where the magnetic splitting is comparable to pumping rate. Using an example of emission and absorption lines of SII ion we  demonstrate how polarimetric studies can probe magnetic fields in circumstellar regions and accretion disks. In addition, we show that atomic alignment induced by anisotropic radiation can induce substantial variations of magnetic dipole transitions within the ground state, thus affecting abundance studies based on this emission. Moreover, we show that the radio emission arising this way is polarized, which provides a new way to study magnetic fields, e.g. at the epoch of Universe reionization. ",
        "introduction": "%In contrast to other properties of astrophysical plasma, our knowledge of astrophysical magnetic field is very limited in compatible with its importance. Lacking of precise measurement techniques, empirical methods are widely in use to infer the magnetic field. Exploring new tools is thus extremely important.  %Recent years have seen growing attention to spectropolarimetry. Indeed unique information can be learned from line polarizations. It gives clue of the geometry and topology of spatially unresolved regions, and most important it releases the information of magnetic field through different various physics phenomena.  Magnetic fields have important or dominant effects in many areas of astrophysics, but have been very difficult to quantify. We believe that atomic alignment is a new promising way of studying magnetic fields in the radiation-dominated environments. In fact, this diagnostics is far more sensitive then those based on the Zeeman effect. The basic idea of atomic alignment is simple: consider {\\em atoms} or {\\em ions} irradiated by a nearby star and embedded in a magnetic field. Anisotropic radiation pump the atoms differentially from different magnetic sublevels, resulting in over- or under-populations of the atomic states of various magnetic quantum numbers, M. The non-LTE populations produce observable polarization in the absorption or emission line involved in the interaction. The sensitivity of aligned atoms to magnetic field arises from atoms precessing in the magnetic field. It is obvious that to be sensitive to weak magnetic fields, the atomic level should be long-lived. The ground and metastable atomic levels with fine and hyperfine structure correspond to this requirement. As a result, these species are the focus of our study.   This paper is the third in the series of papers where we study atomic alignment. Rigorously speaking, ``atomic alignment\" means that the angular momentum of the atoms is not distributed isotropically: in ``fluorescent alignment\" this is due to angular momentum pumping by anisotropic illuminating light; ``magnetic realignment\" occurs when this alignment is modified by precession in a magnetic field. In the paper, however, we shall use, whenever this does not cause confusion, the terms \"alignment\" and \"realignment\" interchangeably. In Yan \\& Lazarian (2006, henceforth YLa) we dealt with polarization of absorption lines of the atoms (ions) having fine splitting, while in Yan \\& Lazarian (2007, henceforth YLb) we dealt with polarization of absorption and emission lines arising from atoms (ions) with hyperfine splitting. Although both papers were focused on developing new techniques for studies of magnetic fields in diffuse media, in both papers we discussed other effects of atomic alignment, e.g. the variations of absorption and emission intensities induced by atomic alignment. We showed that atomic alignment opens wide avenues for studies of weak magnetic fields, i.e. fields $\\lesssim10^{-5}$~G, which makes it a unique tool for magnetic field studies in many diffuse astrophysical environments. In YLab we discussed both cases of interplanetary, interstellar, circumstellar and extragalactic studies employing atomic alignment.  In this paper, we, first of all, fill in the gap and discuss the polarization of emission lines from atoms (ions) with fine splitting. As in our earlier papers, we show that this provides a technique to study the {\\em geometry} of weak magnetic fields. Similar to our earlier paper, henceforth, for the sake of simplicity, we talk about {\\em atomic alignment} and alignment of {\\em atoms} even when we deal with {\\em ions}.  More importantly, we discuss particular regimes that are also sensitive to the magnetic field intensity.  It is well known, that magnetic fields play essential roles in the dynamics of circumstellar discs and have a strong influence on the evolution of pre-main sequence stars, on accretions, outflows, stellar winds etc. Magnetically controlled accretion has been confirmed for the low mass T Tauri stars (see Bertout 1989, Camenzind 1990; K\\\"onigl 1991; Cameron \\& Campbell 1993; Shu et al. 1994; Johns-Krull et al. 1999). For intermediate mass stars, particularly Herbig Ae/Be stars, it is believed that a global magnetic field with a complex configuration is responsible for various peculiar phenomena observed for Herbig stars (e.g., Hubrig et al. 2007). For massive stars, magnetic field affects the stellar wind and causes substantial deviation from the spherical geometry (see Ignace, Nordsieck , \\& Cassinelii 1999). Zeeman effect, including both Zeeman broadening and the circular polarization is the main diagnostics adopted by the community. However, the measurements based on Zeeman effect have serious limitations: first of all, Zeeman effect is only applicable to very strong magnetic field $\\gtrsim$kG; secondly, the Zeeman splitting will be smeared out by various line-broadening processes, rotational, winds, accretion, turbulence, especially in the highly dynamical discs; and moreover Zeeman effect requires the geometry of the field to be particularly simple, as, otherwise, there will be significant cancellations of polarization between the Zeeman components.  On the other hand, the spectropolarimetry based on atomic alignment and Hanle effect does not have these constrains. Hanle effect can measure magnetic field down to 1G. And both direction and strength can be measured\\footnote{Note, that, unlike the weak magnetic field regime, which is uniquely relevant to diffuse medium, the stronger field regime, i.e. Hanle regime, is relevant to Solar studies. The research into Hanle effect resulted in important change of the views on solar chromospheres (see Bommier \\& Sahal-Brechot 1978, Landi Degl'Innocenti 1983, 1984, 1999, Stenflo \\& Keller 1997, Trujillo Bueno 1999). Studies of magnetic fields with the Hanle effect that take into account atomic alignment have been successful from the ground and resulted in calls for space-based mission to study the Sun polarization with more lines (Trujillo Bueno et al. 2005).}.  %The regime of stronger field is much better studied for particular atoms than the regime of the weak field. For us extending our treatment to that regime provides an additional advantage in terms of benchmarking of our results.  However, if Hanle effect affects the ground state, again very {\\em weak} magnetic fields can be studied. In the paper we also study ground state Hanle effect, which takes place when the precession rate gets comparable with the pumping rate of the source. As the result, polarization of both {\\em absorption} and {\\em emission} will be influenced according to the topology and the direction of magnetic field. As far as we know, no {\\em emission} lines from species with fine structure of the ground level have been reported, to be observed to emanate from interstellar or circumstellar gas. {\\em Absorption} lines are readily available, however (see Morton 1975).  The polarization of absorption lines that we study here is thus currently look more promising. Lower-level Hanle effect had been noticed in the analysis of HeI $D_3$ and Na $D_2$ polarization for solar prominences (Bommier 1980, Landolfi \\& Landi Degl'Innocenti 1985). However, given the intense pumping in solar case, the lower level Hanle regime is not well separated from the Hanle regime. For the weaker pumping by more distant source we study here, the lower-level Hanle regime become distinctive. Moreover, we shall demonstrate a new diagnostic through the lower level Hanle effect with the polarization of {\\em absorption} lines.  Atomic alignment in diffuse media demonstrates a real profusion of interesting physical effects, all of which we cannot cover in the paper. Nevertheless, we briefly touch upon an issue of the influence of atomic alignment on the measurements of element abundances, in particular, using magnetic dipole lines. Although physically different, this effect can be calculated with the same formalism as in the rest of the paper, which justifies its discussion within the same publication.  In what follows we outline the different regimes of atomic alignment and Hanle effect in \\S2. The basic formalism for studying the problem is described in \\S3. Following this, we  discuss in \\S4 the magnetic realignment regime where complementary studies to YLa,b, namely, on polarized absorption and emission are provided. The polarizations in both upper level and lower level Hanle regimes are presented in \\S5. The atomic alignment is included there, and has a notable influence to the predictions of polarizations. Example studies for a circumstellar disk with poloidal or toroidal magnetic field are performed in \\S6. In \\S7, we discuss the influence of atomic alignment to radio lines. It is demonstrated that atomic alignment not only modulates the distortion of CMB arising from optically pumped lines, therefore affects, e.g. abundance estimates in early universe, but also causes the line to be polarized, thus suggesting a new promising technique to detect magnetic field in the epoch of reionization. Discussion and summary are provided in, respectively, \\S\\S8 and 9. ",
        "conclusions": "In this paper we extended our studies of atomic alignment in diffuse medium. We provided more calculations of polarization of absorption lines arising from aligned species. We filled the gap of our earlier studies by considering polarization of emission lines arising from aligned species with fine structure. In addition, we considered atomic alignment in the presence of Hanle effect. By studying both ground and upper state Hanle effect, we attempted to give a more comprehensive coverage of the diagnostic of atomic resonant line on weak magnetic field in diffuse circumstellar, interstellar and intergalactic medium.  \\subsection{Previous work on atomic alignment} Let us briefly state some milestones in the studies of atomic alignment. A more detailed discussion is provided in YLa and YLb. First of all, atomic alignment was studied in laboratory in relation with early day maser research (see Hawkins 1955).  In astrophysical context atomic alignment was proposed by Varshalovich (1968). The atomic alignment for a idealized atom with two levels and magnetic field directed along the line of sight was considered by Landolfi \\& Landi Degl'Innocenti (1986). We {\\em benchmarked} our results by their calculations in the restricted geometry of observation. Polarization of emission lines arising from atomic pumping with mixing by magnetic field in the direction perpendicular to the symmetry axis of pumping radiation was considered in Lee, Blandford \\& Western (1994). There radiation field, nevertheless, is treated classically and we suspect that this cannot provide correct results. Originally, for the sake of simplicity, in astro-ph/0412174, we attempted a similar approach motivated by the calculations in Hawkins (1955). However, our further studies proved that the results obtained without full QED treatment are erroneous. As we pointed out in YLa,b, because conservation and transfer of angular momentum is the essence of the problem, it is necessary to quantize the radiation field as the atomic states are quantized. Our results for the polarization of the Na emission in YLb differ from the corresponding results in Hawkins (1955).  Most of the research on the atomic alignment is done in Solar physics domain\\footnote{Interestingly enough, our study shows that atomic alignment effects are more important for the diffuse medium, e.g. around luminous stars, where collisional depopulation is reduced compared to the Sun.}. This research, however, dealt with the atomic alignment in a regime, different from that important for diffuse medium. The differences include the relative role of collisions, magnetic fields and radiation.  In particular,  Sun  has radiation transfer effects that complicate the interpretation of the effects arising from alignment, but are not so important for many circumstellar and interstellar species. For this reason we developed atomic alignment theory applicable for the diffuse astrophysical medium. Some of our results, e.g. polarization of for absorption lines do not have an analogs in Solar research. Because the alignable species and the regimes of alignment are different, we performed an extended program of calculation of alignment for a number of astrophysically important species. However, in the present paper we also benchmarked our codes by reproducing some results that were obtained by the Solar community.  An extensive studies of polarization of absorption lines arising from atoms with fine structure submerged in weak magnetic field of an arbitrary direction are provided in YLa.  The first studies of astrophysically important cases of atomic alignment for atoms with hyperfine structure in weak magnetic field of arbitrary direction are provided in YLb. This paper continues and extend our studies above. Our approach can be applied to other problems, e.g. the ones discussed in Lee et al. (1994).  We stress the importance of taking into account magnetic fields while calculating the polarization. For instance, polarization of absorption arising from the atomic alignment is suggested as the mechanism of the observed polarization of H $\\alpha$ lines (Kuhn et al. 2007). However, the effect of magnetic field is neglected there, while our study present and earlier studies in YLab show that magnetic fields should be accounted for to get correct polarization measures.  \\subsection{Studies of stronger magnetic fields}  In our previous papers, we conducted an extensive study of atomic alignment and demonstrated their diagnostic power for weak magnetic field in diffuse medium (e.g. in ISM). When field gets stronger to 1G or so, the magnetic mixing becomes comparable to the line-width, this is the domain of Hanle regime. Originally defined as modification of the polarization of resonant scattering, Hanle effect has been known for decades. However, until recently (see Landi Degl'Innocenti 1999, Trujillo Bueno 1999), the alignment of the ground level and the effect of magnetic field on the ground level was not considered. Nevertheless, magnetic field is treated as a destructive factor to the ground level alignment. Magnetic field, in fact, realign the atoms towards (either $\\|$ or $\\bot$, see YLa) to the magnetic field in addition to reducing the degree of alignment on the ground state. The possibility of using the effect to diagnose the direction of magnetic field as well as the magnetic field strength is not well explored.  In this paper, using SII as an example, we provided a comparative study of Hanle effect with and without atomic alignment on the ground state. A notable effect of atomic alignment on the Hanle effect has been identified. The atomic alignment carries on additional information of the direction of magnetic field, which would be veiled otherwise. And quantitative results are provided for a generic geometry, which enables a possibility of a precise study of both magnetic field strength and its 3D direction. Indeed there are four parameters even for a known radiation source, $\\theta_0$, $\\theta_B,\\,\\phi_B$ and B field strength.  For the study of magnetic field in circumstellar medium, the geometry of the system is clearer than in ISM, namely, the direction of radiation is well defined. Three lines (they can be from the same species) are thus adequate to get the whole information of magnetic field.  We also consider lower level Hanle regime, which occurs when the inverse of the lifetime on the ground level-the photon pumping rate ($9.8\\times 10^4 (R_*/r)^2$,\\, for pumping of S II $J_l\\rightarrow J_u=5/2$ by an o star), becomes comparable to the magnetic precession rate $\\nu_L\\simeq 88(B/G)$. This can happen for a strong pumping source, e.g., quasars and supernovae, or a relatively weak magnetic field. In this situation, ground state is not aligned in respect to magnetic field. Neither the magnetic field nor the radiation field is a good quantization axis. As the result, coherence appears and the ground state density now resembles the upper state in the Hanle regime. {\\em Absorption} from the ground state is modulated the same way as the emission in the Hanle regime. Both the degree and the direction of the polarization varies with the magnetic field. Unlike the solar case, {\\em absorption} lines are readily available for circumstellar and interstellar medium. Furthermore, the upper level also gets the imprint of the magnetic field through excitation from the ground state. And therefore, the polarization of {\\em emission} also carries the information of the magnetic field. Unfortunately, the polarization of emission in the lower level Hanle regime is not well separated with that in the Hanle regime in the polarization diagram. This entails a complexity for interpreting the observations. However, if we combine emission and absorptions, or different species, we should be able to disentangle the two regimes and obtain the information of the magnetic field and the environment.   {\\em Absorption}  from atoms in the ground Hanle regime are, nevertheless, polarized in a {\\em unique} way which can distinguish itself from the other two regimes. In the magnetic realignment regime or the Hanle regime, polarization of absorption can only be parallel or perpendicular to ${\\bf B}$ in the plane of sky. In the ground Hanle regime, however, the polarization can be in any direction as the case for emission. Thus if the polarizations of multiplets are neither parallel nor perpendicular to each other, the region in study must be in the ground Hanle regime.  The polarization of {\\em absorption} lines is thus more informative.   Note that in this paper we consider the region where radiation source is far enough to be treated as a point source. There is thus a substantial dilution of the radiation field so that the dilution factor $W\\ll 1$. Thus $B{\\bar J}^0_0/A=[J_u]/[J_l]W/(e^{h\\nu/k_BT}-1)\\ll 1$, meaning that the radiative pumping rate is thus much smaller  than the spontaneous emission rate. The two regimes, namely, Hanle regime ($\\Gamma'\\simeq 1$) and ground Hanle regime ($\\Gamma'\\simeq 1$) are well separated physically. For the same reason, we don't include forward scattering when treating the absorptions since the scattering intensity $I_{em}\\propto \\rho_u\\propto {\\bar J}^0_0$ is much smaller.  \\subsection{Radio Emission Lines Influenced by Optical Pumping}  A new development of the theory here, which is considered neither in YLab, nor in Solar research, is the calculation of the magnetic dipole transitions between the sublevels of the ground state, provided that the atoms are subject to {\\it anisotropic} optical pumping. We provided calculations for OI transitions, but our approach is applicable to various atoms with hyperfine or fine splitting of levels. We showed that the emission arising from such atoms is polarized, which provides a new way of studying magnetic fields. For instance, in YLb we studied the variations of 21-cm transmission arising from the alignment of HI atoms. Our present paper predicts that the transmission within fine splitting of levels will be polarized. Apart from apparent galactic and extragalactic applications, this may be an interesting process to study magnetic fields at the epoch of reionization, which hopefully will be available with the instruments are currently under construction.  As for our study of OI pumping, we obtained, first of all, a factor of 4/3 correction for the $\\Delta T/T_{cmb}$ and $\\Delta I_\\nu/B_\\nu(T_{cmb})$ in the case of isotopic pumping of OI 63.2$\\mu{\\rm m}$ compared to early study by Hern\\'andez-Monteagudo et al. (2007). We also showed in the realistic pumping by OI Balmer $\\alpha$ lines, the anisotropy can introduce a factor of 2 variations in the resulting $\\Delta T/T_{cmb}$ and $\\Delta I_\\nu/B_\\nu(T_{cmb})$. Only when the angle between the line of sight and the incident radiation is $54.7^\\circ$ (so-called Van Vleck angle), the results are coincident with those from isotropic pumping. These are important changes to be taken in modeling of the corresponding processes, as OI is being used to study processes in the early Universe. More importantly, we predicted that the arising emission is polarized with high degree of polarization (up to 30 percent). This provides a unique opportunity to trace magnetic fields in the early Universe and also gauge the degree of reionization that arises from anisotropic optical pumping.   \\subsection{Wide extend of alignment effects}  In our studies we were focused on new ways to study magnetic field that atomic alignment of the atoms/ions with fine and hyperfine structure provides. Being alignable in their ground state, these species can be realigned in weak magnetic fields, which is extremely good news for the studies of weak magnetic fields in astrophysical diffuse gas. The latter studies are currently very limited, with polarimetry based on grain alignment  being the most widely used technique (see Whittet 2005 and ref. therein). However, in spite of the progress of grain alignment theory (see Lazarian 2007 for a review), the quantitative studies of magnetic fields with aligned atoms are not always possible. Atoms, unlike dust grains, have much better defined properties, which allows, for instance, tomography of magnetic fields by using different species\\footnote{As we discussed in YLa long-lived alignable metastable states that are present for some atomic species between upper and ground states may act as proxies of ground states. The life time of the metastable level may determine the distance from the source over which the atoms are aligned being on metastable level. Absorption from such metastable levels can be used as diagnostics for magnetic field in the star vicinity.}, which would be differentially aligned at different distances from the source. In addition, ions can trace magnetic fields in the environments in which grains cannot survive. In fact, rather than compare advantages and disadvantages of different ways of studying magnetic fields, we would  stress the complementary nature of different ways of magnetic field studies. The subject is starved for both data and new approaches to getting the data. Note, that atomic realignment happens on the Larmor precession time, which potentially allows to study the dynamics of fast variations of magnetic fields, e.g., related to MHD turbulence.  An incomplete list of objects where effects of alignment should be accounted for arises from our studies, which include this paper, as well as YLab. These include diffuse medium in the early Universe, quasars, AGNs, reflection nebulae, high and low density ISM, circumstellar regions, accretion disks and comets. One can easily add more astrophysical objects to this list. For instance, Io sodium tail can be studied the same way as sodium tail of comets. In general, in all environment when optical pumping is fast compared with the collisional processes we expect to see effects of atomic alignment and magnetic realignment of atoms. The wide variety of atoms with fine and hyperfine structure of levels ensures multiple ways that the information can be obtained. Comparing information obtained through different species one can get deep insights into the physics of different astrophysical objects. If the implications of atomic alignment influenced the understanding of particular features of the Solar spectrum, then the studies of atomic alignment in a diffuse astrophysical media can provide much deeper and yet unforeseen changes in our understanding of a wide variety of physical processes.  We note, that it is also important that some atoms do not demonstrate the effects of alignment. A comparison of the properties demonstrated by these atoms with those that demonstrate the alignment allows to gauge the importance of the effects of alignment. In addition, we would like to stress that the effects of atomic alignment are not limited by tracing all important astrophysical magnetic fields. As we discussed in YLab and this paper, the variations of the optical/UV absorptions and emissions induced by atomic alignment may be really appreciable and may substantially influence the interpretation of the absorption lines in terms of element abundances. Our present study of effects of alignment for the radio lines further extends the range of the effects for which a key role can be played by atomic alignment. This proves that further study of the effects of atomic alignment are necessary if we want to correctly interpret numerous pieces of observational data.      %\\subsection{Benchmarking our results}  %Hanle regime is also the regime most studied in connection of solar physics. So, this allows us to see the correspondence of our calculations to those provided in the corresponding literature....."
    },
    "0711/0711.2090_arXiv.txt": {
        "abstract": "We study the effect of an imposed  magnetic field on the motion of charged dust particles in  magnetically active regions of a protoplanetary disc. Assuming a power law structure for the vertical and the toroidal components of the magnetic field for the regions beyond magnetically dead region of the disc, the radial and the vertical velocities of the charged particles, in the asymptotic case of small particles, are calculated analytically. While grains with radii smaller than a critical radius significantly are affected by the magnetic force, motion of the particles with larger radii is independent of the magnetic field. The critical radius depends on the magnetic geometry and the charge of the grains. Assuming that a grain particle has one elementary charge and the physical properties of the disc correspond to a minimum-mass solar nebula, we show that only micron-sized grains are affected by the magnetic force. Also, charge polarity determines direction of the radial velocity. For such small particles, both the radial and the vertical velocities increase  due to the magnetic force. ",
        "introduction": "Physical structures of an accretion disk are determined by various physical factors such as gravity of the central star, turbulence, magnetic fields. Understanding the time evolution of dust particles in circumstellar disc-like structures around protostars and young stellar objects is crucial for explaining the formation of planetary systems. Due to various complex physical processes that are playing significant roles in dust and gas discs, it is difficult to present a complete and self-consistent model for such systems. Because of the gas drag force, the motion of the dust particles is different from that of gas (e.g., Weidenschilling 1977). The gas drag force takes angular momentum from particles, resulting in their inward migration (Whipple 1972). Garaud, Barri$\\grave{e}$re-Fouchet \\& Lin (2004) found that small particles move smoothly to the midplane, while large bodies oscillate about the midplane with decreasing amplitude.  Takeuchi \\& Lin (2002) (here after TL) studied grains motion in optically thick disc, using the assumption that radiation pressure on the dust is negligible. They showed that as the gas accretion proceeds, small dust particles are left behind in the disc, and the dust to gas ratio increases. In a subsequent paper, Takeuchi \\& Lin (2003) investigated the outflow of dust particles on the surface layers of optically thick discs. Klahr \\& Lin (2001) studied the motion of grains in optically thin disc, where the dust experiences both gas drag and radiation pressure. In a similar analysis, Takeuchi \\& Artymowicz (2001)  suggested that ring-like features which have been observed in some circumstellar debris discs (e.g. Schneider et al. 1999; Marsh et al. 2005), arise naturally in discs passing through the transitional phase from gas-dominated to dust-dominated discs. The dust particles migrate radially until they are in co-rotation with the gas.   To our knowledge, most of the previous studies are assuming that the motion of particles in the gas and dust accretion discs exclusively caused by gravitation, gas drag force and radiation pressure of the central protostar. However, collisional dust grain electrification in the early solar nebula has long been discussed to be a cause for lightning, which could explain the melting of meteoritic chondrule precursors. Poppe, Blum $\\&$ Henning (2000) studied collisional grain charging experimentally. They showed that collisional grain charging is stronger than previously discussed with respect to preplanetary grains and should be considered concerning the preplanetary dust aggregation and the formation of lightning in the solar nebula. It also has influences on other physical conditions in gas and dust discs, because charged particles can interact with the magnetic field. Although the strength and direction of magnetic fields in accretion discs are rather unknown, the observation of winds and outflows is explained by magnetic acceleration mechanisms that require the existence of a magnetic field with strength of $\\approx 10^{-4}$ T at 1 AU solar distance (K$\\ddot{o}$nigl $\\&$ Ruden 1993).    The possible effects of magnetic fields on the structure of accretion discs have been studied both numerically or analytically using simplified models (e.g. Lovelace, Romanova \\& Newman  1994; Bisnovatyi-Kogan \\& Lovelace 2000; Igumenshchev, Narayan \\& Abramowicz 2003). These studies, generally show that magnetic fields may have significant effects on the structure of the discs. So, in gas and dust discs, one may expect such type of behaviors due to the magnetic fields.  However, the ionization structure of protoplanetary discs which determines the coupling between the magnetic field and the grains is likely to be inhomogeneous and complex. The ionization fraction at different locations, and even different times, are not easy to calculate (e.g. Ilgner \\& Nelson 2006). In fact, the idea of a magnetically dead-zone was first proposed by Gammie (1996). He argued that there is a poorly ionized region where the growth of magnetorotational instability against ohmic dissipation can not be sustained. The dead zone typically stretching out to around 10 AU to 15 AU (e.g. Reyes-Ruiz 2001; Matsumura \\& Pudritz 2006). Wardle (2007) studied the ionization equilibrium and magnetic diffusivity as a function of height from the disc midplane at radii 1 and 5 AU.  He showed that the dead zone may not extend over the entire vertical dimension of the disc, because not only the ionization fraction but the gas density are  strong functions of height above the midplane and so the magnetic coupling is likely to increase away from the midplane. Moreover, the ionization fraction itself is a strong function of the presence (and size) of the dust particles mixed with the gas and recent calculations have shown that once the grains have settled, the ionisation level is sufficient to couple the magnetic field to the gas at 1 AU, even at the midplane (Wardle 2007). But the disc becomes magnetically active beyond the dead zone (e.g. Reyes-Ruiz 2001). In this study, we are applying a simplified and phenomenological model for describing the gas component of a gas and dust disc for regions beyond the dead zone.    We are assuming only the dust particles which are influenced by an imposed magnetic field.  The properties of any external magnetic fields threading a dust and gas disc are not well-known. However, it is very unlikely that the dipolar field of the central protostar is important at several AU from the protostar where the planets are believed to form. More probably such an external field comes from the the molecular cloud core out of which the gas and dust disc formed.   These considerations show that magnetic forces on micron-sized particles in gas and dust disks will be worth further investigation. In fact, the structure of magnetic fields within protostellar discs may be studied via polarimetry provided that grains are aligned in respect to magnetic field within the discs. In this study, we assume that the charged particles are collisionless and indestructible spheres in a gaseous disc and, to keep the problem tractable, we also neglect the feedback of the particles' drag on the motion of the gas.  We mainly follow the approach of TL in analyzing the dynamics of grains, but in an imposed magnetic field. We show that a vertical imposed magnetic field affects the radial and the vertical motions of charged dust particles. The basic equations of the model are described in section 2. In sections 3 we solve for individual particles' velocities, which are charged and small with sizes smaller than the mean free path of the gas molecules. In the final section, we summarize our results. ",
        "conclusions": "We presented a model for the dynamics of charged grains in a magnetized, protoplanetary disc. Properties of the gas component correspond to a minimum-mass solar nebula and the toroidal and the poloidal components of the magnetic field are assumed to be power-law. The vertical and the radial velocities of the charged dust particles depend on the geometry of the magnetic field, charge and mass of the grains. We assumed that each particle is charged one elementary charge.   Our model illustrates the importance of charge of dust grains in their motions beyond the dead zone. Observations of protoplanetary discs, at least in the optical and infrared, mostly probe the upper disc layer. For charged grains  with radii smaller than $s_{\\rm c}$,  effect of the magnetic force on their motion is stronger in the upper layers comparing to the midplane region. Since magnetic force prevents sedimentation of such grains to the midplane of the disc, one may expect grains with radii larger than a critical radius in midplane region, unless smaller grains are neutral. In a disc resembling to the minimum-mass solar nebula, and assuming that all grains are charged, our model predicts that grains in midplane of the disc are mostly larger than $s_{\\rm c}=1$ $\\mu$m. Dust sedimentation is known to affect the infrared spectra and the images of discs. In particular the far-infrared emission may be reduced by strong sedimentation. Although we did not investigate  spectra of the disc, it is interesting to study  possible modifications to the spectra due to the charge of the grains and the magnetic force. However, in a more realistic model not only the coupling between the field lines and the charged grains should be considered but also the significant effects of magnetic fields on the gas component.  Takeuchi \\& Lin (2003) studied the outflow of dust particles in the upper layers of optically disc. They showed that grains in surface layers are moving outward due to the stellar radiation. The outward mass flux is carried primarily by particles of size $0.1$ $\\mu$m according to their calculations. Although we neglected the stellar radiation in our model, we expect strong outflows of sub-micron sized particles, if they are charged positively.  In model of Takeuchi \\& Lin (2003), radiation pressure  is the main physical factor in making surface outflows of grains. But in our model, it is the magnetic force  which drives positively charged grains smaller than $s_{\\rm c}$ to move outward. The actual optical depth along the radial direction in the midplane of the disk is huge, and thus, the radiation pressure from the central star onto the dust grains should be negligible. But neutral grains smaller than the critical radius in the midplane region can move outward once they charged positively. When they reach to the surface layers, not the magnetic force but the radiation pressure contribute to their outward motion. It will be interesting to study surface outflows under combined action of gas drag,  the radiation pressure, and the magnetic force in order to understand their relative importance in making surface outflows.    We can summarize our results:  (i) Motion of particles with radii smaller than $s_{\\rm c}$ are affected by the magnetic field, but larger particles move independent of their charge and the magnetic force. For input parameters corresponding to a minimum-mass solar nebula, the critical radius is around 1 $\\mu$m and so, sub-micron sized particles are affected by the field lines.  (ii) However, if the exponent of the magnetic field decreases or the charge of the grains increases, the critical radius $s_{\\rm c}$ increases.  (iii) In our model, magnetic force increases the radial migration of charged particles comparing to when they are neutral. Also,  vertical velocity increases due to the magnetic force.  (iv) Although neutral dust grains tend to settle towards the midplane of the disc, magnetic force prevents charged grains with radii smaller than the critical radius from settling down.  (v) Direction of the radial motion is determined by the polarity of charged grains."
    },
    "0711/0711.4096_arXiv.txt": {
        "abstract": "The paucity of reliable achromatic breaks in Gamma-Ray Burst afterglow light curves motivates independent measurements of the jet aperture. Serendipitous searches of afterglows, especially at radio wavelengths, have long been the classic alternative. These survey data have been interpreted assuming a uniformly emitting jet with sharp edges (``top-hat'' jet), in which case the ratio of weakly relativistically beamed afterglows to GRBs scales with the jet solid angle. In this paper, we consider, instead,  a very wide outflow with a luminosity that decreases across the emitting surface. In particular, we adopt the  universal structured jet (USJ) model,   that is an alternative to the top-hat model for the structure of the jet. However, the interpretation of the survey data is very different: in the USJ model we only observe the emission within the jet aperture and the observed ratio of prompt emission rate to afterglow rate should solely depend on selection effects. We compute the number and rate of afterglows expected in all-sky snapshot observations as a function of the survey sensitivity. We find that the current (negative) results for OA searches are in agreement with our expectations. In radio and X-ray bands this was mainly due to the low sensitivity of the surveys, while in the optical band the sky-coverage was not sufficient. In general we find that X-ray surveys are poor tools for OA searches, if the jet is structured. On the other hand, the FIRST radio survey and future instruments like the Allen Telescope Array (in the radio band) and especially GAIA, Pan-Starrs  and LSST (in the optical band) will have chances to detect afterglows. ",
        "introduction": "\\label{sec:intro}  Surveys for transient sources may detect Gamma-Ray Burst (GRB) afterglows. In this paper, we call an ``orphan'' afterglow {\\it any} afterglow associated with such serendipitous searches, as opposed to ``triggered'' GRB afterglows, localized through the preceding prompt $\\gamma$-ray emission. Rhodes (1997) suggested that these surveys could be used to put constraints on the geometrical beaming angle of the GRB jets in the ``top-hat'' (thereafter TH) model.  In this model, GRBs are assumed to be uniformly emitting within a cone of angle $\\theta_{\\rm jet}$ with sharp edges, where the luminosity drops suddenly to an undetectable level.   The suggestion by Rhodes (1997) is based on the fact that the prompt $\\gamma$-ray emission is relativistically beamed within an angle $\\theta_{\\rm jet}+1/\\Gamma$, where $1/\\Gamma << \\theta_{\\rm jet} $.  Thus, if the line of sight lies outside this angle, the GRB is unlikely to be detected. However, as the outflow slows down in the afterglow phase, the visible region increases to eventually encompass the observer's line of sight.  At late times, $\\Gamma\\sim 1$ and the emission is roughly isotropic. This behaviour would suggest that more long-wavelength transients than $\\gamma$-ray ones {\\rm may} be expected. Detections of transients at different wavelengths may thus be used to constrain the beaming factor $b \\propto \\theta_{\\rm jet}^{-2}$. A measure of this quantity would be of great importance as it would allow one to calibrate $\\theta_{\\rm jet}$ and then estimate the true GRB rates ($\\propto b$) and energetics ($\\propto b^{-1}$).  Searches for OAs have been performed at various wavelengths, but none of them has yielded a firm detection. Several authors have used observations in the radio band to constrain the GRB rate and energetics (e.g. Perna \\& Loeb 1998; Woods \\& Loeb 1999; Paczyski 2001; Levinson 2002; Gal-Yam et al. 2006), since late time (i.e. nearly isotropic) afterglow emission peaks in this band.  The simplest and most common assumption is that the predicted number of OAs in a snapshot observation is {\\it proportional} to the beaming factor. Perna \\& Loeb (1998) used the lack of detections to set an upper limit of $b\\la 10^3$.  More recently, Levinson et al. (2002) and Gal-Yam et al. (2006) compared their data with a more detailed model for radio afterglows.  They showed that the number of expected OAs in a flux limited survey is {\\it inversely} proportional to $b$ and they place a lower limit of $b \\gsim 60$.   X-ray survey data have been used to search for OAs by Grindlay (1999) and by Greiner et al. (2000), while optical searches have been more numerous (Schaefer et al. 2002; Vanden Berk et al. 2002; Becker et al. 2004; Rykoff et al. 2005; Rau et al. 2006, Malacrino et al. 2007). Using shorter wavelengths than radio to constrain the beaming factor necessarily requires a careful comparison with theoretical predictions (e.g. Totani \\& Panaitescu 2002; Nakar, Piran \\& Granot 2002), since the emission is likely to be still relativistically beamed when OAs are detected in those bands.  Malacrino et al. (2007, 2007b) put the tightest optical constraints so far (see their fig.3). Their non-detection resulted in an upper limit for the number of OAs on the sky that is marginally consistent with the predictions of Totani \\& Panaitescu (2002, thereafter TP02) and consistent with Nakar et al. (2002, thereafter N02) and Zou et al. (2007, thereafter Z07).   The purpose of this paper is to predict results for searches of afterglows in surveys, for a different GRB jet structure. Currently, OA surveys generally aim at constraining the jet angle. However, if GRB outflows are {\\it not} geometrically beamed but they rather have an anisotropic luminosity distribution, the interpretation of data within the TH scenario would be misleading.   Numerical simulations of collapsing massive stars (e.g. MacFadyen \\& Woosley 1999) show that the jet emerges from the star with an energy distribution $E(\\theta)$ and Lorentz factor $\\Gamma(\\theta)$ that vary as a function of the angle $\\theta$ from the jet axis. It has been shown (Rossi, Lazzati \\& Rees 2002; Zhang \\& \\Mesz ~2002) that, if $E(\\theta) \\propto \\theta^{-2}$ the diversity of afterglow light curves can be ascribed to different viewing angles within the context of an universal structured jet (USJ). In the USJ model, the outflow is geometrically wide.  In this paper, we will postulate that for each GRB two simultaneous and opposite jets with $\\theta_{\\rm jet}=90^{\\circ}$ are produced. In this model $b=1$. The feature of having emission into $4\\pi$ solid angle is attractive since it can explain the lack of OA detections, while the non-uniform energy distribution allows one to avoid the huge energy requirement, demanded by GRBs with an isotropic equivalent energy $\\gsim 10^{54}$ ergs (e.g.~GRB 990123). An important consequence of the energy distribution law $E(\\theta) \\propto \\theta^{-2}$ is that it establishes a unique relation between the viewing angle and the observed luminosity, once a radiation efficiency law with angle is assumed. Thus, unlike in the TH model, the observed luminosity function is not a free parameter. Consequently, the uncertainties in the predicted OA rates in the USJ scenario are smaller than in the TH model (see \\S~\\ref{sec:USJvsUJ}).    In this paper, we use the USJ framework to compute the expected number of transients in an all-sky snapshot  and their rate as a function of the survey sensitivity. The aim is pursued by means of Monte Carlo simulations of the afterglow properties, as observed by X-ray, optical and radio surveys.  The procedure is described in \\S~\\ref{sec:simule}. We show prospects for OA detections with current and future surveys in \\S~\\ref{sec:results}. This allows us to identify the best survey characteristics to increase the chance of detection and single out the most promising future missions. A comparison of our results with the top-hat predictions is performed in \\S~\\ref{sec:USJvsUJ}. Finally,  we discuss and conclude our work in \\S~\\ref{sec:fine}.  Throughout the paper, we assume a flat Universe with $H_{0}=73\\ \\mathrm{km ~s^{-1} Mpc^{-1}}$, $\\Omega_\\mathrm{m}=0.3$ and $\\Omega_\\mathrm{\\Lambda}=0.7$. ",
        "conclusions": "\\label{sec:fine}  The realization that GRBs may be jetted has triggered studies of orphan afterglows. Early studies assumed that the prompt GRB emission comes from a sharp jet, in which case the ratio of afterglows to that of GRBs yields a constraint on the GRB beaming fraction (or equivalently the opening angle of the GRB jet). This parameter is of importance for a proper assessment of GRB rates and energetics.  The structured jet model has offered an equivalent explanation for afterglow phenomenology. However, its interpretation of the lightcurve breaks is different: they would arise from viewing angle effects and not from geometrical collimation.  In fact, if GRB jets are indeed structured, most, if not all, afterglows should be generally preceded by a prompt emission pointing towards the observer. Therefore, even if in practice the relative number of detections in various bands depends on the survey strategy, the ratio should tend to unity ($b=1$) if events are detectable at arbitrarily low fluxes.   In this paper, we have investigated the detection prospects of afterglows for flux limited surveys, in the USJ framework. We conclude that large sky coverage is essential in all bands. In addition, X-ray and radio instruments should push their flux limit below $10^{-14}$ erg s$^{-1}$ cm$^{-2}$ (at 1 keV) and $\\sim 1$ mJy (at 5 GHz) respectively. Current and planned X-ray surveys are thus not suited for OAs searches, if the jet is structured. The FIRST and the future ATA projects could be successful in detecting radio OAs. The potential is even better for future optical all-sky surveys, such as GAIA and Pan-Starrs. We also note that it would be worthwhile to exploit the great sensitivity of the Suprime-Cam, for which 5-10\\% of the whole sky is sufficient for positive detections. Certainly, a combination of X-ray, optical and radio observations would yield the most of information and will help understanding whether the USJ describes correctly the structure of the jet."
    },
    "0711/0711.4575_arXiv.txt": {
        "abstract": "{Precession and nutation of the Earth depend on the Earth's dynamical flattening, $H$, which is closely related to the second degree zonal coefficient, $J_2$ of the geopotential. A small secular decrease as well as seasonal variations of this coefficient have been detected by precise measurements of artificial satellites (Nerem et al. 1993, Cazenave et al. 1995) which have to be taken into account for modelling precession and nutation at a microarcsecond accuracy in order to be in agreement with the accuracy of current VLBI determinations of the Earth orientation parameters. However, the large uncertainties in the theoretical models for these $J_2$ variations (for example a recent change in the observed secular trend) is one of the most important causes of why the accuracy of the precession-nutation models is limited (Williams 1994, Capitaine et al. 2003). We have investigated in this paper how the use of the variations of $J_2$ observed by space geodetic techniques can influence the theoretical expressions for precession and nutation. We have used time series of $J_2$ obtained by the ``Groupe de Recherches en G\\'eod\\'esie spatiale'' (GRGS) from the precise orbit determination of several artificial satellites from 1985 to 2002 to evaluate the effect of the corresponding constant, secular and periodic parts of $H$ and we have discussed the best way of taking the observed variations into account. We have concluded that, although a realistic estimation of the $J_2$ rate must rely not only on space geodetic observations over a limited period but also on other kinds of observations, the monitoring of periodic variations in $J_2$ could be used for predicting the effects on the periodic part of the precession-nutation motion. ",
        "introduction": "Expressions for the precession of the equator rely on values for the precession rate in longitude that have been derived from astronomical observations (i.e. observations that were based upon optical astrometry until the IAU~1976 precession, and then on Very Long Baseline Interferometry (VLBI) observations for more recent models). The IAU~2000 precession-nutation model provided by Mathews et~al. (2002) (denoted MHB~2000 in the following), that was adopted by the IAU beginning on 1~January~2003, includes a new nutation series for a non-rigid Earth and corrections to the precession rates in longitude and obliquity that were estimated from VLBI observations during a 20-year period. The precession in longitude for the equator being a function of the Earth's dynamical flattening $H$, observed values for this precession quantity are classically used to derive a realistic value for $H$. Such a global dynamical parameter of the Earth is generally considered as a constant, except in a few recent models for precession (Williams 1994, Capitaine et al. 2003) or nutation (Souchay \\& Folgueira 1999, Mathews et~al. 2002, Lambert \\& Capitaine 2004) in which either the secular or the zonal variations of this coefficient are explicitly considered through simplified models.  The recent implementation of the IAU~2000A precession-nutation model guarantees an accuracy of about 200 $\\mu$as in the nutation angles, and all the predictable effects that have amplitudes of the order of 10~$\\mu$as have therefore to be considered. One of these effects is the influence of the variations ($\\Delta H$) in the Earth's dynamical flattening, which are not explicitly considered in the IAU~2000A precession-nutation model. Furthermore, the IAU~2000 precession is based on an improvement of the precession rates values derived from recent VLBI measurements, but it does not improve the higher degree terms in the polynomials for the precession angles $\\psi_A$, $\\omega_A$ of the equator (see Fig. \\ref{fig:Prec_Nut_Angles}). This precession model is not dynamically consistent because the higher degree precession terms are actually dependent on the precession rates (Capitaine et~al.,~2003) and need to be improved, even though VLBI observations are unable to discriminate between recent solutions due to the limited span of the available data (Capitaine et al., 2004). One alternative way for such an improvement is to improve the model for the geophysical contributions to the precession angles and especially the influence of $\\Delta H$ (or equivalently $\\Delta J_2$).  The $H$ parameter is linked to the dynamical form-factor, $J_2$ for the Earth (i.e.~the $C_{20}$ harmonic coefficient of the geopotential) which is determined by space geodetic techniques on a regular basis. Owing to the accuracy now reached by these techniques, the temporal variation of a few Earth gravity field coefficients, especially $\\Delta C_{20}$, can be determined (for early studies, see for example Nerem et al.~(1993), Cazenave et al.~(1995) or Bianco et al.~(1998)). They are due to Earth oceanic and solid tides, as well as mass displacements of geophysical reservoirs and post-glacial rebound for $\\Delta C_{20}$. This coefficient $C_{20}$ can be related to the Earth's orientation parameters and more particularly to the Earth precession-nutation, through $H$. The purpose of this paper is to use space geodetic determination of the geopotential to estimate $\\Delta H$, in order to investigate its influence on the precession-nutation model. The $C_{20}$ data used in this study have been obtained from the positioning of several satellites between 1985 and 2002. We estimate also the constant part of $H$, based on such space geodetic measurements, and compare its value and influence on precession results with respect to those based on VLBI determinations.  In Sect.~2 we recall the equations expressing the equatorial precession angles as a function of the dynamical flattening $H$. We provide the numerical values implemented in our model, compare the values obtained for $H$ by various studies and discuss the methods on which they rely. In Sect.~3 the relationship between $\\Delta H$ and $\\Delta C_{20}$ is discussed, depending on the method implemented. We explain how these geodetic data are taken into account in Sect.~4. We present our results in Sect.~5, and discuss them in the last part. We investigate how the use of a geodetic determination of the variable geopotential can influence the precession-nutation results, considering first the precession alone, and second the periodic contribution.  In the whole study, the time scale for $t$ is TT Julian centuries since J2000, which will be denoted cy. ",
        "conclusions": "This study was based on new considerations: the use of a geodetic determination of the variable geopotential to investigate its influence on the developments of the precession angles. The major effect on the precession is due to the $J_2$ secular trend which implies an acceleration of the $\\psi_A$ precession angle. But for the moment, the available time span for $J_2$ satellite series is not as long as we need to determine a reliable $\\dot{J}_2$ value. The $J_2$ secular trend estimation based on our $C_{20}$ residuals series from 1985 to 1998 is: $\\dot{J}_2 = -2.5 \\times 10^{-9}$/cy. The accuracy of the precession expression is limited to about $1.5$ mas/cy$^2$ due to the uncertainty in this $J_2$ rate value.  Then, we can notice that the main periodic effect is due to the $18.6$-yr periodic term in $\\Delta C_20$ due to solid Earth tides. But we must say that computing the $\\Delta C_{20}$ with satellite positioning observations requires making some assumptions on the geophysical contributions to $\\Delta C_{20}$, for instance from atmospheric pressure, and oceanic or solid Earth tides. Actually, models are used, but they are not perfect and we may have some errors. So the $\\Delta C_{20}$ residuals obtained may be affected by these errors, which is why the total $\\Delta C_{20}$ contributions (residuals observed + models assumed) constitute a better series to evaluate the effects on the precession angles. This introduces Fourier terms into the $\\psi_A$ development ($4~\\mu$as and $105~\\mu$as in cosine and sine respectively; see Table \\ref{tab:Periodic_results}) that we should compare to the MHB2000-nutations. Indeed, the different terms of the total $\\Delta H$ (or $\\Delta C_{20}$) contributions have same periods as the ($\\Delta \\psi$, $\\Delta \\epsilon$) nutations. This implies that there is some coupling between the observed $\\Delta H$ effects and the nutations, which may not have been included in the MHB2000-nutations.  In the future, we will be able to compare the $J_2$ data with geophysical models and data, in order to have better ideas on the different contributions and on the secular trend. We will also be able to proceed to numerical study of this problem, and to implement a refined and more realistic Earth model."
    },
    "0711/0711.3952_arXiv.txt": {
        "abstract": "We present multiwavelength observations of a large-amplitude oscillation of a polar crown filament on 15 October 2002, which has been reported by Isobe and Tripathi ({\\it Astron. Astrophys.} {\\bf 449}, L17, 2006). The oscillation occurred during the slow rise ($\\approx$ 1 km s$^{-1}$) of the filament. It completed three cycles before sudden acceleration and eruption. The oscillation and following eruption were clearly seen in observations recorded by the Extreme-Ultraviolet Imaging Telescope (EIT) onboard the Solar and Heliospheric Observatory (SOHO). The oscillation was seen only in a part of the filament, and it appears to be a standing oscillation rather than a propagating wave. The amplitudes of velocity and spatial displacement of the oscillation in the plane of the sky were about 5 km s$^{-1}$ and 15,000 km, respectively. The period of oscillation was about two hours and did not change significantly during the oscillation. The oscillation was also observed in H$\\alpha$ by the Flare Monitoring Telescope at Hida Observatory. We determine the three-dimensional motion of the oscillation from the H$\\alpha$ wing images. The maximum line-of-sight velocity was estimated to be a few tens of km s$^{-1}$, though the uncertainty is large owing to the lack of the line-profile information. Furthermore, we also identified the spatial displacement of the oscillation in 17 GHz microwave images from Nobeyama Radio Heliograph (NoRH). The filament oscillation seems to be triggered by magnetic reconnection between a filament barb and nearby emerging magnetic flux as was evident from the MDI magnetogram observations. No flare was observed to be associated with the onset of the oscillation. We also discuss possible implications of the oscillation as a diagnostic tool for the eruption mechanisms. We suggest that in the early phase of eruption a part of the filament lost its equilibrium first, while the remaining part was still in an equilibrium and oscillated. ",
        "introduction": "\\label{Introduction}  A variety of oscillatory motions of prominences and filaments have been observed and used to infer their structure and physical parameters (see the review by Ballester, 2006). In particular, global oscillations of filaments initiated by flare-associated disturbances (Ramsey and Smith, 1966) may provide useful information on the properties of filaments as oscillators, and hence on their stability.  Isobe and Tripathi (2006; hereafter paper I) reported an event in which a polar-crown filament exhibited large amplitude oscillation in its pre-eruptive slow rise phase. This event provides a rare opportunity to use so-called coronal seismology or prominence seismology in order to study the eruption mechanism of filaments. In paper I we reported that the oscillation repeated for about three cycles without any significant damping, until the onset of rapid acceleration and eruption of the filament. The period and the velocity amplitude measured in the plane of sky of the oscillation were about 2 hours and 4.2 km s$^{-1}$, respectively. It was suggested that the large-amplitude oscillation in the filament was evidence for it retaining a non-linear stable equilibrium in its pre-eruptive, slow-rise phase. The transition from this equilibrium to eruption, \\ie, loss of equilibrium (LOE) or instability, was triggered by rapid ($\\approx$ Alfv\\'{e}n time scale) mechanism such as magnetic reconnection.  In this paper we present further detailed analyses of this event using multiwavelength observations. In addition to the detection of the oscillation in extreme ultraviolet and H$\\alpha$, we discovered that the oscillation was seen also in microwave. To the best of our knowledge, this is the first evidence of spatially resolved, large-amplitude oscillation of a filament in microwave imaging. One of the fundamental issues that remained unknown in paper I was the physical cause behind excitation of the oscillation. It is known that large amplitude filament oscillations are usually excited by flare-associated disturbances (see \\eg, Ramsey and Smith,  1966; Eto \\ietal,  2002; Jing \\ietal, 2003, 2006; Okamoto \\ietal, 2004; Vr\\v{s}nak \\ietal, 2007). However, no flare associated with the onset of oscillation in this event was observed. In this paper we examine the magnetic-field data to study the mechanism of the excitation of the oscillation.  The paper is organized as follows: In Section 2 we briefly summarize the observational data. In Section 3 the oscillation seen in various wavelengths is presented. We discuss the possible excitation mechanisms responsible for the filament oscillation and their implication to eruption mechanisms in Section 4. Conclusions are summarized in Section 5. ",
        "conclusions": "In this paper we presented the observations of a large amplitude oscillation of the polar crown filament in its pre-eruption, slowly rising phase. The oscillation was seen only in a part of the filament, and seems to be a standing oscillation rather than a propagating wave. The spatial displacement of the oscillation was also identified in 17 GHz microwave images from NoRH. Such an oscillation has not been reported during the more than ten years of NoRH observations (although we have not done a systematic survey for such events, so this does not mean that ours is the only such event), and we believe that this is the first observation of spatially resolved filament oscillation by microwave imaging.  In Contrast to previous observations of large-amplitude filament oscillations, there were no flares or even microflares associated with the excitation of the filament. However, we found that a filament barb connected near an emerging flux disappeared and seems to be ejected at the onset of the oscillation. We believe that reconnection between the emerging flux and the barb caused the ejection that blew the filament to excite the oscillation.  The large-amplitude oscillation occurred in the slow-rise phase of the filament, and only a part of the filament exhibited the oscillation. The oscillation is evidence that the oscillating part of the filament still retained a nonlinearly stable equilibrium during the slow rise. It supports the model by Chifor \\ietal (2006) that at first a part of filament loses equilibrium and starts to rise, and the slow rise induces LOE of the remaining part, leading to the destabilization of the whole filament.  Filament and prominence seismology is a powerful tool, not only for measuring the physical parameters, but also for diagnosing the stability and hence the trigger mechanism of filament eruption. It is worthwhile to search for more events in which filaments exhibit large-amplitude oscillation in their pre-eruption phase. From a theoretical point of view, there are many analytical and numerical models of pre-eruption filaments or flux ropes. Investigation of the behavior of these model filaments as oscillators may provide some physical insight into the mechanisms of filament eruptions."
    },
    "0711/0711.4119_arXiv.txt": {
        "abstract": "WMAP observations have accurately determined the position of the first two peaks and dips in the CMB temperature power spectrum. These encode information on the ratio of the distance to the last scattering surface to the sound horizon at decoupling. However pre-recombination processes can contaminate this distance information. In order to assess the amplitude of these effects we use the WMAP data and evaluate the relative differences of the CMB peaks and dips multipoles. We find that the position of the first peak is largely displaced with the respect to the expected position of the sound horizon scale at decoupling. In contrast the relative spacings of the higher extrema are statistically consistent with those expected from perfect harmonic oscillations. This provides evidence for a scale dependent phase shift of the CMB oscillations which is caused by gravitational driving forces affecting the propagation of sound waves before recombination. By accounting for these effects we have performed a MCMC likelihood analysis of the location of WMAP extrema to constrain in combination with recent BAO data a constant dark energy equation of state parameter $w$. For a flat universe we find a strong $2\\sigma$ upper limit $w<-1.10$, and including the HST prior we obtain $w<-1.14$, which are only marginally consistent with limits derived from the supernova SNLS sample. On the other hand we infer larger limits for non-flat cosmologies. From the full CMB likelihood analysis we also estimate the values of the shift parameter $R$ and the multipole $l_a$ of the acoustic horizon at decoupling for several cosmologies to test their dependence on model assumptions. Although the analysis of the full CMB spectra should be always preferred, using the position of the CMB peaks and dips provide a simple and consistent method for combining CMB constraints with other datasets. ",
        "introduction": "Cosmic Microwave Background (CMB) observations have provided crucial insights into the origin and evolution of present structures in the universe \\cite{CMBobs1,CMBobs2,Hinshaw}. Physical processes occurred before, during and after recombination have left distinctive signatures on the CMB. The most prominent feature is a sequence of peaks and dips in the anisotropy power spectrum, the remnant imprints of acoustic waves propagating in the primordial photon-baryon plasma at the time of decoupling \\cite{CMBTh1,CMBTh2,CMBTh3}. This oscillatory pattern carries specific information on several cosmological parameters \\cite{PeaksTh}. As an example the angular scale at which these oscillations are observed provides a distance measurement of the last scattering surface to the sound horizon at decoupling, hence a clean test of cosmic curvature \\cite{Ktest}.  WMAP observations have accurately detected the peak structure of the CMB power spectrum. These data have constrained the geometry of the universe to be nearly flat and have precisely determined other cosmological parameters \\cite{Spergel}. On the other hand constraints on dark energy are less stringent, this is because its late time effects leave a weaker imprint of the CMB which is diluted by degeneracies with other parameters. Indeed other cosmological tests can be more sensitive to the signature of dark energy, nonetheless they still require additional information from CMB to break the parameter degeneracies. As an example CMB constraints are usually combined with those from SN Ia luminosity distance data. Alternatively the CMB can be used in combination with measurements of the baryon acoustic oscillations (BAO) in the galaxy power spectrum \\cite{Eisenstein}. In fact the same acoustic signature present in the CMB is also imprinted in the large scale distribution of galaxy, thus providing a complementary probe of cosmic distances at lower redshifts.  A likelihood analysis of the CMB spectra is certainly the more robust approach to implement CMB constraints with those from other datasets. This can be very time consuming, henceforth one can try to compress the CMB information in few measurable and easily computable quantities. Recent literature has focused on the use of the shift parameter $R$, and the multipole of the acoustic scale at decoupling $l_a$ \\cite{Elgaroy,Wang}. However these quantities are not directly measured by CMB observations, they are inferred as secondary parameters from the cosmological constraints obtained from the full CMB likelihood analysis. Consequently their use as {\\it data} can potentially lead to results which suffer of model dependencies as well as prior parameter assumptions made in the analysis from which the values of $R$ ($l_a$) have been inferred in the first place. In contrast the multipole location of the CMB extrema can be directly determined from the observed temperature power spectrum through model-independent curve fitting. These measurements can then be used to constrain cosmological parameters provide that pre-recombination corrections are properly taken into account.  In this paper we analyse in detail the cosmological information encoded in the position of the CMB extrema as measured by WMAP. Our aim is to provide a simple and unbiased method for incorporating CMB constraints into other datasets which is alternative to that of using $R$ and/or $l_a$ \\cite{Elgaroy,Wang}. Firstly we estimate the amplitude of pre-recombination mechanisms that can displace the location of the CMB extrema with the respect to the angular scale of the sound horizon at decoupling. In particular we show that the WMAP location of the first peak is strongly affected by such mechanisms, while the displacements induced on the higher peaks and dips are smaller. By accounting for these effects we perform a cosmological parameter analysis and infer constraints on dark energy under different prior assumptions, including the cosmic curvature. We then combine these results with measurements of BAO from SDSS and 2dF data \\cite{Percival}, and confront the inferred constraints with those obtained using SN Ia data from the Supernova Legagy Survey \\cite{Astier}. Finally we test for potential model dependencies of $R$ (and $l_a$) by performing a full likelihood analysis of the WMAP spectra for different sets of cosmological parameters.   The paper is organized as follows: in Section~\\ref{Peaks} we review the physics of the CMB acoustic oscillations. In Section~\\ref{WMAP} we discuss the relative shifts of the multipoles of the WMAP peaks and dips. In Section~\\ref{Cosmo} we present the results of the cosmological parameter inference using the location of the CMB extrema in combination with BAO. In Section~\\ref{sndata} we confront the results with the SN Ia likelihood analysis from the SNLS sample. We discuss the results on the shift parameter in Section~\\ref{shift} and present our conclusions in Section~\\ref{Conclusions}. ",
        "conclusions": "The multipoles of the CMB extrema can be directly measured from the WMAP spectra and used to combine CMB information with other cosmological datasets. Corrections to the location of the CMB peaks and dips from pre-recombination effects need to be taken into account for an unbiased parameter inference. Here we have shown that the position of the first peak as measured by WMAP-3yrs data is strongly displaced with the respect to the actual location of the acoustic horizon at recombination. This displacement is caused by gravitational driving forces affecting the propagation of sound waves before recombination. These effects are smaller on higher harmonics, indicating the presence of a scale dependent phase shift which becomes negligible on scales well inside the horizon.  We have performed a cosmological parameter inference using the position of the WMAP peaks and dips in combination with recent BAO measurements and derived constraints on a constant dark energy equation of state under different model parameter assumptions.  The method we have presented here is alternative to using the shift parameter $R$ and/or the multipole of the acoustic horizon at decoupling $l_a$. We have tested for potential model dependencies of $R$ and $l_a$ by running a full CMB spectra likelihood analysis for different class of models. Indeed for simple dark energy models the inferred constraints on $R$ and $l_a$ do not differ from those inferred assuming the vanilla $\\Lambda$CDM. Nevertheless we have suggested caution in using these secondary parameters as {\\it data}, since hidden assumptions may lead to biased results particularly when testing models which greatly depart from the $\\Lambda$CDM cosmology.  Indeed we do advocate the use of the full CMB spectra, particularly for constraining the properties of dark energy. In fact more information on dark energy is encoded in the full CMB spectrum than just in the distance to the last scattering surface. Nevertheless we think that using the location of the CMB extrema provide a fast and self-consistent approach for combining in a friendly user way the CMB information with complementary cosmological data."
    },
    "0711/0711.3025_arXiv.txt": {
        "abstract": "Since 2002, the number of detected blazars at gamma-ray energies above 100 GeV has more than doubled. I study 17 blazars currently known to emit $E>100$~GeV gamma rays. Their intrinsic energy spectra are reconstructed by removing extragalactic background light attenuation effects. Luminosity and spectral slope in the $E>100$~GeV region are then compared and correlated among each other, with X-ray, optical and radio data, and with the estimated black hole (BH) masses of the respective host galaxies.  According to expectations from synchrotron self-Compton emission models, a correlation on the 3.6-$\\sigma$ significance level between gamma-ray and X-ray fluxes is found, while correlations between gamma-ray and optical/radio fluxes are less pronounced. Further, a general hardening of the blazar spectra in the $E>100$~GeV region with increasing gamma-ray luminosity is observed, both for the full 17-source sample and for those sources which have been detected at distinct flux levels. This goes in line with a correlation of the gamma-ray luminosity and the synchrotron peak frequency, which is also seen. Tests for possible selection effects reveal a hardening of the spectra with increasing redshift. The blazar gamma-ray emission might depend on the mass of the central BH. The blazars under study do, however, show no correlation of the BH masses with the spectral index and the luminosity in the $E>100$~GeV region.  I also consider temporal properties of the X-ray and $E>100$~GeV gamma-ray flux. No general trends are found, except for the observation that the blazars with the most massive BHs do not show particularly high duty cycles. These blazars include Mkn\\,501 and PKS\\,2155-304, for which recently very fast flares have been reported. In general, VHE flare time-scales are not found to scale with the BH mass.  As a specific application of the luminosity study, a constraint for the still undetermined redshift of the blazar PG\\,1553+113 is discussed. ",
        "introduction": "All but one of the detected extragalactic very high energy (VHE, defined by $E>100\\,\\mathrm{GeV}$) gamma ($\\gamma$) ray sources so far are blazars. Within the unified scheme \\citep[e.g.][]{UrryPadovani} of active galactic nuclei (AGN), blazars comprise the rare and extreme subclasses of BL~Lac objects and flat spectrum radio quasars (FSRQs). These are characterised by high apparent luminosities, short variability time-scales, and apparent superluminal motion of jet components. These observations can be explained by highly relativistic, beamed plasma outflows (jets) closely aligned to the observer's line of sight \\citep{1979ApJ...232...34B} powered by central supermassive black holes accreting at sub-Eddington rates \\citep{1969Natur.223..690L,1978Natur.275..516R}. The prime scientific interest in VHE $\\gamma$-ray emitting blazars (in the following, `VHE blazars') is twofold: (1) To understand the particle acceleration and $\\gamma$-ray production mechanisms, assumed to take place in the jets and to be linked to the central supermassive black hole (BH). Knowledge of the VHE emission process will also contribute to the further understanding of the accretion processes in AGN, jet formation processes, and the jet structure. (2) To use the VHE $\\gamma$-rays as a probe of the extragalactic background light \\citep[EBL; e.g.][]{hauser,2005PhR...409..361K} spectrum in the wavelength range between about $0.3$ to $30~\\mu\\rm m$. Determining the EBL spectrum in this wavelength range may allow to constrain the star formation rate (convolved with the initial mass function) in the early Universe. In order to assess both issues, it is essential to have a large sample of VHE $\\gamma$-ray blazars at hand. Ideally it should encompass a wide range in redshift for EBL studies and at the same time include groups of sources at similar distances in order to probe and compare properties of the individual sources without possible systematic uncertainties caused by the EBL de-absorption. \\begin{figure*} \\center{\\includegraphics[width=0.65\\linewidth]{fig01} \\caption{Currently known VHE $\\gamma$-ray blazars along with the identified (66 objects) and tentatively identified AGN (27 objects) in the 3rd EGRET catalogue of $\\gamma$-ray sources (solid grey dots: identified AGN; open grey dots: tentatively identified AGN). EGRET data from \\citet{hartman}. The sources are shown in a galactic coordinate system.\\label{fig:comp:sp}}} \\end{figure*}  The preconditions for such studies have much improved recently: Before 2004, only a few nearby extragalactic sources had been established as VHE $\\gamma$-ray emitters \\citep[e.g.][]{Moriicrc} and provided hardly enough data to perform comparative studies. Around 2004, the third generation of imaging air Cerenkov telescopes (IACTs, e.g. MAGIC, \\citealt{Baixeras,Cortina} and H.E.S.S., \\citealt{HESStech}), the most successful tools so far to explore VHE $\\gamma$-rays, started to deliver scientific results. To date, the VHE blazar sample with available spectral information in the VHE region comprises 17 BL Lac objects, % among them one LBL object, BL~Lacertae. Furthermore, now the redshifts of the known VHE blazars reach up to $z=0.212$ -- or even to $z=0.536$, considering the recently announced discovery of the first FSRQ in VHE $\\gamma$-rays, 3C~279 \\citep{279}. For 3C~279, however, no VHE spectrum has been published yet, therefore the VHE luminosity and spectral slope of 3C~279 are not yet available for this study. M87, a FR~I radio galaxy also detected in VHE $\\gamma$-rays, is not included in the study, as its VHE $\\gamma$-ray production mechanism may differ from that in blazars \\citep{hess87n}. Fig.~\\ref{fig:comp:sp} shows the sky positions of the known VHE blazars in galactic coordinates along with the AGN identified in the 3rd EGRET catalogue.   The electromagnetic continuum spectra of blazars extend over many orders of magnitude from radio frequencies to sometimes multi-TeV energies and are dominated by non-thermal emission that consists in a $\\nu F_\\nu$ representation of two pronounced peaks. The low-energy peak, located between the IR and hard X-rays, is thought to arise from synchrotron emission of ultrarelativistic electrons, accelerated by shocks moving along the jets at relativistic bulk speed. Depending on the location of the low-energy peak, BL~Lac objects are often referred to as high-frequency peaked (HBL; in the UV to X-ray domain) or low-frequency peaked (LBL; in the near-IR to optical) BL~Lac objects \\citep{fossati}, although the transition is smooth rather than dichotomic. The origin of the high-energy peak at MeV to TeV energies is still debated. It is commonly explained by inverse Compton upscattering of low-energy photons by electrons. The seed photons may originate from synchrotron radiation produced by the same electron population \\citep*[synchrotron-self Compton (SSC) models; e.g.][]{maraschi,coppi2} or belong to ambient thermal photon fields \\citep*[external inverse Compton models; e.g.][]{mk89,1994ApJ...421..153S,1994ApJS...90..945D}. In hadronic models, which can also explain the observed features, interactions of a highly relativistic jet outflow with ambient matter \\citep{1997ApJ...478L...5D,1993ApJ...402L..29B}, proton-induced cascades \\citep{MannheimProton}, synchrotron radiation by protons \\citep{AharonianHad,MueckeProtheroe}, or curvature radiation, are responsible for the high energy photons. The AGN identified in the EGRET data are predominantly powerful FSRQs and quasars with SEDs peaking at rather low frequencies, and thus only few of these (Mkn 421, PKS 2155-304, BL~Lacertae and 3C~279) were also detected in the VHE range.  Knowing the variability time-scales of the VHE $\\gamma$-ray emission and the form of the two-bump spectral energy distribution (SED) enables the derivation of all input parameters of one-zone SSC models \\citep{Tavecchio1998}, which describe the observed emission in BL~Lac objects reasonably well. Strictly simultaneous and temporally-resolved measurements of the SED, however, are only rarely possible and more often than not are also severely restricted by the (temporal) instrumental resolution. Generally, detailed spectral studies particularly in low-emission states are rather demanding. In addition, the determination of the location of the high-energy peak $\\nu^\\mathrm{IC}_\\mathrm{peak}$ generally would require complementary satellite detector coverage of the SED between some hundred MeV and $\\approx 50$~GeV. This region of the SED, however, is difficult to access due to the low fluxes expected (up to $\\approx 100$ MeV) from extragalactic sources in between the two bumps and due to insufficient instrumental sensitivity for energies exceeding some GeVs.  In this paper, for the first time studies of the VHE emission properties of the complete set of all currently known VHE blazars (cf. Tab.~\\ref{tab:TeVSources}) are performed. First, the detected VHE blazars are brought into context with the AGN searches conducted by IACTs so far and the expected $\\gamma$-ray attenuation by the EBL in Sect.~\\ref{sect:grhorsect}. After a study of the black hole mass distribution of the VHE blazars in Sect.~\\ref{sect:comp:bhmavg}, I infer intrinsic emission properties in the VHE $\\gamma$-ray regime in Sect.~\\ref{sect:intrvhe}. Because measurements of the location and shape of the high-energy bump are elusive at present for almost all VHE blazars, the observed $\\gamma$-ray luminosity and spectral slope in the VHE region are used as auxiliary observables to characterise the VHE $\\gamma$-ray emission. The main part of the paper (Sect.~\\ref{sect:corstudies1}--\\ref{sect:intrinsicspectra}) is devoted to the search for correlations of these observables with the X-ray emission properties, the optical and the radio luminosity, and with the black hole mass estimations. In Sect.~\\ref{sect:comp:501lumilim} the VHE blazar luminosity distribution is used to address the specific problem of the unknown redshift of PG\\,1553+113 by deriving an upper redshift for this VHE blazar. Finally, Sect.~\\ref{sect:dcsect} turns to the study of X-ray and VHE $\\gamma$-ray timing properties. Sect.~\\ref{sect:conc} summarises the main conclusions of the studies.  \\begin{table} \\caption{Extragalactic VHE $\\gamma$-ray sources, listed in chronological order of their discovery.} \\label{tab:TeVSources} \\begin{tabular}{lccl} \\hline Source & Type & Redshift $z$ & Discovery reference  \\\\ \\hline Mkn\\,421       & HBL & 0.030\\phantom{0} & \\citet{Punch1992}  \\\\ Mkn\\,501       & HBL & 0.034\\phantom{0} & \\citet{Quinn1996} \\\\ 1ES\\,2344+514  & HBL & 0.044\\phantom{0} & \\citet{Catanese1998}  \\\\ 1ES\\,1959+650  & HBL & 0.047\\phantom{0} & \\citet{Nishiyama1999}  \\\\ PKS\\,2155-304  & HBL & 0.116\\phantom{0} & \\citet{Chadwick}  \\\\ H\\,1426+428    & HBL & 0.129\\phantom{0} & \\citet{Horan}  \\\\ M87$^a$            & FR~I  & 0.0044           & \\citet{Tluczykont}\\\\ \\hline PKS\\,2005-489  & HBL &  0.071\\phantom{0}& \\citet{HESS2005}   \\\\ 1ES\\,1218+304 & HBL & 0.182\\phantom{0}  & \\citet{MAGIC1218} \\\\ H\\,2356-309   & HBL & 0.165\\phantom{0} & \\citet{HESSAGNNature}   \\\\ 1ES\\,1101-232  & HBL & 0.186\\phantom{0} & \\citet{HESSAGNNature}  \\\\ PG\\,1553+113   & HBL & $^b$            & \\citet{HESS1553}, \\\\ &     &               & \\citet{MAGIC1553published}  \\\\ Mkn\\,180       & HBL & 0.045\\phantom{0} & \\citet{MAGIC180}   \\\\ PKS\\,0548-322  & HBL & 0.069\\phantom{0} & \\citet{0548icrc} \\\\ BL~Lacertae    & LBL & 0.069\\phantom{0} & \\citet{MAGICBLLac} \\\\ 1ES\\,1011+496  & HBL & 0.212\\phantom{0} & \\citet{1011} \\\\ 1ES\\,0229+200  & HBL & 0.139\\phantom{0} & \\citet{0229} \\\\ 1ES\\,0347-121  & HBL & 0.188\\phantom{0} & \\citet{0347} \\\\ 3C~279$^c$     & FSRQ& 0.536\\phantom{0} & \\citet{279} \\\\ \\hline \\end{tabular}  \\medskip The upper part of the table shows the confirmed sources prior to the advent of new generation instruments like MAGIC and H.E.S.S., while the lower panel summarises the sources discovered after 2002. HBL: High-frequency peaked BL~Lac object, LBL: Low-frequency peaked BL~Lac object, FSRQ: Flat spectrum radio quasar, FR: Fanaroff--Riley galaxy. $^a$ M87 is not included in the present study. $^b$ This redshift is currently under discussion, cf. Sect. \\ref{sect:comp:501lumilim}. $^c$ No spectrum of 3C~279 has been published yet. \\end{table} ",
        "conclusions": "\\label{sect:conc} Before the new generation of IACT became operational, only six firmly detected extragalactic VHE $\\gamma$-ray sources were known \\citep[e.g.,][]{Moriicrc}; not for all of them differential energy spectra had been inferred. To date, there are 17 BL Lac objects and the FSRQ 3C~279 known to emit VHE $\\gamma$-rays. This substantially enlarged set  of blazars called for a synoptic study. I collected and derived intrinsic properties of the VHE $\\gamma$-ray emission (luminosity, spectral hardness, temporal properties) and further included X-ray, optical and radio emission properties. As an accessible property of the BHs in the centres of the blazar host galaxies, the BH mass was also used. The studies yield the following results. \\begin{itemize} \\item[(i)] So far, only blazars with $M_\\bullet \\ga 10^8 \\mathrm{M}_{\\sun}$ show VHE $\\gamma$ emission. Whether this experimental finding constitutes a selection effect or will reveal interesting physics, remains to be seen. \\item[(ii)] The VHE luminosity of the blazars under study and the corresponding X-ray luminosity show a hint of a correlation, as expected from leptonic acceleration/SSC models. \\item[(iii)] A correlation between the spectral slope in the VHE region and the peak location of the synchrotron peak is found. Such a correlation is also expected from SSC models. \\item[(iv)] There were no correlations found between the $\\gamma$ emission properties and the $M_\\bullet$ of the galaxies that host the blazars. Also, no correlation of $M_\\bullet$ could be observed with the flare duty cycles and the flare time-scales. Thus, VHE $\\gamma$-ray emission properties may not dominantly depend on $M_\\bullet$. Other possibly interesting BH parameters are not yet within instrumental reach. \\item[(v)] There is an indication that the VHE $\\gamma$ luminosity is correlated with the spectral hardness. This correlation can be formulated as a decrease of $\\Delta \\Gamma \\approx 0.82$ per decade of luminosity. \\item[(vi)] This behaviour is also found for some individual blazars that were observed in different emission states: There are indications that in some variable sources the observed spectra become harder with increasing luminosity, while in others no hardening is found. \\end{itemize}  Investigations of the temporal properties of the X-ray emission show that the blazars with the most massive BHs have rather low duty cycles. The temporal behaviour of the VHE $\\gamma$-ray emission is not well studied for most of the VHE blazars, but the present data do not seem to support a scaling of the flux doubling time-scales with the BH masses.  Obviously we are still dealing with a low number of sources and certainly with an incomplete source sample, which leaves regions in the parameter space empty. Nevertheless, the rather impressive number of 17 VHE blazars has permitted first comparative studies. It also allows conclusions concerning EBL models: A hint at a marginal correlation between the intrinsic spectral hardness and the source distance is likely due to an EBL over-prediction. This result was also recently quantified by \\citet{HESSAGNNature} using the two distant blazars 1ES\\,1101-232 and H~2356-309. In clarifying the situation, more distant sources particularly with intrinsically hard spectra (as e.g. 1ES\\,0229+200) will play a decisive role; PG\\,1553+113 with its yet undetermined distance might also turn out to be a good candidate for such analyses once a measurement of its redshift succeeds. In the meantime, the observed luminosity distribution of the studied VHE $\\gamma$ blazars was used to constrain the unknown distance of PG\\,1553+113 by assuming that the properties of this blazar are not too different from the most extreme objects in the blazar sample. Conversely, a large distance of PG\\,1553+113 implies an unusually high luminosity or an unusually high Doppler factor.  It can be said that the era of blazar astronomy has been entered -- astronomy being understood as the study of generic properties of a given class of objects. With the currently known VHE blazars and the given instrumental sensitivity, VHE blazar astronomy starts to become less biased: Not only blazars with hard spectra or during outbursts have been detected, but also low (quiescent?) emission states in the VHE range are now being observed and studied.  With continuing discoveries of new sources, it will become easier to disentangle EBL absorption effects and intrinsic absorption effects from the measured spectra. Already currently, the VHE blazar sample contains groups of objects at very similar distances, e.g. 1ES\\,1218+304 and 1ES\\,1101-232 ($\\Delta z=0.004$), or the triplet 1ES\\,2344+514, Mkn~180 and 1ES\\,1959+650 ($\\Delta z=0.003$). Such groups are ideal for direct comparisons between the respective individual objects, as they are subject to very similar EBL attenuation.  Since 2003, on average three new blazars have been detected per year. When comparing this rate to the ten years it took to discover the first six VHE blazars, one can with good reason hope to accumulate many more sources in the near future, which will refine synoptic blazar studies as the pioneering one conducted here. In addition, the first observation of new object classes like LBLs and FSRQs will certainly help to sharpen the VHE view of a possible blazar sequence and the underlying physics in the near future."
    },
    "0711/0711.3739_arXiv.txt": {
        "abstract": "{} {We aim to investigate the ability of simple spectral models to describe the GRB early afterglow emission.} {We performed a time resolved spectral analysis of a bright GRB sample detected by the Swift Burst Alert Telescope and promptly observed by the Swift X--ray Telescope,with spectroscopically measured redshift in the period April 2005 -- January 2007.  The sample consists of 22 GRBs and a total of 214 spectra. We restricted our analysis to the softest spectra sub--sample which consists of 13 spectra with photon index $>$ 3.} {In this sample we found that four spectra, belonging to GRB060502A, GRB060729, GRB060904B, GRB061110A prompt--afterglow transition phase, cannot be modeled neither by a single power law nor by the Band model. Instead we find that the data present high energy ($>$ 3 keV, in the observer frame) excesses with respect to these models.  We estimated the joint statistical significance of these excesses at the level of 4.3 $\\sigma$.  In all four cases, the deviations can be modeled well by adding either a second power law or a blackbody component to the usual synchrotron power law spectrum. The additional power law would be explained by the emerging of the afterglow, while the blackbody could be interpreted as the photospheric emission from X--ray flares or as the shock breakout emission. In one case these models leave a 2.2$\\sigma$ excess which can be fit by a Gaussian line at the energy the highly ionized Nickel recombination.} {Although the data do not allow an unequivocal interpretation, the importance of this analysis consists in the fact that we show that a simple power law model or a Band model are insufficient to describe the X-ray spectra of a small homogeneous sample of GRBs at the end of their prompt phase.} ",
        "introduction": "The X-ray telescope (XRT, \\cite{Burrows05}) on board the Swift satellite (\\cite{Gehrels04}) allows to perform time resolved spectroscopy of a large number of gamma ray burst (GRB) afterglows in the 0.3-10 keV energy band.  The vast majority of the spectra can be modeled by a single absorbed power law (SPL) model.  In most of the afterglows a strong spectral evolution is observed in the early phases with spectral indexes varying in the range 0.5-5 (\\cite{Obrien06},\\cite{Butler07}, \\cite{Zhang07}).  These observations are consistent with the classical fireball model which describes the burst and afterglow emission due to synchrotron radiation. \\begin{figure*} \\begin{tabular}{cc} \\includegraphics[width=8.0cm]{f1_a.ps}  & \\includegraphics[width=8.0cm] {f1_h.ps}  \\\\ \\includegraphics[width=8.0cm]{f1_b.ps}   & \\includegraphics[width=8.0cm]{f1_j.ps}  \\\\ \\includegraphics[width=8.0cm]{f1_c.ps}   & \\includegraphics[width=8.0cm]{f1_k.ps}  \\\\ \\includegraphics[width=8.0cm]{f1_d.ps}   & \\includegraphics[width=8.0cm]{f1_i.ps}  \\\\ \\includegraphics[width=8.0cm]{f1_e.ps}   & \\includegraphics[width=8.0cm]{f1_l.ps}  \\\\ \\includegraphics[width=8.0cm]{f1_f.ps}   & \\includegraphics[width=8.0cm]{f1_m.ps}  \\\\ \\includegraphics[width=8.0cm]{f1_g.ps}   &                                         \\\\ \\end{tabular} \\caption{The 13 very soft spectra are shown together with their best SPL fit. They are sorted in time from the top-left corner.  The observed time interval and the photon index are indicated in the figures. The asterisks denote the four spectra for which we found significant departures from SPL model.} \\label{fig:f1} \\end{figure*} It is possible that detection of different spectral features, signature of other emission mechanisms, like thermal components or recombination lines might give some insight about GRB progenitors, chemical composition, physical conditions and geometry of the GRB environment. Before the Swift mission, several X--ray line detections were reported in GRB late ($>$ 10 hr) afterglow observations. The statistical significance of these detections has been questioned by \\cite{Sako05}, who concluded that there were no credible X--ray features in any GRB afterglow. Butler et al. (2005), however, showed that for GRB011211 the different estimates of the statistical significance can be explained by the different approach in the continuum modelling.  In the Swift afterglow observations only a few deviations from SPL have been reported.  Most of them have been explained by the curvature of the synchrotron spectrum and the presence of the $\\nu F_{\\nu}$ peak (E$_{\\rm{peak}}$) within the XRT band (Falcone et al. 2006, Butler \\& Kocevski 2007, Goad et al. 2007, Mangano et al. 2007, Godet 2007a).  Moreover Butler (2007) found anomalous soft X-ray emission in the spectra of four GRBs which can be interpreted as multiple emission lines due to K shell transition in light metals as well as thermal emission from a blackbody with temperature $\\sim$0.1 keV. \\cite{Grupe07} and Godet et al. (2007a) found that early afterglow data of GRB060729 and GRB050822, respectively, can be fitted with a SPL plus a blackbody with decreasing temperature in the first few hundreds seconds from the beginning of the prompt emission. Campana et al. (2006) found a cooling thermal component in the spectrum of SN2006aj/XRF060218, interpreted as due to the supernova shock break-out (this interpretation has been subsequently questioned by Li 2007 and Ghisellini et al. 2007).  Starting from the idea that any deviation from a SPL spectral model, if present, would be a faint signal mostly covered by the high level ``noise'' of synchrotron emission, we searched for the best observational conditions to detect it. In particular, we searched for high energy excesses with respect to the SPL when the spectrum is steepest and, at least in the hard part of the energy band, could be the non dominating component.  Throughout this paper, all errors are quoted at 68\\% confidence level for one parameter of interest, unless otherwise specified. The reduced $\\chi^2$ will be denoted as $\\chi^2_\\nu$ and the number of degrees of freedom with the abbreviation ``d.o.f.''.  We follow the convention $F_\\nu(\\nu,t) \\propto t^{-\\alpha}\\nu^{-\\beta}$, where $\\alpha$ and $\\beta$ are the temporal decay slope and the spectral index, respectively. As time origin, we will adopt the Burst Alert Telescope (BAT) trigger (T0).  The photon index is $\\Gamma = 1 + \\beta$.  Last, we adopt the standard ``concordance'' cosmology parameters, $\\Omega_{\\rm m} = 0.27$, $\\Omega_\\Lambda = 0.73$, $h_0 = 0.71$. ",
        "conclusions": "Our most solid result is that we found a small homogeneous and fairly defined sample of afterglow spectra (the soft sample) for which deviations from the SPL spectral model are highly probable.  We started from an homogeneous sample of bright GRB afterglows with known redshift and we studied their spectral evolution. We split the data in different time slices and we focused on the softest spectra. In this sub-sample at least 4 cases out of total of 13 present highly significant deviations from the SPL spectral model during the prompt--afterglow transition phase.  We could firmly exclude that these excesses can be explained as statistical fluctuations of a SPL spectrum or some instrumental effects.  We also excluded that data can be fitted by a Band or a cutoff power law model. We fitted these spectra adding one of three different trial components to the SPL model. We did not try to discriminate among these different models on a purely statistical basis, and we discussed them using the component time variability and energetics.  In a very recent paper Yonetoku et al. (2007) show that the very soft spectrum of the early afterglow of GRB 060904A presents a feature which is very similar to what we described in the present work. They selected and stacked the data in the time intervals of the GRB 060904A early afterglow, where the spectral photon index is larger than 4.0 (we note that this GRB is not included in our sample because its redshift is not known). In a very similar way to our results, they found that, in this spectrum, the data show a clear hardening break around 2 keV. This feature leaves a significant excess with respect to the double power law model above 2 keV.  They conclude that this spectrum consists of two emission components, the second being consistent with the spectrum of the late afterglow (our DPL model). These two (independent) studies represent a direct piece of evidence that the emission observed in the early phases of the afterglow is composed by more than one component. Differently from Yonetoku et al. (2007), in our sample, we showed that, if the second component were the emerging afterglow emission, at early time this should be much more luminous than the expectation from the classical afterglow model. We showed that spectral studies of the prompt--afterglow transition phase can be the starting point to seperate different emitting components and could provide useful information in order to better understand the afterglow light curve complexity."
    },
    "0711/0711.4605_arXiv.txt": {
        "abstract": "% We review the evidence for a constant star formation rate per unit mass in dense molecular gas in the Milky Way and the extragalactic correlations of $L_{IR}$ with L$^{\\prime}$ from observations of dense molecular gas. We discuss the connection between the constant SFR/M interpretation in dense gas and the global Schmidt-Kennicutt star formation law. ",
        "introduction": "The connection between star formation studied in the Milky Way and in galaxies is important for understanding the global evolution of stars in the universe.  Lessons learned locally, where angular resolution permits detailed studies of high-mass star forming regions, may be applied to star formation regions that are unresolved in other galaxies. Star formation occurs in dense molecular gas (Evans 1999);  therefore, it is necessary to study the properties of dense molecular gas to understand the conditions in the material actively involved in forming protostars.  An important aspect of this study is the efficiency of star formation or how effectively dense gas is converted into stars.  One method for determining the star formation efficiency in Milky Way molecular clouds is to calculate the ratio of the bolometric luminosity to the mass of a star-forming clump within a giant molecular cloud ($L/M$). Within high-mass star forming regions in the Milky Way, submillimeter single-dish telescopes can resolve individual cluster-forming clumps ($\\theta_{mb} < \\approx 30^{\\prime\\prime}$).  Higher resolution interferometric observations resolve the clumps into individual star-forming cores (e.g. Brogan et al. 2007).  Only a few high-mass clumps have been observed at high resolution. Since several systematic single-dish surveys of clumps have been made (e.g., Plume et al. 1997; Zinchenko et al. 2000; Sridharan et al. 2002), we focus on the properties of the cluster-forming clumps. If we assume a universal stellar Initial Mass Function (IMF) for all star-forming regions within the Milky Way, then the bolometric luminosity is directly proportion to the star formation rate ($L_{bol} \\propto SFR$). The virial mass of a molecular clump is calculated from \\begin{equation} M_{\\rm{virial}} = \\frac{5 R \\Delta v^2}{G \\ln 2} \\frac{a_{\\rm{density}}}{a_{shape}} \\; , \\end{equation} where R is the size of the clump, $\\Delta v$ is the FWHM linewidth, $a_{density} = (1 - 2p/5)/(1 - p/3)$ (for $p < 2.5$) is the correction factor for a power-law density of the form $n \\propto r^{-p}$, and $a_{shape}$ is the correction factor for an ellipsoidal shape (Bertoldi \\& McKee 1989). The size and shape are determined from the FWHM intensity contour of the clump (see Shirley et al. 2003).  An optically thin linewidth must be determined from observations of isotopomers (e.g. H$^{13}$CN, C$^{34}$S, H$^{13}$CO$^+$) since the linewidth of common dense gas tracers can be significantly broadened due to large optical depths (Philips et al. 1979). Radiative transfer modeling of optically thin dust continuum emission indicates that the average large scale structure of dense molecular clumps follows a single power-law density $n \\propto r^{-p}$ (Mueller et al. 2002; Williams et al. 2005), for which the correction to $M_{vir}$ can be substantial ($a_{density} = 0.6$ for a $r^{-2}$ power-law).  \\begin{figure}[!ht] \\plotone{sfrm_plot.eps} \\caption{$L_{bol}/M_{vir}$ is plotted for high-mass star forming clumps observed in various dense molecular gas tracers (top panel).  The average $\\log L/M$ is plotted versus the critical density of the molecular tracer calculated at $40$ K (bottom panel).  The arrows in density correspond to the effective density for which a $1$ K line may be observed in a typical clump (see Evans 1999).  Notice that $\\langle \\log L_{bol}/M_{vir} \\rangle$ is nearly constant for all molecular tracers above $10^4$ cm$^{-3}$.} \\end{figure}  The ratio of $L_{bol}/M_{vir}$ versus $M_{vir}$ is plotted in Figure 1 for mapping surveys of CS J=2-1 and J=5-4 (Wu et al. 2008, in prep.; Plume et al. 1992; Shirley et al. 2003), HCN J=3-2 (Wu \\& Evans 2005), and HCO$^+$ J=3-2 (Shirley et al. 2008, in prep.).  All of these dense gas tracers display a remarkable lack of correlation.  When all of these tracers are considered together, this result is consistent with a constant SFR per unit mass of dense gas.  The average star forming efficiency traced by dense gas is constant among high-mass star forming regions in the Milky Way and this result appears to be independent of the dense gas tracer observed (bottom panel, Figure 1). The critical density of a transition ($n_{crit} = A_{ul}/\\gamma_{ul}$) is crude approximation for determining the density at which a particular transition is excited.  In reality, radiative transfer effects can significantly lower the density at which a particular line may be easily detected ($n_{eff} < n_{crit}$) by up to an order of magnitude (see Evans 1999). The observed $\\langle \\log(L_{bol}/M_{vir}) \\rangle$ is nearly constant for all tracers with $n_{eff} > 10^4$ cm$^{-3}$.  The result from molecular line observations is confirmed by making $L_{bol}/M_{dust}$ comparisons from dust continuum observations.  Since the dust emission on large scales is optically thin at (sub)millimeter wavelengths, the dust mass may be calculated with assumptions of the average dust temperature and the dust opacity (Hildebrand 1983).  For instance, the $350$$\\mu$m radiative transfer modeling survey of Mueller et al. (2002) find a constant $\\log(L_{bol}/M_{dust})$ that is within a factor of two of the $\\log(L_{bol}/M_{vir})$ calculated from dense molecular gas (Shirley et al. 2003). A factor of two is within the uncertainty in dust opacity at long wavelengths in high-mass star forming regions. Thus, there is strong evidence that, regardless of tracer, the SFR/M is constant in dense molecular gas within the Milky Way. ",
        "conclusions": "Galactic observations indicate that the SFR per unit mass traced by dense molecular transitions is constant, independent of tracer above $n_{eff} > 10^4$ cm$^{-3}$. In contrast, extragalactic observations of HI, CO, and optical and near-infrared emission lines indicate that the SFR follows a super-linear Schmidt-Kennicutt law with $SFR \\propto n^{1.5}$.  Recent radiative transfer models of molecular emission on galactic scales indicate that the observed linear correlations observed between $L_{IR}$ and $L^{\\prime}$ for dense gas tracers may be related to an underlying Schmidt-Kennicutt law. Indeed, within the Milky Way, accounting of the gas content (HI + CO) and cluster populations within 1 kpc of the sun reveal that the Schmidt-Kennicutt relationship predicts the observed star formation surface density (Evans 2007).  How do we reconcile the two possibilities for the SFR law (constant vs. super-linear or both)?  In the absence of direct observations of molecular clumps in other galaxies, the problem presents a difficult theoretical hydrodynamic-radiative transfer challenge which we have only begun to explore.  Within the Milky Way, it is still a theoretical challenge to understand which process mitigate the observed constant SFR/M in dense gas (see article by Mac Low in this volume). On galactic scales, one intriguing possibility is that the local Schmidt-Kennicutt law may flatten at high densities such that the SFR $\\propto n$.  This possibility should be explored theoretically. Also, observation of higher excitation lines at high spatial resolution with interferometers (e.g., PdBI, CARMA, and ALMA) are needed to resolve issues of excitation and abundance variations within galaxies.  The synthesis of current observational and theoretical studies of dense molecular gas in the Milky Way and other galaxies is beginning to reveal the global properties of star formation in the gas that is actively involved in star formation."
    },
    "0711/0711.2215_arXiv.txt": {
        "abstract": "Different observation data for cool star---Procyon (Obs\\_IDs of 63, 1461 and 1224) available from {\\it Chandra Data Public Archive} were co-added and analyzed. Emissivities of emission lines of highly charged silicon ions (Si~VII--Si~XII) were calculated over temperatures by adopting the published data of Liang et al. (2007, {\\it Atom. Data and Nucl. Data Tables}, {\\bf 93}, 375). Using the emission measure derived by Raassen et al. (2002, A\\&A, {\\bf 389}, 228), the theoretical line fluxes are predicted, and the theoretical spectra are constructed by assuming the Gaussian profile with instrumental broadening (0.06~\\AA\\,). By detailed comparison between observation and predictions, several emissions lines are identified firstly such as emissions at 43.663~\\AA\\, (Si~XI), 45.550~\\AA\\, (Si~XII), 46.179~\\AA\\, (Si~VIII), 50.874~\\AA\\, (Si~X), 64.668~\\AA\\, (Si~IX), and 73.189~\\AA\\, (Si~VII) etc. Several emission lines are re-assigned in this work, such as the emission line at 52.594~\\AA\\, to Si~X (52.612~\\AA\\,), at 69.641~\\AA\\, to the blending of Si~VII (69.632~\\AA\\,) and Si~VIII (69.664~\\AA\\,) lines, as well as at 70.050~\\AA\\, to Si~VII (70.027~\\AA\\,). The prediction reveals the large discrepancies between the $3s$--$2p$ line (63.715~\\AA\\,) {\\it versus} $3d$--$2p$ line (61.012~\\AA\\,) again for lower charge stage Si~VIII. Solar flare observation is also added for the assessment of present calculation. Different assignments for some lines between Procyon and solar flare, have been found to be due to the hotter emitter in Sun than in Procyon coronae. ",
        "introduction": "Since the launch of {\\it Chandra} and {\\it XMM}-Newton X-ray missions, a large quantity of high quality X-ray spectra with high-resolution has been obtained for nearly all classes of astrophysical X-ray sources~\\cite{BGK00,CHD00,ABG01,BBG01,NMS01,RBK01,FCD04}, which allows detailed plasma diagnostics to be performed for a wide range of celestial objects.  In the spectra of cool stars, rich emission lines of highly charged iron ions were detected due to its high abundance and high effective collecting area (200~cm$^2$ at 8.5~\\AA\\, for HEG/MEG+ACIS-S instrument; 35~cm$^2$ at 10.0~\\AA\\, for LETG+HRC-S instrument\\footnote{http://cxc.harvard.edu/ciao/manuals.html}) between 6---18~\\AA\\,. In this wavelength region, a good agreement between predictions and observations is achieved so far~\\cite{BCK01}, which profits from the Iron-Project~\\cite{HBE93} and earlier running program--Opacity project~\\cite{SYM94}. Moreover, many laboratory experiments have been performed for iron by adopting grating and crystal spectrometers with high-resolution. The laboratory platform can be early vacuum spark, tokamak used for magnetic confinement fusion study, intense laser used for inertial confinement fusion research, and newly developed facility--electron beam ion trap (EBIT, Beiersdorfer 2003). Though the difference of $3s$--$2p$ line intensity {\\it versus} $3d$--$2p$ line intensity of Fe~XVII is a debating problem, a good theoretical performance is obtained when compared with spectra of highly charged sulfur and argon, as revealed by experimental measurements at Livermore EBIT~\\cite{LBB03,LBB05a}.  Besides rich emission lines of highly charged Fe ions, some emission lines of L-shell calcium, argon, sulfur and silicon have been identified for cool stars. Yet detailed analysis for spectra of these ions receive a scant attention due to the low effective area over the wavelength range of 35---100~\\AA\\, spanned by lines of L-shell Si, Ar and Ca ions, and the absence of accurate atomic data of these ions. Several literatures~\\cite{KPW98,DKR99,KST00,LZ06,LZS06c} also reveal the great potentials of L-shell silicon ions. Such as, the line intensity ratio $I$(50.524)/$I$(50.691) of Si~X, $I$(52.306)/$I$(46.391) of Si~XI are good diagnostic methods for the electron density; the line ratio $I$(52.306)/$I$(43.743) of Si~XI shows a good performance for the electron temperature ($T_{\\rm e}$) determination. Similar characteristics were explored for S~X, Ar~XIV and Ca XVI etc.~\\shortcite{KCF93,KPM00,KAW01,KAK03}. Furthermore, spatial information of coronae could be assessed indirectly using correlation ($EM=n_e^2V$) of the electron density with emission measure.  A laboratory measurement for silicon has been performed at the Livermore EBIT-II. Unfortunately, only the spectra between 80---90~\\AA\\, is shown~\\cite{LBB05b}, yet it is enough to demonstrate the large differences between the measurement and theory. Recently, accurate atomic data including energy levels and spontaneous radiative decay rates were calculated for Si~IX---Si~XI by consideration of large configuration interaction and relativistic effects~\\cite{LZZ07}. A self-consistent electron impact excitation rates over a large temperature grids are available, which are based on relativistic distort-wave (RDW) method. A laboratory measurement of soft X-ray spectra (40---180~\\AA\\,) of highly charge silicon ions were also performed by irradiating the silicon target by intense laser beam at Institute of Physics, Chinese Academy of Sciences, which satisfactorily benchmarks the theoretical modelling at high density~\\cite{LZZL07}.  In this work, we calculate the line emissivities of highly charged silicon ions (Si~VI---Si~XII) by adopting new available and accurate atomic data. The available observation data for cool star--Procyon (Obs\\_IDs 63, 1461 and 1224) are co-added and analyzed. By detailed comparison between the predictions and the co-added spectrum. Emissions lines of highly charged silicon are investigated. The observation of solar flare is also summarized for the assessment of the calculation. ",
        "conclusions": "In summary, available observation data (Obs\\_IDs of 63, 1461 and 1224) for cool star--Procyon are co-added and analyzed. By adopting our published atomic data and some unpublished data, line emissivities of Si~VI--Si~XII ions are calculated at an electron density of 3.0$\\times10^8$~cm$^{-3}$ (typical value for cool stars) and temperatures over 0.1--5.0~MK. In this work, energy levels of 878, 312, 560, 320, 350 and 40 have been included for Si~VII--Si~XII, respectively. Moreover, decay rates of transitions (including E1, M1, E2 and M2) up to tens of thousand are included. Some electron impact excitation data are replaced by available $R-$matrix data to take resonant effects into account as far as possible.  Based upon the EM derived by Raassen et al.~\\shortcite{RMA02} and the calculated emissivities, we estimated the theoretical line fluxes of highly charged Si~VII--Si~XII ions. By detailed comparison between the observations and predictions, several emissions lines [with bold font in second or forth column of Table 2] are identified firstly to our best knowledge. The identification is assessed by comparing with solar spectrum with higher resolution (0.02~\\AA\\,). The prediction indicates the large discrepancies between the $3s$--$2p$ line (69.641~\\AA\\,) {\\it versus} $3d$--$2p$ line (61.012~\\AA\\,) again for lower charge stage Si~VIII. This reveals that more accurate calculation of cross section in electron-ion interactions is very necessary. For emission lines at 52.594~\\AA\\, (Si~X), 55.078~\\AA\\, (Si~IX) and 70.027 (Si~VII), this work indicates that the contamination from Ni~XVIII (52.615~\\AA\\,), Mg~IX (55.060~\\AA\\,) and Fe~XII (70.010 and 70.054~\\AA\\, of Fe~XV), respectively, can be negligible for Procyon. However, in solar flare, the contaminations from Mg, Fe and Ni ions become dominant. This is due to the solar flare is hotter than the Procyon coronae. For the peak around 69.641~\\AA\\,, Si~VII line has the comparable contribution with the Si~VIII lines identified by Raassen et al.~\\shortcite{RMA02}. However, their total contribution is less than 25\\% in solar flare. The discrepancies and the identification strongly suggest the benchmark from laboratory measurements at low-density plasmas."
    },
    "0711/0711.0026_arXiv.txt": {
        "abstract": "A mayor problem that arises in the computation of stellar atmosphere models is the self consistent determination of the temperature distribution via the constraint of energy conservation. The energy balance includes the gains due to the absorption of radiation: $\\int \\chi(\\nu) J(\\nu) d\\nu$, and the losses due to emission: $\\int \\chi(\\nu) S(\\nu) d\\nu\\ .$ It is well known that within each one of the two above integrals the part corresponding to spectral ranges whose opacity $\\chi(\\nu)$ is huge can overcome by many orders of magnitude the part that corresponds to the remaining frequencies. On the other hand, at those frequencies where $\\chi(\\nu)$ is very large, the mean intensity $J(\\nu)$ of the radiation field shall be equal, up to many significant digits, to the source function $S(\\nu)$ and consequently to the Planck function $B(\\nu,T)$. Then their net share to the energy balance shall be null, albeit separately their contributions to the gain and loss integrals are the most important numerically. Thus the spectral range whose physical contribution to the overall balance is null will dominate numerically both sides of the energy balance equation, and consequently the errors on the determination of $J(\\nu)$ and $S(\\nu)$ at these frequencies will falsify the balance. It is possible to circumvent the numerical problem brought about by the foregoing circumstances by solving the radiative transfer equation for the variable $I({\\bf{n}},\\nu) - S(\\nu)$, instead of the customary intensity $I({\\bf{n}},\\nu)$.  We present here a novel iterative algorithm, based on iteration  factors already employed by us with success, which makes it possible a fast correction of the temperature by computing directly the difference between the radiative losses and gains at each step of the iterations. ",
        "introduction": "\\label{sec:1}  The distinctive feature of our algorithm for an iterative temperature correction procedure comes from the organization into two fundamental macro-blocks of the ensemble of the mathematical equations which represent the physical phenomena that occur in a stellar atmosphere.  Within each block a self-consistent physical problem is solved by assuming that the current values of some input variables are known, so that the updated values of the output variables of the block can be obtained.  The first macro-block (the {\\it{structural}} one) consists of the equations that express the conservation of mass and momentum. The latter, together with the equation of state, are the backbone of the {\\it{''hard ''}} structure of the model. If we assume that the temperature distribution $T(r)$ is given, we can deal with the equations of this block in a quite easy way, and get the distribution of the pressure $P(r)$ and density $\\rho(r)$, hence the values of all the relevant thermodynamical and transport coefficients. All these values depend of course on the initial temperature distribution assumed.  Once we know the hard structure, namely $P(r)$ and $\\rho(r)$, as well as the ensemble of the thermodynamical and transport coefficients at each point of the physical system, we can tackle the radiative transfer (RT) equations, coupled through the constraint of the conservation of energy, in order to obtain with relative ease a new temperature distribution. These equations, together with the energy conservation constraint, form the second macro-block (the {\\it{energetic}} one). The two macro-blocks are coupled both physically and mathematically {\\it{via}} the equation of state. However this coupling is not very strong, so each block can be treated mathematically in an independent way by means of a sequential procedure that is iterated till the required numerical convergence. According to our experience, such an iterative procedure quickly converges (in about ten iterations), and numerical difficulties never show up. An exhaustive presentation of the algorithm can be found in \\cite{Simonneau88}. ",
        "conclusions": "\\label{sec:4}  We have presented here a novel method for building stellar atmosphere models, more precisely, for the calculations required by the determination of the temperature at each point. Because of radiative transfer one must take into account the coupling among the physical conditions at each and every point of the system. In particular, the temperature at each point is linked with those at all the others. We ideally remove the coupling, and proceed sequentially to get the physical conditions at each point under the assumption that those at all the other points be known. In order to achieve the consistency - {\\it{i.e.}} the coupling - we developed an iterative procedure. Such a treatment constitutes a numerical simulation of the physical processes that are at the origin of the corresponding equations. We tackle and solve those problems that arise when the fluxes to be determined are the difference between quantities (densities) that are very close, so that the difference between their values is less than the numerical accuracy of the computations.  As a proof of the quality of the method proposed, we claim that the integrated radiative flux $H^{\\ i}$, computed at each step of iterations, quickly converges (in a few iterations) to the actual value $H^{\\ast}$. The inaccuracies on the converged numerical values of $H^{\\ i}$ are due to the quadrature of $H^{\\ i}(\\nu)$ over frequencies, inaccuracies that we estimate to be more close to $0.01\\%$ than to $0.1\\%$. As the problem with the numerical quadratures over frequencies is the same as above, we impose that the equality between the two terms of the transcendental equation (\\ref{eq:14}) be numerically accurate to one order of magnitude less, namely to $0.001\\%$, in order not to increase the unavoidable indetermination already brought about by the numerical treatment of the data."
    },
    "0711/0711.2509_arXiv.txt": {
        "abstract": "We seek to improve estimates of the power spectrum covariance matrix from a limited number of simulations by employing a novel statistical technique known as shrinkage estimation. The shrinkage technique optimally combines an empirical estimate of the covariance with a model (the {\\it target}) to minimize the {\\it total} mean squared error compared to the true underlying covariance. We test this technique on N-body simulations and evaluate its performance by estimating cosmological parameters. Using a simple diagonal target, we show that the shrinkage estimator significantly outperforms both the empirical covariance and the target individually when using a small number of simulations. We find that reducing noise in the covariance estimate is essential for properly estimating the values of cosmological parameters as well as their confidence intervals. We extend our method to the jackknife covariance estimator and again find significant improvement, though simulations give better results. Even for thousands of simulations we still find evidence that our method improves estimation of the covariance matrix. Because our method is simple, requires negligible additional numerical effort, and produces superior results, we always advocate shrinkage estimation for the covariance of the power spectrum and other large-scale structure measurements when purely theoretical modeling of the covariance is insufficient. ",
        "introduction": "\\label{sec:intro}  Large-scale structure statistics, especially power spectra, provide precise constraints on cosmological theories. Accurate measurements are now possible with large-volume surveys and advancing computational power. However, the measured power spectrum is not the only required ingredient for estimating cosmological parameters; the covariance matrix also carries a great deal of information that is vital for properly estimating parameter values and their confidence intervals. Observational effects such as the survey geometry, redshift-space distortions, and non-linear clustering make theoretical modeling of the covariance matrix difficult, and often simulations are used to study them in detail. Covariance matrices estimated from a finite number of simulations will contain noise. Cosmological parameter estimation requires the inverse of the covariance matrix to properly weight the measurements. Matrix inversion is an inherently non-linear operation that is sensitive to the noise of all the elements. \\citet{cs06} showed that when the off-diagonal elements of a covariance matrix are excessively noisy it is better for parameter estimation to use a diagonal approximation of the covariance. This reduces the effects of noise, but ignores important information in the covariance.  Covariance matrices for large-scale structure measurements are often estimated using the unbiased empirical covariance matrix, $\\mathbfss{S}$ (see equation~\\ref{eq:cov}), a close relative of the maximum-likelihood estimator, $\\mathbfss{S}^{(ML)} = \\frac {n-1} {n} \\mathbfss{S}$. These estimators work well in the regime where the number of repeat observations, $n$, is much greater than the number of parameters measured for each observation, $p$.  However, in the regimes where $n \\sim p$ or $n \\ll p$ the covariance matrix estimates become ill-conditioned and unstable during inversion, which is necessary for optimal weighting of the data. This is an indication that these estimators do not produce good approximations of the true underlying covariance matrix in these regimes. \\citet{efron82} provides some insight into the difference between maximum-likelihood as a {\\it summarizer} and as an {\\it estimator}.  Maximum-likelihood is an excellent summarizer of data in the sense of trying to represent the important statistical information about a dataset in a small set of numbers. Though maximum-likelihood is asymptotically optimal for estimation in the limit of infinite data, the use of this summary of information for the purpose of making estimates with a finite set of data is not always the best option. \\citet{stein56} proved that one can construct estimators in high-dimensional ($d \\geq 3$) inference problems that outperform maximum-likelihood estimators in the sense of minimizing the {\\it total} mean squared error. Maximum-likelihood produces the best estimates of individual parameters, but the alternatives can often reduce the error on many of the parameters while only slightly increasing the error on a few, resulting in an overall improvement. \\citet{stein56} also showed that the maximum likelihood estimator has the best performance among estimators that transform correctly under translation, implying that any estimator that outperforms maximum-likelihood will necessarily involve an arbitrary choice.  \\citet{ss05} employ a method known as {\\it shrinkage estimation} to construct covariance matrices for functional genomics measurements in the $n \\ll p$ regime.  Their technique optimally combines a high-dimensional estimate that has little or no bias with a low-dimensional estimate that may be biased but has much less variance. The result minimizes the total mean squared error, which is the sum of bias (squared) and variance. They argue that their method can also perform some amount of regularization, resulting in a covariance matrix that has a full set of positive-definite eigenvalues and is well-conditioned (i.e., the ratio of the largest to smallest eigenvalue is not so large that inversion becomes unstable). They employ a lemma from \\citet{lw03} to analytically calculate the optimal linear combination of the low and high dimensional estimates.  In this paper our goal is to provide a simple recipe for using shrinkage estimation to improve the covariance matrix of the matter power spectrum from a limited number of simulations over the ubiquitous sample covariance estimator. Our method aims to reduce the total noise while retaining as much information about real covariance in the simulations as possible. Shrinkage estimation achieves this by optimally combining a theoretical model with the empirical estimate. We will assess the improvements our method offers by examining the performance of the covariance matrices through inversion and use in cosmological parameter estimation. Although we focus on the matter power spectrum, the shrinkage technique is relevant for many studies in large-scale structure and cosmology.  The outline of this paper is as follows: in Section~\\ref{sec:shrink} we introduce shrinkage estimation and describe its application to covariance matrices. Section~\\ref{sec:toy} applies the shrinkage technique to several toy problems before moving to a more complicated example involving galaxy clustering. We describe our technique for measuring matter power spectra from N-body simulations in section~\\ref{sec:sims}. In Section~\\ref{sec:results} we construct several estimates of the power spectrum covariance matrix and compare their performance by estimating cosmological parameters. Finally we review our results, make recommendations, and discuss future directions of this project in Section~\\ref{sec:disc}. ",
        "conclusions": "\\label{sec:disc}  We have introduced shrinkage as a technique for improving estimates of the covariance matrix for power spectrum measurements. We tested our methods on dark matter simulations and showed improvement over the empirically estimated covariance matrix from a limited number of simulations or jackknife resamplings. In order to clearly assess the potential improvement from using shrinkage estimation, we chose an intentionally difficult scenario where traditional methods of estimating the covariance were unlikely to yield satisfactory results. All of these methods would perform better if we allowed ourselves more simulations per Monte Carlo estimate or if we did not push as far into the non-linear clustering regime.  The shrinkage technique would still outperform the other methods, but perhaps the differences would be less obvious.  A good estimate of the covariance matrix of a power spectrum measurement is essential for extracting cosmological information via parameter fitting.  Including the covariance between different bins is a good step towards properly estimating the confidence intervals on cosmological parameters.  However, the increased number of free parameters of a full covariance estimate (as opposed to a diagonal approximation) can cause the covariance estimate to be noisy if only a relatively small number of simulations are available.  This noise can adversely affect the estimate of the parameter itself.  A diagonal approximation to the covariance can be more easily constrained with a limited number of simulations, leading to better estimates of the parameter values.  However, the confidence intervals can be severely underestimated if actual covariance is ignored.  Neither alternative is appealing.  If a similar measurement was performed with the two-point correlation function, the Fourier dual of the power spectrum, a full covariance matrix is especially important as bins will be strongly correlated, even in the linear clustering regime. Realistic survey geometries will also cause additional covariance on large scales for the power spectrum.  Shrinkage estimation is an optimal way of combining a model with many degrees of freedom and a model with few degrees of freedom to minimize the total error on the covariance estimate. In our example the shrinkage versions of the Monte Carlo and jackknife estimators clearly outperformed their counterparts without shrinkage, with the shrinkage version of the Monte Carlo estimator producing the best results. The lemma of \\citet{lw03} as employed by \\citet{ss05} allows a mathematically and numerically simple way of calculating the optimal shrinkage intensity. This means that there is minimal addition work required to use a shrinkage version of a covariance estimator. Shrinkage estimation can result in a massive improvement in the limit of a small number of simulations and will not adversely affect the covariance estimate in the limit of a large number of simulations. For these reasons we always recommend the use of the shrinkage versions of covariance estimators in all regimes.  We briefly investigated the effects of shrinkage estimation in the limit of a large (though not infinite) number of simulations. We applied shrinkage estimation to our reference covariance matrix estimated from all 4096 sub-volumes using the target from equation~\\ref{eq:target} and found an optimal shrinkage intensity $\\hat{\\lambda}^{\\star} = 0.0096$.  This number is the same order as the relative noise we expect in each element of the matrix, $1/\\sqrt{4096} = 0.0156$.  We then calculated the eigensystems of both matrices.  The dot products of the corresponding eigenvectors always exceeded $0.996$, indicating that they are essentially identical.  The (sorted) eigenvalue spectra are shown in Fig.~\\ref{fig:eigen}.  The eigenvalues are the same to within $1\\%$ for the first 10 eigenmodes.  After the tenth eigenmode the eigenvalues from the reference matrix become increasingly smaller compared to the shrinkage version.  By the final eigenmode the difference is $\\sim 50\\%$.  The shrinkage version of the reference matrix should be a more accurate estimate of the true underlying covariance matrix. The non-linear nature of matrix inversion can cause errors $\\gg 1\\%$ even when individual elements of the covariance matrix are estimated to $\\sim 1\\%$. We ran our parameter estimation test using the shrinkage version of the reference matrix and found that $\\langle \\hat{\\sigma}_8 \\rangle$ moved by less than $0.5\\%$. This is small compared to the width of the distribution, which is $\\sim 5\\%$. The average minimum $\\chi^2$ did improve from $52.4$ to $41.1$ with the shrinkage version of the covariance matrix, though this is still large for $18 - 1 = 17$ degrees of freedom. The remaining discrepancy is dominated by bias from problems with modeling the power spectrum into the non-linear regime or power loss in the simulation at smaller scales due to low resolution, not a grossly inaccurate estimate of the variances. The amplitude is mainly sensitive to smooth eigenmodes, which have large eigenvalues, so there is little change in the estimated value. Parameters that are more sensitive to the shape of the power spectrum may be more sensitive to the lower eigenvalue modes and show more than a $1\\%$ change. The impact of these differences could be estimated with a study of the information content of the power spectrum covariance in terms of cosmological parameter confidences (i.e., \\citealt{ns07}), but this is beyond the scope of this paper.    \\figeigen    We employed a very simple diagonal target matrix in this paper, but better targets can clearly improve the efficiency of the shrinkage technique.  A much more realistic model for covariance on small scales could be constructed using the halo model.  For realistic measurements it may also be advantageous to model some of the effects of survey geometries, redshift-space distortions, and clustering bias. Targets that depend on a small number of free parameters may be very useful for some of these effects (e.g., clustering bias). Targets can also be developed for a wide range of large-scale structure measurements in addition to the power spectrum. The exploration of more sophisticated targets is beyond the scope of this paper and is left to future studies.  Ultimately we would like to develop more diagnostics of the performance of our covariance estimates.  Fig.~\\ref{fig:lambda} shows the estimated error bar, $\\Delta$, as a function of the shrinkage intensity, $\\hat{\\lambda}^{\\star}$, for the shrinkage versions of the Monte Carlo and jackknife estimators.  There is clearly some correlation for the shrinkage version of the Monte Carlo estimator, so knowledge of $\\hat{\\lambda}^{\\star}$ could help one gauge how much the error bars are underestimated. The exploration of such diagnostics should proceed as better targets are developed.    \\figlambda    The difficulties in estimating the power spectrum covariance matrix in the context of making precision cosmological measurements are of even greater concern for higher-order clustering measurements.  Higher-order clustering measurements have a configuration space with more degrees of freedom than the power spectrum (or two-point correlation function).  Even a lower resolution measurement will have more bins and a much larger covariance matrix, and noise will cause larger deviations in the inverse matrix.  Theoretical modeling of the covariance matrix for an N-point correlation function generally involves correlations up to the 2N-point \\citep[e.g.,][]{higher}, making the models more uncertain. The ability to optimally combine simulations and a theoretical model with a small number of free parameters will make dramatic improvements.  Shrinkage estimators could also be used for covariance matrices of measurements outside of large-scale structure, including the cosmic microwave background power spectrum. Finally, we note that Section~\\ref{sub:shrink} makes no specific references to covariance matrices and that shrinkage is a general estimation technique. We are studying additional applications of shrinkage estimation for cosmological measurements."
    },
    "0711/0711.2989_arXiv.txt": {
        "abstract": "Pairwise forces between particles in cosmological N-body simulations are generally softened to avoid hard collisions.  Physically, this softening corresponds to treating the particles as diffuse clouds rather than point masses. For particles of unequal mass (and hence unequal softening length), computing the softened force involves a nontrivial double integral over the volumes of the two particles.  We show that Plummer force softening is consistent with this interpretation of softening while spline softening is not.  We provide closed-form expressions and numerical implementation for pairwise gravitational force laws for pairs of particles of general softening scales $\\varepsilon_1$ and $\\varepsilon_2$ assuming the commonly used cloud profiles: NGP, CIC, TSC, and PQS. Similarly, we generalize Plummer force law into pairs of particles of general softenings.  We relate our expressions to the gaussian, Plummer and spline force softenings known from literature.  Our expressions allow possible inclusions of pointlike particles such as stars or supermassive black holes. ",
        "introduction": "Gravitational softening is a building block in the foundation of any modern cosmological N-body simulation; at the same time it has traditionally been given relatively little coverage in literature. In this paper we give this problem a share of our dedicated attention.  Any practical numerical simulation will have finite dynamic range, limiting the resolution with which the system of interest may be studied.  Oftentimes, the resolution requirements for simulations are not completely uniform across the simulation volume: there may be special regions of interest, which we would like to study at high resolution, whereas the remainder of the simulation volume need not be simulated in great detail.  For example, we might be interested in simulating individual dark matter halos embedded within realistic, time-varying cosmological environments (e.g.  \\cite{Diemand05}, \\cite{Diemand07}).  Given finite computational resources, an efficient strategy is to employ adaptive resolution, i.e.\\ to use high resolution in the regions of interest and low resolution elsewhere \\cite{Bertschinger01}.  Two types of resolution that arise in N-body simulations are mass resolution and force resolution.  The mass resolution is simply limited by the finite number of particles; increasing the particle number within some region increases its mass resolution.  Force resolution is limited by softening of pairwise forces between particles.  For simulations of collisionless systems, pairwise forces must be softened at short distances to avoid artificial collisionality due to finite particle number (e.g. \\cite{white78}, \\cite{power03}, \\cite{lukic07} and references therein).  For example, Plummer softening corresponds to calculating the force between two particles separated by $\\vec{r}$ as \\begin{equation}\\label{eq_plummerforce} \\vec{F}(\\vec{r})=-\\frac{m_1 m_2}{(|r|^2 + \\eps_P^2)^{3/2}}\\vec{r}, \\label{plummer} \\end{equation} where $\\eps_P$ here is the Plummer softening length. Other softening laws are commonly used in the literature, with their own corresponding softening lengths. Increasing the force resolution therefore means decreasing the force softening length.  This is allowed within regions of high mass resolution where the particle number is enhanced; elsewhere, the force softening length must remain large to avoid collisionality.  One common choice is to scale the softening length as the cube root of the mass, $\\eps\\propto m^{1/3}$, which holds fixed the maximal density of all particles.  In Eqn.~(\\ref{plummer}), one subtlety that arises when computing the pairwise force between particles of different mass is that the appropriate choice of softening length becomes uncertain: should the force be softened by the larger $\\eps$, the smaller, or some average of the two?   The chief requirements are that the pairwise forces be symmetric in order to conserve momentum (i.e.\\ the same $\\eps$ is used for $F_{12}$ as for $F_{21}$); and that the softening length is not too small, in order to avoid hard collisions. For example, taking $\\eps={\\rm max}(\\eps_1,\\eps_2)$ satisfies both of these conditions.  This choice is clearly not optimal, however, since the effective force resolution is degraded.  In addition to basic conservation laws, the consistency of computed forces with gravity is also important. Consider a set of point masses connected by massless springs for example: such system clearly satisfies both energy and momentum conservation laws but its evolution is inconsistent with gravity. This illustrates that in setting up the effective softening one should also ensure the consistency of computed forces with gravity.  The uncertainty in the appropriate softening method may be resolved by resorting to a well-known physical interpretation of force softening.  The appropriate physical model is that simulation particles represent not point masses, but rather spatially extended clouds of finite density.  This picture naturally leads to finite maximal pairwise forces : if $\\rho\\rightarrow\\rho_{\\rm max}$ towards the center of a cloud, then the force exerted on a test particle will reach a maximum and eventually vanish as the test particle approaches the interior and then the center of the cloud. Of course, to be self-consistent within this physical picture, we should not compute pairwise forces between particles treating one particle as an extended cloud and the other as a point particle, but rather treating both particles as extended.  The pairwise interaction potential then becomes \\begin{equation}\\label{correct} U_{12} = -\\int \\frac{\\rho_1(\\vec{x_1})\\rho_2(\\vec{x_2})}{|\\vec{x_1}-\\vec{x_2}|}\\; d^3x_1 d^3x_2, \\end{equation} and the force is obtained by differentiation of the potential with respect to separation between the centroids of clouds~1 and 2.   Equation~(\\ref{correct}), while clearly the correct choice of force softening, may appear daunting to evaluate.  Pairwise force calculation is often the limiting step in N-body simulations, and so an expression for the pairwise force that requires few floating point operations in its evaluation is an absolute requirement.  The 6-dimensional integral in Eqn.~(\\ref{correct}) would hardly appear promising in this regard, and so a simpler choice of softening, although less physically self-consistent, may seem more appealing. However, we show in this paper that (perhaps surprisingly), the potential in Eqn.~(\\ref{correct}) is in fact analytic for commonly used softening kernels. For the popular ``spline'' density kernel \\cite{Monaghan85}, the potential in Eqn.~(\\ref{correct}) may be expressed with rational functions, while for Plummer force kernel the integral can be done numerically. This allows the efficient calculation of self-consistent forces between particles of unequal mass using our proposed W-shape and the extended Plummer force laws.  Our force profiles can also be effectively used for particle splitting methods \\cite{mes05} applied to pure gravity, or simulation of dark matter including pointlike objects such as super-massive black holes or stars.  Section~\\ref{sec:review} gives general expressions for pairwise gravitational potential and force given the specified density shape for particles in the pair. A numeric integration is necessary for Plummer softening discussed in section~\\ref{sec_plummer}. Closed form analytic solution for standard $W_n$ shapes for~$n=1,2,3,4$ is presented in section~\\ref{sec:softening}. As $n\\rightarrow \\infty$, this solution asymptotically approaches gaussian softening discussed in section~\\ref{sec_gaussian}.  In section~\\ref{sec_relat} we compare all the discussed softening methods and relate them to each other. We conclude in section~\\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  An N-body simulation is a simulation of massive particles under the influence of gravity.  In cosmological simulation interparticle force is softened to avoid hard collisions. Simulation particles have often been viewed as particles of fixed density shapes in previous literature (e.g. expression for $\\rho(r)$ in Eqn.~(5) of \\cite{Springel05}).  To ensure that the nature of interaction between particles in simulations remains gravitational we adpot this interpretation in which force softening corresponds to interaction between cloudlike particles of a fixed density shape.   We prove that the Plummer interparticle force softening is consistent with this force softening model and generalize the Plummer force law into the case of unequal softening scales $\\eps_1$ and~$\\eps_2$.  Interestingly, we mathematically prove that the previosuly used in literature (e.g.  \\cite{Springel05}) method of {\\it spline softening} is inconsistent with gravitational interaction in the sense that there is no possible solution for the density shape~$\\rho(r)$ that for a pair of identical particles of this shape would under gravitational interaction give us the force law used for interparticle force law in spline softening.  We provide the closed form solution for potential and force laws between widely known $W_n$-cloud shapes of general softenings~$\\eps_1$ and $\\eps_2$.  Our generalized interparticle force laws should be useful for N-body simulations with adaptive mass refinement, in which particles of different mass interact gravitationally.  Examples include simulations of the first collapsed objects in the universe \\cite{Diemand05}, the ``Via Lactea'' simulation \\cite{Diemand07}, pure gravity extension of particle splitting~\\cite{mes05}, and simulations of dark matter that includes pointlike objects, e.g.  stars or supermassive black holes.  As an easy way to extend the force law~$f(\\eps,r)$ into pairs of particles of unequal softenings~$\\eps_1$ and $\\eps_2$ one would plug their symmetric combination of ~$\\eps_1$ and $\\eps_2$ as the softening scale~$\\eps$. We found that using the simple gaussian motivated recipe~$\\eps = \\sqrt{(\\eps_1^2+\\eps_2^2)/2}$ for simulations with adaptive mass resolution generally leads to significant systematic uncertainties for the pairs of particles of unequal softenings. The overall effect of these errors can be established by numerical convergence tests similar to \\cite{Heitmann05}. However the problem is immediately resolved by using our proposed profiles.  For convenience, we also made a numeric implementation of our analytic expressions.  The most efficient way to use these laws in an N-body simulation would be to pre-compute the interaction law once before running the simulation, and then to interpolate by look-up from the precomputed table. The {\\tt GRACOS} package \\cite{sb05} will implement these profiles for adaptive softening."
    },
    "0711/0711.2171_arXiv.txt": {
        "abstract": "{ I present here our project, the Bologna Open Clusters Chemical Evolution (BOCCE) project, aimed at using Open Clusters as tracers of the disk properties and their evolution with time. We are collecting and homogeneously analyzing data, both photometric and spectroscopic, on a large sample of open clusters, representative of the old cluster population, and I show here results obtained on a subset of our clusters. ",
        "introduction": " ",
        "conclusions": "To briefly summarize : \\begin{itemize}  \\item we are studying (old) open clusters as tracers of the disk properties;  \\item to this end we are deriving ages, distances, reddenings, metallicities and detailed abundances for a large sample of old OCs (about 20 already analyzed photometrically, and about 10 spectroscopically, up to now);  \\item in doing so, we try to maintain the maximum homogeneity of methodology, to reach really sound and significant results;  \\item if we compare them to field stars, the open clusters do indeed appear to be good tracers of the general disk abundances, confirming that we may use them to trace the disk properties;  \\item abundances of open and globular clusters and field stars  studied by our group will be on a common scale, ensuring meaningful comparisons between different stellar populations; \\item finally, our sample can be used to test stellar evoutionary models of different ages and metallicities, being complementary to the older, metal-poorer globular clusters.  \\end{itemize}"
    },
    "0711/0711.0342_arXiv.txt": {
        "abstract": "% We use PostgreSQL DBMS for storing XML metadata, described by the IVOA Characterisation Data Model. Initial XML type support in the PostgreSQL has recently been implemented. We make heavy use of this feature in order to provide comprehensive search over Characterisation metadata tree. We built a prototype of the Characterisation metadata query service, implementing two access methods: (1) HTTP-GET/POST based interface implements almost direct translation of the query parameter name into XPath of the data model element in the XML serialisation; (2) Web-Service based interface to receive XQuery which is also directly translated into XPath. This service will be used in the ASPID-SR archive, containing science-ready data obtained with the Russian 6-m telescope. ",
        "introduction": "Storing and querying structured metadata is an important point for building astronomical archives containing heterogeneous datasets coming from various telescopes and/or instruments. IVOA Characterisation Data Model (McDowell et al., in prep.) allows to describe a position of any astronomical dataset in the multidimensional space of physical parameters. Thus, building query interface on top of the characterisation metadata, will provide enough capabilities for elaborated queries often needed for sophisticated scientific usage of the resource. ",
        "conclusions": ""
    },
    "0711/0711.0518_arXiv.txt": {
        "abstract": "The Cosmic Microwave Background fluctuations provide a powerful probe of the dark ages of the universe through the imprint of the secondary anisotropies associated with the reionisation of the universe and the growth of structure. We review the relation between the secondary anisotropies and and the primary anisotropies that are directly generated by quantum fluctuations in the very early universe. The physics of secondary fluctuations is described, with emphasis on the ionisation history and the evolution of structure. We discuss the different signatures arising from the secondary effects in terms of their induced temperature fluctuations, polarisation and statistics. The secondary anisotropies are being actively pursued at present, and we review the future and current observational status. ",
        "introduction": "In the post WMAP era for Cosmic Microwave Background (CMB) measurements and in preparation of Planck and the post-Planck era, attention is now shifting towards small angular scales of the order of a few arc-minutes or even smaller. At these scales, CMB temperature and polarisation fluctuations are no longer dominated by primary effects at the surface of last scattering but rather by the so-called secondary effects induced by the interaction of CMB photons with the matter in the line of sight.  Current and future CMB experiments have two main goals: i) measuring small angular scale temperature fluctuations (below a few arc minutes), and ii) measuring the CMB polarisation power spectrum. These goals are fundamental for our understanding of the universe.  The small-scale anisotropies are directly related to the presence of structures in the universe whereas the two types of polarisation ($E$ and $B$-modes, which we discuss later) probe both the reionisation of the universe, i.e. the formation of the first emitting objects, and the inflationary potential. In this review, we focus on the end of the dark ages and the astrophysical probes of reionisation.  There have been rapid and important advances in the recent past. We already have on the one hand measurements, by ACBAR, CBI, BIMA, VSA, of the temperature power spectrum for $2000<\\ell<4000$ with CBI and BIMA data showing an excess of power as compared with the predicted damping tail of the CMB (Figure ~\\ref{fig:high-l}). On the other hand, DASI, Archeops, Boomerang, Maxipol, CBI, QUaD and WMAP have direct measurements of the $E$-mode polarisation. The situation will change even more in the near future with anticipated results from experiments currently taking data or in preparation (QUaD, BICEP, EBEX, CLOVER, QUIET, SPIDER, Planck).  \\begin{figure} \\epsfxsize=12cm \\epsfysize=10cm \\hspace{2cm} \\epsffile{high-l.eps} \\caption{From Spergel et al. (2006): The compilation of the small scale CMB measurements from ground-based and balloon experiments (Ruhl et al.  2003, Abroe et al. 2004, Kuo et al. 2004, Readhead et al. 2004, Dickinson et al. 2004).  The red, dark orange and light orange lines represent the predictions from the $\\Lambda$CDM model fit to the WMAP data for the best fit, the 68\\% and 95\\% confidence levels respectively. Excess of power is seen at the largest $l$ values.} \\label{fig:high-l} \\end{figure}   All of this experimental activity is motivated by what now amounts to the standard model of cosmology. The CMB temperature fluctuations which are generated prior to decoupling are measured on scales from 90 degrees to several arc minutes. This has led to a model of precision cosmology. The basic infrastructure is the Friedmann-Lemaitre model with zero curvature, a cosmological constant (or dark energy), a baryonic content and non-baryonic dominant cold dark matter component (with fractions given by the recent WMAP data (Spergel et al. 2007)). Superimposed on the cosmological background are the primordial adiabatic density fluctuations, described by a nearly scale-invariant power spectrum $|\\delta_k|^2 \\propto k^{n-1},$ at horizon crossing (in the comoving gauge), that generated the large-scale structure via gravitational instability of the cold matter component. However it has become increasingly apparent that to further refine these parameters, and to face the more intriguing challenge of establishing possible deviations from the concordance model one has to address the degeneracies between cosmological parameters with those from the secondary anisotropies as well as the extragalactic astrophysical foregrounds.  The primary CMB anisotropies are due to the gravitational redshift at large angular scales (Sachs \\& Wolfe 1967) and to the evolution of the primordial photon-baryon fluid evolution under gravity and Compton scattering at lower scales (Silk 1967, Peebles \\& Yu 1970, Sunyaev \\& Zel'dovich 1970) to which one adds photon diffusion damping at small scales (Silk 1967). Primary fluctuations have provided us with an unparalleled probe of the primordial density fluctuations that seeded large-scale structure formation.  Indeed on large angular scales, greater than the angular scale subtended by the sound horizon at recombination, one can directly view the approximately scale-invariant spectrum of primordial quantum fluctuations.  On their way towards us, the photons interact with cosmic structures and their frequency, energy or direction of propagation are affected. These are the secondary effects that involve the density and velocity fields and incorporate Compton scattering off electrons. This review is devoted to a study of these secondary anisotropies.  The CMB photons we observe today have traversed the universe from the last scattering surface to us and have thus interacted with matter along their path through the universe. These interactions generate the secondary anisotropies that arise from two major families of interactions. The first family includes the gravitational effects (Figure \\ref{fig:sectot} panel a), including gravitational lensing, the Rees-Sciama effect (RS), moving lenses and decaying potentials usually referred to as the integrated Sachs-Wolfe effect (ISW). These anisotropies arise from the interactions of the photons with gravitational potential wells. The second family incorporates the effects of scattering between CMB photons and free electrons (Figure \\ref{fig:sectot} panel b) such as inverse Compton interaction (the Sunyaev-Zel'dovich (SZ) effect) and velocity-induced scatterings such as the Ostriker-Vishniac (OV) effect and inhomogeneous reionisation.   We define secondary anisotropies in the CMB to include all temperature fluctuations generated since the epoch of matter-radiation decoupling at $z\\sim 1100.$ The following contributions may be distinguished.  \\begin{enumerate} \\item {} The integrated Sachs-Wolfe (ISW) effect is due to CMB photons traversing a time-varying linear gravitational potential. The relevant scale is the curvature scale freeze-out  in concordance cosmology: the horizon at $1+z\\sim (\\Omega_\\Lambda/\\Omega_{\\rm m})^{1/3}$. This corresponds to an angular scale of about 10$^\\circ$. \\item {} The Rees-Sciama (RS) effect is due to CMB photons traversing a non-linear gravitational potential,  usually associated with   gravitational collapse. The relevant scales are those of galaxy clusters and superclusters, corresponding to  angular scales of 5-10 arc minutes. \\item {} Gravitational lensing of the CMB by intervening large-scale structure does not change the total power in fluctuations, but power is redistributed preferentially towards smaller scales. The effects are significant only below a few arc minutes. Its effects may be significant on large scales when the observable of interest is the $B$-mode power spectrum. \\item {} The Sunyaev-Zel'dovich (SZ) effect from hot gas in clusters is due to the first order correction for energy transfer in Thomson scattering.  It is on the scale of galaxy clusters and superclusters, although it may be produced on very small scales by the first stars in the universe.  There is a spectral distortion, energy being transferred from photons in the Rayleigh-Jeans tail of the cosmic blackbody radiation to the Wien tail. \\item {} The kinetic Sunyaev-Zel'dovich effect is the Doppler effect due to the motion of hot gas in clusters that scatters the CMB. It causes no spectral distortion. \\item {} The Ostriker-Vishniac (OV linear) effect is also due to Doppler boosting. It is the linear version of the kinetic Sunyaev-Zel'dovich effect. It is proportional to the product of $\\Delta n_e$ and $\\Delta v,$. This is effective on the scale of order 1 arc minute. \\item {} Discrete sources provide an appreciable foreground, especially at lower frequencies for radio sources and high frequencies for infra-red and submillimetre sources. \\item {} Polarisation is primarily a secondary phenomenon.  The primary effect from last scattering is induced by out-of-phase velocity perturbations and provides evidence for the acausal nature of the fluctuations. The secondary polarisation is associated with the reionisation of the universe and is on large scales corresponding to the horizon at reionisation. Inhomogeneous reionisation and scattering at the galaxy cluster scale leads to smaller scale polarisation. The reionisation signal is weak, amounting to no more than 10 percent of the primary signal. \\item {} $B$-mode polarisation can be induced by shear perturbations. One source is gravitational lensing of primary CMB fluctuations.  A second is relic gravity waves from inflation. These are pure $B$-modes, and fall off rapidly on scales smaller than the horizon at recombination, corresponding to about half a degree. Mixing by Faraday rotation in the intracluster medium also contributes to $B$-mode generation on small angular scales. The $B$-mode polarisation amplitude only amounts to about a percent of the primary signal, and its discovery will pose the major challenge for future experiments.  \\end{enumerate}  \\begin{figure} \\hspace{-2cm} \\epsfxsize=16.cm \\epsfysize=9cm \\hspace{2.0cm} \\epsffile{second-grav-scat-hudod.eps} \\caption{From Hu \\& Dodelson (2002): The power spectrum of the secondary temperature anisotropies arising from gravitational effects (panel a) and scattering effects (panel b). The power spectrum of primary anisotropies is shown for comparison. The calculations use a flat universe with $\\Omega_\\Lambda = 0.67, \\Omega_{\\rm b}h^2 = 0.02, \\Omega_{\\rm m}h^2 = 0.16, n=1$. Acronyms are defined in the text. $\\delta$- and i-mod refer to density and ionisation fraction modulation respectively (Sect. \\ref{sec:reion}). ``Suppression'' and ``Doppler'' refer to the damping and anisotropy generation at reionisation (Sect. \\ref{sec:reion}). } \\label{fig:sectot} \\end{figure} ",
        "conclusions": ""
    },
    "0711/0711.0204_arXiv.txt": {
        "abstract": " ",
        "introduction": "The philosophy of UNIT is to develop an instrument allowing to obtain observations through the Internet from a PC at any location. UNIT is primarily supposed for professional applications. It can also be employed for educational aims by students via a www gateway. The foundation of this project must be the ideas of high technology, innovative methods of observations and data processing, know-how. Just such approaches allow the project will be able to survive in the conditions of high competition in the world scientific community.  A typical station of the Network is as follows: standard equipment + trained operator + Internet. Correspondents and participants of the Network can be organized in any country of the world, including the Ukrainian station in Antarctic Continent. ",
        "conclusions": ""
    },
    "0711/0711.3800_arXiv.txt": {
        "abstract": "We investigate the effect of X-ray echo emission in gamma-ray bursts (GRBs). We find that the echo emission can provide an alternative way of understanding X-ray shallow decays and jet breaks. In particular, a shallow decay followed by a ``normal\" decay and a further rapid decay of X-ray afterglows can be together explained as being due to the echo from prompt X-ray emission scattered by dust grains in a massive wind bubble around a GRB progenitor. We also introduce an extra temporal break in the X-ray echo emission. By fitting the afterglow light curves, we can measure the locations of the massive wind bubbles, which will bring us closer to finding the mass loss rate, wind velocity, and the age of the progenitors prior to the GRB explosions. ",
        "introduction": "\\label{sec:intro}  Gamma-ray bursts (GRBs) are the most luminous and intriguing explosions in the universe. A number of physical interpretations for GRBs have been proposed since they were discovered about four decades ago. The most successful model for GRBs and their afterglows is the fireball shock model \\citep[and references therein]{m06}. Although the interpretation for the prompt gamma-ray emission is still controversial even in this standard scenario, the long-wavelength afterglows have been widely accepted to be the emission from relativistic shocks sweeping up an external medium \\citep[and references therein]{m06,z07}.  In spite of the success of the standard external shock model for interpreting most features of afterglows, there are still a few issues recently raised by the \\textit{Swift} X-ray Telescope (XRT) \\citep{g04}. The first issue is the ubiquitously detected shallow decay of early X-ray afterglows \\citep{n06,z06}. In general, to account for the shallow decay, a ``refreshed shock\" scenario is assumed, either by a continuous activity of the GRB progenitor \\citep{dl98,zm01,d04}, or by a power-law distribution of the Lorentz factors in the ejecta \\citep{rm98,sm00}. However, although the ``refreshed shock\" model is generally able to explain the flatter temporal X-ray slopes from current analysis, it may not be able to explain self-consistently both the X-ray and the optical afterglows of some well-monitored GRBs, such as GRB 050319 and GRB 050401 \\citep{fp06}.  The second issue is the lack of jet breaks that are supposed to be detected simultaneously in X-ray and optical light curves \\citep{p06,s07,br07}. This raises a serious concern about our understanding of the afterglow emission. Currently, possible solutions to this problem are additional mechanisms taking place in the early X-ray emission \\citep{fp06}, X-ray emission arising from different outflows with the optical emission \\citep{p06}, or alternatively, the temporal break in X-rays is masked by some additional source of X-ray emission \\citep{s07}, so that a jet break at the expected time could still be seen in the optical band. The last interpretation may bring renaissance to the physical meaning of a jet break \\citep{r99} and $E_{\\rm jet}$ \\citep{f01}, and therefore revive the relation between $E_{\\rm jet}$ and $E_{\\rm peak}$ \\citep{g04a}.  Recently \\cite{sd07} realized that the impulsive prompt X-ray emission scattering off the surrounding medium could give rise to a long-term X-ray flux, which is comparable to that of most observed afterglows. The scattered X-ray emission (which is called echo emission herein) could be an additional contributor to the early X-ray afterglow, which was originally put forward by \\cite{k98} and then further discussed by \\cite{mg00} and \\cite{ss03}. \\cite{sd07} also worked out detailed light curves of the echo emission that could nicely reproduce most of the temporal features of observed X-ray afterglows, i.e., a shallow decay followed by a ``normal\" decay and a further rapid decay \\citep{z06,n06}. This scenario provides a self-consistent solution to both the shallow decay and the jet break problem. By fitting the X-ray light curves of GRB 060813 and GRB 060814, \\cite{sd07} concluded that the scattering medium, i.e., a dusty shell should be at a distance of about tens of parsecs from the GRB source.  Coincidentally, this conjectured dusty shell by \\cite{sd07} is consistent with the massive wind bubble around a Wolf-Rayet (WR) star, which has been proposed as the progenitor of a GRB. The study of stellar structure and evolution has provided a mass-loss history of a massive star, starting with the Main Sequence (MS) stage as an O star, which will evolve into a WR star and then produce a final supernova and/or a GRB. By the computation of stellar wind and atmosphere models for a O-type star of 60 ${\\rm M_\\sun}$ mass, \\cite{l94} concluded that the stellar mass lost to the wind bubble is about 32 ${\\rm M_\\sun}$ in the MS stage, 8 ${\\rm M_\\sun}$ in the Luminous Blue Variable (LBV) stage, and 16 ${\\rm M_\\sun}$ in the WR stage. Such a massive wind bubble will produce a significant shell mainly swept up by the MS wind at the time of the supernova or GRB (see Fig.\\ref{fig:bubble}), with a distance from the progenitor of about $R\\sim$ tens of parsecs \\citep{g96a,m03,d07}. If the clumps formed during the LBV and WR stages survive long enough, they will be hitting the MS shell roughly at the time of the supernova or GRB \\citep{g96a}.  In this paper, we apply the echo emission model to a sample of 36 GRBs with well-monitored X-ray light curves by the XRT onboard \\textit{Swift}. We find that most of the light curves without significant flares can be reproduced by our model. In this sample, 14 bursts have available redshifts by observations. By fitting their X-ray light curves, we find the distance $R$ of the wind bubbles from GRB progenitors. A measurement of $R $ might bring us closer to finding the mass loss rate, density and wind velocity of the progenitor prior to the explosion.  Our paper is structured as follows: In $\\S$\\ref{sec:bubble}, we first discuss the existence of massive stellar wind bubbles around GRBs when they explodes and the subsequent X-ray echo emission. We formulate the theoretical light curves of X-ray echo emission in $\\S$\\ref{sec:theory}, and we introduce the temporal break in echo emission and fit the observational light curves of 36 GRBs in $\\S$\\ref{sec:break}. Discussion and conclusions follow in $\\S$\\ref{sec:discuss}. ",
        "conclusions": " \\begin{enumerate} \\item The shallow decay phase can still be understood without additional contributions from the central engine \\citep{z06,l07a}. Interestingly, the feature of a shallow decay followed by a ``normal\" decay and a further rapid decay can be reproduced by the echo from the prompt X-ray emission scattered by the massive wind bubble around the progenitor \\citep{g96a,m03,d07}. In general, dust formation seems to be prevailing in carbon-rich WR stars of the WC subtype and there is where echo emission would be most efficient.  \\item Echo emission can be an additional X-ray source that masked the X-ray jet breaks \\citep{s07}. Therefore, the lack of jet breaks \\citep{p06,br07,p07,l07b} might also be understood via competing echo emission. In addition, some results in the pre-\\textit{Swift} era, say, some correlations relevant for optical breaks might still be valid \\citep{f01,g04a,lz05} .  \\item There is an extra temporal break in X-ray echo emission introduced by $t_{\\rm c}$ above. By fitting the light curves, we can indirectly measure the locations of the massive wind bubbles,  which will bring us closer to finding the mass loss rate, wind velocity, and the age of the progenitors prior to the GRB explosions \\citep{g96a,d07}.  \\item Even though some features in X-ray afterglows would be better understood if the echo emission is taken into account, the external shock model is still the dominant scenario for optical/infrared afterglows, since echo emission is negligible in the later. In addition, the echo emission model does not completely rule out other interpretations (energy injection, etc.) in the framework of the fireball-shock model. Technically, echo emission takes place at a distance of $\\sim$ tens of parsecs, while the internal or/and external shocks are believed to take place within a distance of $\\sim 10^{17}$ cm. It is more likely that both of them are taking place in the X-ray afterglow, which potentially complicates the light curves. However, this may raise the problem about GRB radiative efficiencies again. Even by assuming that the X-ray emission is the forward shock emission, the GRB efficiency is already too high to be accommodated by the leading internal shock models \\citep[and references therein]{z07a}. Assuming an echo-dominated X-ray emission would exacerbate the situation. \\end{enumerate}"
    },
    "0711/0711.4757_arXiv.txt": {
        "abstract": "We present new high-resolution near-IR spectroscopy and OH maser observations to investigate the population of cool luminous stars of the young massive Galactic cluster RSGC1. Using the 2.293\\micron\\ CO-bandhead feature, we make high-precision radial velocity measurements of 16 of the 17 candidate Red Supergiants (RSGs) identified by Figer et al. We show that F16 and F17 are foreground stars, while we confirm that the rest are indeed physically-associated RSGs. We determine that Star F15, also associated with the cluster, is a Yellow Hypergiant based on its luminosity and spectroscopic similarity to $\\rho$~Cas. Using the cluster's radial velocity, we have derived the kinematic distance to the cluster and revisited the stars' temperatures and luminosities. We find a larger spread of luminosities than in the discovery paper, consistent with a cluster age 30\\% older than previously thought (12$\\pm$2Myr), and a total initial mass of $(3\\pm1) \\times 10^{4}$\\msun. The spatial coincidence of the OH maser with F13, combined with similar radial velocities, is compelling evidence that the two are related. Combining our results with recent SiO and H$_2$O maser observations, we find that those stars with maser emission are the most luminous in the cluster.  From this we suggest that the maser-active phase is associated with the end of the RSG stage, when the luminosity-mass ratios are at their highest. ",
        "introduction": "\\label{sec:intro} The Red Supergiants (RSGs) represent a key evolutionary phase in the life-cycle of stars with initial masses of $\\sim$8--30\\msun\\ \\citep[e.g.][]{Mey-Mae00}. Though comparatively brief, the mass-loss rate in this stage can be many orders of magnitude greater than on the main-sequence \\citep[e.g.][]{Repolust04,v-L05}, and the mass lost in the RSG phase can determine the terminal mass of the star, the appearance of the supernova (SN) explosion, and nature of the stellar remnant \\citep[e.g.][]{Heger03}.  The study of RSGs is hampered by low-number statistics -- until recently only $\\sim$200 were known in the Galaxy, and around 40 in the Large Magellanic Cloud (LMC). Further, the majority of these stars are isolated, and hence have a variety of ages, initial masses and metallicities. Ideally, we would like to study large numbers of RSGs {\\it coeval clusters}, where we can be confident that the variables of metallicity and initial stellar mass are fixed.  Two recent discoveries now present us with the opportunity to study statistically-significant numbers of RSGs in such clusters. In \\citet[][hereafter FMR06]{Figer06} the discovery of an apparent cluster of 14 RSGs was presented; while \\citet[][hereafter DFK07]{RSGC2paper} report on the remarkable stellar population of a second cluster in the same region, which it was shown contains 26 RSGs. In order to use these two clusters as testbeds with which to study the pre-SN evolution of massive stars, we must first quantitatively study the properties of the stars within each.  In the case of the second cluster, hereafter RSGC2 (also known as Stephenson~2), DFK07 used high-resolution spectroscopy to obtain accurate radial velocities of many stars in the region of the cluster, and were able to separate proper cluster-members from background/foreground stars. Further, from the radial velocities they made a quantitative discussion of the kinematic distance to the cluster, thus enabling them to determine the stars' luminosities, temperatures, ages and initial masses.  The details of the stellar properties of the first cluster, hereafter RSGC1, are less well constrained. In the discovery paper FMR06 showed with low-resolution spectra that the stars were all late-type; but with no radial velocity information, they relied on the stars' similar near-IR colours to argue that the stars were all at the same distance. They determined this distance by associating a nearby OH maser, detected by \\citet{Blommaert94}, with the RSGs. However this maser was {\\it singly-peaked} -- such masers are formed in the winds of RSGs and hence are typically {\\it doubly-peaked}, with separation twice the outflow speed. If the second peak was missed in the original observations, this would lead to an incorrect radial velocity for the cluster, and hence kinematic distance. Consequently, the cluster membership, as well as the properties of the luminous cool stars of RSGC1 are poorly-constrained when compared to RSGC2.  Here, we study the cool supergiants of RSGC1 with data of similar quality to that presented in DFK07 for RSGC2. We present high resolution spectroscopy of the CO bandhead feature of 16 of the 17 $K$-bright stars in the field identified in FMR06, allowing us to determine accurate radial velocities of the stars, and establish the cluster membership. We also present new observations of the OH maser source OH25.25-0.16, showing that it is indeed doubly-peaked, and that the central radial velocity of the profile is consistent with the average stellar radial velocity. We use this value, in conjunction with contemporary Galactic rotation curve parameters, to reappraise the distance to the cluster, the stars' temperatures and luminosities, and the age and intial mass of the cluster.  We begin in Sect.\\ \\ref{sec:obs} with a description of the observations and data reduction steps, and we describe the results and analysis of the data in Sect.\\ \\ref{sec:res}. In Sect.\\ \\ref{sec:disc} we derive the cluster's age and mass, and compare with similar analyses of RSGC2. Finally, we use the stellar population of RSGC1 to investigate the maser-active phase of RSGs. ",
        "conclusions": "We have presented a comprehensive investigation into the physical properties of the luminous cool stars in the massive cluster RSGC1. Using high-resolution spectroscopy we have accurately measured the cluster's radial velocity, derived its kinematic distance, and reappraised the stars' temperatures and luminosities. We find a larger luminosity spread than in the discovery paper, which is well fitted by a cluster age of 12Myr and cluster mass of $(3 \\pm 1) \\times 10^4$\\msun. The mass is similar to that of the nearby cluster RSGC2, and we suggest that the difference in the number of RSGs in each is due to the separation in cluster ages, with RSGC2 being somewhat older. This implies that the initial masses of the RSGs in each cluster are different, which we determine to be $\\sim$18\\msun\\ for RSGC1 and $\\sim$14\\msun\\ for RSGC2. The clusters therefore allow the study of RSG evolution as a function of initial mass, while constraining the variable of metallicity. Finally, with new 1612MHz radio observations we find compelling evidence that the OH maser is associated with star F13, and collating recent maser observations of the cluster we argue that the maser-active phase is associated with stars in the latter stages of RSG evolution."
    },
    "0711/0711.3946_arXiv.txt": {
        "abstract": "The Solar Optical Telescope (SOT) on board Hinode satellite observed an X3.4 class flare on 2006 December 13. Typical two-ribbon structure was observed, not only in the chromospheric Ca\\emissiontype{II} H line but also in G-band and Fe\\emissiontype{I} 6302 \\AA \\, line. The high-resolution, seeing-free images achieved by SOT revealed, for the first time, the sub-arcsec fine structures of the ``white light\" flare. The G-band flare ribbons on sunspot umbrae showed a sharp leading edge followed by a diffuse inside, as well as previously known core-halo structure. The underlying structures such as umbral dots, penumbral filaments and granules were visible in the flare ribbons. Assuming that the sharp leading edge was directly heated by particle beam and the diffuse parts were heated by radiative back-warming, we estimate the depth of the diffuse flare emission using the intensity profile of the flare ribbon. We found that the depth of the diffuse emission is about 100 km or less from the height of the source of radiative back-warming. The flare ribbons were also visible in the Stokes-V images of Fe\\emissiontype{I} 6302 \\AA \\, as a transient polarity reversal. This is probably related to \"magnetic transient\" reported in the literature. The intensity increase in Stokes-I images indicates that the Fe\\emissiontype{I} 6302 \\AA \\, line was significantly deformed by the flare, which may cause such a magnetic transient. ",
        "introduction": "Solar flares that are visible in optical continuum are called white light flares (WLFs). They have been of interest because the radiative loss by the WL emission is significant, they are similar to stellar flares in many aspects, and their energy transport and emission mechanisms have been challenging problems \\citep{neidig89}. Observers used to believe that WL emissions are seen only in the most energetic flares \\citep{hiei82, neidig89}. However, seeing-free observation from spacecrafts have revealed that even C-class flares show enhancement in optical continuum \\citep{matthews03, hudson06}.  Close correlation in space and time between WL and hard X-ray (HXR)  emissions have been reported in many events, strongly suggesting that WLFs are energized by nonthermal electron beams \\citep{rust75, hudson92, neidigkane93, chending06}. Recently, several authors studied the energetics of WLFs, examining whether nonthermal electron beam had enough energy to power the WLFs \\citep{metcalf03, chending06, fletcher07}. Their results indicate that, for most of WLFs, the energy required to power the WL emission was comparable to the total energy carried by the electron beam, and therefore the WLFs were energized by the low energy electrons,  about 20 keV or less, which carried most of the energy. Such low energy electrons cannot penetrate into the photosphere or lower chromosphere, so they deposit the energy in the upper chromosphere \\citep{hudson72}. Hence, the WL emission may also come from the upper chromosphere, or the energy is transported by radiation into deeper layer at which the WL emission occurs (so called back-warming effect; \\citealt{machado89, metcalf90}).  Morphology also gives some clues to examine the mechanisms of WLFs. High-resolution images  often show separation of a  WLF kernel into a bright inner core and a fainter outer region (halo) \\citep{neidig93, hudson06}. Using high-cadence ground based observation, \\citet{xu06} examined the timing of WL and HXR emission in the core and halo separately. They found that the peak of WL in the halo showed a delay from that of HXR, but they were almost simultaneous in the core. This suggests direct heating by electron beams in the core and indirect heating  by back-warming (or something else) in the halo.  The high-resolution, seeing-free images from the Solar Optical Telescope (SOT; Tsuneta et al. 2007) on board Hinode (Kosugi et al. 2007) provide a qualitatively new data set for studying WLFs. After its launch in 2006 September, Hinode has successfully observed several flares. The X3.4 flare on 2006 December 13 was of particular interest from the view point of WLF, because the SOT observed clear two ribbons not only in the chromospheric Ca\\emissiontype{II} H images but also in the G-band images and  the Stokes I and V images of Fe\\emissiontype{I} 6302 \\AA \\, line. G-band and the Fe\\emissiontype{I} 6302 line are not ``white light\" in a strict sense, but they both originate in the photosphere. Moreover, the flare emission seen in Stokes V images may be related to so called magnetic transients (magnetic anomalies), the transient signals in magnetograms associated with WLFs, whose origin remains uncertain \\citep{patterson81, zirin81, kosovichev01, qiu03}. The aim of this paper is to present the initial SOT observation of the flare ribbons of this event and discuss their possible indications for the mechanisms of WLFs. ",
        "conclusions": "The high-resolution and seeing-free images from the SOT has revealed the very fine structures in a WLF (more precisely, G-band and Fe\\emissiontype{I} 6302 flare). The flare ribbons in G-band and 6302I clearly show two distinct features: bright and sharp leading edges (kernels) and following (surrounding) diffuse part. In the diffuse part, the pre-flare structure such as umbral dots, penumbral filaments and granules remain visible, giving an impression that the diffuse part are optically thin. The maximum contrast in the flare ribbons is about 600\\% for G-band and about 50\\% for 6302I.  Assuming that the leading edge (kernels) is heated directly by nonthermal particles beams and the diffuse parts are heating by radiative back-warming, we made a simple model to derive the depth of the layer of diffuse emission using the intensity profile of the flare ribbon. We found that the layer of diffuse emission is about 100 km below the height of the irradiation source. If the height of initial energy deposition and hence that of the irradiation source is in the upper chromosphere as suggested by previous studies \\citep{metcalf03, chending06, fletcher07}, our result indicates that the height of the diffuse component, which we believe due to the radiative back-warming, is also in the upper or middle chromosphere.  The 6302V images (longitudinal magnetogram) show transient sign reversal, which is probably the same phenomenon known as magnetic transient. The sign reversal was co-spatial and co-temporal with the flare ribbon, and probably due to a significant distortion of the Fe I 6302 line during the flare, though we cannot be conclusive because of the lack of spectroscopic data.  \\bigskip  Hinode is a Japanese mission developed and launched by ISAS/JAXA, collaborating with NAOJ as a domestic partner, NASA and STFC (UK) as international partners. Scientific operation of the Hinode mission is conducted by the Hinode science team organized at ISAS/JAXA. This team mainly consists of scientists from institutes in the partner countries. Support for the post-launch operation is provided by JAXA and NAOJ (Japan), STFC (U.K.), NASA, ESA, and NSC (Norway). This work was carried out at the NAOJ Hinode Science Center, which is supported by the Grant-in-Aid for Creative Scientific Research ``The Basic Study of Space Weather Prediction\" from MEXT, Japan (Head Investigator: K. Shibata), generous donations from Sun Microsystems, and NAOJ internal funding. H.I. is supported by a Research Fellowship from the Japan Society of the Promotion of Science for Young Scientists. He also appreciates the fruitful comments from H. Hudson, S. Krucker, T. Yokoyama, and the anonymous referee."
    },
    "0711/0711.1491_arXiv.txt": {
        "abstract": "{{}We present spatially-resolved, rest-frame optical spectroscopy of a z$\\sim 3$ Lyman-break galaxy (LBG), Q0347-383~C5, obtained with SINFONI on the VLT. This galaxy, among the $\\sim 10$\\% brightest LBGs, is only the second z$\\sim 3$ LBG observed with an integral-field spectrograph. It was first described by Pettini et al. (2001), who obtained WFPC2 F702W imaging and longslit spectroscopy in the K-band. We find that the emission line morphology is dominated by two unresolved blobs at a projected distance of $\\sim 5$ kpc with a velocity offset of $\\sim 33$ km s$^{-1}$. Velocity dispersions suggest that each blob has a mass of $\\sim 10^{10}$ M$_{\\odot}$. Unlike Pettini et al. (2001), our spectra are deep enough to detect H$\\beta$, and we derive star-formation rates of $\\sim 20-40$ M$_{\\odot}$ yr$^{-1}$, and use the H$\\beta$/[OIII] ratio to crudely estimate an oxygen abundance $12+[O/H]=7.9-8.6$, which is in the range typically observed for LBGs.  We compare the properties of Q0347-383~C5 with what is found for other LBGs, including the gravitationally lensed ``arc$+$core'' galaxy (Nesvadba et al. 2006), and discuss possible scenarios for the nature of the source, namely disk rotation, a starburst-driven wind, disk fragmentation, and merging of two LBGs. We favor the merging interpretation for bright, extended LBGs like Q0347-383~C5, in broad agreement with predicted merger rates from hierarchical models.} ",
        "introduction": "\\label{introduction} Our understanding of high-redshift galaxies is growing at a rapid pace. By combining ground and space-based observations, all of the electromagnetic spectrum from the X-rays to the radio is now being used to select and study galaxies in the early universe.  Among the ever-growing number of high-redshift galaxy populations, Lyman Break Galaxies (LBGs) play a dominant role, from an astrophysical and from an observational point of view. Recently, \\citet{reddy07} argued that blue, star-forming galaxies like LBGs may have had the largest contribution to the bolometric energy output of galaxies at redshifts z$\\sim 2-3$, emphasizing their importance for our understanding of the cosmic star-formation history.  Observationally, more than 1000 LBGs have spectroscopic redshifts at z$\\sim 3$ alone \\citep[e.g.,][]{steidel03}, and their ensemble properties, like their spatial clustering, luminosity function, and average rest-frame UV spectral properties are known at unprecedented precision \\citep[e.g.,][]{adelberger05,shapley01,papovich01,shapley03}. Very recently, \\citet{law07} presented adaptive-optics assisted integral-field spectroscopy of the $z= 3.3$ LBG DSF2237a-C2 using OSIRIS on the Keck telescope.  The picture that emerges from these observations is far from simple. \\citet{law07} point out that at resolutions of $\\sim 1$ kpc reached with adaptive optics, UV-selected z$\\sim 2-3$ galaxies have irregular kinematics, which are likely not dominated by large-scale gravitational motion, but perhaps are more related to merging or gas-cooling. \\citet{genzel06,nmfs06} however argue that at least a subsample of blue, star-forming galaxies at somewhat lower redshifts, z$\\sim 2$, may show the signs of large, spatially-extended, rotating disks. Distinguishing between the two scenarios is difficult, due to the low spatial resolution of the data relative to the size of the targets \\citep[see also][]{kronberger07}.  Here we present a study of one of the first z$\\sim 3$ LBGs described in the literature, Q0347-383~C5, which was initially described by \\citet{steidel96b}.  With ${\\cal R}=$23.82 mag, Q0347-383~C5 is within the tail of the $\\sim 10$\\% brightest LBGs.  WFPC2 imaging shows a relatively complex morphology, extending over $\\sim 1$\\arcsec down to the faint surface brightness detection limit of the image (18000s through the F702W filter) and a half light radius about a few tenths of an arc second \\citet{pettini01}.  Extents this large are not rare among the bright LBG population \\citep{conselice03} and its half-light radius is also rather typical \\citep{ferguson04}.  \\citet{pettini01} obtained longslit spectroscopy in the K-band for a small sample of z$\\sim 3$ LBGs, including Q0347-383~C5, which was one out of two of their sources with spatially-extended spectra and a velocity gradient of $\\sim 70$ km s$^{-1}$ in the [OIII]$\\lambda\\lambda$4959,5007 emission line. They placed a 3$\\sigma$ limit on H$\\beta$ of $F(H\\beta) < 1.7 \\times 10^{-17}$ erg s$^{-1}$ cm$^{-2}$.  Using the near-infrared spectrograph SINFONI on the VLT, we obtained deep, spatially-resolved spectroscopy of the [OIII]$\\lambda\\lambda$4959,5007, and H$\\beta$ emission line. This is only the second unlensed LBG with integral-field spectroscopy in the literature with such observations \\citep[][discuss the z=3.2 LBG DSF2238a-C2, for which they obtained rest-frame optical integral-field spectroscopy using OSIRIS on the Keck with a laser guide star.]{law07} Given the small number of LBGs with integral-field spectroscopy, and the large observational expense of such observations, even the study of a single object is already a significant step forward. Such observations overcome many uncertainties related to longslit spectroscopy, such as slit-losses, and allow us to trace the emission line morphology and kinematics across the two-dimensional surface of the target. This makes them particularly suited to disentangle the often complex emission line morphology of high-redshift galaxies.  Throughout the paper we adopt a flat H$_0 =$70 km s$^{-1}$ Mpc$^{-3}$ concordance cosmology with $\\Omega_{\\Lambda} = 0.7$ and $\\Omega_{M} = 0.3$. In this cosmology, D$_L=$ 28 Gpc and D$_A =$ 1.5 Gpc at z=3.23. The size scale is 7.5 kpc arcsec$^{-1}$. The age of the universe at this redshift and cosmological model is 1.9 Gyr. ",
        "conclusions": "\\label{sec:summary} We presented an analysis of rest-frame optical integral-field spectroscopy of the $z=3.23$ Lyman-Break Galaxy Q0347-383~C5 in the K band. This galaxy is one of the largest known LBGs, and in particular large enough for seeing-limited observations. Q0347-383~C5 was first described by \\citet{pettini01}, who obtained F702W HST continuum imaging and longslit spectroscopy in the K-band  We detect the [OIII]$\\lambda\\lambda$4959,5007 doublet with line properties that are similar to those discussed in \\citet{pettini01}, but in addition, we also identify H$\\beta$ with a flux of $9\\times10^{-18}$ erg s$^{-1}$ cm$^{-2}$.  The [OIII]/H$\\beta$ line ratio is high, but not too high for a low-metalicity star-forming galaxy, and corresponds to an oxygen abundance within the range of metallicities of LBGs measured by \\citet{pettini01}. The observations do not suggest that the optical spectrum of Q0347-383~C5 is dominated by an AGN.  The [OIII]$\\lambda$5007 line image shows two knots at a projected distance $\\sim 0.7$\\arcsec\\ (5.4 kpc) with a small relative velocity of 33 km s$^{-1}$. Line morphology and kinematics do not resemble those expected for an outflow or a rotating disk, and more likely originate from a merger of either two intermediate-mass galaxies with a dynamical mass of $\\le 10^{10} M_{\\odot}$ each, or perhaps massive sub-clumps of a fragmented disk as postulated by \\citet{immeli04,bournaud07}. The large masses of individual knots make it more likely that we see the merging of two galaxies each tracing its individual dark matter halo or subhalo, although this is a very difficult distinction to make with present day data. The density of similarly luminous z$\\sim 3$ LBGs is consistent with predictions from recent models of the cosmic evolution of the merger rate. Star-formation rates estimated from the observed H$\\beta$ flux correspond to $\\sim 20-40$ M$_{\\odot}$ in each clump, which is not unusual for LBGs generally.  Most z$\\sim 3$ LBGs are significantly more compact than Q0347-388~C5, with typical half-light radii of r$_e\\sim 0.3$\\arcsec. Such scales are difficult to resolve with 10-m class telescopes, even with adaptive optics assisted observations. From such observations \\citet{law07} find that DSF2237a-C2, their only target at z$>$3, has a velocity gradient and velocity dispersions of the same magnitude as the shear.  While superficially these characteristics could be suggestive of a rotating disk, \\citet{law07}, from a comparison of their data to a simple exponential rotating disk model, emphasize that this source is unlikely to be a thin, rotationally-supported disk. Both galaxies are among the largest LBGs and are comparably bright, which sheds doubts as to whether the properties of the overall population of $z\\sim 3$ LBGs are well described by the properties of its largest members. \\citet{nesvadba06} found evidence for rotation on sub-kpc scales in a strongly-lensed LBG at z$=3.24$, but such scales are well beyond reach for generic LBGs even with adaptive optics. While adaptive optics-assisted observations allow to probe the dynamics of high-redshift galaxies at sub-kpc resolution, they must concentrate on galaxies with particularly bright line emission, to ensure reasonable observing times as pointed out by \\citet{law07}.  This will inevitably lead to biases between observed LBG samples and the parent population of LBGs, and is a reason why studies of gravitationally lensed are not superceded, but are rather complemented, by high angular resolution observations of LBGs with adaptive optics, in spite of uncertainties related to the gravitational magnification.  More positively, observing galaxies with bright line emission will plausibly provide information about particularly rapid phases of star-formation and galaxy growth, whatever mechanism is responsible for initiating such phases. Prudence and caution however are certainly justified when generalizing the results of high redshift galaxies given the current limitation in astronomical instrumentation and the small sample sizes with detailed 3-dimensional spectroscopy observations."
    },
    "0711/0711.3494_arXiv.txt": {
        "abstract": " ",
        "introduction": "A lot of dust particles are produced by small bodies in the solar system. The dust located within about 2 AU from the Earth is seen as the zodiacal light. There are various points of view on the contributions of asteroidal, cometary, and trans-Neptunian dust to the zodiacal cloud. The estimates of the contributions made in several works are summarized in Table 1. These estimates were based on the Infrared Astronomical Satellite ({\\it IRAS}) and {\\it COBE/DIRBE} observations, on cratering rates, shape of microcraters, etc. In the present paper, for estimates of the contributions we analyzed some of these observations using our studies of migration of dust particles produced by different small bodies. We considered a wide range of particle masses, whereas some other scientists used results of calculations for one or two sizes of particles, e.g., Liou et al. (1995) considered 9 $\\mu$m diameter dust particles, and studies by Gorkavyi et al. (2000a,b) and Ozernoy (2001) were based on 1 $\\mu$m and 5$\\mu$m particles modeling. In our analysis for the first time we use the observations of velocities of zodiacal dust particles obtained by Reynolds et al. (2004) with the use of the Wisconsin H-Alpha Mapper (WHAM) spectrometer.   To set the stage for our work we first review some % published estimates of asteroid/cometa\\-ry contributions to the zodiacal cloud in more detail than in Table 1.    % A significant fraction of cometary dust in the near-Earth space was proposed by Southworth (1964), Liou et al. (1995), and Zook (2001). Based on cratering rates from an ensemble of Earth- and Lunar-orbiting satellites, Zook (2001) estimated that the cometary contribution to the near-Earth flux of particles is $\\sim$75\\%. His conclusion was based on (1) the comparison of the meteoroid penetration rates of the 25-$\\mu$m thick cells of the Earth-orbiting {\\it Explorer 16} and {\\it 23} satellites with the penetration rate of the five {\\it Lunar Orbiter} satellites that had nearly identical cells of the same thickness and on (2) the studies of the crater production rate on the leading edge of the Earth-orbiting Long Duration Exposure Facility ({\\it LDEF}) satellite, as compared to that on the trailing edge. For the estimates of the cometary contribution, Zook also used (1) Humes's (1993) result that it takes a mean Earth-entry velocity of about 17 km s$^{-1}$ to give agreement with the {\\it LDEF} observations and (2) Jackson and Zook's (1992) numerical modeling, which showed that meteoroids originating in the main belt of asteroids will strike the top of the Earth's atmosphere with a mean velocity between 12 and 13 km s$^{-1}$.   $<$ Place for Table 1 $>$ \\centerline{}  Grogan et al. (2001), Dermott et al. (2001), and Wyatt (2005) suggested that at least 30\\% of zodiacal dust comes from the break-up of asteroids in order to explain formation of dust bands (i.e., excesses of dust at elliptic latitudes $\\le$10$^\\circ$ (Kelsall et al. 1998)), as dust bands alone supply as much as 30\\%. Kortenkamp and Dermott (1998) suggested that the Earth predominately accretes asteroidal dust. Dermott et al. (2002) concluded that most of zodiacal dust particles can be of asteroidal origin and have eccentricities $e$$<$0.1. Gr\\\"un (1994) and Gr\\\"un et al. (2001) considered that zodiacal particles orbit the Sun at low inclination ($i$$<$30$^{\\circ}$) and moderate eccentricity ($e$$<$0.6) orbits. Our studies presented in Section 4 % are in accordance with Gr\\\"un's estimates.  Nesvorn\\'y et al. (2006) compared the {\\it IRAS} observations with their computer model of the thermal emission of the Karin and Veritas family particles. Their best-fit model results suggest that the Karin and Veritas family particles contribute by 5-9\\% in 10-60-$\\mu$m wavelengths to the zodiacal brightness within 50$^\\circ$ latitudes around the ecliptic, and by 9-15\\% within 10$^\\circ$ latitudes. The high brightness of the zodiacal cloud at large latitudes suggests that it is mainly produced by particles with higher orbital inclinations than what would be expected for asteroidal particles produced by sources in the main asteroid belt. Based on these results, Nesvorn\\'y et al. infer that asteroidal dust represents a smaller fraction of the zodiacal cloud than previously thought (e.g., by Dermott et al., 2001). They hypothesize that up to $\\approx$50\\% of interplanetary dust particles measured by the {\\it LDEF} may be made up of particles species from the Veritas and Karin families. Based on their modelling, the disproportional contribution of Karin/Veritas particles to the zodiacal cloud (only 5-9\\%) and to the terrestrial accretion rate (30-50\\%) suggests that the effects of gravitational focusing by the Earth enhance the accretion rate of Karin/Veritas particles relative to those in the background zodiacal cloud. Nesvorn\\'y et al. (2006) noted that the size distribution of asteroidal particles can be a strong function of the heliocentric distance. They infer that the zodiacal cloud emission may be dominated by high-speed cometary particles, while the terrestrial impactor flux contains a major contribution from asteroidal sources. Hahn et al. (2002) concluded that, though about 80\\% of the dust particles in the sample of dust collected in the Earth's stratosphere by U2 aircraft (Brownlee et al. 1993) have low entry velocities consistent with asteroidal orbits, the dust released from low-inclination Jupiter-family comets can also  have low entry velocities, and the Earth's gravitational focusing naturally selects for low-velocity dust over  all dust.   Sykes et al. (2004) infer that the zodiacal cloud scale height is not a good discriminator between asteroids and comets as the main supply of dust because inclinations of some Jupiter-family comets are not very different from those for asteroids and the half-width of the distribution of asteroidal  orbital inclinations (12$^\\circ$-16$^\\circ$) does not differ much from the half-width (14$^{\\circ}$) at half-maximum number density for the Kelsall model. Some models based on {\\it in situ} particle detections suggest that the inclination distribution may have a half-width as wide as 20$^\\circ$-30$^{\\circ}$ (Dikarev et al. 2002). If the latter estimate is true, then the fraction of zodiacal dust produced by long-period comets can exceed 10-20\\%.  Landgraf et al. (2002) concluded that Comet 29P/Schwassmann-Wachmann 1 itself is able to provide a major fraction of the solar system dust that is currently found between 6 and 8 AU. Comparison of the number density of dust particles produced by different small bodies with the observed constant density of $\\sim$10 $\\mu$m particles between 3 and 18 AU (Humes 1980, Gr\\\"un 1994) showed (Ozernoy 2001; Landgraf et al. 2002; Ipatov and Mather 2006) that a considerable fraction of dust at such distances is produced by comets. The conclusions on a considerable fraction of cometary dust are also in agreement with earlier studies of the dynamics of Jupiter-family comets (Ipatov and Mather 2003, 2004a-b, 2006), which showed that some former cometary objects could get high eccentric orbits located entirely inside Jupiter's orbit and stay in these orbits for a long time. Some of these objects could disintegrate producing a substantial amount of dust.  Gr\\\"un et al. (1985) proposed that the main contribution to the zodiacal light is from particles that range from 20 to 200 $\\mu$m in diameter (for silicate particles, this range corresponds to the ratio of the Sun's radiation pressure force to gravitational force $\\beta$$\\sim$0.002-0.02). The cratering record on {\\it LDEF} showed (Love and Brownlee 1993) that the cross-sectional area distribution of particles accreted by the Earth reaches maximum at a particle diameter $d$$\\sim$100 $\\mu$m. Smaller particles have a larger surface per unit of mass of particles. For a fragmentation power law $n(r_p)dr_p$=$n_{\\circ}r_p^{-q}dr_p$, the brightness of produced particles is proportional to $r_p^{2-q}$. Dynamical lifetimes  of particles are usually smaller for smaller particles, and many particles with diameter $d$$\\sim$1 $\\mu$m are relatively quickly removed from the solar system. If we consider that a lifetime of a particle $t_{dl}$$\\sim$$d$ (at 1$<$$d$$<$1000 $\\mu$m, this assumption is in a general agreement with dynamical lifetimes of particles obtained by Ipatov and Mather 2006; destruction of particles is not considered) and $q$=3.5, then the brightness of migrating particles with diameter $d$ is proportional to $d$$^{-1/2}$ and the total brightness of particles with diameters between $d_1$ and $d_2$ is proportional to $\\sqrt{d_2} - \\sqrt{d_1}$. These our estimates are in accordance with that particles with $d$$<$20 $\\mu$m are not main contributors to the zodiacal light.  Dermott et al. (2001, 2002) noted that the cross-sectional area of material in the asteroid belt should be concentrated in particles with $d$$\\sim$1000 $\\mu$m and many of these large particles are broken up by collisions before they reach a distance from the Sun $R$=1 AU. Gr\\\"un et al. (1985) considered that the Poynting-Robertson drag lifetime is comparable to the collisional lifetime for zodiacal particles greater than 100 $\\mu$m. Liou et al. (1996) and Moro-Martin and Malhotra (2002) studied collisional lifetimes of interplanetary particles destroyed by interstellar dust particles, and the sublimation temperature of particles at several distances to the Sun. They found that collisional destruction is most important for trans-Neptunian particles with $d$ between 6 and 50 $\\mu$m.   Ground-based spectroscopic observations allow one to study Doppler shifts of scattered solar Fraunhofer lines in the zodiacal light (e.g., James 1969; Hicks et al. 1974; Fried 1978; East and Reay 1984; Clarke 1996; Clarke et al. 1996).  Analysis of these shifts thus provides an opportunity to explore velocities of interplanetary dust in the inner solar system. Reynolds et al. (2004) were the first to obtain accurate measurements of the centroid velocities and line profiles of the scattered solar Mg I $\\lambda$5184 absorption line in the zodiacal light, both along the ecliptic equator and at high ecliptic longitudes.  The main goal of our paper is to compare the WHAM observations with our models of the zodiacal dust cloud based on our calculations of the migration of dust particles produced by different small bodies. The models for migration of dust particles and calculations of the radial velocity profile of the scattered Mg~I line are discussed in Section 2. Earlier we (Ipatov et al. 2005, 2006; Ipatov and Mather 2006; Madsen et al. 2007) compared velocities corresponding to shifts of the Mg~I line obtained in our models with the WHAM observations in a short form for only a few sizes of particles. Our present studies of the velocities (Section 3) are based on analysis of many data obtained for a wide range of sizes of particles and for various sources of particles. In the papers by other authors, spectroscopic observations of the zodiacal light have been compared with analytical models, but this is the first comparison of the models based numerical integrations of migration of particles with the observations. We also compare in detail the mean width of the line (the end of Section 3) and the variation of a  number density with distance from the Sun $R$$\\le$3 AU (Section 5.1) obtained in our model for different sources and sizes of particles with the results of observations, study the typical eccentricities and inclinations of zodiacal particles that better fit the WHAM observations (Section 4), and discuss the fractions of asteroidal, cometary, and trans-Neptunian particles in the zodiacal cloud which better satisfy various observations (Section 5). ",
        "conclusions": "Our study of velocities corresponding to Doppler shifts and widths of the Mg~I line in the zodiacal light is based on the distributions of positions and velocities of migrating dust particles originating from various solar system sources. These distributions were obtained from our integrations of the orbital evolution of particles produced by asteroids, comets, and trans-Neptunian objects. At the elongations of the fields observed by WHAM, the model curves of the characteristic velocity of the line vs. the solar elongation (`velocity-elongation' curves) were close to each other  for several scattering functions considered. The differences between the curves for several sources of dust reached its maximum at elongation between 90$^\\circ$ and 120$^\\circ$. Therefore it is important for the future zodiacal light Doppler shifts measurements  to pay a particular attention  to observing at this elongation range, since this is the elongation range that allows to have the best discrimination between the different dust sources.  The comparison of `velocity-elongation' curves and the line width averaged over the elongations obtained at observations made by Reynolds et al. (2004) with the corresponding % curves and mean widths obtained in our models shows that asteroidal dust particles alone cannot explain these observations, and that particles produced by comets, including high-eccentricity comets (such as Comet 2P/Encke and long-period comets), are needed. The conclusion that a considerable fraction of zodiacal dust is cometary particles is also supported by the comparison of the variations of a number density with a distance from the Sun obtained in our models with the spacecraft observations.  Cometary particles originating inside Jupiter's orbit and particles produced beyond Jupiter's orbit (including trans-Neptunian dust particles) can contribute to zodiacal dust about 1/3 each, with a possible deviation from 1/3 up to 0.1-0.2. The fraction of asteroidal dust is estimated to be $\\sim$0.3-0.5. The estimated contribution of particles produced by long-period and Halley-type comets to zodiacal dust does not exceed 0.1-0.15. The same conclusion can be made for particles originating from Encke-type comets (with $e$$\\sim$0.8-0.9).  The velocity amplitudes of `velocity-elongation' curves are greater for greater mean eccentricities and inclinations, but they depend also on distributions of particles over their orbital elements. The mean eccentricities of zodiacal particles located at 1-2 AU from the Sun that better fit the WHAM observations are between 0.2 and 0.5, with a more probable value of about 0.3."
    },
    "0711/0711.4561_arXiv.txt": {
        "abstract": "We conducted a survey for infrared excess emission from 16 nearby main sequence shell stars using the Multiband Imaging Photometer for \\si\\ (MIPS) on the \\emph{Spitzer Space Telescope}. Shell stars are early-type stars with narrow absorption lines in their spectra that appear to arise from circumstellar (CS) gas. Four of the 16 stars in our survey showed excess emission at 24~\\um\\ and 70~\\um\\ characteristic of cool CS dust and are likely to be edge-on debris disks. Including previously known disks, it appears that the fraction of protoplanetary and debris disks among the main sequence shell stars is at least $48\\% \\pm 14\\%$. While dust in debris disks has been extensively studied, relatively little is known about their gas content. In the case of $\\beta$~Pictoris, extensive observations of gaseous species have provided insights into the dynamics of the CS material and surprises about the composition of the CS gas coming from young planetesimals \\citep[e.g.][]{Roberge:2006}. To understand the co-evolution of gas and dust through the terrestrial planet formation phase, we need to study the gas in additional debris disks. The new debris disk candidates from this \\si\\ survey double the number of systems in which the gas can be observed right now with sensitive line of sight absorption spectroscopy. ",
        "introduction": "\\label{sec:intro}  Protoplanetary disks may be classified by their relative abundance of primordial gas and dust left over from star formation. As the amount of this material decreases, disks evolve from primordial disks to transitional disks to debris disks. Debris disks, which are found around main sequence stars, contain little or no primordial material. Instead, they are composed of recently produced secondary material generated by collisions between and evaporation of planetesimals, analogs of Solar System asteroids and comets. Terrestrial planet formation may be occurring in the younger debris disks, while the older ones seem to correspond to the ``clearing out'' stage early in the history of the Solar System, during which most left-over planetesimals are removed from the system. Impacts by water-rich planetesimals probably delivered most of the Earth's surface volatiles during the clearing phase \\citep[e.g.][]{Morbidelli:2000}.  Gas in debris disks is generally hard to detect. Only one bona fide debris disk, 49~Ceti, shows any trace of sub-mm CO emission \\citep{Zuckerman:1995, Dent:2005}, indicating that debris disk gas abundances are low relative to those in primordial and transitional disks. It appears that primordial molecular gas left over from star formation has largely dissipated by the debris disk phase. However, sensitive UV/optical absorption spectroscopy has revealed small amounts of secondary gas in a few debris disks: $\\beta$~Pictoris \\citep[e.g.][]{Lagrange:1998}, 51~Ophiuchi \\citep[e.g.][]{Roberge:2002}, $\\sigma$~Herculis \\citep{Chen:2003}, and HD~32297 \\citep{Redfield:2007a}. The gas seen is primarily atomic. The primary production mechanisms for secondary gas in debris disks are currently unknown, but may include photon-stimulated desorption from circumstellar (CS) dust grains \\citep{Chen:2007} and/or grain-grain collisions \\citep{Czechowski:2007}.  \\object[HD39060]{$\\beta$~Pic} is the only debris disk whose gas is well-characterized. In this disk, most measured gaseous elements have roughly solar composition relative to each other \\citep{Roberge:2006}. The exception is carbon, which is extremely overabundant relative to every other measured element (e.g.\\ C$/$O $= 18 \\times$ solar). This discovery was surprising, since the central star has solar metallicity \\citep{Holweger:1997}. The carbon overabundance suggests that a previously unsuspected -- and currently unknown -- process is operating in the disk, during which the planetesimals preferentially lose volatile carbonaceous material. However, the lack of similar abundance information for other debris disks renders interpretation of the \\bp\\ results difficult.  Detections of gas in debris disks to date have almost exclusively used optical or UV absorption spectroscopy of edge-on disks. Such observations are sensitive to very small amounts of cold gas. Unfortunately, the geometric constraint that the line of sight to the central star must pass through the disk severely limits the number of disks that may be probed for gas in this way. There is a strong need for additional debris disks in which gas can be observed, in order to study the evolution of gas abundance and composition throughout the entire planet formation phase.  \\begin{deluxetable*}{lcccccr} \\centering \\tablecolumns{7} \\tablewidth{0pt} \\tablecaption{Target Stars \\label{tab:targets}} \\tablehead{% \\multicolumn{1}{l}{ID} & \\colhead{Other Name} & \\colhead{Spectral Type} & \\colhead{Distance} & \\colhead{Inside the} & \\multicolumn{2}{c}{Exp.\\ Times (sec)} \\\\ \\colhead{}   & \\colhead{}           &  \\colhead{} & \\colhead{(pc)}     & \\colhead{LIB?} & \\colhead{24~$\\mu \\mathrm{m}$} & \\colhead{70~$\\mu \\mathrm{m}$} } \\startdata \\multicolumn{7}{l}{\\textit{Main Sequence Shell Stars from \\citet{Hauck:2000}}} \\vspace*{0.5ex} \\\\ \\object{HD21620} \\ \\tablenotemark{a} &  \t& A0Vn & 143 & y & 37 & 6965 \\\\ \\object{HD24863}\t\t     &  \t& A4V  & 107 & y & 37 & 3105 \\\\ \\object{HD39283}       \t\t     & $\\xi$~Aur & A2V &  74 & y & 37 & 671 \\\\ \\object{HD77190} \\ \\tablenotemark{a} & 67~Cnc \t& A8Vn &  59 & y & 37 & 2685 \\\\ \\object{HD98353} \\ \\tablenotemark{a} & 55~UMa \t& A2V  &  56 & y & 37 & 420 \\\\ \\object{HD118232} \t    \t     & 24~Cvn \t& A5V  &  58 & y & 37 & 294 \\\\ \\object{HD142926}\t \t     & 4~Her \t& B9pV & 148 & y & 37 & 2853 \\\\ \\object{HD158352}\t\t     &  \t& A8V  &  63 & y & 37 & 797 \\\\ \\object{HD196724}\t\t     & 29~Vul  \t& A0V  &  65 & y & 37 & 671 \\\\ \\object{HD199603} \\ \\tablenotemark{a} &  \t& A9V  &  85 & y & 37 & 2685 \\\\ \\object{HD223884} \t              &  \t& A5V  &  92 & y & 37 & 3483 \\\\ \\object{HD224463} \t   \t      &  \t& F2V  & 108 & y & 37 & 4280 \\\\ \\vspace*{-1.5ex} \\\\ \\multicolumn{7}{l}{\\textit{Shell Stars with Time-Variable Infalling Gas}} \\vspace*{0.5ex} \\\\ \\object{HD42111} \\ \\tablenotemark{a}  &  \t& A3Vn & 187 & n & 37 & 3609 \\\\ \\object{HD50241}\t              & $\\alpha$~Pic   & A7IV &  30 & y & 37 & 168 \\\\ \\object{HD148283} \\ \\tablenotemark{a} & 25~Her \t& A5V  &  79 & y & 37 & 671 \\\\ \\object{HD217782} \\ \\tablenotemark{a} & 2~And \t& A3Vn & 107 & y & 37 & 545 \\enddata \\tablenotetext{a}{ \\ Binary system.} \\end{deluxetable*}  A class of peculiar stars called shell stars may be the avenue toward finding edge-on debris disks containing gas. These are stars whose spectra show narrow absorption lines arising from line of sight gas at the velocity of the star. Most shell stars are B- or A-type. In some cases, the narrow lines may arise from interstellar (IS) material, although the fact that the gas is at the velocity of the star makes this interpretation less likely. In addition, the absorption lines are generally stronger than can be readily explained by IS gas \\citep[e.g.][]{Abt:1973}. The shell stars have high projected rotational velocities compared to non-shell early-type stars, showing that gas is more likely to be detected when the line of sight intersects the rotational plane of the system \\citep[e.g.][]{Holweger:1999}. This indicates that the gas is usually confined to that plane and strengthens the case for a CS rather than IS origin of most shell absorption lines.  Many shell stars appear to be rapidly rotating evolved stars that have recently ejected mass from their equatorial planes. However, \\citet{Hauck:2000} evaluated the luminosity classes of 57 shell stars and found that 40\\% of them are main sequence stars. The nature of the main sequence shell stars is not fully understood. Some are rapidly rotating classical Be stars with hot excretion disks. We note, however, that 6 out of the 23 main sequence stars in the \\citet{Hauck:2000} sample have protoplanetary or debris disks. Among the six are the well-known disk systems \\bp, HD~163296, and MWC~480. In fact, \\bp\\ was classified as a shell star a decade before discovery of its dust disk \\citep{Slettebak:1975}.  In this paper, we present a study aimed at finding additional debris disks hiding among the shell stars. We used the Multiband Imaging Photometer for \\si\\ (MIPS) on the \\emph{Spitzer Space Telescope} to measure the fluxes from 16 nearby main sequence shell stars at 24 and 70 \\um. The observations were designed to probe sensitively for excess flux arising from CS dust. The goal is to assemble a unique and valuable set of debris disk systems where the gas may be studied in as much detail as the dust. ",
        "conclusions": "\\label{sec:conc}  We have found four debris disk candidates in our \\si\\ survey of 16 main sequence shell stars. These stars show IR excess emission characteristic of moderate amounts of cool CS dust. We have begun follow-up observations of all of our shell stars, to more fully characterize their CS material.  Based on their SEDs, it seems likely that our candidates will be confirmed as debris disks. If so, they are likely to be edge-on, which is the most advantageous orientation for coronagraphic imaging of a disk. These disk candidates will likely double the number of debris disks with currently detectable CS gas. With this sample, we may begin a program to determine why some debris disks have both gas and dust and how these different components are related.  As mentioned in the Introduction, the nature of the main sequence shell stars has been something of a puzzle. Our \\si\\ study suggests that at least half of the stars in this class have protoplanetary or debris disks. Surveys for shell stars using ground-based optical spectroscopy, followed by IR photometry to look for CS dust, may be an effective way of finding protoplanetary and debris disks."
    },
    "0711/0711.0424_arXiv.txt": {
        "abstract": "Radiative transfer equation in an accretion disk wind is examined analytically and numerically under the plane-parallel approximation in the subrelativistic regime of $(v/c)^1$, where $v$ is the wind vertical velocity. Emergent intensity is analytically obtained for the case of a large optical depth, where the flow speed and the source function are almost constant. The usual limb-darkening effect, which depends on the direction cosine at the zero-optical depth surface, does not appear, since the source function is constant. Because of the vertical motion of winds, however, the emergent intensity exhibits the {\\it velocity-dependent} limb-darkening effect, which comes from the Doppler and aberration effects. Radiative moments and emergent intensity are also numerically obtained. When the flow speed is small ($v \\leq 0.1c$), the radiative structure resembles to that of the static atmosphere, where the source function is proportional to the optical depth, and the usual limb-darkening effect exists. When the flow speed becomes large, on the other hand, the flow speed attains the constant terminal one, and the velocity-dependent limb-darkening effect appears. We thus carefully treat and estimate the wind luminosity and limb-darkening effect, when we observe an accretion disk wind. ",
        "introduction": "Accretion disks are now widely believed to be energy sources in various active phenomena in the universe: in protoplanetary nebulae around young stellar objects (YSOs), in cataclysmic variables (CVs) and supersoft X-ray sources (SSXSs), in galactic X-ray binaries and microquasars ($\\mu$QSOs), and in active galaxies (ANGs) and quasars (QSOs). Accretion-disk models have been extensively studies during these three decades (see Kato et al. 1998 for a review). Besides the traditional standard disk by Shakura and Sunayev (1973), new type disks, such as advection-dominated accretion flows (ADAF) or radiatively-inefficient accretion flows (RIAF) for the very small mass-accretion rate (e.g., Narayan, Yi 1994), and supercritical accretion disks or so-called slim disks for the very large mass-accretion rate (e.g., Abramowicz et al. 1988).  Accretion disk winds have been also extensively examined in relation to astrophysical jets and outflows: in bipolar outflows from YSOs, in mass outflows from CVs and SSXSs, in relativistic jets from $\\mu$QSOs, AGNs, QSOs, and in gamma-ray bursts (GRBs). In particular, intense radiation fields of luminous supercritical accretion disks may be responsible for relativistic jets from super-Eddington sources, such as luminous $\\mu$QSOs, GRS~1915$+$105 and SS~433, luminous QSOs, 3C~273, and energetic GRBs (see, e.g., Fukue 2004 for references).  In such circumstances, radiative transfer in accretion disk winds as well as accretion disks becomes more and more important.  Radiative transfer in the standard disk has been investigated in relation to the structure of a static disk atmosphere and the spectral energy distribution from the disk surface (e.g., Meyer, Meyer-Hofmeister 1982; Cannizzo, Wheeler 1984). Furthermore, gray and non-gray models of accretion disks were constructed under numerical treatments (K\\v ri\\v z and Hubeny 1986; Shaviv and Wehrse 1986; Adam et al. 1988; Mineshige, Wood 1990; Ross et al. 1992; Shimura and Takahara 1993; Hubeny, Hubeny 1997, 1998; Hubeny et al. 2000, 2001; Davis et al. 2005; Hui et al. 2005) and under analytical ones (Hubeny 1990; Artemova et al. 1996; Fukue, Akizuki 2006a).  Radiative transfer in the accretion disk wind, on the other hand, has not been well considered both in the non-relativistic and relativistic regimes. Recently, radiative transfer in a moving disk atmosphere was firstly investigated in the subrelativistic regime (Fukue 2005a, 2006a), and in the relativistic regime (Fukue 2005b, 2006b; Fukue, Akizuki 2006b). In contrast to the static atmosphere, in the moving atmosphere the boundary condition at the surface of zero optical depth should be modified (Fukue 2005a, b). Moreover, the usual Eddington approximation violates in the highly relativistic flow (Fukue 2005b; see also Turolla, Nobili 1988; Nobili et al. 1991; Turolla et al. 1995; Dullemond 1999), and the velocity-dependent variable Eddington factor was proposed (Fukue 2006b for a plane-parallel case; Akizuki, Fukue 2007 for a spherical case).  Radiation hydrodynamical (RHD) simulations were also performed for radiation-dominated supercritical disks with winds by several researchers (Eggum et al. 1985, 1988; Okuda et al. 1997, 2005; Okuda, Fujita 2000; Okuda 2002; Ohsuga et al. 2005; Ohsuga 2006). In these current studies of RHD simulations for disks and winds, they were done in the subrelativistic regime up to the order of $(v/c)^1$, using the moment formalism and the flux-limited diffusion (FLD) approximation (Levermore, Pomraning 1981). The flux-limited diffusion method provides good approximations to the exact solutions but only if they are derived from transfer equations in which terms of the order of $(v/c)^2$ or higher have been retained (Yin, Miller 1995).   Radiative transfer problems on accretion disks and winds are not well understood yet, in particular for the relativistic cases. Hence, in order clarify the physics, in addition to RHD simulations, we must treat the simplified problem in the analytical way.  In this paper, we thus examine radiative transfer in the accretion disk wind, which is assumed to blow off from the luminous disk in the vertical direction (plane-parallel approximation), and analytically and numerically obtain the flow solutions for the case without internal heating. In the previous studies, radiative flows in the vertically moving atmosphere were solved in the subrelativistic regime (Fukue 2005a, 2006a), and in the relativistic regime (Fukue 2005b, 2006b; Fukue, Akizuki 2006b), where only radiative moments were obtained. In the present paper, we further obtain the radiation intensity as well as radiative moments.  In the next section we describe the basic equations. In section 3, we show analytical solutions, while we present numerical solutions in section 4. The final section is devoted to concluding remarks. ",
        "conclusions": "In this paper we have examined the radiative transfer problem in an accretion disk wind under the plane-parallel approximation in the subrelativistic regime of $(v/c)^1$. The flow velocity, the radiation pressure distribution, and other quantities are analytically and numerically solved as a function of the optical depth for the case without heating. Furthermore, the emergent intensity from the wind top with zero optical depth is also analytically obtained under the assumption of constant flow velocity, while numerically calculated.  When the source function is constant, te usual limb-darkening effect does not appear. However, the emergent intensity exhibits the {\\it velocity-dependent} limb-darkening effect, which originates from the Doppler and aberration effects, associating with the vertical subrelativistic motion of winds. As a result, a wind luminosity would be overestimated by a pole-on observer and underestimated by an edge-on observer, when we observe an optically thick accretion disk wind.   It should be noted that the {\\it apparent} optical depth in the relativistically moving media. Abramowicz et al. (1991) pointed out that the optical depth in the relativistic flow decreases as $\\gamma(1-\\beta\\mu)\\tau$ toward the downstream direction, due to the Doppler and aberration effects. Inspecting equation (\\ref{II}) or solution (\\ref{i_sol10}), we find that, in the present subrelativistic flow, the optical depth $\\tau$ is apparently replaced by $(1-\\beta\\mu)\\tau$. This is just consistent with the results by Abramowicz et al. (1991) within the order of $(v/c)^1$.  In the present modelized plane-parallel geometry, the gas density at the wind top is finite, since the wind velocity there is finite. Rigorously speaking, this situation is inconsistent with the definition of the wind top, where the optical depth is effectively zero. In more realistic cases, the disk wind would geometrically diverge, and the density quickly drops. As a result, the surface of a moving photosphere exists, even if the wind terminal velocity is finite. The present model mimics such a realistic disk wind, where the velocity-dependent limb darkening effect would also appear. In other words, as long as the acceleration of wind takes place in a small height, and then the wind diverges, the present assumption would be valid.  The radiative transfer problem investigated in the present paper must be quite {\\it fundamental problems} for accretion disk physics and astrophysical jet formation. In this paper, for simplicity, we only considered the subrelativistic case of $(v/c)^1$. The fully relativistic case of $(v/c)^2$ will be explored in the future.    \\vspace*{1pc}  The author would like to thank an anonymous referee for valuable comments. This work has been supported in part by a Grant-in-Aid for Scientific Research (18540240 J.F.) of the Ministry of Education, Culture, Sports, Science and Technology."
    },
    "0711/0711.1614_arXiv.txt": {
        "abstract": "% Microlensing is potentially sensitive to multiple-planet systems containing analogs of all the solar system planets except Mercury, as well as to free floating planets. I review the landscape of microlensing planet searches, beginning with an outline of the method itself, and continuing with an overview of the results that have been obtained to date. Four planets have been detected with microlensing. I discuss what these detections have taught us about the frequency of terrestrial and giant planets with separations beyond the ``snow line.''  I then discuss the near and long-term prospects for microlensing planet searches, and in particular speculate on the expected returns of next-generation microlensing experiments both from the ground and from space. When combined with the results from other complementary surveys, next generation microlensing surveys can yield an accurate and complete census of the frequency and properties of essentially all planets with masses greater than that of Mars. ",
        "introduction": "%  \\centerline{{\\it``I don't understand. You're looking for planets you can't}} \\centerline{{\\it see around stars you can't see.''} -- Debra Fischer, c.\\ 2000} \\centerline{} \\centerline{{\\it``Microlensing is a cult.''} -- Dave Koerner, c.\\ 2000} \\centerline{}  At the time when these words were uttered, microlensing planet searches had not yet made any definitive detections, despite having been in production in earnest for over five years, and despite the fact that numerous talks had been given vociferously extolling the virtues of the method.  Furthermore, it was a well-known `fact' (in part perpetuated by those giving said talks [including myself]) that the primary virtue of the microlensing method was its ability to provide statistics on large-separation and low-mass planets.  Individual detections would be of little interest, it was thought, because there would be no opportunity for follow-up and the  properties of the host stars would remain unknown, due primarily to their very large distances from Earth.  It is now seven years later, and the landscape and perception of microlensing planet searches has changed dramatically.  Four detections have been published, and an additional $\\sim$half a dozen planets have been detected and will be announced in the not-to-distant future.  The four published detections have already provided new and important information about the frequency of ``Super Earths'' and Jovian planets with separations beyond the ``snow line.'' Furthermore, it has become clear that the original assessments of the (im)possibility of follow-up and characterization of the planet host stars were overly pessimistic.  Indeed, for three of the four detected planets, reasonably precise ($\\sim 20\\%$) estimates of the mass of the host star have been made or will soon be possible. ",
        "conclusions": ""
    },
    "0711/0711.1752_arXiv.txt": {
        "abstract": "%  This is the written account of the first of two panel discussions, on {\\it Star formation in galaxies: how do we continue?} The chair of the panel was Phil James, and panel members were John Beckman, Torsten B\\\"oker, Daniela Calzetti, Angeles D\\'\\i az, and Rob Kennicutt. The panel and audience discussed the four following questions: 1) What are the most critically needed techniques to give accurate measurements of total rates and efficiencies of star formation?  2) Do we understand the form of the initial mass function and its variation as a function of redshift and environment?  3) Are there multiple modes of star formation in galaxies (bulge vs. disk, burst vs. continuous) or does the Schmidt law explain everything?  4) How do we bring together our understanding of star formation in our Galaxy and in external systems? ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.2028_arXiv.txt": {
        "abstract": "We report the discovery of five new millisecond pulsars in the globular cluster NGC~6440 and three new ones in NGC~6441; each cluster has one previously known pulsar. Four of the new pulsars are found in binary systems. One of the new pulsars, PSR J1748$-$2021B in NGC 6440, is notable for its eccentric ($e = 0.57$) and wide ($P_b = 20.5$ days) orbit. If the rate of advance of periastron is due solely to general relativity, we can estimate of the total mass of this binary system: $2.92 \\pm 0.20 M_{\\odot}$. This would imply an anomalously large mass for this pulsar, which could introduce important constraints in the study of the equation of state for cold neutron matter. ",
        "introduction": "The recent discovery of 30 millisecond pulsars (MSPs) in Terzan 5 \\cite{rhs+05,hrs+06}, made using the GBT's S-band receiver and the pulsar spigot \\cite{kel+05} shows that this observing system (henceforth GBT/S/PS) has an unprecedented sensitivity to MSPs outside the Arecibo sky. This is only partly due to the factor of 3 higher gain of the GBT as compared to the Parkes telescope. The higher observing frequency ($\\sim$2\\,GHz), large observing bandwidth ($600$\\,MHz), and relatively fine frequency resolution ($\\sim$0.78\\,MHz) provide enough frequency and time resolution to detect faint MSPs at high dispersion measures (DM). An impressive demonstration of this is the discovery of PSR~J1748$-$2446ad, the fastest spinning pulsar known ($P\\,=\\,1.396\\,$ms). This was found in Terzan~5, where the average DM of the pulsars is 239\\,cm$^{-3}$\\,pc. This motivated us to search for MSPs in other promising GCs with this system.  For the first phase of the survey, we selected the globular clusters (GCs) NGC~6388, NGC~6440 and NGC~6441 based on their high stellar interaction rates at the core, $\\Gamma_c$ \\cite{frb+07}. Despite the large $\\Gamma_c$s of these GCs, only two pulsars were known in them: PSR~B1745$-$20A (NGC6440A, \\cite{lmd96}) and PSR~J1750$-$37A (NGC~6441A, \\cite{pcm+06}). The spin periods and DMs of NGC~6440A and NGC~6441A are 288.6\\,ms and 111.6\\,ms and 220 and 234~cm$^{-3}$\\,pc, respectively.  We assumed that the lack of other known pulsars in these high-DM clusters was largely due to a bias of the earlier surveys against the detection of MSPs. The GBT/S/PS system can greatly reduce this bias.  With full track observations made in 2005, we discovered eight new MSPs, five of them in NGC~6440, three in NGC~6441 and none in NGC~6388. Preliminary results on these objects were presented in B\\'egin (2006)\\cite{beg06}. Their pulse profiles, together with those of the previously known objects, are presented in Figure~\\ref{fig:profiles}. These pulse profiles were obtained by adding {\\em all} detections of the pulsars, carefully excluding all sub-integrations that are corrupted by RFI.  \\begin{figure} \\label{fig:profiles} \\includegraphics[height=.23\\textheight]{fig1.eps} \\caption{Average 1950 MHz pulse profiles for the 10 pulsars known in the GCs NGC~6440 and NGC~6441. These profiles cover one full rotation. The pulsars in binary systems have an asterisk after their names, and the newly discovered pulsars have their names in boldface. The horizontal width of the rectangles indicates the system's total time resolution, including the effects of dispersive smearing, relative to each pulsar's spin period. The dip in power after the main pulse of NGC~6440A might be a SPIGOT artifact.} \\end{figure}  \\begin{figure}[htp] \\label{fig:fluxes} \\includegraphics[height=.33\\textheight]{fig2.eps} \\caption{Cumulative distribution of pseudo luminosities at 1950 MHz for the pulsars in NGC~6440, NGC~6441 and Terzan 5, assuming the most recent distance estimates.} \\end{figure} ",
        "conclusions": ""
    },
    "0711/0711.3168_arXiv.txt": {
        "abstract": "\\sloppy{We explore the relationship among three coronal mass ejections (CMEs), observed on 28 October 2003, 7 November 2004, and 20 January 2005, the type II burst-associated shock waves in the corona and solar wind, as well as the arrival of their related shock waves and magnetic clouds at 1 AU}. Using six different coronal/interplanetary density models, we calculate the speeds of shocks from the frequency drifts observed in metric and decametric radio wave data.  We compare these speeds with the velocity of the CMEs as observed in the plane-of-the-sky white-light observations and calculated with a cone model for the 7 November 2004 event. We then follow the propagation of the ejecta using Interplanetary Scintillation (IPS) measurements, which were available for the \\mbox{7 November 2004} and 20 January 2005 events. Finally, we calculate the travel time of the interplanetary (IP) shocks between the Sun and Earth and discuss the velocities obtained from the different data. This study highlights the difficulties in making velocity estimates that cover the full CME propagation time. ",
        "introduction": "In order to deduce reliable values of arrival time at Earth for CME-related ejecta, it is first necessary to determine the CME departure speed from the Sun where the phases of slow rise, acceleration, constant velocity, and deceleration all need to be considered. The problem is made more difficult by the fact that Earth-directed events are so-called halo or partial-halo CMEs that have a quasi-symmetric appearance around the solar disc when viewed in white light. Thus it is more difficult to obtain true earthward velocities at the departure times of these events.  A range of tools exists for the estimation of these speeds that are mostly applicable from the lower corona to distances up to $r \\sim$ 10 \\rsun. These include the observation of radio bursts in the dynamic spectra (principally type II bursts) and the analysis of white-light images obtained with coronagraphs. For values \\mbox{$r >$ 50 \\rsun}, when the faster CMEs are usually decelerating, the Interplanetary Scintillation (IPS) technique can be used. While it is possible to track the expansion of EUV or X-ray emitting plasmas by using appropriate imagers or by directly measuring velocities spectroscopically, these methods have been applied to CMEs in just a few isolated cases. While we expect a greater use of both imaging and spectroscopy in the future, through data obtained by the {\\it Hinode} and STEREO missions, neither of these latter methods has been employed for the events discussed here.  Since it is important to understand the strengths and limitations of these methods before applying them to CME speed estimation, it is necessary to briefly describe the underlying physics. We therefore present a short survey in the next section that outlines the basis for use of radio-burst dynamic spectra and white-light images for speed measurement. For $r >$ 50 \\rsun, use of the IPS technique is most appropriate, and so we outline the basic features of that method. We then describe the application of these techniques to the events that had been selected for the Sun-Earth Connection Workshop (in this Topical Issue), namely those of 28 October 2003, 7 November 2004, and 20 January 2005, and the results obtained from the analysis. We conclude with a general discussion of all three halo CMEs and the uncertainties in speed estimation. ",
        "conclusions": "The halo CMEs on 28 October 2003 and 7 November 2004 were both preceded by a slower CME. Both halo CMEs were found to be decelerating after initiation. During the 20 January 2005 event no evident CME ``interaction'' was observed, and there is some evidence to suggest that this CME was accelerating at low atmospheric heights. In all three events deceleration from the initial CME speed is evident from the arrival times of the magnetic clouds and the measured speeds near Earth, see Table \\ref{table6}. The transit speeds of the 7 November 2004 CME/magnetic cloud are discussed in more detail in Culhane {\\it et al.} (\\citeyear{culhane06}).  As all three CMEs were faster than the solar wind, the drag force exerted by the wind slowed them down considerably. Radio scintillation observations at heliocentric distances $>$50 {\\rsun} on 7\\,--\\,9 November 2004 and 20\\,--\\,21 January 2005 also record CME deceleration, at 3 and 30 \\mbox{m s$^{-2}$} respectively, but where this deceleration begins is not so clear. On 7 November 2004 the white-light CME was observed to decelerate at 20 m s$^{-2}$, which is consistent with the expectation that the CME should decelerate more rapidly near the Sun where the density is high, since the drag force depends on the density and velocity. The observed speeds near Earth from IPS measurements are in agreement with the measured speeds of the IP shocks and magnetic clouds.  Wang {\\it et al.} (\\citeyear{wang05}) have made simulations of CME interactions, in a case where a slow CME is overtaken by a faster one, with the result that the faster CME is slowed down significantly by the blocking action of the preceding magnetic cloud. This implies that the travel time is dominated by the slower cloud. Therefore it is possible that interaction with earlier, slower CMEs affected the speeds in the first two events. The simulations made by Lugaz, Manchester, and Gombosi (\\citeyear{lugaz05}) indicate that a CME can accelerate when it first meets the rear edge of a preceding CME (due to the density drop), and only decelerates later when it passes through the center and front of the earlier launched CME. Radio observations made at the time of CME acceleration would then give too high an instantaneous speed value. This could possibly explain the higher radio-source (shock) velocity compared to the CME velocity on the 7 November 2004 event, where the CME and radio source heights are in good agreement.   \\begin{table}[] \\caption{Estimated speeds for the magnetic clouds (MCs) and the IPS velocity \\mbox{estimates near Earth} } \\label{table6} \\begin{tabular}{llcccl} \\hline Launch         & Arrival  & Average & MC speed & Sun       & Notes \\\\ time           & time     & transit & near     & to        & \\\\ &          & speed   & Earth    & Earth     & \\\\ &          & (km/s)  & (km/s)   & (AU)      & \\\\ \\hline 28 Oct. 2003   & 29 Oct.  & 2140      & 1200 & 0.993 & CME launch \\\\ 11:00 UT       & 05:58 UT &           &      &       & to shock \\\\ &          &           &      &       & arrival\\\\ \\hline 07 Nov. 2004   & 09 Nov.  & 1190\\,--\\,1200 & ~800 & 0.990 & CME launch \\\\ 15:53-16:20 UT & 02:00 UT  &           &      &       & to shock  \\\\ &           &           &      &       & arrival \\\\ 07 Nov. 2004   & 09 Nov.   & 960\\,--\\,970  &  &       & CME launch \\\\ 15:53-16:20 UT & 10:00 UT  &           &      &       & to shock \\\\ &           &           &      &       & arrival \\\\ 07 Nov. 2004   & 09 Nov.   & 790\\,--\\,820  &  &       & to shock/\\\\ 15:53-16:20 UT & 17:55\\,--\\,18:55 UT &     &  &       & magnetic cloud \\\\ &           &           &      &       & arrival \\\\ IPS near Earth &           &           & ~700  &       & \\\\ \\hline 20 Jan. 2005   & 21 Jan.   & 1180      & ~675  & 0.984 & flare start \\\\ 06:36 UT       & 16:50 UT  &           &      &       & to first  \\\\ &           &           &      &       & shock\\\\ 20 Jan. 2005   & 21 Jan.   & 1130      &      &       & flare start\\\\ 06:36 UT       & 18:20 UT  &           &      &       & to second  \\\\ &           &           &      &       & shock \\\\ IPS near Earth &           &           & 1000 &       & \\\\ \\hline \\end{tabular} \\end{table}  The estimated radio source heights are also in agreement with the CME heights on 28 October 2003. The calculated shock velocity from the DH type II burst at 11:16\\,--\\,11:21 UT (3600 km s$^{-1}$) indicates a shock speed only somewhat higher than the CME speed observed after 11:30 UT (2900 km s$^{-1}$). The white-light CME was also observed to decelerate. Since the high-speed full-halo CME was propagating in the wake of the earlier partial-halo CME, the high momentary speed could reflect temporary acceleration or abrupt changes in the local density. Estimation of the continuation of the DH type-II-burst lane suggests that the shock was much slower (2600 km s$^{-1}$) later on.  The source locations at decimetric-metric wavelengths on 28 October 2003 have recently been analysed by Pick {\\it et al.} (\\citeyear{pick05}), who discovered that the radio sources outline an H$\\alpha$-Moreton wave, which is a direct signature of a propagating shock wave at the photospheric level (see {\\it e.g.} Uchida, \\citeyear{uchida74}). The average speed of the Moreton wave was around 2000 km s$^{-1}$. The authors also report type III bursts where the burst envelope moves with a speed of 2500 km s$^{-1}$, in the same direction (southwest) as the most compact (brightest) CME structure. Therefore it is plausible that the burst envelope of the type III bursts and the DH type II burst were closely associated with the propagating CME.  In the case of 20 January 2005 we see several type-II-burst lanes in the dynamic spectra at decimetric-metric and DH wavelengths, and it is possible that several shocks were formed during the intense event. The shock speeds derived from the decimetric observations (650\\,--\\,1000 km s$^{-1}$) are not too far from the estimated initial CME speed, before acceleration (as was shown in Figure~\\ref{fig4e})."
    },
    "0711/0711.4923_arXiv.txt": {
        "abstract": "We have obtained time-resolved spectroscopic observations for a sample of 10 cool Ap stars using the ultra-stable spectrograph HARPS at the ESO 3.6-m telescope. The aim of our study was to search for low-amplitude oscillations in Ap stars with no or inconclusive evidence of pulsational variability. Here we report initial results of our investigation. We confirm the presence of $\\approx$\\,16-min period pulsations in $\\beta$~CrB (HD~137909) and demonstrate multiperiodic character of oscillations in this star. Furthermore, we discovered very low amplitude 9-min pulsations in HD~75445 -- an object spectroscopically very similar to known roAp stars. ",
        "introduction": "Many recent time-resolved spectroscopic studies of the rapidly oscillating Ap (roAp) stars (e.g., Kurtz et al. \\cite{kc06}, Ryabchikova et al. \\cite{tr07}) have focused on the known, high-amplitude pulsators and did not address the fundamental question of what distinguishes pulsating and non-oscillating Ap (noAp) stars. Ryabchikova et al. (\\cite{tr04}) suggested that there might exist no real dichotomy between roAp and noAp stars, and weak pulsational variability is present in all cool Ap stars.  With the aim to test this hypothesis and to obtain the first complete picture of the incidence of pulsations in Ap stars, we used the ultra-stable spectrograph HARPS at ESO to obtain time-resolved spectra for a sample of 10 noAp and weakly oscillating Ap stars. Here we summarize preliminary results of the analysis of HARPS spectra of HD\\,75445 and HD\\,137909 ($\\beta$~CrB).  \\begin{figure}[!t] \\includegraphics[width=6cm]{P14_f1a.eps} \\includegraphics[width=6.0cm]{P14_f1b.eps} \\caption{{\\it Left panel:} Periodogram showing discovery of 9-min pulsational signal in the combined RVs of 29 Nd~{\\sc iii} lines in HD\\,75445. {\\it Right panel:} Mean RV variation of 40 Ce~{\\sc ii} lines in $\\beta$~CrB. Multiperiodic character of oscillation is clearly evident. } \\label{fig1} \\end{figure} ",
        "conclusions": ""
    },
    "0711/0711.4112_arXiv.txt": {
        "abstract": "We estimate the power of relativistic, extragalactic jets by modelling the spectral energy distribution of a large number of blazars.  We adopt a simple one--zone, homogeneous, leptonic synchrotron and inverse Compton model, taking into account seed photons originating both locally in the jet and externally.  The blazars under study have an often dominant high energy component, which, if interpreted as due to inverse Compton radiation, limits the value of the magnetic field within the emission region.  As a consequence, the corresponding Poynting flux cannot be energetically dominant.  Also the bulk kinetic power in relativistic leptons is often smaller than the dissipated luminosity.  This suggests that the typical jet should comprise an energetically dominant proton component.  If there is one proton per relativistic electrons, jets radiate around 2--10 per cent of their power in high power blazars and 3--30 per cent in less powerful BL Lacs. ",
        "introduction": "The radiation observed from blazars is dominated by the emission from relativistic jets (Blandford \\& Rees 1978) which transport energy and momentum to large scales. As the energy content on such scales already implies in some sources jet powers comparable with that which can be produced by the central engine (e.g. Rawlings \\& Saunders 1991), only a relatively small fraction of it can be radiatively dissipated on the `blazar' (inner) scales.  However we still do not know the actual power budget in jets and in which form such energy is transported, namely whether it is mostly ordered kinetic energy of the plasma and/or Poynting flux.  In addition, the predominance of one or the other form can change during their propagation. These of course are crucial pieces of information for the understanding on how jets are formed and for quantifying the energy deposition on large scales.  In principle the observed radiation can -- via the modelling of the radiative dissipation mechanism -- set constraints on the minimum jet power and can even lead to estimates of the relative contribution of particles (and the corresponding bulk kinetic power), radiation and magnetic fields.  The modelling depends of course on the available spectral information and conditions on the various jet scales (i.e. distances from the central power--house).  Attempts in this direction include the work by Rawlings \\& Saunders (1991), who considered the energy contained in the extended radio lobes of radio--galaxies and radio--loud quasars.  By estimating their life--times they could calculate the average power needed to sustain the emission from the lobe themselves (Burbidge 1959).  At the scale of hundreds of kpc the {\\it Chandra} satellite observations, if interpreted as inverse Compton scattering on the Cosmic Microwave Background (Tavecchio et al. 2000; Celotti, Ghisellini \\& Chiaberge 2001), indicate that jets of powerful blazars are still relativistic.  This allowed Ghisellini \\& Celotti (2001) to estimate a minimum power at these distances for PKS 0637--752, the first source whose large scale X--ray jet was detected by {\\it Chandra}.  Several other blazars were studied by Tavecchio et al. (2004), Sambruna et al. (2006) and Tavecchio et al. (2007) who found that the estimated powers at large scales were comparable (within factors of order unity) with those inferred at much smaller blazar scales.  Celotti \\& Fabian (1993) considered the core of jets, as observed by VLBI techniques, to derive a size of the emitting volume and the number of emitting electrons needed to account for the observed radio luminosity.  The bulk Lorentz factor, which affects the quantitative modelling, was estimated from the relativistic beaming factor required not to overproduce, by synchrotron self--Compton emission, the observed X--ray flux (e.g. Celotti 1997).  A great advance in our understanding of blazars came however with the discovery that they are powerful $\\gamma$--ray emitters (Hartman et al. 1999).  Their $\\gamma$--ray luminosity often dominates (in the powerful flat spectrum radio--loud quasars, FSRQs) the radiative power, and its variability implies a compact emitting region.  The better determined overall spectral energy distribution (SED) and total observed luminosity of blazars constrain -- via pair opacity arguments (Ghisellini \\& Madau 1996) -- the location in the jet where most of the dissipation occurs.  For a given radiation mechanism the modelling of the SED also allows to estimate the power requirements and the physical conditions of this emitting region.  Currently the models proposed to interpret the emission in blazars fall into two broad classes.  The so-called `hadronic' models invoke the presence of highly relativistic protons, directly emitting via synchrotron or inducing electron--positron (e$^{\\pm}$) pair cascades following proton-proton or proton--photon interactions (e.g., Mannheim 1993, Aharonian 2000, M{\\\"u}cke et al. 2003, Atoyan \\& Dermer 2003).  The alternative class of models assumes the direct emission from relativistic electrons or e$^{\\pm}$ pairs, radiating via the synchrotron and inverse Compton mechanism. Different scenarios are mainly characterised by the different nature of the bulk of the seed photons which are Compton scattered. These photons can be produced both locally via the synchrotron process (SSC models, Maraschi, Ghisellini \\& Celotti 2002), and outside the jet (External Compton models, EC) by e.g. the gas clouds within the Broad Line Region (BLR; Sikora, Begelman \\& Rees 1994, Sikora et al. 1997) reprocessing $\\sim$10 per cent of the disk luminosity.  Other contributions may comprise synchrotron radiation scattered back by free electrons in the BLR and/or around the walls of the jet (mirror models, Ghisellini \\& Madau 1996), and radiation directly from the accretion disk (Dermer \\& Schlickeiser 1993; Celotti, Ghisellini \\& Fabian, 2007).  Some problems suffered by hadronic scenarios (such as pair reprocessing, Ghisellini 2004) make us favour the latter class of models. By reproducing the broad band properties of a sample of $\\gamma$--ray emitting blazars via the SSC and EC mechanisms, Fossati et al. (1998) and Ghisellini et al. (1998; hereafter G98) constrained the physical parameters of a (homogeneous) emitting source.  A few interesting clues emerged.  The luminosity and SED of the sources appear to be connected, and a spectral sequence in which the energy of the two spectral components and the relative intensity decrease with source power seems to characterise blazars, from low power BL Lacs to powerful FSRQs (opposite claims have been put forward by Giommi et al. 2007; see also Padovani 2006). This SED sequence translates into an (inverse) correlation between the energy of particles emitting at the spectral peaks and the energy density in magnetic and radiation fields (Ghisellini, Celotti \\& Costamante 2002; hereafter G02).  An interpretation of such findings is possible within the internal shock scenario (Ghisellini 1999; Spada et al. 2001; Guetta et al 2004), which could account for the radiative efficiency, location of the dissipative region and spectral trend if the particle acceleration process is balanced by the radiative cooling. In such a scenario the energetics on scales of 10$^2-10^3$ Schwarzschild radii is dominated by the power associated to the bulk motion of plasma. This is in contrast with an electromagnetically dominated flow (Blandford 2002; Lyutikov \\& Blandford 2003).  Within the frame of the same SSC and EC emission model, in this work we consider the implications on the jet energetics, the form in which the energy is transported and possibly the plasma composition.  In particular we estimate the (minimum) power which is carried by the emitting plasma in electromagnetic and kinetic form in a significant sample of blazars at the scale where the $\\gamma$--ray emission -- and hence most of the luminosity -- is produced.  Such scale corresponds to a distance from the black hole of the order of $10^{17}$ cm (Ghisellini \\& Madau 1996), a factor 10--100 smaller than the VLBI one.  The found energetics are lower limits as they only consider the particles required to produce the observed radiation, and neglect (cold) electrons not contributing to the emission.  In Section 2 the sample of sources is presented. In Section 3 we describe how the power in particles and field have been estimated, and the main assumptions of the radiative model adopted.  The results are reported in Section 4 and discussed in Section 5.  Preliminary and partial results concerning the power of blazar jets were presented in conference proceedings (see e.g. Ghisellini 1999, Celotti 2001, % Ghisellini 2004). %  We adopt a concordance cosmology with $H_0=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_\\Lambda=0.7$ and $\\Omega_{\\rm M}=0.3$. ",
        "conclusions": "The first important result emerging from this work is that the power of extragalactic jets is large in comparison to that emitted via accretion.  This result is rather robust, since the uncertainties related to the particular model adopted are not crucial: the finding follows from a comparison with the emitted luminosity, which is a rather model--independent quantity, relying only on the estimate of the bulk Lorentz factor.  The findings about the kinetic and Poynting powers instead depend on the specific modelling of the blazar SEDs as synchrotron and Inverse Compton emission from a one--zone homogeneous region.  Hadronic models may yield different results.  Furthermore, the estimated power associated to the proton bulk motion relies also on the amount of `cold' (non emitting) electron--positron pairs in the jet.  We have argued that if pairs had to be dynamically relevant their density at the jet base would make annihilation unavoidable. However the presence of a few pairs per proton cannot be excluded.  If there were no electron--positron pairs, the inferred jet powers are 10--100 times larger than the disk accretion luminosity, in agreement with earlier claims based on individual sources or smaller blazar samples (Ghisellini 1999; Celotti 2001; Maraschi \\& Tavecchio 2003; Sambruna et al. 2006).  Such large powers are needed in order to energise the emitting leptons at the ($\\gamma$--ray) jet scale and the radio--lobes hundreds of kpc away.   The finding that blazar jets are not magnetically dominated is also quite robust, but only in the context of the (widely accepted) framework of the synchrotron--inverse Compton emission model.  In this scenario the dominance of the high energy (inverse Compton) component with respect to the synchrotron one limits the magnetic field.  This is at odd with magnetically driven jet acceleration, though this appears to be the most viable possibility. In blazars thermally driven acceleration, as invoked in Gamma--Ray Bursts, does not appear to be possible. In Gamma-Ray Bursts the initial fireball is highly opaque to electron scattering and this allows the conversion of the trapped radiation energy into bulk motion (see e.g. Meszaros 2006 for a recent review).  In blazars the scattering optical depths at the base of the jet are around unity at most, and even invoking the presence of electron--positron pairs to increase the opacity is limited by the fact that they quickly annihilate.  Thus, if magnetic fields play a crucial role our results would require that magnetic acceleration must be rapid, since at the scale of a few hundreds Schwarzschild radii, where most of radiation is produced, the Poynting flux is no longer energetically dominant (confirming the results by Sikora et al. 2005). However models of magnetically accelerated flows indicate that the process is actually relatively slow (e.g. Li, Chiueh \\& Begelman 1992, Begelman \\& Li 1994). Apparently the only possibility is that the jet structure is more complex than what assumed and a possibly large scale, stronger field does not pervade the dissipation region, as also postulated in pure electromagnetic scenarios (see e.g. Blandford 2002, Lyutikov \\& Blandford 2003).  The third relevant results refers the difference between FSRQs and BL Lacs. This concerns not only their jet powers but also the relative role of protons in their jets. BL Lacs would be more dissipative and therefore their jets should decelerate.  This inference depends on assuming one proton per emitting lepton also in these sources, and this is rather uncertain (i.e. there could be more than one proton per relativistic, emitting electron).  If true, it can provide an explanation to why VLBI knots of low power BL Lacs are moving sub-luminally and in turn account for the different radio morphology of FR I and FR II radio--galaxies, since low power BL Lacs are associated to FR I sources."
    },
    "0711/0711.0782_arXiv.txt": {
        "abstract": "With more than 260 extrasolar planetary systems discovered to-date, the search for habitable planets has found new grounds. Unlike our solar system, the stars of many of these planets are hosts to eccentric or close-in giant bodies. Several of these stars are also members of moderately close ($<$40 AU) binary or multi-star systems. The formation of terrestrial objects in these \"extreme\" environments is strongly affected by the dynamics of their giant planets and/or their stellar companions. These objects have profound effects on the chemical structure of the disk of planetesimals and the radial mixing of these bodies in the terrestrial regions of their host stars. For many years, it was believed that such effects would be so destructive that binary stars and also systems with close-in giant planets would not be able to form and harbor habitable bodies. Recent simulations have, however, proven otherwise. I will review the results of the simulations of the formation and long-term stability of Earth-like objects in the habitable zones of such \"extreme\" planetary systems, and discuss the possibility of the formation of terrestrial planets, with significant amounts of water, in systems with hot Jupiters, and also around the primaries of moderately eccentric close binary stars. ",
        "introduction": "An analysis of the orbital properties of the currently known extrasolar planets indicates that many of these objects have dynamical characteristics that are profoundly different from those of the planets in our solar system. While around the Sun, planets keep nearly circular orbits with smaller planets at closer distances and the larger ones farther way, from the 263 exoplanets that were discovered at the time of the writing of this article, 68 were at distances smaller than 0.1 AU from their host stars, and 126 revolve in orbits with eccentricities larger than 0.2 (figure 1). A deeper look at the parent stars of these planets indicates that the differences between their planetary systems and the Solar System are not limited to the semimajor axes and eccentricities of their planetary companions. Several of these stars (approximately 20\\% of currently known extrasolar planet-hosting stars) are members of binaries or multi-star systems.  Since the masses of the majority of extrasolar planets are within the range of a few Neptune- to several Jupiter-masses, the eccentric and close-in orbits of these bodies, and the fact that several of them exist around the components of binary stars, raise questions on the formation of such {\\it extreme} planetary systems, and also on the possibility of the existence of smaller planets, such as terrestrial-class objects, around their host stars, and particularly in their systems' habitable zones. In this paper, I review the current status of research on the formation of habitable planets in these extreme planetary environments, and discuss the results of the simulations of Earth-like planet formation in systems with close-in giant planets and binary companions.   \\begin{figure} \\plottwo {f1a.eps}{f1b.eps} \\caption{Graphs of the mass and eccentricity of some of the extrasolar planets. The graph on the left shows close-in giant planets. The graph on the right shows planets with large eccentricities. The two encircled areas depict extreme planetary systems. Graphs from exoplanet.eu.} \\label{fig1} \\end{figure} ",
        "conclusions": ""
    },
    "0711/0711.2264_arXiv.txt": {
        "abstract": "{} {We aim to provide atmospheric parameters and rotational velocities of a large sample of O- and early B-type stars, analysed in a homogeneous and consistent manner, for use in constraining theoretical models.} {Atmospheric parameters, stellar masses and rotational velocities have been estimated for approximately 250 early B-type stars in the Large (LMC) and Small (SMC) Magellanic Clouds from high-resolution VLT-FLAMES data using the non-LTE TLUSTY model atmosphere code.  This data set has been supplemented with our previous analyses of some 50 O-type stars (Mokiem et al. 2006, 2007) and 100 narrow-lined early B-type stars (Hunter et al. 2006, Trundle et al. 2007) from the same survey, providing a sample of $\\sim$400 early-type objects.} {Comparison of the rotational velocities with evolutionary tracks suggest that the end of core hydrogen burning occurs later than currently predicted. We also show that the large number of the luminous blue supergiants observed in the fields are unlikely to have directly evolved from main-sequence massive O-type stars as neither their low rotational velocities or position on the H-R diagram are predicted. We suggest that blue-loops or mass-transfer binary systems may populate the blue supergiant regime. By comparing the rotational velocity distributions of the Magellanic Cloud stars to a similar Galactic sample we find that (at 3$\\sigma$ confidence level) massive stars (above 8M$_{\\rm \\sun}$) in the SMC rotate faster than those in the solar neighbourhood.  However there appears to be no significant difference between the rotational velocity distributions in the Galaxy and the LMC. We find that the $v \\sin i$ distributions in the SMC and LMC can modelled with an intrinsic rotational velocity distribution which is a Gaussian peaking at 175\\,km\\,s$^{-1}$ (SMC) and 100\\,km\\,s$^{-1}$ (LMC) with a $\\frac{1}{e}$ half width of 150\\,km\\,s$^{-1}$. We find that in NGC346 in the SMC, the 10-25\\,M$_{\\rm \\sun}$ main-sequence stars appear to rotate faster than their higher mass counterparts. It is not expected that O-type stars spin down significantly through angular momentum loss via stellar winds at SMC metallicity, and hence this could be a reflection of mass dependent birth spin rates. Recently  Yoon et al. (2006) have determined rates of GRBs by modelling rapidly rotating massive star progenitors. Our measured rotational velocity distribution for the 10-25M$_{\\rm \\sun}$ stars is peaked at slightly higher velocities than they assume, supporting the idea that GRBs could come from rapid rotators with initial masses as low as 14\\,M$_{\\rm \\sun}$ at low metallicities. }{} ",
        "introduction": "\\label{s_intro}  The evolution of massive stars has traditionally been described in terms of mass and metallicity. However in the past 20 years it has become evident that rotational effects are equally important, from star formation through to their deaths in supernovae explosions and associated gamma-ray bursts (Woosley \\& Heger \\cite{woo06}). For example, Heger \\& Langer (\\cite{heg00}) discuss how rotation changes the lifetime and evolution of massive stars on the Hertzsprung-Russell diagram through rotationally induced mixing between the core and envelope.  Rotational effects are also believed to be metallicity dependent. At low metallicity, line-driven winds are weaker and mass-loss rates lower (Kudritzki et al. \\cite{kud87}; Kudritzki \\& Puls \\cite{kp2000}; Vink et al. \\cite{vin01}; Mokiem et al. \\cite{mok06}). Hence a star will lose less angular momentum and thereby maintain higher rotational velocities on the main-sequence. This promotes rotational mixing over a longer period of an object's evolution which in turn extends its lifetime. {\\bf Additionally nucleosynthetically processed material will} be mixed from the core into the photosphere and hence abundance anomalies are expected, with helium and nitrogen enhancements and carbon and oxygen depletions. Recent observations of massive stars have revealed abundance anomalies in O-stars (Bouret et al. \\cite{bou03}, Mokiem et al. \\cite{mok06}, \\cite{mok07}), B-type supergiants (Dufton et al. \\cite{duf05}, Hunter et al. \\cite{hun06}, Lennon et al. \\cite{len03} and Trundle \\& Lennon \\cite{tru05}), B-type main-sequence stars (Hunter et al. \\cite{hun06}, Korn et al. \\cite{kor02}, \\cite{kor05}), A-type supergiants (Venn \\cite{ven99}) and lower mass stars (Smiljanic et al. \\cite{smi06}).  Maeder \\& Meynet (\\cite{mae01}) have included rotation and metallicity effects in stellar evolutionary models to explain both the relative frequency of blue and red supergiants as well as observed abundance trends in A-type supergiants at low metallicity. However much remains uncertain. For example, they adopt initial rotational velocities of 300\\,km\\,s$^{-1}$, which equates to an average main-sequence rotational velocity of 220-240\\,km\\,s$^{-1}$ (there being little change in the rotational velocity over the main-sequence lifetime of a B-type star), but statistically significant observational samples are necessary both to estimate typical rotational velocities as well as to calibrate the effect of metallicity on the rotational velocities. Additionally rotation has been described as the key ingredient in enabling single stars to explode with associated long-duration gamma-ray bursts by Woosley \\& Heger (\\cite{woo06}), although rotational velocities close to 400\\,km\\,s$^{-1}$ are required and hence gamma-ray bursts would be relatively rare events compared to supernovae. The requirement that little angular momentum is lost during their stellar lifetime obviously favours low metallicity regimes with their smaller mass-loss rates.  A large survey of massive stars has been undertaken at the European Southern Observatory using the Fibre Large Array Multi-Element Spectrograph (FLAMES) on the 8.2m Keuyen Very Large Telesope. The observations are discussed in detail in Evans et al. (\\cite{eva05}, hereafter Paper I) and Evans et al. (\\cite{eva06}, hereafter Paper II) and are summarised in Sect.~\\ref{s_obs}. Briefly, some 750 stars have been observed with mainly O and early B spectral types in the Galaxy and the Large (LMC) and Small (SMC) Magellanic Clouds. The analysis of many of these objects is well underway or completed. Mokiem et al. (\\cite{mok06}, \\cite{mok07}) have derived atmospheric parameters and helium abundances for the majority of the O-type stars in the sample and discuss the dependence of mass-loss rates on metallicity and calibrate the wind-momentum luminosity relation. Hunter et al. (\\cite{hun06}) and Trundle et al. (\\cite{tru07}) have derived atmospheric parameters and metal abundances (CNO in particular) for a subset of the narrow-lined B-type main-sequence and supergiant stars finding that significant mixing has occurred in many of the objects. Dufton et al. (\\cite{duf06a}, hereafter Paper III) have estimated atmospheric parameters and rotational velocities of the entire Galactic sample and report that more massive stars have lower rotational velocities (as expected from enhanced mass-loss) and that cluster objects appear to rotate significantly faster than objects in the Galactic field.  In order to examine the effects of rotation, here we perform a similar analysis to Paper III and estimate atmospheric parameters and rotational velocities for the Magellanic Cloud sample of B-type stars. {\\bf The principle intention of performing such an analysis was to provide rotational velocities for a large sample of early-type stars in a homogenous and consistent manner. This dataset can then be used to answer several questions of importance to massive star evolution:  \\begin{itemize} \\item Where does the end of core hydrogen burning occur and how can large populations of post core hydrogen burning objects be explained? \\item Are the rotational velocities of early B-type stars dependent on the metallicity of the host environment? \\item Are mass-loss effects important at low metallicity and how do these affect the predictions of gamma-ray burst rates from massive star progenitors? \\end{itemize} }   In Sect.~\\ref{s_obs} we briefly describe the FLAMES survey and our selection criteria for the objects discussed here. In Sect.~\\ref{s_analysis} we present our methods for deriving the atmospheric parameters and projected rotational velocities. {\\bf In Sect.~\\ref{s_nature} the rotational velocity distributions of the different subsets of the sample (for example, binaries, Be-type stars, high mass stars and supergiants) are compared. In Sect.~\\ref{s_discussion} the relevant subsets are utilised to answer the three questions posed above.} Finally in Sect.~\\ref{s_conclude} we present the principal conclusions from our analysis. ",
        "conclusions": "\\label{s_conclude}  Atmospheric parameters and projected rotational velocities are presented for approximately 400 O- and early B-type stars towards four clusters in the Magellanic Clouds, which combines this work, the analysis of some 50 O-type stars (Mokiem et al. \\cite{mok06}, \\cite{mok07}) and the analysis of 100 narrow lined B-type stars (Hunter et al. \\cite{hun06} and Trundle et al. \\cite{tru07}). This represents the full data set from our ESO Large Programme, the VLT Flames survey of massive stars in the Magellanic clouds and the Galaxy. {\\bf From the subsets of the data we find that supergiants are the slowest rotators in the sample, typically having rotational velocities less than 80\\,km\\,s$^{-1}$. Additionally binaries tend to rotate slower than single objects, but given that binarity is easier to detect at low projected rotational velocities this result may be dominated by selection effects. It is also found that objects with Be phenomena are typically the fastest rotators in our dataset, which is in agreement with previous studies. The number of high mass objects (M$>$25\\,M$_{\\rm \\sun}$) in the combined sample is limited but there is some evidence that the most massive objects rotate slower than their less massive counterparts. We utilise the single core-hydrogen burning objects along with these subsets to address the questions outlined in Sect.~\\ref{s_intro}}.  \\begin{itemize} \\item  {\\bf Current evolutionary models do not predict the large number of giants observed in our samples which may imply that the theoretical tracks need to be extended. Indeed we find that stars classed as giants typically rotate some $\\sim$50\\,km\\,s$^{-1}$ slower than dwarf objects and this is consistent with a modest slowing down of the rotational velocity over the core hydrogen burning lifetime, indicating that giants are still in the hydrogen burning phase. From the bi-model distribution of rotational velocities for supergiant and core-hydrogen burning objects we observe the end of the core-hydrogen burning phase to be at 3.2\\,dex. We also show that supergiants are unlikely to have evolved directly from the hydrogen burning main-sequence, but instead, binarity or blue loops may be a more likely explanation for the observed supergiant population. }  \\item Our sample is dominated by field objects and we compare the rotational velocity distributions with Galactic field objects from the literature, finding that the SMC metallicity stars rotate fastest, in agreement with theoretical predictions. The difference between the SMC stars and Galactic field stars is significant at the 3$\\sigma$ level, suggesting we have detected a real difference. However we find no significant difference between the rotational velocity distributions of the Galactic and LMC stars. By assuming random angles of inclination, we can model the observed $v \\sin i$ distributions with broad Gaussians with the parameters listed in  Table~\\ref{t_mcvels}.  \\begin{table} \\centering \\caption[]{Peak and width of the Magellanic Cloud rotational velocity distributions determined by fitting the distributions with a Gaussian function, see Sect.~\\ref{s_vmet}.} \\label{t_mcvels} \\begin{tabular}{lcc} \\hline \\hline Region & Peak & $\\frac{1}{e}$ half-width \\\\ & km\\,s$^{-1}$ & km\\,s$^{-1}$ \\\\ \\hline \\\\ LMC    & 100  & 150 \\\\ SMC    & 175  & 150 \\\\ \\\\ \\hline \\end{tabular} \\end{table}   \\item {\\bf Objects more massive than 25\\,M$_{\\rm \\sun}$ tend to rotate slower than less massive objects and this} could be due to mass-loss effects, with a very massive star losing angular momentum through strong stellar winds, although this is unexpected in the SMC. An alternative explanation is that the initial birth rotational velocity is mass dependent, perhaps due to protostellar winds.  {\\bf However, the consequence is that Yoon et al. (\\cite{yoo06}) may have underestimated the rotational velocity distribution of 10-20\\,M$_{\\rm \\sun}$ GRB progenitors.} We therefore suggest that the rotational velocity distribution derived here should be employed in future GRB rate calculations, and this may even increase the rate of GRBs expected at low metallicty due to {\\bf the greater prevalence of B-type stars}.  \\end{itemize}"
    },
    "0711/0711.1267_arXiv.txt": {
        "abstract": "{The forthcoming new-generation radio telescope SKA (Square Kilometre Array) and its precursors will provide a rapidly growing number of polarized radio sources.} {Our analysis looks at what can be learned from these sources concerning the structure and evolution of magnetic fields of external galaxies.} { \\emph{Recognition} of magnetic structures is possible from Faraday rotation measures (${\\rm RM}$) towards background sources behind galaxies or a continuous $\\RM$ map obtained from the diffuse polarized emission from the galaxy itself. We constructed models for the ionized gas and magnetic field patterns of different azimuthal symmetries (axisymmetric, bisymmetric and quadrisymmetric spirals, and superpositions) plus a halo magnetic field. $\\RM$ fluctuations with a Kolmogorov spectrum due to turbulent fields and/or fluctuations in ionized gas density are superimposed. Assuming extrapolated number density counts of polarized sources, we generated a sample of $\\RM$ values within the solid angle of the galaxy. Applying various templates, we derived the minimum number of background sources and the minimum quality of the observations. For a large number of sources, \\emph{reconstruction} of the field structure without precognition becomes possible.} {Any large-scale regular component of the magnetic field can be clearly recognized from $\\RM$ data with the help of the $\\chi^2$ criterium. Under favorite conditions, a few dozen polarized sources are enough for a reliable result. A halo field with a vertical component does not affect the results of recognition. The required source number increases for small inclinations of the galaxy's disk and for larger $\\RM$ turbulence. A flat number density distribution of the sources can be overcome by more sensitive observations. Application of the recognition method to the available $\\RM$ data in the region around M~31 indicates that there are significant $\\RM$ contributions intrinsic to the background sources or due to the foreground of the Milky Way. A reliable reconstruction of the field structure needs at least 20 $\\RM$ values on a cut along the projected minor axis.} {Recognition or reconstruction of regular field structures from the $\\RM$ data of polarized background sources is a powerful tool for future radio telescopes. Measuring $\\RM$ at frequencies around 1~GHz with the SKA, simple field structures can be recognized in galaxies up to about 100\\,Mpc distance and will allow to test dynamo against primordial or other models of field origin. The low-frequency SKA array and low-frequency precursor telescopes like LOFAR may also have good $\\RM$ sensitivity if background sources are still significantly polarized at low frequencies.} ",
        "introduction": "\\label{sec:intro}  Magnetic fields play an important role for the structure and evolution of galaxies. Radio synchrotron emission provides the best tools for measuring the strength and structure of galactic magnetic fields \\citep{beck96}. However, the sensitivity of present-day radio observations only allows detailed studies for a couple of nearby galaxies \\citep{beck05}.  Rotation measures ($\\RM$) towards polarized background sources located behind spiral galaxies can trace regular magnetic fields in these galaxies out to large distances, even where the synchrotron emission of the galaxy itself is too weak to be detected \\citep{han98,gaensler05}. However, with the sensitivity of present-day radio telescopes, the number density of polarized background sources is only a few sources per solid angle of a square degree, so that only M~31 and the LMC, the two angularly largest galaxies in the sky, could be investigated so far \\citep{han98,gaensler05}.  Future high-sensitivity radio facilities will be able to resolve the detailed magnetic field structure in galaxies by observing their polarized intensity and $\\RM$ directly. Furthermore, they will be able to observe a huge number of faint radio sources, thus providing the high density background of polarized point sources. This opens the possibility of mapping the Faraday rotation of polarized background sources towards nearby galaxies and study their magnetic fields. The detection of large-scale field patterns (modes) and their superpositions would strengthen dynamo theory for field amplification and ordering, while the lack of $\\RM$ patterns would indicate a primordial origin or fields structured by gas flows \\citep{beck06}. Mapping the $\\RM$ of the diffuse polarized emission of galaxies at high frequencies is restricted to the star-forming regions where cosmic-ray electrons are produced. At low frequencies the extent of synchrotron emission is larger due to the propagation of cosmic ray electrons, but Faraday depolarization increases much faster with wavelength, so that the diffuse polarized emission is weak. On the other hand, $\\RM$ towards polarized background sources are generated by regular fields plus ionized gas, which extend to larger galactic radii, as indicated by the existing data sets \\citep{han98,gaensler05}, and are less affected by Faraday depolarization.  A major step towards a better understanding of galactic magnetism will be achieved by the Square Kilometre Array (SKA, www.skatelescope.org) and its pathfinders, the Allen Telescope Array (ATA) in the US, the Low Frequency Array (LOFAR) in Europe, the MeerKAT in South Africa, and the Australia SKA Pathfinder (ASKAP). The SKA will be a new-generation telescope with a square kilometre collecting area, a frequency range of 70~MHz to 25~GHz with continuous frequency coverage, a bandwidth of at least 25\\%, a field of view of at least 1 square degree at 1.4~GHz, and angular resolution of better than 1~arcsecond at 1.4~GHz. The SKA is planned to consist of three separate arrays: a phased array for low frequencies (70--300~MHz), a phased array for medium frequencies (300--1000~MHz), and an array of single-dish antennas for high frequencies (1--25~GHz).  ``Cosmic magnetism'' is one of six Key Science Projects for the SKA with the plan to measure a grid of more than $10^7$ $\\RM$ data towards polarized sources over the whole sky \\citep{gaensler04}. This will allow measurement of the evolution of magnetic fields in galaxies from the most distant to nearby galaxies and in the Milky Way. Large-scale field structures should be observable via $\\RM$ mapping of the diffuse polarized emission in galaxies \\citep{beck06}. Even more promising are $\\RM$ measurements towards polarized background sources, which are discussed in this paper. ",
        "conclusions": "\\label{sec:concl}    {\\em Recognition} of simple structures of regular magnetic fields in galaxies, characterized by spiral pitch angle, amplitude of regular field strength, and azimuthal phase, can be reliably performed from a limited sample of $\\RM$ measurements towards polarized background sources. Applying templates to the field models, we could show that single and mixed modes can be clearly recognized based on the $\\chi^2$ criterium. The quality of $\\RM$ observations can be characterized by the number $N$ of detected polarized sources within the solid angle of the galaxy, varying with distance $D$ and instrumental noise in polarization $\\sigma_p$ as $N \\propto \\sigma_p^{-\\gamma} D^{-2}$, and the ``observation parameter'' $Q$, which varies as $Q \\propto \\sigma_p D^{2/\\gamma}$ and also depends on wavelength and bandwidth of the observations. Field recognition is successful beyond a minimum source number $N^*$ and below a maximum observation parameter $Q^*$.  The slope $\\gamma$ of the cumulative number density counts of polarized sources with flux densities $P$ ($N \\propto P^{-\\gamma}$) is actually known beyond $P=0.5$~mJy from available source counts at 1.4~GHz. Extrapolation to smaller $P$ yields $0.7 \\le \\gamma \\le 1.1$. $\\gamma$ still needs to be determined, also at lower and higher frequencies. For the optimistic case of $\\gamma=1.1$, significant inclination of the galaxy's disk ($i > 25^{\\circ}$) and weak $\\RM$ turbulence (due to turbulent magnetic fields and/or fluctuations of the ionized gas density), a few dozen polarized sources are already sufficient to reliably recognize a simple field structure of type ASS, BSS, or QSS. Higher modes (BSS and QSS) are easier to recognize; i.e., they need less sources or yield smaller errors of the pitch angle for the same source number compared to the basic ASS mode. The requirement for a telescope's dynamic range in polarization (polarization purity) may exceed 30~dB in case of a flat slope $\\gamma$ of the cumulative number density counts of polarized sources.  Vertical fields do not affect the results of the reconstruction for galaxies with inclinations $i>15^{\\circ}$. The method should not be applied for galaxies at smaller inclinations where $\\RM$ from the disk field is small and a vertical field may dominate $\\RM$.  Strong turbulence or small inclination need more polarized sources for a statistically reliable recognition. In the case of small inclination angles, $\\RM$ turbulence can get comparable to the $\\RM$ amplitude of the regular field and the required number of polarized sources may become unreasonably large. As the required number of polarized sources per solid angle also increases for large inclination angles due to geometry, the recognition method works best for intermediate inclination angles.  Uncertainties in the position of the galaxy's center of a few arcseconds are even acceptable for galaxies at 100~Mpc distance if the field pattern is type ASS. For BSS-type fields accuracies of better than 2\\arcsec\\ are needed beyond 50~Mpc distance, and better than 1\\arcsec\\ beyond 50~Mpc distance for QSS-type fields.  The application of the recognition method to the available $\\RM$ values within the solid angle of the galaxy M~31 is consistent with the analysis of the diffuse polarized emission if significant $\\RM$ contributions intrinsic to the background sources or fluctuating $\\RM$ in the Milky Way foreground are assumed. The regular field of M~31 is type S0, i.e. axisymmetric spiral in the plane and symmetric with respect to the plane.  {\\em Reconstruction} of the field structure without precognition needs a large number of sources. A reliable reconstruction of the field structure is possible for at least 20 $\\RM$ values on a cut along the projected minor axis.  Future RM observations with the SKA at about 1~GHz will extend the possible targets for field recognition to about 100~Mpc distance, allow detailed statistical studies of the frequency of field structures, dependencies on galaxy properties, and test dynamo against primordial models or other models of field generation. About 60 spiral galaxies can be observed within a distance of 10~Mpc and $\\approx 60000$ spirals within 100~Mpc.  Radio telescopes operating at low frequencies (LOFAR, ASKAP, and the low-frequency SKA array) may also be useful instruments for field recognition or reconstruction with the help of $\\RM$, if background sources are still significantly polarized at low frequencies."
    },
    "0711/0711.1860_arXiv.txt": {
        "abstract": "The quadrupole power of cosmic microwave background (CMB) temperature anisotropies seen in the \\wmap data is puzzlingly low. In this paper we demonstrate that Minimum Variance Optimization (MVO), a technique used by many authors (including the \\wmap science team) to separate the CMB from contaminating foregrounds, has the effect of forcing the extracted CMB map to have zero statistical correlation with the foreground emission. Over an ensemble of universes the true CMB and foreground are indeed expected to be uncorrelated, but any particular sky pattern (such as the one we happen to observe) will generate non-zero measured correlations simply by chance. We call this effect ``cosmic covariance'' and it is a possible source of bias in the CMB maps cleaned using the MVO technique. We show that the presence of cosmic covariance is expected to artificially suppress the variance of the Internal Linear Combination (ILC) map obtained via MVO. It also propagates into the multipole expansion of the ILC map, generating a quadrupole deficit with more than 90\\% confidence. Since we do not know the CMB and the foregrounds {\\it a priori}, there is therefore an unknown contribution to the uncertainty in the measured quadrupole power, over and above the usual cosmic variance contribution. Using the MVO on a series of Monte Carlo simulations that assume Gaussian CMB fluctuations, we estimate that the real quadrupole power of the CMB lies in the range $[305.16,400.40] (\\mu {\\rm K}^2)$ (at the $1-\\sigma$ level). ",
        "introduction": "Detailed measurements of the pattern of temperature anisotropies in the cosmic microwave background (CMB) have provided cosmologists with unprecedented opportunities to probe the physics of the early Universe. In particular, the Wilkinson Microwave Anisotropy Probe (WMAP) data \\citep{wmapresults,wmapfg,wmapcl,wmap3ytem} have played a pivotal role in the establishment of the standard ``concordance'' cosmological model and have heavily constrained alternatives to the inflationary paradigm for the origin of large-scale structure.  According to the concordance cosmology, the variations in temperature of the CMB across the celestial sphere should possess a very broad spectrum usually expressed in terms of the amplitude of spherical harmonic modes labelled by the usual angular harmonic frequency $\\l$. The power of the pattern of anisotropies power is expected to be strongest for the quadrupole ($\\l=2$) which correspond to variations on an angular scale of  $90^\\circ$ on the sky. In principle, therefore, the quadrupole should be the easiest harmonic mode to detect. In practice, however, the measured quadrupole anisotropy of the CMB from \\wmap seems to sit rather uncomfortably with the standard cosmology: its amplitude seems too low \\citep{efssig,dtzh}. Of course, the power that resides in each multipole of the CMB pattern (including the quadrupole) is not so much measured as extracted. There are various diffuse foreground emissions to be separated from the ``true'' primordial CMB, such as free-free emission arising from electron-ion scattering, synchrotron emission from cosmic ray electrons accelerating in the galactic magnetic fields, and dust emission. However, although a number of authors have  employed different foreground subtraction schemes, all find the derived quadrupole power to be lower than the Concordance Cosmological Model with a significance level of 5\\% or better \\citep{efssig,lilc,dt,park3y,saha}. This has even prompted suggestions of physics beyond the standard model \\citep{efsspacurv,dodecahedral}.  In this paper we suggest that the quadruple power deficit in \\wmap may be an artifact of the Minimum Variance Optimization method used by many authors to clean CMB maps from foreground contamination. In the next Section, we explain the basics of the MVO method. In Section 3 we show how this method is prone to a bias introduced by the accidental alignment of structures in the CMB and foreground templates leading to an effect we call ``cosmic covariance''. In Section 4 we quantify the likely effect of this bias using Monte Carlo simulations, and we estimate the quadrupole power taking into account the effect of cosmic covariance. We present our conclusions in Section 6. ",
        "conclusions": "We have shown that the MVO method for eliminating foreground contamination from CMB maps is prone to a bias if there exist accidental alignments between foreground features and structures in the true CMB sky. This bias acts in such a way that a suppression of the quadrupole is more likely than an enhancement. It also reduces the overall variance of the recovered map compared to the true CMB sky. This suggests that attempts to attribute the low quadrupole to exotic physics may be premature.  The MVO was originally designed for the separation of Gaussian signals, whose statistical properties are completely characterized by their power spectra. For a linear combination of such signals, the variance is clearly the natural choice for a parameter to be minimized, and in such a case this method generally leads to a good reconstruction of the power in each component. In this paper, we have demonstrated even for the MVO to work on Gaussian signals it requires a statistical ensemble: for a single realization the role of cosmic covariance becomes crucial. The foregrounds relevant to CMB analysis are also highly non--Gaussian, which makes this problem even worse.  For composite foregrounds with varying spectra, the variance after application of the MVO will induce even more foreground contamination in the ILC map because the weighting coefficients will then be tuned not only for the subtraction of cosmic covariance, but for varying morphology among the different foregrounds. It is then possible that the variance deficit generated by the subtraction of cosmic covariance can be compensated by a foreground residual, which can push the quadrupole power back to the value that is consistent with the Concordance Cosmological model. Extra caution, therefore, will have to be applied when employing the MVO to the interpretation of results from the upcoming ESA Planck Experiment.  \\smallskip  {\\it Acknowledgments:} We acknowledge the use of \\healpix \\footnote{{\\tt http://healpix.jpl.nasa.gov/}} package \\citep{healpix} to produce $\\alm$ from the \\wmap data, and the use of \\glesp\\footnote{{\\tt http://www.glesp.nbi.dk/}} package. \\vfill"
    },
    "0711/0711.3112_arXiv.txt": {
        "abstract": "{} { We study the influence of large-scale photospheric motions on the destabilization of an eruptive filament, observed on October 6, 7, and 8, 2004, as part of an international observing campaign (JOP~178).} {Large-scale horizontal flows were invetigated from a series of MDI full-disc Dopplergrams and magnetograms. From the Dopplergrams, we tracked supergranular flow patterns using the local correlation tracking (LCT) technique. We used both LCT and manual tracking of isolated magnetic elements to obtain horizontal velocities from magnetograms.} { We find that the measured flow fields obtained by the different methods are well-correlated on large scales. The topology of the flow field changed significantly during the filament eruptive phase, suggesting a possible coupling between the surface flow field and the coronal magnetic field. We measured an increase in the shear below the point where the eruption  starts and a decrease in shear after the eruption. We find a pattern in the large-scale horizontal flows at the solar surface that interact with differential rotation.} { We conclude that there is probably a link between changes in surface flow and the disappearance of the eruptive filament.} ",
        "introduction": "Dynamic processes on the Sun are linked to the evolution of the magnetic field as it is influenced by the different layers from the convection zone to the solar atmosphere. In the photosphere, magnetic fields are subject to diffusion due to supergranular flows and to the large-scale motions of differential rotation and meridional circulation. The action of these surface motions on magnetic fields plays an important role in the formation of large-scale filaments (Mackay and Gaizauskas 2003). In particular, the magnetic fields that are transported across the solar surface can be sheared by dynamic surface motions, which in turn result in shearing of the coronal field. This corresponds to the formation of coronal flux ropes in models, which can be compared with H$\\alpha$ filament observations (Mackay and van Ballegooijen 2006b). Many theoretical models try to reproduce the basic structure and the stability of filaments by taking  surface motions into account, as quoted above. These models predict that magnetic flux ropes involved in solar filament formation may be stable for many days and then suddenly become unstable, resulting in filament eruption. Observations show that twisting motions are a very common characteristic of eruptive prominences (see for example Patsourakos and Vial 2002). However, it is still unknown whether the magnetic flux ropes emerge already twisted or if it is only the photospheric motion that drives the twisting of the filament magnetic field. The destabilization of the filament can also be linked to oscillations (Pouget 2006). Therefore, it remains uncertain as to the mechanisms that drive filament disappearance. Destabilization can come from the interior of the structure or by means of an outside flare.  In a previous paper (Rondi et al. 2007, hence forth Paper I), local horizontal photospheric flows were measured at high spatial resolution (0.5\\arcsec) in the vicinity of and beneath a filament before and during the filament's eruptive phases (the international JOP178 campaign). It was shown that the disappearance of the filament originates in a filament gap. Both parasitic and normal magnetic polarities were continuously swept into the gap by the diverging supergranular flow. We also observed  the interaction of opposite polarities in the same region, which could be a candidate for initiating the destabilization of the filament by causing a reorganization of the magnetic field.  In this paper we investigate the large-scale photospheric flows at moderate spatial resolution (2\\arcsec) beneath and in the vicinity of the same  eruptive filament. The observation and coalignment between data from various instruments are explained in Section~2. In Section~3, we describe the different methods of determining the flow field on the Sun surface. The large-scale flows associated with the filament are shown Section~4. The properties of these flows before and after the filament eruption are described in Section~5. In Section~6, we investigate the topology of horizontal flows in the filament area over the 3 days around the filament eruption. A discussion of the results and general conclusions can be found in Section~7. ",
        "conclusions": ""
    },
    "0711/0711.3937_arXiv.txt": {
        "abstract": "Searches for statistically significant correlations between arrival directions of ultra-high energy cosmic rays and classes of astrophysical objects are common in astroparticle physics.  We present a method to test potential correlation signals of \\textit{a priori} unknown strength and evaluate their statistical significance sequentially, i.e., after each incoming new event in a running experiment.  The method can be applied to data taken after the test has concluded, allowing for further monitoring of the signal significance.  It adheres to the likelihood principle and rigorously accounts for our ignorance of the signal strength. ",
        "introduction": "\\label{sec:intro}  One of the major goals in astroparticle physics is the identification and the study of sources of ultra-high energy cosmic rays, defined as cosmic rays with energies larger than $10^{18}$\\,eV.  The discovery of discrete sources would answer longstanding questions about how and where particles are accelerated to such energies.  So far, no discrete sources have been positively identified. One major obstacle for the identification of potential sources is the small number of detected events.  Until a few years ago, the published world data set of cosmic rays with energies above $4\\,\\times 10^{19}$\\,eV consisted of little more than 100 events, mainly recorded with the Akeno Giant Air Shower Array (AGASA) in Japan between 1984 and 2003~\\citep{Takeda:1999sg}, and the High Resolution Fly's Eye (HiRes) Experiment in Utah between 1997 and 2006~\\citep{Abbasi:2004ib}.  Nevertheless, the small data set has been subjected to exhaustive searches for deviations from isotropy.  These include searches for point sources; searches for an excess of clustering in the distribution of arrival directions on various angular scales; and searches for correlations with classes of known astrophysical objects that were considered likely sites of cosmic ray acceleration.  Some of these searches resulted in potential signals, but because of the small size of the data set, the statistical significance could not be established in a reliable manner.  Consequently, while the discovery of discrete sources was claimed repeatedly, statistically independent data routinely failed to support earlier claims.  An example is the search for correlations of cosmic ray arrival directions with objects of the BL Lac class~\\citep{Tinyakov:2001nr,Gorbunov:2004bs,Abbasi:2005qy}.  With a new generation of large-aperture astroparticle physics detectors like the Pierre Auger Observatory nearing completion in Malarg\\\"ue, Argentina and the Telescope Array detector under construction in Utah, the amount of ultra-high energy data is now growing at an unprecedented pace.  The Pierre Auger Observatory, for instance, began scientific data taking in January 2004 and has already accumulated over $9\\times10^3\\,\\mathrm{km}^{2}\\,\\mathrm{sr}\\,\\mathrm{yr}$ of integrated exposure, more than any previous experiment.  \\subsection{Basic Search Techniques in Cosmic Ray Physics}  The fact that previous experiments have failed to find statistically significant deviations from isotropy in skymaps of ultra-high energy cosmic rays can be seen as an indication that the sources are weak.  In this case, the most promising correlation searches are not those which aim at finding sources individually, but rather those conducted on a statistical basis; i.e., searches for significant correlations of cosmic ray arrival directions with catalogs of astrophysical objects.  When studying correlations with objects from a source catalog, one tests whether the probability $p$ of a given event to arrive from the direction of an object in the catalog is significantly larger than the probability $p_0$ of the correlation occurring by chance.  These analyses are typically binned, so an event is said to correlate with an object from the catalog if the angle between its arrival direction and the object's position is smaller than some angle $\\theta$.  If the particles are neutral, $\\theta$ could be chosen to reflect the point spread function of the detector.  In the case of cosmic rays, however, the particles are most likely charged and therefore deflected by Galactic and intergalactic magnetic fields of (unknown) strength. Consequently, $\\theta$ is usually chosen to be larger than the resolution of the detector to account for magnetic smearing.  Typically, potential signals are identified after intensive searches using different angular scales, different energy thresholds, different source catalogs, and other parameters that are found to maximize the signal strength. Therefore, an unbiased chance probability for the observed signal can only be established by discarding the data set used to find the signal and testing the signal with statistically independent data.  For the test, the source catalog and all analysis parameters are fixed {\\it a priori} to obtain an unbiased chance probability for the signal.  Once the \\textsl{a priori} analysis parameters are identified, the problem is easily formulated in terms of a classical hypothesis test, in which new data are checked for compatibility with a null hypothesis $\\mathcal{H}_0$ (``the data exhibit no significant correlation'') or an alternative ``signal'' hypothesis $\\mathcal{H}_1$.  There are several ways to perform such a test. For example, one can run the test after the new data set has reached a certain size $n$, or after the experiment has run for a certain fixed amount of time.  Formally, the size of the data set and the acceptance or rejection of the null hypothesis are determined by two probabilities, $\\alpha$ and $\\beta$, which are usually chosen before the start of the test.  These values define the experimenter's tolerance for different sorts of experimental errors: $\\alpha$ is the probability of wrongly rejecting the null hypothesis when $\\mathcal{H}_0$ is true (a type-1 or ``false positive'' error); and $\\beta$ is the probability of wrongly accepting the null hypothesis when $\\mathcal{H}_0$ is false (a type-2 or ``false negative'' error).  In a classical one-sided hypothesis test, where a p-value $P$ is used to estimate the agreement of the data with the null hypothesis, the result $P<\\alpha$ implies rejection of $\\mathcal{H}_0$ at the ``confidence level'' $1-\\alpha$.  Meanwhile, the desired probability of rejecting a false null hypothesis $(1-\\beta)$ fixes the required size of the data set $(n)$.  \\subsection{One-Shot vs. Sequential Testing}  If one chooses to evaluate $P$ after a predefined number of events has been recorded, or a predefined amount of time has elapsed, then the significance of the signal is tested only once.  However, it is often desirable to evaluate and test the signal sequentially, i.e., after each new event, rather than at the end of the test.  This approach allows for the possibility of claiming a statistically significant result earlier than with methods that check the signal only once, a distinct advantage when event rates are quite low.  It also avoids another practical disadvantage of hypothesis tests that arises when the experiment, for one reason or another, has to discontinue data taking before the predefined number of events is taken.  In that case, the ``one-shot'' analysis does not lead to a conclusion.  A sequential analysis can be performed in several ways.  If $P$ is evaluated after every incoming event and not just once after all $n$ events are collected, a ``penalty'' factor has to be inserted to account for the fact that there are now more opportunities to satisfy the test by chance~\\citep{Anscombe:1954,Armitage:1969}.  This penalty factor can be evaluated with simulations and will depend on $n$.  The dependence of $P$ on $n$ is an undesirable feature of the method; rather than depending on the data that were actually recorded, $P$ now depends on the number of events that an observer would have recorded had he decided to perform a ``one-shot'' test.  The interpretation of the data therefore depends on data not actually taken.  This feature of the test violates the likelihood principle~\\citep{Berry:1987}.  In addition, the inclusion of the penalty factor means that data arriving after the test has ended cannot be used to calculate $P$ for the entire data set.  It is therefore not possible to include new data in the calculation of the probability.  In many practical situations, data taking continues after the test has ended, and it is highly desirable to monitor the signal probability with new data.  The classical sequential likelihood ratio test developed by~\\citet{Wald:1945,Wald:1947} avoids the limitations that arise when using the p-value $P$.  Wald defines the likelihood ratio evaluated after the $n^{th}$ event as \\begin{equation} \\mathcal{R}_n=\\frac{P(\\mathcal{D}|\\mathcal{H}_1)}{P(\\mathcal{D}|\\mathcal{H}_0)}~~, \\label{eq:likelihood_ratio} \\end{equation} where the denominator and numerator represent the probability of observing a data set $\\mathcal{D}$ given a null hypothesis (no correlation) and an alternative (correlation).  The ratio $\\mathcal{R}_n$ can be evaluated after each incoming event (i.e. after the $n^{th}$ event) without statistical penalty, and the test stops with the acceptance or rejection of the null hypothesis when $\\mathcal{R}_n$ falls below or exceeds a predefined value (details will be given in Section~\\ref{sec:method}). Moreover, the evaluation of $\\mathcal{R}_n$ can continue after the decision to see whether new data continue to favor or disfavor the selected hypothesis.  The probabilities $P(\\mathcal{D}|\\mathcal{H}_0)$ and $P(\\mathcal{D}|\\mathcal{H}_1)$ in eq.\\,(\\ref{eq:likelihood_ratio}) depend on the expected correlations in case of random coincidences and true signals, respectively.  In correlation studies, the strength of the signal is typically not known before the test is complete; so in the analysis proposed by~\\citet{Wald:1945,Wald:1947}, one simply takes a ``best guess'' at the lower bound of the signal strength.  In this paper, we extend Wald's technique to marginalize the signal strength, which more rigorously accounts for our ignorance of the true signal.  As in the classical likelihood ratio test, this extended test can be applied after each new event without statistical penalty, so that it adheres to the likelihood principle.  It also allows for the evaluation of the significance of the signal after the test has been fulfilled, as well as in cases where the test stops prematurely.  We note that the usefulness of this test is not limited to cosmic ray physics. It can be applied in many other areas of astroparticle physics or astrophysics where event rates are low, for example in searches for discrete sources of high energy neutrinos or $\\gamma$-rays. ",
        "conclusions": "\\label{sec:summary}  We have outlined a sequential analysis technique for testing a point null hypothesis with probability $p_0$ against a signal probability $p$.  The method is based on the sequential analysis proposed in~\\citet{Wald:1945,Wald:1947}, but replacing the likelihood ratio used to evaluate the significance of a signal with one that marginalizes the signal strength.  In many sequential tests, the signal strength is unknown when the test starts.  Typically, the signal probability $p$ can in principle have any value in the interval $\\left[p_0,1\\right]$.  Rather than choosing a fixed threshold for $p$, as suggested in~\\citet{Wald:1945,Wald:1947}, we have argued that, in general, the better alternative is to marginalize $p$ and account for our ignorance exactly.  In the marginalization of the signal likelihood, the integration starts at some value $p_1=p_0+\\delta$, where $\\delta$ is an \\textsl{ad hoc} parameter reflecting the experimenter's belief about the strength of the signal, the capability of his experiment, and other \\textsl{a priori} knowledge.  Because of the integration of the signal likelihood over a range in $p$, the parameters $\\alpha$ and $\\beta$ have lost their intuitive meaning if the method is applied to data sets where $p$ is fixed, as is typically the case for real data.  However, we have shown that for most values of $\\delta$ and $p$ that occur in correlation searches, the type-1 and type-2 error rates of the sequential analysis are consistent with the classical interpretations of the probabilities $\\alpha$ and $\\beta$.  Note that we have run a test with one of two outcomes (i.e., an acceptance or rejection of $\\mathcal{H}_0$), defining $\\alpha$ and $\\beta$, rather than one outcome (say, only a rejection of $\\mathcal{H}_0$) such as in~\\citet{Darling:1968}.  The latter case supposes that we are only concerned about reporting a signal. However, it is important to state a null result at some point in the interest of reducing reporting bias. That is, it is important to ensure that 1\\% of the results that claim an excess of events are indeed a 1\\% effect.  The sequential analysis technique proposed here is efficient, allows the signal significance to be evaluated after the test has been fulfilled, adheres to the likelihood principle, and rigorously accounts for our ignorance of the signal strength."
    },
    "0711/0711.4726_arXiv.txt": {
        "abstract": "We study effects of kinetic helicity fluctuations in a turbulence with large-scale shear using two different approaches: the spectral tau-approximation and the second order correlation approximation (or first-order smoothing approximation). These two approaches demonstrate that homogeneous kinetic helicity fluctuations alone with zero mean value in a sheared homogeneous turbulence cannot cause large-scale dynamo. Mean-field dynamo can be possible when kinetic helicity fluctuations are inhomogeneous which cause a nonzero mean alpha effect in a sheared turbulence. On the other hand, shear-current effect can generate large-scale magnetic field even in a homogeneous nonhelical turbulence with large-scale shear. This effect was investigated previously for large hydrodynamic and magnetic Reynolds numbers. In this study we examine the threshold required for the shear-current dynamo versus Reynolds number. We demonstrate that there is no need for a developed inertial range in order to maintain the shear-current dynamo (e.g., the threshold in the Reynolds number is of the order of 1). ",
        "introduction": "It has been widely recognized that astrophysical large-scale magnetic fields originate due to the mean-field dynamo (see, e.g., \\cite{M78,P79,KR80,ZRS83,RSS88,O03,BS05}). Such dynamo can be driven by the joint action of the mean kinetic helicity of turbulence and large-scale differential rotation. On the other hand, recently performed numerical experiments \\cite{B05,YHS07,BRR07} have demonstrated existence of a nonhelical large-scale dynamo in a turbulence with a large-scale shear whereby mean $\\alpha$ effect vanishes. Note that a sheared turbulence is a universal feature in astrophysical \\cite{O03,BS05,DCG05} and laboratory \\cite{MP07} flows.  One of the possible mechanism of the nonhelical large-scale dynamo in a homogeneous sheared turbulence is a shear-current effect that has been extensively studied during recent years (see \\cite{RK03,RK04,RKL06,RKCL06}). In particular, the deformations of the original nonuniform magnetic field lines are caused by upward and downward turbulent eddies. In a sheared turbulence the inhomogeneity of the original mean magnetic field breaks a symmetry between the influence of the upward and downward turbulent eddies on the mean magnetic field. This creates the mean electric current along the mean magnetic field and results in the nonhelical shear-current dynamo. Indeed, the large-scale velocity shear creates anisotropy of turbulence that produces a contribution to the electromotive force, ${\\bf W} {\\bf \\times} {\\bf J}$, caused by the shear, where ${\\bf W}$ is the background large-scale vorticity due to the shear and  ${\\bf J}$ is the large-scale electric current. Joint effects of the electromotive force ${\\bf W} {\\bf \\times} {\\bf J}$ and stretching of the mean magnetic field due to the large-scale shear motions cause the mean-field dynamo instability. Note also that the numerical experiment with Taylor-Green forcing \\cite{PON05} seems to be another example of a mean-field dynamo produced by a combined effect of a nonhelical turbulence and a complicated large-scale flow.  Other effect that might explain the nonhelical large-scale dynamo is related to  kinetic helicity fluctuations in a sheared turbulence. A problem associated with dynamics of large-scale magnetic field in the presence of kinetic helicity fluctuations in a shear-free turbulence has been formulated for the first time by Kraichnan \\cite{K76} (see also \\cite{M78}). In particular, he assumed that in small scales $l_\\nu \\ll l_{\\rm turb} \\ll l_0$ (and $\\tau_\\nu \\ll \\tau_{\\rm turb} \\ll \\tau_0)$ there is a small-scale turbulence generated by forcing ${\\bf F}^{(u)}$. On the other hand, in the scales $l_0 \\ll l \\ll l_\\chi$ (and $\\tau_0 \\ll \\tau \\ll \\tau_\\chi$) there are kinetic helicity fluctuations (or $\\tilde \\alpha$ fluctuations) with a zero mean generated by forcing ${\\bf F}^{(\\chi)}$. The mean-field effects occur at large scales $L \\gg l_\\chi$ (and times $\\tau_{_{L}} \\gg \\tau_\\chi$), where the large-scale helicity is zero. The large-scale quantities are determined by double averaging over velocity fluctuations and over kinetic helicity fluctuations. It was found in \\cite{K76} (see also \\cite{M78}) that kinetic helicity fluctuations in a shear-free turbulence cause both, a reduction of turbulent magnetic diffusion and a large-scale drift velocity, ${\\bf V}^{(\\alpha)} \\propto \\bec{\\nabla} \\langle \\tilde{\\alpha}^2 \\rangle$, of the mean magnetic field.  In recent time the dynamo problem related to kinetic helicity fluctuations has been extensively studied. In particular, various mathematical aspects of this problem has been discussed in \\cite{SOK97} (see also \\cite{F03}). Some numerical experiments which examine effects of kinetic helicity fluctuations have been performed in \\cite{VB97,BRR07}. In particular, numerical simulations of the magnetic field evolution in accretion discs in \\cite{VB97} have demonstrated that kinetic helicity fluctuations with a zero mean can result in generation of large-scale magnetic field.  It has been pointed out in \\cite{SIL2000} that inhomogeneous kinetic helicity fluctuations in a sheared turbulence can produce a mean-field dynamo. In particular, a joint action of a large-scale shear and a nonzero mean alpha effect caused by the inhomogeneous kinetic helicity fluctuations can result in generation of a large-scale magnetic field, where the mean alpha effect in a sheared turbulence is proportional to ${\\nabla} \\langle \\tilde{\\alpha}^2 \\rangle$. This mean-field dynamo is similar to the $\\alpha\\Omega$-dynamo. On the other hand, it has been suggested in \\cite{PROC07} using phenomenological arguments that homogeneous kinetic helicity fluctuations in a homogeneous turbulence with shear may generate a large-scale magnetic field.  The main goal of this study is to examine a possibility for a nonhelical large-scale dynamo associated with homogeneous kinetic helicity fluctuations with a zero mean in a homogeneous turbulence with a large-scale shear. In a study of a sheared turbulence we use two different approaches, namely, the spectral tau-approximation \\cite{RK04} and second order correlation approximation (SOCA, sometimes refers in astrophysical literature as first-order smoothing approximation (FOSA), see, e.g., \\cite{RAS06,RUK06}). We also investigate dynamo effects associated with inhomogeneous kinetic helicity fluctuations in a sheared homogeneous turbulence.  This paper is organized as follows. In Sec.~II we formulate the governing equations and outline the procedure of derivation based on the $\\tau$-approach, that allows us to determine a contribution to the mean electromotive force caused by a combined action of the sheared turbulence and the kinetic helicity fluctuations. In Sec.~III we study the effects of kinetic helicity fluctuations in a sheared turbulence using the $\\tau$-approach. In Sec.~IV we investigate similar effects using the SOCA-approach. In Sec.~V we discuss the threshold required for the shear-current dynamo versus Reynolds number. In Sec.~VI we draw concluding remarks. Finally, the detailed derivations of the effects of kinetic helicity fluctuations in a sheared turbulence using the $\\tau$-approach and the SOCA-approach have been performed in Appendixes A and B, respectively. ",
        "conclusions": "We study effects of kinetic helicity fluctuations in a homogeneous turbulence with large-scale shear using the spectral tau-approximation and the second order correlation approximation (SOCA). We show that homogeneous kinetic helicity fluctuations  alone with a zero mean cannot cause a large-scale dynamo in a sheared turbulence. This negative result based on SOCA and tau approach, is a quantitative one: the sign of a certain coefficient in the mean electromotive force $(\\propto -\\sigma_{\\alpha})$ turns out to be unfavorable for the mean-field dynamo. This result is in a contradiction to that suggested in \\cite{PROC07} using phenomenological arguments. In order to compare with the results obtained in \\cite{PROC07}, we rewrite Eqs.~(\\ref{E2}) and~(\\ref{E3}) in the following form: \\begin{eqnarray} {\\partial \\bar A \\over \\partial t} &=& \\sigma_{_{S}} \\, S \\, l_0^2 \\, \\bar B'_y + \\tilde \\eta_{_{T}} \\, \\bar A''  \\;, \\label{EE2}\\\\ {\\partial \\bar B_y \\over \\partial t} &=& - S \\, \\bar A' + \\tilde \\eta_{_{T}} \\, \\bar B''_y \\;, \\label{EE3} \\end{eqnarray} where the mean magnetic field, $\\bar{\\bf B} = \\bar B_y(z) \\, {\\bf e}_y + \\bec{\\nabla} {\\bf \\times} [\\bar A(z) \\, {\\bf e}_y]$. Equation~(\\ref{EE3}) is similar to Eq.~(8) derived in \\cite{PROC07}, whereby $\\Omega'$ is replaced by $-S$, while Eq.~(\\ref{EE2}) is similar to Eq.~(9) derived in \\cite{PROC07}.  The first term in the right hand side of Eq.~(\\ref{EE2}) determines two different effects, namely the shear-current effect $\\propto \\sigma_{_{B}} \\, S \\, \\bar B'_y$ and the effect $\\propto -\\sigma_{\\alpha}\\, S \\, \\bar B'_y$ caused by the homogeneous kinetic helicity fluctuations in a sheared turbulence. The shear-current effect  results in the mean-field dynamo, while the second effect $(\\propto -\\sigma_{\\alpha})$ acts against the mean-field dynamo (see Eqs.~(\\ref{BB13})-(\\ref{B15}) and~(\\ref{BB15})). These two effects can be interpreted as the off-diagonal terms in the tensor of turbulent magnetic diffusion, and they cannot be reduced to the standard $\\alpha$ effect.  Note that the first term in the right hand side of Eq.~(\\ref{EE2}) that determines the effect of homogeneous kinetic helicity fluctuations in a sheared turbulence, is similar to the first term in the right hand side of Eq.~(9) derived in \\cite{PROC07} except for it has opposite sign. The latter is crucial for the mean-field dynamo. In this comparison we have not taken into account the shear-current effect determined by the coefficient $\\sigma_{_{B}}$, because this effect has not been considered in \\cite{PROC07}. The results obtained in the present study are derived using the rigorous mean field theory based on SOCA (see Sec.~IV). These results are also in a good agreement with those obtained by the tau-approach (see Sec.~III). On the other hand, the results of derivation performed in \\cite{PROC07} by phenomenological arguments using ad hoc mean-field equations are in a disagreement with our results.  The shear-current effect causes the mean-field dynamo in a homogeneous nonhelical turbulence with imposed large-scale shear. This effect has been studied previously (see \\cite{RK03,RK04,RKL06}) only for large hydrodynamic and magnetic Reynolds numbers. In the present study we determine the threshold required for the shear-current dynamo as a function of Reynolds number and demonstrate that threshold value of the Reynolds number is of the order of 1. This implies that there is no need for a developed inertial range in order to maintain the shear-current dynamo.  In the present study we also recover the results obtained in \\cite{K76} (see also \\cite{M78}) for a shear-free turbulence, whereby a negative contribution of kinetic helicity fluctuations to the turbulent magnetic diffusion, $\\eta^{(\\alpha)}_{_{T}} = - \\tau_\\chi \\, \\langle \\tilde{\\alpha}^2 \\rangle^{(\\alpha)}$, and a large-scale drift velocity of the mean magnetic field, ${\\bf V}^{(\\alpha)} \\propto \\tau_\\chi \\, \\bec{\\nabla} \\langle \\tilde{\\alpha}^2 \\rangle^{(\\alpha)}$, have been found.  On the other hand, we have demonstrated that inhomogeneous kinetic helicity fluctuations in a sheared turbulence cause a nonzero mean alpha effect, $\\bar{\\alpha}^{(S,\\alpha)} \\propto - \\tau_\\chi^2 \\, S \\, {\\nabla}_z \\langle \\tilde{\\alpha}^2 \\rangle^{(\\alpha)}$, where the mean vorticity due to the large-scale shear is ${\\bf W}^{(S)} = S \\, {\\bf e}_z$. The mean alpha effect $\\bar{\\alpha}^{(S,\\alpha)}$ is formed by a combined action of a large-scale shear in turbulent flow and inhomogeneous kinetic helicity fluctuations even when $\\langle \\tilde{\\alpha} \\rangle^{(\\alpha)}=0$. The large-scale shear and the mean alpha effect can cause a mean-field dynamo (see \\cite{SIL2000}) that is similar to the $\\alpha\\Omega$-dynamo.  The discussed effects in this study might be important in astrophysics (e.g., accretion discs, colliding protogalactic clouds, merging protostellar clouds \\cite{RKCL06}) and laboratory dynamo experiments. In particular, non-symmetrical explosions of supernova may produce fluctuations of kinetic helicity located in larger scales than  small-scale turbulence existing in convective zones inside stars. On the other hand, the shear-current dynamo acts together with the $\\alpha$-shear dynamo which is similar to the $\\alpha \\Omega$ dynamo. The shear-current effect does not quenched (see \\cite{RK04,RKL06}) contrary to the quenching of the nonlinear $\\alpha$ effect, the turbulent magnetic diffusion, the effective drift velocity. This implies that the shear-current effect might be the only surviving effect, which can explain the origin of large-scale magnetic fields in sheared astrophysical turbulence."
    },
    "0711/0711.4984_arXiv.txt": {
        "abstract": "We report the results for rotational velocities, $V_{\\rm rot}$~sin~$i$, and macroturbulence dispersions, $\\zeta_{\\rm RT}$, for 12 metal-poor field red giant branch (RGB) stars and 7 metal-poor field red horizontal branch (RHB) stars. The results are based on Fourier transform analyses of absorption line profiles from high-resolution ($R \\approx\\ 120,000$), high-S/N ($\\approx\\ 215$ per pixel; $\\approx\\ 345$ per resolution element) spectra obtained with the Gecko spectrograph at CFHT. The stars were selected from the studies of 20 RHB and 116 RGB stars from Carney et al.\\ (2003, 2007), based primarily on larger-than-average line broadening values. We find that $\\zeta_{\\rm RT}$ values for the metal-poor RGB stars are very similar to those for metal-rich disk giants studied earlier by Gray and his collaborators. Six of the RGB stars have small rotational values, less than 2.0 \\kms, while five show significant rotation/enhanced line broadening, over 3 \\kms. We confirm the rapid rotation rate for RHB star HD~195636, found earlier by Preston (1997). This star's rotation is comparable to that of the fastest known rotating blue horizontal branch (BHB) stars, when allowance is made for differences in radii and moments of inertia. The other six RHB stars have somewhat lower rotation but show a trend to higher values at higher temperatures (lower radii). Comparing our results with those for BHB stars from Kinman et al.\\ (2000), we find that the fraction of rapidly rotating RHB stars is somewhat lower than is found among BHB stars. The number of rapidly rotating RHB stars is also smaller than we would have expected from the observed rotation of the RGB stars. We devise two empirical methods to translate the line broadening results obtained by Carney et al.\\ (2003, 2007) into $V_{\\rm rot}$~sin~$i$ for all the RGB and RHB stars they studied. Binning the RGB stars by luminosity, we find that most metal-poor field RGB stars show no detectable sign, on average, of rotation, which is not surprising given the stars' large radii. However, the most luminous stars, with $M_{\\rm V} \\leq\\ -1.5$, do show net rotation, with mean values of 2 to 4 \\kms, depending on the algorithm employed, and also show signs of radial velocity jitter and mass loss. This ``rotation\" may in fact prove to be due to other line broadening effects, such as shock waves or pulsation. ",
        "introduction": "In previous papers (Carney et al.\\ 2003, Carney et al.\\ 2007; hereafter C2003, C2007) we reported on radial velocities and line broadenings for 136 metal-poor field red giant branch (RGB) and red horizontal branch (RHB) stars, based on 2413 high-resolution, low-S/N spectra. One of the more intriguing results was that the more luminous red giants, as well as many of their evolutionary progeny, red horizontal branch stars, showed significant line broadening. Interpreting the enhanced line broadening as rotation, C2003 explored the possibility that it might have arisen from absorption of one (or more) jovian mass planets that were engulfed only as the red giants swelled to large enough radii.  C2003 suggested several follow-up studies. First, the sample size should be expanded, and C2007 presented the results for 45 stars to complement the original 91-star sample of C2003. Second, high-precision radial velocity monitoring of metal-poor dwarfs and subgiants should be undertaken to explore the frequency of jovian-mass planets with orbital periods of order one year, corresponding to aphelion distances comparable to the maximum radial size of metal-poor RGB stars. The initial results of such a study have been reported (Sozzetti et al.\\ 2006), and it appears that such planetary companions are not sufficiently common to explain the modest frequency of significant line broadening among the most luminous metal-poor red giants.  This paper explores the separate contributions of rotation and macroturbulence, based on a selected subsample of stars studied by C2003 and C2007. For example, if macrotubulence is a strong function of luminosity, the enhanced line broadenings found by C2003 and C2007 among the field red giants might be explained. If macroturbulence is a strong function of temperature, perhaps the line broadening seen in some of the red horizontal branch stars could also be explained. But if rotation is the cause of the enhanced line broadening among the stars with the largest radii, some new explanations must be sought.  To test these hypotheses, we decided to exploit methods developed by Gray (1982), whereby line profiles measured using very high-resolution, high-S/N spectra could be analyzed via Fourier transform methods to distinguish the contributions of rotation and macroturbulence. ",
        "conclusions": "We have obtained high-resolution, high-S/N spectra for 12 metal-poor field RGB stars and 7 metal-poor field RHB stars. Fourier transform analyses have yielded good estimates for the macroturbulence dispersion, $\\zeta_{\\rm RT}$, and the rotation velocity, $V_{\\rm rot}$~sin~$i$, for all the stars. We obtained consistent results for HD~29574, which was observed during both observing runs, and for $\\eta$~Ser, which had been studied previously by Gray \\& Pallavicini (1989). It is good to recall that we selected our stars from the C2003 and C2007 samples of 116 RGB stars and 20 RHB stars, and with a bias toward stars with larger line broadening values (referred to as rotational velocities by C2003). The RGB stars appear to show very similar macroturbulence behavior as a function of luminosity, gravity, and temperature as do disk giants. The twelve RGB stars studied here, however, show a larger range in rotational velocities, with half showing values of 2.0 \\kms\\ or less, and five having values in excess of 3.0 \\kms. For the seven field RHB stars, we confirm the rapid rotation of HD~195636 discovered by Preston (1997). When allowance is made for the star's larger radius compared to BHB stars, its rotational angular momentum is comparable to the largest seen in field BHB stars.  We have explored the use of two empirical methods, neither justified physically, to attempt to exploit the more extensive data on line broadening from C2003 and C2007. The derived rotational velocities of the 13 field RHB stars are, as expected from the CfA results, modest. To compare our results with the much lower resolution BHB data from Kinman et al.\\ (2000), we consider the percentages of BHB stars whose rotation rates are comparable to the spectral resolution, and then make allowances for the differences in radii between the BHB stars and our sample of RHB stars. We find that the RHB stars have fewer rapidly rotating stars than does the BHB sample studied by Kinman et al.\\ (2000).  Both algorithms were applied to binned samples of the field RGB stars, so that statistical fluctuations would be diminished. All but the bin containing the most luminous stars ($M_{\\rm V} \\leq\\ -1.5$) showed nearly zero mean rotation, as expected for the large radii. It is clear that the most luminous metal-poor field RGB stars show enhanced rotation or some other source of line broadening, as found initially by C2003. Unlike the results from C2003, however, the line broadening is relatively modest, on average being 2 \\kms\\ or 4 \\kms, depending on which algorithm is employed. Our CFHT observations did not extend to the lower luminosities, but are consistent with this result. The twelve RGB stars studied had $<M_{\\rm rot}>$ = $-2.0$ and $<V_{\\rm rot}$~sin~$i>$ = 2.3 \\kms.  We draw attention to the fact that the transition in luminosity between negligible and significent rotation occurs at about the same $M_{\\rm V}$ value as the appearance of velocity jitter and mass loss. This is highly suggestive of a common underlying physical origin, which may be a sign of shock waves and acoustic heating of chromospheres (Cuntz et al.\\ 1994) or pulsation (Smith \\& Dupree 1988). We recommend a dedicated monitoring program of a few key stars to compare timescales and the presence of line broadening, velocity variations, and mass loss.  We acknowledge financial support from the National Science Foundation to the University of North Carolina through grant AST0305431 and to Bowling Green State University through grant AST0307340.  \\clearpage  \\begin{deluxetable}{lrrrrrrrl} \\tablewidth{0pc} \\tablenum{1} \\footnotesize \\tablecaption{Observational Data \\label{tab1}} \\tablehead{ \\colhead{Star} & \\colhead{$\\lambda$} & \\colhead{HJD$-$2,450,000} & \\colhead{$V$} & \\colhead{Exp} & \\colhead{S/N\\tablenotemark{a}} & \\colhead{$V_{\\rm rad}$} & \\colhead{$\\sigma$} & \\colhead{Comments} } \\startdata HD~3008    & 6150 & 3366.7716 & 9.70 & 40 & 175 & $-81.83$ & 0.34 & RGB; CM~Cet; jitter \\nl BD$-$18~271 & 6150 & 4015.8939  & 9.85 & 90 & 145 & $-210.54$ & 0.19 & RGB; jitter \\nl CD$-36$~1052 & 5430 & 3366.8083 & 10.00 & 70 & 150 & $+304.36$ & 0.32 & RHB \\nl HD~23798   & 6150 & 3366.8681 & 8.32 & 25 & 210 & $+88.83$ & 0.54 & RGB \\nl HD~25532   & 5430 & 4016.9720 & 8.24 & 80 & 175 & $-111.88$ & 0.34 & RHB \\nl BD+6~648   & 6150 & 4015.1177 & 9.09 & 270 & 285 & $-142.41$ & 0.31 & RGB \\nl HD~29574   & 6150 & 3366.8934 & 8.38 & 60 & 155 & $+17.86$ & 0.37 & RGB; HP~Eri; jitter \\nl & 6150 & 4015.0694 &      & 35 & 250 & $+17.67$ & 0.48 & \\nl HD~82590   & 5430 & 3366.9454 & 9.42 & 75 & 180 & $+214.31$ & 0.37 & RHB; NSV~4526 \\nl BD+22~2411 & 6150 & 3367.0135 & 9.95 & 90 & 160 & $+35.05$ & 0.26 & RGB; jitter \\nl HD~106373  & 5430 & 3367.1146 & 8.91 & 75 & 190 & $+83.68$ & 0.21 & RHB \\nl HD~110281  & 6150 & 3367.0801 & 9.39 & 45 & 170 & $+139.90$ & 0.51 & RGB; KR~Vir; jitter \\nl HD~165195  & 6150 & 4011.7779 & 7.34 & 120 & 380 & $+0.50$ & 0.40 & RGB; V2564~Oph; jitter \\nl HD~184266  & 5430 & 4016.7436 & 7.57 & 50 & 250 & $-349.20$ & 0.39 & RHB \\nl HD~187111  & 6150 & 4015.8079 & 7.75 & 40 & 395 & $-186.16$ & 0.15 & RGB \\nl HD~195636  & 5430 & 4016.7877 & 9.57 & 140 & 130 & $-258.34$ & 1.62 & RHB \\nl HD~214925  & 6150 & 4015.7595 & 9.30 & 50 & 215 & $-327.26$ & 0.60 & RGB; jitter \\nl HD~214362  & 5430 & 4016.9199 & 9.10 & 60  & 105 & $-92.48$  & 0.14 & RHB \\nl HD~218732  & 6150 & 4015.8478 & 8.47 & 40 & 270 & $-294.24$ & 0.25 & RGB; LS~Aqr; SLSB; jitter \\nl HD~221170  & 6150 & 4009.8693 & 7.71 & 100 & 260 & $-121.69$ & 0.40 & RGB; NSV~14589 \\nl $\\eta$~Ser & 6150 & 3981.7138 & 3.26 & 20 & 300 & $+12.24$ & 0.15 & RGB; check; NSV~10675 \\nl \\tablevspace{4pt} \\enddata \\tablenotetext{a}{S/N values are per pixel. A typical resolution element is 2.3 pixels.} \\end{deluxetable}  \\clearpage  \\begin{deluxetable}{lrrrrrrrl} \\tablewidth{0pc} \\tablenum{2} \\footnotesize \\tablecaption{Line Broadening Results \\label{tab2}} \\tablehead{ \\colhead{Star} & \\colhead{$M_{\\rm V}$} & \\colhead{$T_{\\rm eff}$} & \\colhead{log $g$} & \\colhead{[Fe/H]} & \\colhead{$V_{\\rm broad}$} & \\colhead{$V_{\\rm rot}$ sin $i$} & \\colhead{$\\zeta_{\\rm RT}$} & \\colhead{Comments} } \\startdata HD~3008 & $-1.5$ & 4140 & 1.00 & $-1.43$ & 9.2 & $4.4 \\pm 0.8$ & $6.9 \\pm 0.6$ & RGB \\nl BD$-18$~271 & $-2.1$ & 4150 & 0.70 & $-1.98$ & 7.3 & $0.0 \\pm 1.5$ & $7.2 \\pm 1.0$ & RGB \\nl CD$-36$~1052 & $+0.62$ & 5890 & 2.50 & $-2.00$ & 14.4 & $8.8 \\pm 0.8$ & $8.9 \\pm 0.8$ & RHB \\nl HD~23798 & $-1.8$ & 4310 & 1.00 & $-1.90$ & 5.0 & $0.0 \\pm 1.0$ & $6.7 \\pm 0.6$ & RGB \\nl HD~25532 & $+0.79$ & 5320 & 2.54 & $-1.33$ & 8.5 & $4.8 \\pm 1.0$ & $7.6 \\pm 0.7$ & RHB \\nl BD+6~648 & $-1.79$ & 4160 & 0.87 & $-1.82$ & 6.1 & $1.2 \\pm 1.5$ & $6.6 \\pm 1.0$ & RGB \\nl HD~29574  & $-2.11$ & 3960 & 0.57 & $-2.11$ & 10.2 & $3.7 \\pm 1.0$ & $7.4 \\pm 0.6$ & RGB \\nl & &      &      &         &      & $4.2 \\pm 0.7$ & $7.3 \\pm 0.5$ & \\nl HD~82590  & $+0.7$ & 5960 & 2.70 & $-1.85$ & 13.0 & $7.7 \\pm 0.6$ & $11.0 \\pm 0.5$ & RHB \\nl BD+22~2411 & $-1.7$ & 4320 & 1.00 & $-1.95$ & 7.3 & $0.0 \\pm 2.0$ & $7.5 \\pm 0.7$ & RGB \\nl HD~106373 & $+0.57$ & 6160 & 2.70 & $-2.48$ & 13.5 & $10.8 \\pm 0.7$ & $6.5 \\pm 1.5$ & RHB \\nl HD~110281 & $-2.6$ & 3850 & 0.20 & $-1.75$ & 11.5 & $5.5 \\pm 1.0$ & $6.2 \\pm 1.0$ & RGB \\nl HD~165195 & $-2.14$ & 4200 & 0.76 & $-2.16$ & 7.6 & $1.8 \\pm 0.7$ & $6.4 \\pm 0.5$ & RGB \\nl HD~184266 & $+0.7$ & 5490 & 2.60 & $-1.87$ & 11.7 & $5.0 \\pm 0.5$ & $11.5 \\pm 0.3$ & RHB \\nl HD~187111 & $-1.54$ & 4260 & 1.04 & $-1.65$ & 5.2 & $2.4 \\pm 0.5$ & $5.5 \\pm 0.7$ & RGB \\nl HD~195636 & $+0.5$ & 5370 & 2.40 & $-2.40$ & 21.5 & $22.2 \\pm 1.0$ & $10.0 \\pm 1.5$ & RHB \\nl HD~214925 & $-2.5$ & 3890 & 0.30 & $-2.14$ & 9.7 & $4.5 \\pm 0.7$ & $8.4 \\pm 0.4$ & RGB \\nl HD~214362 & $+0.6$ & 5700 & 2.60 & $-2.20$ & 11.1 & $7.5 \\pm 1.0$ & $8.5 \\pm 0.7$ & RHB \\nl HD~218732 & $-2.8$ & 3900 & 0.20 & $-2.00$ & 11.1 & $3.1 \\pm 0.5$ & $6.4  \\pm 0.7$ & RGB \\nl HD~221170 & $-1.7$ & 4410 & 1.10 & $-1.56$ & 7.4 & $1.0 \\pm 1.0$ & $6.4 \\pm 0.5$ & RGB \\nl $\\eta$~Ser & $+1.87$ & 4890 & 3.21 & $-0.42$ & 4.0 & $1.0 \\pm 0.8$ & $4.1 \\pm 0.5$ & disk RGB \\nl \\tablevspace{4pt} \\enddata \\end{deluxetable}  \\clearpage  \\begin{deluxetable}{lllrrrrrr} \\tablewidth{0pc} \\tablenum{3} \\footnotesize \\tablecaption{Disk Giants \\label{tab3}} \\tablehead{ \\colhead{Star} & \\colhead{HR} & \\colhead{Sp Type} & \\colhead{[Fe/H]} & \\colhead{$T_{\\rm eff}$} & \\colhead{log $g$} & \\colhead{$M_{\\rm V}$} & \\colhead{$V_{\\rm rot}$ sin $i$} & \\colhead{$\\zeta_{\\rm RT}$} } \\startdata $\\xi$ Cyg & 8079 & K4.5 Ib & $-0.45$ & 4090 & 1.42 & $-4.60_{-0.37}^{+0.45}$ & $1.6 \\pm 2.0$ & $10.1 \\pm 1.0$ \\nl 58 Per & 1454 & G8 II & $-0.29$ & 4260 & 2.21 & $-2.22_{-0.38}^{+0.47}$ & $6.3 \\pm 1.4$ & $10.2 \\pm 0.9$ \\nl $\\theta$ Lyr & 7314 & K0 II & $-0.01$ & 4500 & 1.93 & $-2.76_{-0.24}^{+0.27}$ & $3.6 \\pm 1.4$ & $5.6 \\pm 0.9$ \\nl $\\theta$ Her & 6695 & K1 II & $-0.24$ & 4330 & 1.28 & $-2.71_{-0.23}^{+0.26}$ & $3.4 \\pm 0.6$ & $7.9 \\pm 0.2$ \\nl $\\lambda$ Lyr & 7192 & K2.5 II & $-0.02$ & 4220 & 2.21 & $-3.75_{-0.50}^{+0.65}$ & $3.2 \\pm 1.0$ & $6.3 \\pm 0.5$ \\nl $\\pi$ Her & 6418 & K3 II & $-0.18$ & 4100 & 1.68 & $-2.10_{-0.12}^{+0.13}$ & $3.7 \\pm 0.1$ & $4.9 \\pm 0.6$ \\nl $\\beta$ Cyg & 7417 & K3 II & $-0.17$ & 4270 & 1.79 & $-2.27_{-0.14}^{+0.15}$ & $3.0 \\pm 0.3$ & $6.1 \\pm 0.2$ \\nl $\\gamma$ Aql & 7525 & K3 II & $-0.29$ & 4210 & 1.63 & $-3.38_{-0.22}^{+0.24}$ & $3.2 \\pm 0.5$ & $7.0 \\pm 0.4$ \\nl 1 Lac & 8498 & K3 II-III & $-0.12$ & 4350 & 1.75 & $-2.61_{-0.24}^{+0.27}$ & $3.6 \\pm 1.4$ & $5.6 \\pm 0.9$ \\nl $\\alpha$ Hya & 3748 & K3 II-III & $-0.12$ & 4120 & 1.77 & $-1.68_{-0.09}^{+0.09}$ & $0.0 \\pm 1.4$ & $6.5 \\pm 0.9$ \\nl $\\beta$ Sge & 7488 & G8 IIIa & $-0.03$ & 4850 & 2.79 & $-1.39_{-0.20}^{+0.22} $ & $9.1 \\pm 0.7$ & $8.2 \\pm 0.7$ \\nl $\\lambda$ Peg & 8667 & G8 IIIa & $-0.10$ & 4800 & 3.20 & $-1.46_{-0.18}^{+0.19}$ & $7.8 \\pm 0.4$ & $7.8 \\pm 0.4$ \\nl $\\beta$ Her & 6148 & G8 III & $-0.27$ & 4920 & 2.62 & $-0.49_{-0.10}^{+0.10}$ & $3.4 \\pm 1.0$ & $6.8 \\pm 1.0$ \\nl $\\eta$ Dra & 6132 & G8 III & $-0.21$ & 4940 & 3.10 & $+0.59_{-0.03}^{+0.03}$ & $2.2 \\pm 1.0$ & $5.5 \\pm 1.0$ \\nl $\\alpha^{2}$ Cap & 7754 & G8 IIIb & $-0.18$ & 5000 & 3.05 & $+0.98_{-0.06}^{+0.07}$ & $3.2 \\pm 0.4$ & $4.6 \\pm 0.3$ \\nl $\\alpha$ Cas & 168 & K0 III & $-0.09$ & 4610 & 2.71 & $-1.98_{-0.09}^{+0.09}$ & $4.9 \\pm 0.4$ & $6.2 \\pm 0.3$ \\nl $\\delta$ Tau & 1373 & K0 III & 0.00 & 4940 & 2.85 & $+0.40_{-0.09}^{+0.10}$ & $2.5 \\pm 1.0$ & $6.2 \\pm 1.0$ \\nl $\\gamma$ Tau & 1346 & K0 III & $-0.02$ & 4930 & 2.90 & $+0.28_{-0.12}^{+0.12} $ & $2.4 \\pm 1.0$ & $5.9 \\pm 1.0$ \\nl $\\beta$ Cet & 188 & K0 III & $-0.09$ & 4820 & 2.87 & $-0.30_{-0.05}^{+0.05}$ & $3.0 \\pm 1.0$ & $5.9 \\pm 1.0$ \\nl $\\theta^{1}$ Tau & 1411 & K0 III & $+0.04$ & 4960 & 3.17 & $+0.43_{-0.09}^{+0.09}$ & $3.4 \\pm 1.0$ & $4.9 \\pm 1.0$ \\nl $\\beta$ Gem & 2990 & K0 III & $-0.04$ & 4850 & 2.75 & $+1.08_{-0.02}^{+0.02}$ & $2.5 \\pm 1.0$ & $4.2 \\pm 1.0$ \\nl $\\epsilon$ Cyg & 7949 & K0 III & $-0.27$ & 4730 & 2.89 & $+0.78_{-0.03}^{+0.03}$ & $3.0 \\pm 1.0$ & $4.2 \\pm 1.0$ \\nl $\\alpha$ UMa & 4301 & K0 III & $-0.20$ & 4660 & 2.46 & $-1.10_{-0.04}^{+0.04}$ & $2.6 \\pm 1.0$ & $5.2 \\pm 1.0$ \\nl $\\gamma^{1}$ Leo & 4057 & K2 III & $-0.49$ & 4470 & 2.35 & $-0.32_{-0.07}^{+0.07}$ & $2.6 \\pm 1.0$ & $5.2 \\pm 1.0$ \\nl $\\alpha$ Ari & 617 & K2 III & $-0.25$ & 4480 & 2.57 & $+0.47_{-0.04}^{+0.04}$ & $3.1 \\pm 1.0$ & $3.9 \\pm 1.0$ \\nl $\\beta$ Oph & 6603 & K2 III & $+0.02$ & 4550 & 2.63 & $+0.77_{-0.04}^{+0.04}$ & $1.6 \\pm 1.0$ & $4.0 \\pm 1.0$ \\nl $\\alpha$ Ser & 5854 & K2 III & $+0.03$ & 4530 & 2.76 & $+0.88_{-0.03}^{+0.03}$ & $0.0 \\pm 1.0$ & $4.8 \\pm 1.0$ \\nl $\\alpha$ Boo & 5340 & K2 III & $-0.60$ & 4200 & 2.19 & $-0.30_{-0.02}^{+0.02}$ & $2.4 \\pm 1.0$ & $5.2 \\pm 1.0$ \\nl $\\eta$ Ser & 6869 & K2 III & $-0.42$ & 4890 & 3.21 & $+1.87_{-0.03}^{+0.03}$ & $2.0 \\pm 0.5$ & $4.0 \\pm 0.5$ \\nl $\\epsilon$ Sco & 6241 & K2.5 III & $-0.17$ & 4560 & 2.49 & $+0.78_{-0.04}^{+0.04}$ & $2.6 \\pm 0.5$ & $4.3 \\pm 0.5$ \\nl $\\epsilon$ Crv & 4630 &  K2.5 III & $+0.13$ & 4320 & 2.16 & $-1.82_{-0.14}^{+0.15}$ & $2.6 \\pm 0.5$ & $5.8 \\pm 0.5$ \\nl 26 Aql & 7333 & G8 III-IV & $-0.21$ & 4900 & 3.31 & $+1.63_{-0.08}^{+0.08}$ & $2.8 \\pm 0.5$ & $5.0 \\pm 0.5$ \\nl \\tablevspace{4pt} \\enddata \\end{deluxetable}  \\clearpage  \\begin{deluxetable}{lrrrrrrrrr} \\tablewidth{0pc} \\tablenum{4} \\footnotesize \\tablecaption{The Other 13 Red Horizontal Branch Stars \\label{tab4}} \\tablehead{ \\colhead{Star} & \\colhead{$\\alpha$ (J2000)} & \\colhead{$\\delta$ (J2000)} & \\colhead{[Fe/H]} & \\colhead{$M_{V}$} & \\colhead{$T_{eff}$} & \\colhead{log~$g$} & \\colhead{$R$/$R_{\\odot}$} & \\colhead{$V_{\\rm broad}$} & \\colhead{$V_{\\rm rot}$\\tablenotemark{a}}} \\startdata HD 20 & 00:05:15.3 & $-27$:16:18 & $-1.66$ & 0.7 & 5350 & 2.5 & 8.2 & 5.9 & 0.8 \\nl CD$-23$ 72 & 00:16:16.5 & $-22$:34:40 & $-1.12$ & 0.8 & 5270 & 2.5 & 8.0 & 8.9 & 3.1 \\nl HD 3179 & 00:34:50.6 & $-21$:52:56 & $-0.92$ & 0.9 & 5280 & 2.6 & 7.8 & 5.2 & 0.4 \\nl BD$+44$ 493 & 02:26:49.7 & +44:57:46 & $-2.71$ & 0.8 & 5510 & 2.6 & 7.4 & 3.9 & 0.0 \\nl HD 108317 (binary) & 12:24:04.4 & +05:34:46 & $-2.48$ & 0.5 & 5230 & 2.4 & 9.6 & 5.1 & 0.4 \\nl HD 110885 & 12:45:19.2 & +01:03:20 & $-1.59$ & 0.7 & 5330 & 2.5 & 8.2 & 8.2 & 2.5 \\nl HD 119516 & 13:43:26.7 & +15:34:29 & $-2.49$ & 0.6 & 5440 & 2.5 & 8.6 & 9.1 & 3.3 \\nl BD$+9$ 2860 & 14:13:19.7 & +08:36:40 & $-1.67$ & 0.7 & 5240 & 2.5 & 8.6  & 3.9 & 0.0 \\nl BD$+11$ 2998 & 16:30:16.7 & +10:59:51 & $-1.46$ & 0.8 & 5360 & 2.5 & 8.0 & 6.8 & 1.4 \\nl BD$+9$ 3223 & 16:33:35.6 & +09:06:17 & $-2.41$ & 0.6 & 5310 & 2.4 & 8.9 & 4.8 & 0.2 \\nl BD+17 3248   & 17:28:14.4 & $+17$:30:35  & $-2.07$ & 0.65    & 5240 & 2.44 &  9.0 & 4.3 & 0.0 \\nl BD+25 3410   & 18:02:03.2 & $+25$:00:41  & $-1.37$ & 0.79    & 5740 & 2.69 &  6.7 & 9.7 & 3.9 \\nl BD$-3$ 5215 & 21:28:01.3 & $-03$:07:40 & $-1.64$ & 0.7 & 5420 & 2.6 & 7.9 & 7.3 & 1.8 \\nl \\tablevspace{4pt} \\tablenotetext{a}{These values are derived using Equation~\\ref{eq:vbroadcorr}.} \\enddata \\end{deluxetable}  \\clearpage  \\begin{deluxetable}{lrrrrrrrrrrr} \\tablewidth{0pc} \\tablenum{5} \\footnotesize \\tablecaption{CfA Data Binned by $M_{\\rm V}$ \\label{tab5}} \\tablehead{ \\colhead{$<M_{\\rm V}>$} & \\colhead{N} & \\colhead{$\\sigma$} & \\colhead{$\\sigma_{\\mu}$} & \\colhead{$<R>$} & \\colhead{$\\sigma$} & \\colhead{$\\sigma_{\\mu}$} & \\colhead{$<V_{\\rm broad}>$} & \\colhead{$\\sigma$} & \\colhead{$\\sigma_{\\mu}$} & \\multicolumn{2}{r}{$<V_{\\rm rot}$ sin $i>$} \\\\ \\colhead{} & \\colhead{} & \\colhead{} & \\colhead{} & \\colhead{($R_{\\odot}$)} & \\colhead{} & \\colhead{} & \\colhead{} & \\colhead{} & \\colhead{} & \\colhead{Eq. 1} & \\colhead{Eq. 2} } \\startdata $-2.01$ & 20 & 0.52 & 0.12 & 64.3 & 21.7 & 4.8 & 7.7 & 2.3 & 0.5 & 4.4 & 2.1 \\nl $-1.28$ & 20 & 0.37 & 0.08 & 34.9 & 5.8 & 1.3 & 5.2 & 1.7 & 0.4 & 0.0 & 0.4 \\nl $-0.67$ & 20 & 0.38 & 0.09 & 22.2 & 3.4 & 0.8 & 5.5 & 1.7 & 0.4 & 0.0 & 0.6 \\nl $+0.07$ & 20 & 0.32 & 0.07 & 14.7 & 2.3 & 0.5 & 4.2 & 1.4 & 0.3 & 0.0 & 0.0 \\nl $+0.83$ & 20 & 0.50 & 0.11 & 9.4 & 1.0 & 0.2 & 3.8 & 2.2 & 0.5 & 0.0 & 0.0 \\nl $+1.89$ & 13 & 0.68 & 0.19 & 5.4 & 1.0 & 0.3 & 2.9 & 2.5 & 0.7 & 0.0 & 0.0 \\nl \\tablevspace{4pt} \\enddata \\end{deluxetable}  \\clearpage  \\begin{deluxetable}{lrrrrrrrr} \\tablewidth{0pc} \\tablenum{6} \\footnotesize \\tablecaption{Comparison of Derived Rotational Velocities with Mg~II and Ca~II Emission line Asymmetries \\label{tab6}} \\tablehead{ \\colhead{Star} & \\colhead{$M_{\\rm V}$} & \\colhead{$V_{\\rm rad}$} & \\colhead{$V_{\\rm rot}$ sin $i$} & \\colhead{$V_{\\rm broad}$(CfA)} & \\colhead{$V_{\\rm rot}$(Eq.\\ref{eq:vbroad})} & \\colhead{$V_{\\rm rot}$(Eq.\\ref{eq:vbroadcorr})} & \\colhead{Mg~II} & \\colhead{Ca~II} } \\startdata HD 2796  & $-0.81$ & $ -61.0$ & \\nodata & 7.0 & 4.1 & 1.6 & \\nodata & V$>$R \\nl HD 6755\\tablenotemark{a}  &  1.50 & $-319.2$ & \\nodata &  3.3 & 0.0 & 0.0 & V$>$R & V$>$R  \\nl HD 6833  & $-0.40$ & $-245.0$ & \\nodata & 7.4 & 5.0 & 1.9 & V$<$R & V=R \\nl HD 8724  & $-1.11$ & $-113.2$ & \\nodata & 6.0 & 1.4 & 0.9 & \\nodata &  V$>$R \\nl HD 21022 & $-1.17$ & 122.3 & \\nodata & 5.0 & 0.0 & 0.3 & \\nodata & V$>$R \\nl HD 23798 & $-1.80$ & 89.4 & 0.0 & 5.0 & 0.0 & 0.3 & \\nodata & V=R \\nl HD 26297\\tablenotemark{b} & $-1.48$ & 14.8 & \\nodata & 0.5 & 0.0 & 0.0 & \\nodata & V$>$R  \\nl HD 29574\\tablenotemark{c} & $-2.11$ & 19.8 & 4.0 & 10.1 & 7.8 & 4.3 & \\nodata & V$<$R \\nl HD 36702 & $-1.90$ & 122.7 & \\nodata & 6.5 & 1.5 & 1.2 & \\nodata & V$>$R \\nl HD 63791 &  0.30 & $-108.4$ & \\nodata & 3.7 & 0.0 & 0.0 & \\nodata & V$>$R \\nl HD 83212 & $-0.90$ & 109.1 & \\nodata & 7.3 & 4.5 & 1.8 & \\nodata & V$<$R \\nl HD 88609 & $-1.42$ & $-38.5$ & \\nodata & 4.0 & 0.0 & 0.0 & \\nodata & V$>$R \\nl HD 110281 & $-2.60$ & 141.9 & 5.5 & 11.5 & 9.3 & 5.8 & \\nodata & V=R \\nl HD 118055  & $-1.50$ & $-100.7$ & \\nodata & 7.4 & 4.2 & 1.9 & \\nodata & V$<$R \\nl HD 122956\\tablenotemark{d} & $-0.70$ & 166.0 & \\nodata & 6.3 & 2.9 & 1.1 & V$<$R & V$>$R \\nl HD 126587  & $-0.60$ & 149.0 & \\nodata & 4.3 & 0.0 & 0.0 & V$>$R & \\nodata \\nl BD+1 2916  & $-1.76$ & $-12.1$ & \\nodata & 8.2 & 5.3 & 2.5 & \\nodata & V$>$R \\nl HD 166161  & 0.79 & 68.3 & \\nodata & 4.0 & 0.0 & 0.0 & \\nodata & V$>$R \\nl HD 175305  & 1.80 & $-181.0$ & \\nodata & 7.1 & 5.8 & 1.7 & V$>$R & V=R \\nl HD 184711\\tablenotemark{e}  & $-2.35$ & 102.2 & \\nodata & 7.7 & 4.0 & 2.1 & \\nodata & V$\\leq$R \\nl HDE 232078 & $-2.15$ & $-387.2$ & \\nodata & 10.9 & 8.8 & 5.2 & \\nodata & V=R \\nl HD 187111\\tablenotemark{f} & $-1.54$ & $-186.5$ & 2.4 & 5.2 & 0.0 & 0.4 & \\nodata & V$\\leq$R \\nl HD 204543\\tablenotemark{g}  & $-1.09$ & $-98.4 $ & \\nodata & 5.0 & 0.0 & 0.3 & \\nodata &  V$\\geq$R \\nl HD 216143  & $-1.4$ & $-116.0$ & \\nodata & 5.5 & 0.0 & 0.6 & V$<$R &  V$>$R \\nl HD 218732\\tablenotemark{h} & $-2.80$ & $-312.2$ & 3.1 & 11.1 & 8.7 & 5.4 & \\nodata & V$<$R \\nl HD 221170  & $-1.70$ & $-119.0$ & 1.0 & 7.4 & 4.0 & 1.9 & V$<$R & V$>$R \\nl HD 222434\\tablenotemark{i}  & $-1.20$ & 8.8 & \\nodata & 5.0 & 0.0 & 0.3 & \\nodata & V=R \\nl \\tablevspace{4pt} \\enddata \\tablenotetext{a}{We employ the $\\gamma$ velocity for this binary system.} \\tablenotetext{b}{Dupree \\& Smith 1995 cite $V_{\\rm rad}$ = +135 \\kms, which we believe to be incorrect.} \\tablenotetext{c}{Smith et al.\\ 1992 and Dupree \\& Smith 1995 both found V$>$R.} \\tablenotetext{d}{Smith et al.\\ 1992 and Dupree \\& Smith 1995 both found V$>$R.} \\tablenotetext{e}{Smith et al.\\ 1992 found V$<$R while Dupree \\& Smith 1995 found it to be uncertain. We adopt the former value.} \\tablenotetext{f}{Smith et al.\\ 1992 found V$<$R while Dupree \\& Smith  1995 found V=R.} \\tablenotetext{g}{Smith et al.\\ 1992 found V$>$R and Dupree \\& Smith 1995 found V$\\approx$R. Both suggested $M_{\\rm V}$ = $-0.3$, but we recommend $-1.09$.} \\tablenotetext{h}{We employ the $\\gamma$ velocity for this binary system.} \\tablenotetext{i}{We employ the $\\gamma$ velocity for this binary system.} \\end{deluxetable}  \\clearpage"
    },
    "0711/0711.0105_arXiv.txt": {
        "abstract": "  In the present study, considering the physical conditions that are relevant interactions between supernova remnants (SNRs) and dense molecular clouds for triggering star formation we have built a diagram of  SNR radius versus cloud density in which the constraints above delineate a shaded zone where star formation is allowed. We have also performed fully 3-D radiatively cooling numerical simulations of the impact between SNRs and  clouds under different initial conditions in order to follow the initial steps of these interactions. We determine the conditions that may lead either to cloud collapse and star formation or to complete cloud destruction and find that the numerical results are consistent with those of the SNR-cloud density diagram. Finally, we have applied the results above to the $\\beta-$ Pictoris stellar association which is composed of low mass Post-T Tauri stars with an age of 11 Myr. It has been recently suggested  that its formation could have been triggered by the shock wave produced by a SN explosion localized at a distance of about 62 pc that may have occurred either in the Lower Centaurus Crux (LCC) or in the Upper Centaurus Lupus (UCL) which are both nearby older subgroups of that association (Ortega and co-workers).Using the results of the analysis above we have shown that the suggested origin for the young association at the proposed distance is plausible only for a very restricted range of initial conditions for the parent molecular cloud, i.e., a cloud with a radius of the order of 10 pc, a density of the order of 10$-$20 cm$^{-3}$, and a temperature of the order of 10$-$100 K }. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.2100_arXiv.txt": {
        "abstract": "We present spatially resolved kinematics and stellar population parameters for three low-luminosity galaxies in the NGC~128 group obtained by means of 3D spectroscopy. We briefly discuss their evolutionary scenarii. ",
        "introduction": "The NGC~128 group (d=57~Mpc) includes (spectroscopically confirmed): giant gas-rich S0 NGC~128 with a peanut-shaped bulge, giant Sa NGC~125, 3 low-luminosity S0 closer than 100~kpc in projection to NGC~128: NGC~126 (M$_B$=-18.5, 61 kpc), NGC~127 (M$_B$=-18.6, 14 kpc), and NGC~130 (M$_B$=-19.0, 16 kpc), and several late-type spirals beyond 300~kpc. NGC~128 is known to have a gaseous disc counter-rotating to the stars (Emsellem \\& Arsenault, 1997). ",
        "conclusions": "While NGC~126 and 130 look similar to cluster dS0/S0's, NGC~127 is a star-forming galaxy with unusual kinematics of gas and a bridge connecting it to NGC~128. We propose the following scenario: NGC~127 has recently passed its pericentre, and now we observe an infall of gas from NGC~128 onto NGC~127. After gas removal and several Gyr of passive evolution it will fade down and would be indistinctive from ''normal'' dE galaxies by morphology and luminosity."
    },
    "0711/0711.1515_arXiv.txt": {
        "abstract": "We present new \\textit{Spitzer Space Telescope} observations of the region NGC 2467, and use these observations to determine how the environment of an \\hii~region affects the process of star formation. Our observations comprise IRAC (3.6, 4.5, 5.8, and 8.0 $\\mu$m) and MIPS (24 $\\mu$m) maps of the region, covering approximately 400 square arcminutes.  The images show a region of ionized gas pushing out into the surrounding molecular cloud, powered by an O6V star and two clusters of massive stars in the region.  We have identified as candidate Young Stellar Objects (YSOs) 45 sources in NGC 2467 with infrared excesses in at least two mid-infrared colors.  We have constructed color-color diagrams of these sources and have quantified their spatial distribution within the region.  We find that the YSOs are not randomly distributed in NGC 2467; rather, over 75\\% of the sources are distributed at the edge of the \\hii~region, along ionization fronts driven by the nearby massive stars.  The high fraction of YSOs in NGC 2467 that are found in proximity to gas that has been compressed by ionization fronts supports the hypothesis that a significant fraction of the star formation in NGC 2467 is triggered by the massive stars and the expansion of the \\hii~region.  At the current rate of star formation, we estimate at least 25-50\\% of the total population of YSOs formed by this process. ",
        "introduction": "The process of star formation is still not well understood, especially in high-mass-star-forming regions.  One of the outstanding mysteries is the degree to which the stars already formed in such regions, especially the massive O or B stars, can affect the future formation of stars in the region.  In some high-mass-star-forming regions there is strong evidence for sequential star formation.  The Scorpius-Centaurus (Sco-Cen) OB association is one extensively studied example.  Stellar populations in the different subgroups of Sco-Cen have measurably different ages that suggest a sequence of star formation, with the oldest subgroup (Upper Centaurus-Lupus) located between two younger subgroups (Lower Centaurus-Crux and Upper Scorpius).  The implication is that the oldest subgroup formed in the middle of a molecular cloud after some delay triggered the formation of the two younger subgroups on either side of it (de Geus et al.\\ 1989; Preibisch \\& Zinnecker 1999, 2001).  Other recent evidence for sequential star formation on smaller scales has been observed in bright-rimmed clouds (BRCs) such as BRC 37 by (Ikeda et al.\\ 2008) and in the Orion OB association and Lac OB1 association (Lee et al.\\ 2005; Lee $\\&$ Chen 2007).  In these last two associations new protostars are observed to line up between the massive stars and the parent molecular clouds in the Orion and Lac OB associations, with no young stars found embedded in the molecular clouds far behind the ionization fronts, suggesting the stars form only {\\it after} passage of an ionization front launched by the massive stars (Lee et al.\\ 2005).  Thus evidence exists that massive stars can affect or perhaps even cause the subsequent formation of stars in a region.  An alternative interpretation is that even in high-mass-star-forming regions all star formation is coeval (no delay between the formation of high-mass stars and any other stars in the region).  In an extreme form of this interpretation, massive stars can not modify subsequent star formation because there is none; it has already happened. Hillenbrand et al.\\ (2008; White \\& Hillenbrand 2004) claim that star formation in \\hii~regions environments is coeval and suggest that estimates of ages for low-mass stars have been underestimated and estimates of high-mass stars overestimated.  Hillenbrand et al.\\ (2008) use isochronal fitting to estimate the ages of sources in many star-forming regions and conclude there is no strong evidence for moderate age spreads in either young star-forming regions or young open clusters.  If we are to understand star formation generally, this difference of interpretation must be resolved.  In their complete census of embedded clusters (out to 2 kpc), Lada $\\&$ Lada (2003) have shown that the vast majority (70-90\\%) of all (low-mass) stars formed in rich clusters; roughly 50\\% of all stars formed in clusters rich enough to contain at least one O or B star capable of shaping its environment. If a large fraction (say, one third) of stars in such clusters are {\\it triggered} to form or otherwise affected by the presence of nearby massive stars, then a significant fraction (say, one sixth) of all stars will have been affected by this process.  Given that our own Solar System probably formed in a massive cluster (Hester et al.\\ 2004; Hester \\& Desch 2005), understanding star formation in high-mass-star-forming regions is needed to put our own formation in context.  Accurate predictions of the initial mass function (IMF), as well as a quantification of the timing and conditions of protoplanetary disk growth, rely on identifying processes by which massive stars can affect subsequent star formation and quantifying the extent to which this has occurred in high-mass-star-forming regions. Is star formation in \\hii~region environments coeval?  If it is not, is later star formation physically connected to the presence of high-mass stars?  If so, by what mechanism?  A number of mechanisms by which massive stars can affect the subsequent star formation in a region have been proposed.  Here we review three.  One proposed mechanism for triggered star formation is the ``radiatively-driven implosion\" (RDI) model, which invokes photoionizing ultraviolet (UV) radiation from massive stars feeding back into the environment of the nearby molecular cloud and interstellar medium (ISM).  In this model, pre-existing molecular clumps in the cloud become surrounded on all sides by high-pressure ionized gas heated by the photoionizing UV.  The pressure that is exerted on the surface of a molecular clump leads to the implosion of the clump and to the formation of a cometary globule (Bertoldi 1989; Bertoldi $\\&$ McKee 1990; Larosa 1983).  The collapse phase for this scenario is relatively short (i.e., less than the typical free-fall timescale $t_{\\rm ff} \\, \\sim \\, 3 \\times 10^{5} \\, {\\rm yr}$), and during this phase a shock front (traveling behind the ionization front) will progress into the clump, leading to the formation of a dense core within it (Deharveng et al.\\ 2005).  One observational signature of this mode of triggered star formation would be cometary globules having dense heads and a tail extending away from the ionizing source.  A key prediction of this model is that there should be no newly forming stars located ``ahead'' of the ionization front, in gas that has not yet been radiatively heated.  A second proposed mechanism is the ``collect and collapse'' (C\\&C) model first proposed by Elmegreen \\& Lada (1977).  This model invokes the standard picture of a slow-moving D-type ionization front with associated shock front that proceeds the ionization front (see Osterbrock 1989).  Dense gas tends to pile up in the space between the shock and ionization fronts, with the surface density of this layer increasing over time.  The ``collect and collapse\" model hypothesizes that at some critical threshold (as the layer gets denser and as it cools) the dense layer becomes gravitationally unstable and fragments into dense clumps that then go on to form stars.  The attraction of this model is that both large numbers of stars can form in this layer, as well as massive stars (Whitworth et al.\\ 1994).  Because fragmentation in the dense layer, driven by gravitational instability, grows fastest for long wavelengths, an observational test of the C\\&C scenario is that large fragments typically will be evenly spaced (and detectable in mid-infrared [mid-IR] and millimeter emission) within a shell that surrounds the \\hii~region (Deharveng et al.\\ 2005).  A third proposed mechanism, proposed by Hester et al.\\ (1996) in the context of M16, also invokes a slow-moving D-type ionization front with associated shock, like the C\\&C scenario.  In this scenario, clumps are triggered to form and collapse by the high pressures in the shocked gas ahead of the ionization front.  Dense clumps were probably not pre-existing in the gas, but were either formed from marginally stable overdense regions in the molecular cloud, or created by a fragmentation instability as in the C\\&C model.  As in the RDI model, the high pressures of the compressed gas lead to rapid (i.e., $< t_{\\rm ff} \\approx 3 \\times 10^5 \\, {\\rm yr}$) collapse of the clumps (Garc\\'{i}a-Segura $\\&$ Franco 1996).  As the ionization front advances through the region, the gas surrounding clumps is exposed, leading to objects like cometary globules, which Hester et al.\\ (1996) termed ``evaporating gaseous globules\", or EGGs.  Later, as more gas is photoevaporated, these evaporated clumps are observed as ``proplyds\" or protoplanetary disk.  Because the clumps are triggered to collapse by the high pressures in the post-shock region, a key observational test of this scenario is that Young Stellar Objects (YSOs) will be spatially correlated with ionization fronts within \\hii~region; unlike in the C\\&C scenario, though, protostars will not necessarily be regularly spaced.  Another key test of this scenario is that the evolutionary stage (age) of clumps and subsequent YSOs and disks should increase with distance from the shock front.  While each of the scenarios makes distinct observational tests, it is difficult to constrain which scenario is operating.  For example, proplyds (as predicted by the third scenario) are observed in a variety of regions, including M16 (Hester et al.\\ 1996), W3/W4 (Oey et al.\\ 2005), RCW 49 (Whitney et al.\\ 2004a), and M20 (Cernicharo et al.\\ 1998; Hester et al.\\ 1999).  But the existence of proplyds is not unique to this scenario and would arise in the C\\&C model as well (and probably the RDI, too).  A sequence of star formation has been observed by Lee \\& Chen (2007) in the Ori OB1 and Lac OB1 associations, and a recent study of massive YSOs and ongoing star formation in the LMC by Book et al.\\ (2009) shows evidence for triggered star formation of massive YSOs; but the C\\&C and RDI scenarios are not distinguished.  Further progress must rely on finer observational tools.  As the above discussion indicates, an important diagnostic tool for testing predictions of triggered star formation models is a measure of the ages of detected protostellar objects.  YSOs have infrared excesses (above the stellar photospheric emission) due to a large infalling envelope in the early stages of formation, and from an accretion disk in the later stages.  Traditionally, YSOs have been separated into the ``Class I, II, III'' evolutionary stages, first defined by Adams, Lada, and Shu (1987), based on the slopes of their spectral energy distributions (SEDs) in the near-IR; however, selecting YSO sources from their \\textit{Spitzer} mid-IR colors may be a more robust method (Allen et al.\\ 2004; Whitney et al.\\ 2003a, 2003b \\& 2004b), mainly because not all YSOs have an excess in the near-IR. Recent studies, including Poulton et al.\\ (2008), Indebetouw et al.\\ (2007), and Simon et al.\\ (2007), have also successfully used an SED-fitting tool from Robitaille et al.\\ (2007) to identify and classify YSOs in multiple regions.  The online SED fitter uses a grid of YSO models based on radiative transfer codes from Whitney et al.\\ (2003a, 2003b,$\\&$ 2004b).  In this paper, we use the online SED fitting tool to confirm the candidate YSOs and to estimate their physical properties, specially their ages.  The \\hii~region NGC 2467, also known as Sharpless 311, is located at a distance of 4.1 kpc (Feinstein $\\&$ V\\'{a}zquez 1989).  This region is dominated by one O6 Vn star, HD 64315.  There are also two stellar clusters in the area, Haffner 19 (H19) and Haffner 18ab (H18ab), that contain one later type O and additional B stars, but most ($\\sim 70\\%$) of the ionizing radiation comes from the O6 star.  Using UBVRI broad-band data, Moreno-Corral et al.\\ (2002) determined that there are 34 B stars in the cluster H19, the most massive being a B1V star. They estimate the age of H19 to be $\\approx \\, 2 \\, {\\rm Myr}$. Fitzgerald $\\&$ Moffat (1974) estimate the age of H18ab to only be 1 Myr, and Munari et al.\\ (1998) obtained a best-fit age for H18ab of 2 Myr. Pismis and Moreno (1976) claim that the entire \\hii~region complex is very young, with an age around 2 - 3 Myr.  We will assume an age for NGC 2467 of 2 Myr.  Recently, De Marco et al.\\ (2006) used \\textit{Hubble Space Telescope} (\\textit{HST}) Advanced Camera for Surveys (ACS) data to identify a large number of brightened ridges and cloud fragments in NGC 2467.  The ionization front appears to be very near to many of these cloud fragments, suggesting that they have recently broken off from the molecular cloud as they are uncovered by the advancing ionization front.  In the process these fragments are being photoevaporated, and thus NGC 2467 is an excellent candidate for a study of star formation in an \\hii~region environment, and to test the predictions of the triggered star formation scenarios outlined above.  In this paper, we use \\textit{Spitzer} observations to locate protostars in NGC 2467.  We then use their SEDs to confirm the candidates as protostellar objects, and to determine their physical properties, especially their age.  These observations allow us to constrain the timescales for their formation and the processes at play during their formation.  We also quantify the overall distribution of these objects and their spatial correlation with tracers of star formation with and massive OB stars in the region.  We then test whether the formation of stars in this \\hii~region was coeval or at some level triggered, and attempt to distinguish between the various mechanisms proposed for triggered star formation. ",
        "conclusions": "The mid-IR data from \\textit{Spitzer} show that NGC 2467 is a region of active star formation.  We detected a large number of sources with mid-IR excesses, which is evidence of the youth of the region. 45 YSOs were detected based on their \\textit{Spitzer} colors, and 27 of them also showed an infrared excess in their 2MASS colors.  One of the main questions about star formation that we wanted to address is whether or not there is evidence for triggered star formation or if the population is coeval.  Triggered star formation by \\hii~region expansion appears to be occurring in NGC 2467.  The current population of protostars is spatially correlated with the ionization fronts.  The new protostars appear to be forming in a coherent structure which is associated with the ionization fronts.  We have also shown, based on SED fitting, that this population is very young, with the majority having ages less than a few$\\times$ 10$^{5}$ yr, much younger than the age of the region and of the ages of the OB stars.  We therefore conclude that the star formation in NGC 2467 has been ongoing, and it is not coeval, with strong evidence for triggering from the expansion of the \\hii~region created by the massive OB stars.  We favor an ionization front with a preceding shock as the most likely triggering mechanism, but whichever mechanism, we can estimate the amount of triggered star formation that is occurring.  The estimated current total star formation rate in NGC 2467 is 1.75$\\times$10$^{-3}$ M\\sol~yr$^{-1}$, with a triggered star formation rate equal to 0.42 - 0.90$\\times$10$^{-3}$ M\\sol~yr$^{-1}$. Based on the distribution and age estimates of the current population of candidate YSOs, we estimate that between 25 - 50$\\%$ of the YSOs forming in this region are due to triggering from the advance of ionization fronts created by nearby massive stars.  It is probable that some of the low-mass stars in NGC 2467 have formed spontaneously throughout the lifetime of the region, but we are also seeing a population of very young objects that are currently forming in compressed gas from the \\hii~region expansion. While the data do not suggest that triggering is the only mode of low-mass star formation in this region, it does suggest that triggering is occurring in NGC 2467 and has contributed a substantial fraction to the overall star formation rate.  This work has implications for other star forming regions and even for possible mechanisms of our own Sun's formation.  Although this is only one region, we can attempt to use the results from this region to predict what fraction of all low-mass stars may be triggered by \\hii~region expansion.  Lada and Lada (2003) have shown that as much as 70-90$\\%$ of all low-mass stars form in rich clusters also containing massive stars, therefore if 25-50$\\%$ of low-mass stars in rich clusters are triggered, then as many as 18-45$\\%$ of all low-mass stars may be formed by this mechanism. Possibly including our own Sun and Solar System, where meteoric evidence for live $^{60}$Fe, a short-lived radionucleide, in the early solar system suggests that a supernova event occurred shortly before the formation of the planets (Tachibana $\\&$ Huss 2003; Tachibana et al.\\ 2006).  The environment created by the massive stars in NGC 2467 is having a definite impact on the formation of further generations of stars. We find that newly forming YSOs are mostly found in areas where the shock front driven in advance of the ionization front is compressing the molecular gas.  The sources we are detecting with mid-IR excesses are very young sources, with the majority having ages less than a few $\\times$ 10$^{5}$ yr. Thus the sources we are finding are not coeval with the OB stars. The 45 YSOs detected and the distribution of these sources suggest that NGC 2467 is a prime example of triggered star formation due to \\hii~region expansion."
    },
    "0711/0711.4588.txt": {
        "abstract": "The halo of M31 shows a wealth of substructures, some of which are consistent with the assembly from satellite accretion. Here we report on kinematic and abundance results from Keck/DEIMOS spectroscopy in the near-infrared calcium triplet region of  over 3500 red giant star candidates along the minor axis and in off-axis  spheroid fields of M31. These data reach out to large radial distances of about  160 kpc. % The derived radial velocity distributions show an indication of a kinematically cold substructure around $\\sim 17$ kpc, which has been reported before. We devise a new improved method to measure spectroscopic metallicities from the calcium triplet in low signal-to-noise spectra using a weighted coaddition of the individual lines. The resulting distribution (accurate to $\\sim$0.3 dex down to signal to noise ratios of 5) leads us to note an even stronger gradient in the abundance distribution along M31's minor axis and in particular towards the outer halo fields than previously detected. The mean metallicity in the outer fields reaches below $-2$ dex, with individual values as low as $\\la -2.6$ dex. This is the first time such a metal poor halo has been detected in M31. In the fields towards the inner spheroid we find a sharp decline of $\\sim$0.5 dex in metallicity in a region at  $\\sim$20 kpc, which roughly coincides with the edge of an  extended disk, previously detected from star count maps. A large fraction of red giants that we detect in the most distant fields are likely members of M33's overlapping halo. % A comparison of our velocities with those predicted by new N-body simulations argues that the event responsible for the giant Stream is most likely not responsible for the full population of the inner halo. We show further that the abundance distribution of the Stream is different from that of the inner halo, from which it becomes  evident, in turn, that the merger event that formed the Stream and the outer halo cannot have contributed any significant material to the inner spheroid. All these severe structure changes in the halo suggest a  high degree of infall and stochastic abundance accretion governing the build-up of M31's inner and outer halo. ",
        "introduction": "With the discovery of the Giant Stream (Ibata et al. 2001), the mapping of complex structures in the ``halo'' of M31 (Ferguson et al. 2002; Gilbert et al. 2007; Ibata et al. 2007), and the isolation of an extended kinematic disk structure (Ibata et al. 2005) the idea that the radially more distant populations of M31 originate in accretion events has become established.   Even a subset of M31 satellites might relate to the breakup of a massive progenitor, based on the polar and planar alignment of a number of its early-type satellite galaxies (Koch \\& Grebel 2006).  Prior to this paradigm shift in the description of the M31 halo, relatively shallow Hubble Space Telescope (HST) imagery revealed what appeared to be a mostly metal rich M31 halo (e.g. Rich et al. 1996; Bellazini et al. 2003). Mould \\& Kristian (1986) were the first to find a metal rich (47 Tuc-like) halo population, using ground-based imaging. The widespread presence of this metal rich population as well as the descending red giant branch (RGB; in contrast to the metal poor globular clusters) was noted by Bellazzini et al. (2003).  Ground-based minor axis star counts appear to show a smooth $r^{1/4}$ spheroid (Pritchet \\& van den Bergh 1994), and some studies have argued that the metal rich halo population is an extension of the metal rich bulge (e.g. Mould \\& Kristian 1986; Guhathakurta et al. 2006). But alternatively, the case has been made that the metal rich stars more closely coincide with the perturbed regions (Ferguson et al. 2002).  Metal poor stars around M31 were already indicated in photometric studies of its halo at projected distances of 7 kpc (Mould \\& Kristian 1986) to 20 kpc (Durrell 2001). Subsequent spectroscopic surveys then fully revealed a complex picture of the halo composition: Using the well-established near-infrared calcium triplet (CaT) as a metallicity indicator, Reitzel \\& Guhathakurta (2002) find metal poor stars at 19 kpc on the minor axis.  Benefitting from much larger samples, Chapman et al. (2006) and Kalirai et al. (2006a) argue based on kinematics and metallicity that the expected metal poor halo population is in fact present. Giants with the radial velocity of M31 are claimed as members of the halo to distances in excess of 100 kpc (Ostheimer 2003; Gilbert et al. 2006) and an overall metallicity gradient of the M31 halo is proposed by Kalirai et al. (2006a).   Yet, over much of this region, Ibata et al. (2007) find clear evidence of density enhancements associated with accretion.   What fraction of the outer halo is then comprised of such recently accreted material? In particular the chemical composition of the outer stars remains unknown, and depending on the adopted [$\\alpha$/Fe] ratio,  Kalirai et al. (2006a) estimate a mean metallicity in the outermost field $\\sim$0.3 dex higher than the mean Milky Way metallicity (for Solar scaled abundances), or  a mean that is comparable to the Milky Way halo under the assumption of strong $\\alpha-$enhancement.  In fact, pencil beam ultradeep imaging using the Advanced Camera for Surveys (ACS) on board the HST, offers a complementary picture of the complexity present in the halo populations.   The placement of these deep imaging fields has benefitted from the starcount maps and Keck/DEIMOS kinematic studies. Brown et al. (2003) first demonstrated that a minor axis field projected at 11 kpc contains an indisputable range in age and abundance, extending to nearly-Solar, and a predominant age range from 6--10 Gyr.  Comparing 3 fields, in the inner spheroid (at 11\\,kpc),  the M31 disk and on the Ibata et al. (2001) debris stream, Brown et al. (2006) find the disk to be younger and more metal rich, and lacking old stars. They find the stream and spheroid fields to be indistinguishable based on their color-magnitude diagrams (CMDs), leading to the conclusion that only one progenitor is responsible for the debris field in the inner halo region.  However, if dynamical mixing were efficient in these regions, it could also erase the signatures from different sources. Brown et al. (2007) investigate a field at 21 kpc and find evidence that its population  is marginally older and more metal poor than the inner halo field.  The present-day pencil beam surveys  have found clear evidence of an age range in every field studied and appear to support other evidence of a gradient in age and abundance.  What has been lacking to date has been a survey of abundances and kinematics that ties these fields together and provides a context for the interpretation of these deep fields.  This is one aim of this paper, and a natural extension of our systematic survey of the structure and kinematics of M31 along its southwest minor axis.  While it is attractive to seek {\\it one} massive progenitor for the inner debris field there are several arguments against this position.  First, the existence of both the giant stream and the extended disk suggest at least two very different sources for the debris field at 11 kpc. Second, the dramatic variation of the extended spatial structure of the debris field as a function of metallicity and age (Ferguson et al. 2002) is best understood by invoking  multiple events involving different accretors. Moreover, the distant rotating disk-like population (Ibata et al. 2005) is superposed on other, likely unrelated, structures that are suggestive of shells associated with the giant Stream merger (see also Ibata et al. 2007). It is then an intriguing question, whether a corresponding measurable abundance change occurs at the point where this field ends.  In the context of CDM models (e.g. Bullock \\& Johnston 2005) halos are thought, in general, to accrete from the debris of lower mass satellites.   Yet Mouhcine et al. (2005) find a correlation between parent galaxy luminosity and halo metallicity. Nominally,   M31 has a high metallicity halo and falls in this relationship alongside galaxies with more prominent bulges. There is a paradox: how can the halo of M31 be dominated by stochastic accretion events, yet still have an $r^{1/4}$ profile and still appear to follow trends set by luminosities of the host galaxies?  We report here the culmination of an observational campaign begun in 2002; Table~1 lists the observing run details by Principal Investigator. % This Paper is organized as follows: In \\textsection 2 we present our observations and the standard reduction steps taken, while \\textsection 3 describes our radial velocity measurements.  The dwarf/giant separation is discussed in \\textsection 4 and in \\textsection 5 we devise a new technique to measure spectroscopic metallicities from the CaT. The following sections are then dedicated to the analysis of kinematic (\\textsection 6) and abundance (\\textsection 7) substructures and gradients in M31's halo. Finally, \\textsection 8  summarizes our findings. ",
        "conclusions": "% From a spectroscopic analysis of 1316 confirmed red giant stars along the minor axis on M31 and in spheroid fields out to 160 kpc we  find and confirm the following structural specific features in M31's inner and outer halo: % \\begin{enumerate} % \\item There is evidence of an abundance substructure in the sense that  the mean stellar metallicity strongly declines at $\\sim$20 kpc, where the abundance range within this radius is typically $-0.5$ to $-1$ dex and falls towards $-1.4$ dex at 20--40 kpc. The latter values are consistent with those detected within a smooth underlying M31 halo (Ostheimer 2003; Guhathakurta et al. 2006; Kalirai et al. 2006a; Chapman et al. 2006). Interestingly, the location of the break in the metallicity profile coincides with the edge of the metal rich, extended rotating disk reported by  Ibata et al. (2005), but it cannot be entirely ruled out that there is also a contribution from the kinematic substructure (Gilbert et al. 2007), which would bias  the inner regions towards marginally higher metallicities. % A metallicity gradient has been detected by Kalirai et al. (2006a), but here we show that the decline in our spectroscopic measurements appears to proceed even stronger towards the outermost fields at 160 kpc, where we find mean values around and even below $-2$ dex. % \\item In particular, there is a considerable fraction of metal poor stars below [Fe/H]$_{\\rm CaT}\\la -2$ found at almost all radii. Their presence is not utterly surprising, and red giants that metal poor have been claimed to exist in M31's halo before (e.g., Reitzel \\& Guhathakurta 2002). Furthermore, their established prominence in the Milky Way halo (Carney et al. 1996; Chiba \\& Beers 2000; Carollo et al. 2007) raises the question, why there should not be a comparable distribution present in the M31 halo, if both systems had experienced a similar formation and accretion history. % \\item A considerable fraction of stars in the outermost fields beyond $\\sim$100 kpc (which corresponds to about 100 kpc projected distance to M33) exhibit velocities and metallicities consistent with those of M33 halo stars. This confirms earlier findings from star count maps (Ibata et al. 2007), according to which the stellar halos of these major Local Group spirals overlap to a large extent. % \\item We confirm the earlier detection of a kinematically cold  substructure (Gilbert et al. 2007), located at a radial distance of  15--20 kpc. By comparison of new N-body simulations (Mori \\& Rich 2008), we show that such a substructure is consistent with having originated in the merger event that produced M31's Giant Stellar Stream (Ibata et al. 2001, 2004). Nevertheless, the full radial velocity distribution along the minor axis is difficult to reconcile with a single collisional event of this kind, and it is more likely that a wealth of accretions occurred and formed the halo, which is concordant with the progressive detections of stream-like substructures in star count maps (Ibata et al. 2007). % \\item There is a considerable contribution of stars from a  genuine, ancient M31 halo, potentially ejected during such merger events,  to the velocity distribution at any radius. Moreover, neither stars from the pristine bulge or disk components are likely to be found in the outer or even inner spheroids. % \\end{enumerate} % This leaves us with a picture in which the progenitor that produced the giant Stream and thus presumably donated a major part of M31's halo  cannot be single-handedly responsible for all the substructures seen in our and previous studies -- outside the present-day Stream's sphere of influence,  one sees a predominant occurrence of minor substructures or streams from many a past mergers. A large number of  accretion events is also thought to have contributed to  the formation of M31's disk (e.g., Pe\\~narrubia et al. 2006), as this would produce the proper, observed mix of metallicities. % Ibata et al. (2007) note that the radial metallicity gradient is a mere reflection of Stream debris and numerous  other substructures in the {\\em inner} regions, which is in concordance with model predictions.  In this context the metal rich, genuine Stream material is more centrally concentrated, leaving the impression of a more extended underlying metal poor halo (cf. the stream distribution in Fig.~13). It is unclear at present, whether the gradient is intrinsic to the inner substructures becoming more metal poor with radius, or whether the underlying halo, with its larger dominance in the outer regions, is becoming more metal poor. % Note that Brown et al. (2007) find that an HST field at 25 kpc ($\\sim 1.5\\degr$) is both more metal poor and older than the innermost HST field projected at 11 kpc. Hence, also the HST pointings appear not to be dominated by the canonical Stream's debris. % Moreover,  we find that the innermost deep HST field (Brown et al. 2003) is situated in a region that has some potential contribution from tidal debris, but appears to be dominated by a more metal rich \"inner halo\" population with a {\\it smooth}  velocity distribution, free of the  structures at the velocity extremes that are predicted to be present from interaction models (see left middle panel of Fig.~13). The Brown et al. (2003) 11 kpc field appears to be free of major contamination from an infall event.  The 21 kpc field lies  just outside of the metal rich inner region; Brown et al. (2007) suggest a lower metallicity for this field, which we confirm from our CaT analysis.  The field at 35 kpc apears to be yet more metal poor and may be genuinely representative of the outer halo. % We note that all the HST fields at 11, 21 and 35 kpc have very different circular orbital periods (derived from a simple mass model and the assumption of circular orbits) of $\\sim$0.2, 0.5 and 1\\,Gyr, respectively. The characteristics of the HST deep fields will be discussed in a forthcoming paper (Rich et al.  in prep).  Using the same arguments, we find large orbital periods of $\\sim$6\\,Gyr at 160 kpc, which presupposes them to be bound to M31 (see also Majewski et al. 2007). It is hard to reconcile these time scales with a scenario, in which these distant M31 members are  pressure supported. It is rather an attractive notion that these distant stars might be ejecta of collisions, not yet having completed a single orbit around M31.  The fact that the new metallicity scale we framed in this work is in good agreement with the moderately metal poor  character of the dSph satellites of M31 (Reitzel \\& Guhathakurta 2002; McConnachie et al. 2005; Martin et al. 2006b; Majewski et al. 2007) further support the idea that the outer halo might plausibly emanate from a population similar to the dSphs. % Moreover, the presence of the more metal rich stars beyond 100 kpc suggest an origin of these stars in collisions rather than in a primordial halo. Yet it is challenging to reconcile this scenario with the number count maps of Ibata et al. (2007), which are smooth to within the sensitivity of the MegaCam survey in these regions and with the presence of very metal poor stars: there are no very metal poor stars detected in the Galactic dSphs (Koch et al. 2006; Helmi et al. 2006), nor in the faintest M31 satellites (Martin et al. 2006b, 2007), so that it remains unclear whether systems like the  present-day observed dwarf satellites are responsible for the bulk of the halo.  While our observed radial metallicity gradient stretches 1.5--2 dex over the full extent of our data of 160 kpc, simulations of the hierarchical assembly of Galactic halos do not reproduce any considerable gradients. In this context, the study of Font et al. (2006) predicts gradients in [Fe/H] {\\em averaged} over each full simulated halo of at most 0.5 dex over a few tens of kpc.  Our findings lend striking support to paralleling the M31 and Galactic halos: In a recent work  Carollo et al. (2007) detected a dichotomy in the Milky Way's stellar halo.  In this sense, there exists a clear kinematic and chemical separation into an inner halo (with stars on more eccentric orbits, metallicities around $-1.6$ dex, a flattened density distribution and no considerable rotation) and an outer, independent halo component, which is more spherical in shape, shows evidence of rotation {\\em and is more metal poor} on average. In particular, the outer Galactic halo exhibits a peak metallicity of $-2.2$ dex. The presence of two distinct halos is well explicable in the context of  cosmological structure formation models and incorporates the continuous, chaotic accretion of distinct subhalos that follow an earlier stage of the dissipative merging of massive, yet sub-galactic fragments (see Carollo et al. 2007 for a  discussion of the detailed formation scenario). If structure formation proceeded in analogy to form M31 -- and in the light of $\\Lambda$CDM this is likely -- then the clear distinctness of at least two halo components visible in the MDFs in our M31 study is a natural outcome of the accompanying stochastic abundance accretion.   As a concluding, independent remark, we note that a further benefit of the present study is the derivation of a new improved method to measure metallicities from the calcium triplet, which allowed us, and should encourage future works,  to extract reasonably accurate information from low $S/N$ spectra."
    },
    "0711/0711.1386_arXiv.txt": {
        "abstract": "Accretion disk winds browing off perpendicular to a luminous disk are examined in the framework of fully special relativistic radiation hydrodynamics. The wind is assumed to be steady, vertical, and isothermal. Using a velocity-dependent variable Eddington factor, we can solve the rigorous equations of relativistic radiative hydrodynamics, and can obtain radiatively driven winds accelerated up to the {\\it relativistic} speed. For less luminous cases, disk winds are transonic types passing through saddle type critical points, and the final speed of winds increases as the disk flux and/or the isothermal sound speed increase. For luminous cases, on the other hand, disk winds are always supersonic, since critical points disappear due to the characteristic nature of the disk gravitational fields. The boundary between the transonic and supersonic types is located at around $\\hat{F}_{\\rm c} \\sim 0.1 (\\varepsilon+p)/(\\rho c^2)/\\gamma_{\\rm c}$, where $\\hat{F}_{\\rm c}$ is the radiative flux at the critical point normalized by the local Eddington luminosity, $(\\varepsilon+p)/(\\rho c^2)$ is the enthalpy of the gas divided by the rest mass energy, and $\\gamma_{\\rm c}$ is the Lorentz factor of the wind velocity at the critical point. In the transonic winds, the final speed becomes 0.4--0.8$c$ for typical parameters, while it can reach $\\sim c$ in the supersonic winds. ",
        "introduction": "Mass outflows -- accretion disk winds -- from a luminous disk are observed in various active objects, such as cataclysmic variables (CVs), supersoft X-ray sources (SSXSs), microquasars ($\\mu$QSOs), broad absorption line quasars (BAL QSOs), and so on. Radiatively driven wind emanating from an accretion disk is a clue to the formation mechanism of astrophysical jets and winds in these objects (see Kato et al. 1998, 2007 for a review of accretion disks).  So far, radiatively driven outflows from a luminous disk have been extensively studied by many researchers (Bisnovatyi-Kogan, Blinnikov 1977; Katz 1980; Icke 1980; Melia, K\\\"onigl 1989; Misra, Melia 1993; Tajima, Fukue 1996, 1998; Watarai, Fukue 1999; Hirai, Fukue 2001; Fukue et al. 2001; Orihara, Fukue 2003), and by numerical simulations (Eggum et al. 1985, 1988; Kley 1989; Okuda et al. 1997; Kley, Lin 1999; Okuda, Fujita 2000; Okuda 2002; Okuda et al. 2005; Ohsuga et al. 2005; Ohsuga 2006). In almost all of these studies, however, the luminous disk was treated as an external radiation source (optically thin approximation), and the radiation transfer in the flow was not solved. Although radiation hydrodynamical equations were solved in the numerical simulations, the methods were rather limited; e.g., the equilibrium between gas and radiation was assumed, the flux-limited diffusion approximation was adopted, the flow velocity was subrelativistic on the order of $0.1 c$, or the optically thick to thin transition was not properly treated. Up to now, in relation to accretion disk winds no one solved the fully relativistic radiation hydrodynamical equations.     Recently, radiation hydrodynamical mass outflows have been examined for the first time in the fully relativistic cases for the plane-parallel case (Fukue 2005, 2006; Fukue, Akizuki 2006; see also Akizuki and Fukue 2007 for the spherical case). In these studies, however, the gas pressure was ignored for simplicity, although the relativistic radiation hydrodynamical equations were solved. In this paper, we thus take into account the gas pressure, and obtain the relativistic radiation hydrodynamical {\\it winds} browing off from a luminous flat disk within the framework of a fully special relativistic regime.   In the next section we describe the basic equations in the vertical direction, and examine critical points. In section 3 we obtain transonic winds as well as supersonic ones. The final section is devoted to concluding remarks. ",
        "conclusions": "In this paper we have examined the relativistic radiation hydrodynamical winds from a luminous accretion disk in the relativistic regime of $(v/c)^2$. The wind is assumed to be steady, vertical, and isothermal. Using a velocity-dependent variable Eddington factor, the basic equations can be numerically solved without meeting the pathological singular point at $v=c/\\sqrt{3}$ to reach the relativistic regime.  For less luminous cases, vertical disk winds are transonic types passing through saddle-type critical points, and the wind final speeds are 0.4--0.8$~c$ for typical parameters. For luminous cases, on the other hand, disk winds become supersonic types without passing through any critical points, and the wind final speeds becomes on the order of $c$. The boundary between the transonic and supersonic types is located at around $\\hat{F}_{\\rm c} \\sim 0.1 (\\varepsilon+p)/(\\rho c^2)/\\gamma_{\\rm c}$.  In usual standard accretion disks, the local luminosity is sub-Eddington. Hence, for proton-electron normal plasmas the above condition is safficiently fulfilled, and transonic winds driven by radiation and gas pressures would blow off. For electron-positron pair plasmas, on the other hand, the local luminosity becomes super-Eddington, and pair winds would supersonically blow off.   The maximum attainable velocity of the wind can be roughly estimated by the energy conservation. In the Schwarzschild space-time, without any energy source except for gravity, the total energy of the radiation hydrodynamical flow along the streamline, \\begin{equation} Jc^2 \\frac{\\varepsilon+p}{\\rho c^2} \\gamma \\sqrt{g_{00}} + g_{00} F = \\dot{E} ~({\\rm const.}), \\end{equation} where $g_{00} = \\sqrt{1-r_{\\rm g}/R}$, is conserved. If all of the thermal energy of the gas and the radiative energy in the initial state is converted to the bulk energy of the gas, the final velocity would become its upper limit: \\begin{equation} Jc^2 \\gamma_\\infty= \\left. Jc^2 \\frac{\\varepsilon+p}{\\rho c^2} \\gamma \\sqrt{g_{00}} + g_{00} F \\right|_{z=0}. \\end{equation} If, further, the gas is at rest in the initial state and in the virial state, $[(\\varepsilon+p)/(\\rho c^2)]\\sqrt{g_{00}} |_{z=0} = 1$, then we finally have the maximum attainable velocity as \\begin{equation} \\gamma_\\infty = 1+\\frac{g_{00}F |_0}{Jc^2} = 1+\\left( 1-\\frac{1}{\\hat{r}} \\right) \\frac{\\hat{F}_0}{\\hat{J}}. \\end{equation} For example, in the present typical case of $\\hat{r}=3$, $\\hat{F} \\sim 0.1$, $\\hat{J}=0.1$, we have $\\gamma_\\infty=1.666$, or $\\beta_\\infty=0.8$.  Hence, from the view point of energetics, the final velocity becomes higher and higher, as the ratio $\\hat{F}_0/\\hat{J}$ becomes large; the radiative flux is so high or the mass-loss rate is so low. In other words, if the small amount of the gas gains the large amount of the radiation energy, the final bulk velocity becomes high. In general, however, it may be difficult for such a re-distribution of energy to take place, and the ratio $\\hat{F}_0/\\hat{J}$ would be on the order of unity.  Thus, we conclude that the maximum attainable velocity of accretion disk winds emanating from the disk inner region would be $0.8~c$ or so, as long as it consists of normal plasmas. This conclusion is consistent with the current theoretical works refered in the introduction, those explain mildly relativistic jets of $0.26~c$ ($\\gamma=1.04$) in SS~433 to highly relativistic jets of $0.92~c$ ($\\gamma=2.55$) in several microquasars.  For ultra-relativistic jets of $0.99~c$ ($\\gamma=10$) supposed in several active galactic nuclei or extremely relativistic jets of $0.9999~c$ ($\\gamma=100$) expected to gamma-ray bursts, it may be necessary some other processes, including an extra-ordinary re-distribution of energy, pair dominant plasmas, an energy deposition from other energy sources, such as a nuclear one via neutrino, and so on.     \\vspace*{1pc}  This work has been supported in part by a Grant-in-Aid for the Scientific Research (18540240 J.F.) of the Ministry of Education, Culture, Sports, Science and Technology."
    },
    "0711/0711.3660_arXiv.txt": {
        "abstract": "{} {We report the identification of cyclical changes in the orbital period of the eclipsing cataclysmic variable HT~Cas.} {We measured new white dwarf mid-eclipse timings and combined with published measurements to construct an observed-minus-calculated diagram covering 29 years of observations.} {The data present cyclical variations that can be fitted by a linear plus sinusoidal function with period $36$~yr and semi-amplitude $\\sim 40$~s.  The statistical significance of this period by an F-test is larger than 99.9 per cent.} {We combine our results with those in the literature to revisit the issue of cyclical period changes in cataclysmic variables and their interpretation in terms of a solar-type magnetic activity cycle in the secondary star. A diagram of fractional period change ($\\Delta P/P$) versus the angular velocity of the active star ($\\Omega$) for cataclysmic variables, RS~CVn, W~UMa and Algols reveal that close binaries with periods above the gap (secondaries with convective envelopes) satisfy a relationship $\\Delta P/P \\propto \\Omega^{-0.7 \\pm 0.1}$.  Cataclysmic variables below the period gap (with fully convective secondaries) deviate from this relationship by more than 3-$\\sigma$, with average fractional period changes $\\simeq 5$ times smaller than those of the systems above the gap.} ",
        "introduction": "\\label{sec:intr}  HT Cassiopeiae (HT~Cas) is a short-period ($P_{{\\rm orb}} = 1.77$ hr) eclipsing cataclysmic variable (CV). In these binaries, a late-type star (the secondary) overfills its Roche lobe and transfers matter to a companion white dwarf (Warner 1995). The evolution of CVs is mainly driven by two ingredients: angular momentum loss, to sustain the mass transfer process, and the response of the secondary to the mass loss. Two mechanisms were proposed for angular momentum loss in CVs. The first is gravitational radiation, which is effective only for short orbital periods (Patterson 1984). The second is angular momentum carried away by a stellar wind magnetically coupled to the secondary surface (magnetic braking mechanism; Rappaport et~al. 1983; King 1988). As CVs are tidally locked binaries, any momentum lost by the secondary is also subtracted from the total orbital momentum of the system, causing a secular decrease of the orbital period.  In the disrupted braking model of CV evolution (Rappaport et~al. 1983; Hameury et~al. 1991), the observed dearth of systems with periods in the range 2-3~hr (known as ``period gap''; Knigge 2006) is explained by a sudden drop in the efficiency of the magnetic braking mechanism when the secondary, evolving from longer orbital periods, reach $P_{{\\rm orb}} \\sim 3$~hr and becomes fully convective ($M_{2} \\sim 0.3\\,M_{\\odot}$). However, there are some key assumptions in the standard model, most notably concerning angular momentum loss, that are inconsistent with both the observed spin-down of young, low-mass stars and theoretical developments in our understanding of stellar winds (Andronov et~al.  2003).  The secular evolution of the binary can in principle be detected by measuring the changes in the orbital period of eclipsing CVs. Eclipses provide a fiducial mark in time and can usually be used to determine the orbital period (and its derivative) with high precision. However, attempts to measure the long-term orbital period decrease in CVs have been disappointing: none of the studied stars show the expected rate of orbital period decrease. Instead, most of the well observed eclipsing CVs\\footnote{i.e., those with well-sampled observed-minus-calculated (O$-$C) eclipse timings diagram covering more than a decade of observations.} show cyclical period changes (Baptista et~al. 2003 and references therein). Cyclical orbital period variations are also observed in other close binaries with late-type components -- Algols, RS CVn and W UMa systems (Lanza \\& Rodon\\`{o} 1999). The most promising explanation of this effect seems to be the existence of a solar-type (quasi- and/or multi-periodic) magnetic activity cycle in the secondary star. A number of mechanisms have been proposed which are capable of producing a modulation of the orbital period on time scales of decades, induced by a variable magnetic field in the convective zone of the late-type component (Matese \\& Whitmire 1983; Applegate \\& Patterson 1987; Warner 1988; Applegate 1992; Richman et~al. 1994; Lanza et~al. 1998; Lanza 2006a). The relatively large amplitude of these cyclical period changes probably contributes to mask the low amplitude, secular period decrease.  This paper reports the results of an investigation of orbital period changes in HT~Cas. The observations and data analysis are described in section \\ref{sec:observ}. A discussion of cyclical orbital period changes in CVs, above and below the period gap, are presented in section \\ref{sec:discuss}. ",
        "conclusions": "\\label{sec:discuss}  Our results reveals that the orbital period of HT~Cas shows conspicuous period changes of semi-amplitude $\\sim$~40~s which seems to repeat on a time-scale of about 36~yr. The present work increases the sample of eclipsing CVs in which orbital period modulations were observed and motivated us to update the comparison of cyclical period changes of CVs above and below the period gap performed by Baptista et~al. (2003).  \\subsection{Orbital period modulations in CVs} \\label{sub:mod}  This section reviews the current observational picture on the detection of cyclical period changes in eclipsing CVs. We first address the observational requirements needed to allow detection of cyclical period modulations and then discuss the observational scenario which emerges when a complete sample is constructed based on these requirements. Cyclical orbital period changes are seen in many eclipsing CVs (see Baptista et~al. 2003). The cycle periods range from 5~yr in IP~Peg (Wolf et~al 1993) to about 36~yr in HT~Cas, whereas the amplitudes are in the range $10-10^{2}$~s.  A successful detection of these cyclical period changes demands an (O$-$C) diagram covering at least one cycle of the modulation (i.e., at least about a decade of observations) and that the uncertainty in the (annually averaged) eclipse timings is smaller than the amplitude of the period modulation to allow a clean detection of the latter. Therefore, decade-long time coverage and high precision eclipse timings (better than 10~s and 20~s, respectively for systems below and above the period gap) are basic requirements. A third key aspect concerns the time sampling of the observations. An (O$-$C) diagram constructed from sparse and infrequent eclipse timing measurements may easily fail to reveal a cyclical period change. Figure~\\ref{fig:simul} illustrates this argument. It shows synthetic (O$-$C) diagrams constructed from a period modulation of 20~yr and amplitude 50~s (dotted line). Gaussian noise of amplitude 5~s was added to the annual timings to simulate the typical uncertainties of a real data set. The best-fit ephemeris is indicated by a solid curve/line in each case. The upper panel shows the case of poor data sampling. The gaps around eclipse cycles $\\sim 30000$ and $60000-100000$ mask the period modulation and the avaliable data (solid circles) is best fit by a linear ephemeris (solid line). HT~Cas itself is a good example of the poor sampling case. The revised linear ephemeris of Feline et~al. (2005) was based on a sparsely sampled (O$-$C) diagram. If the timings of Wood et~al.  (1995), Mukai et~al. (1997) and Ioannou et~al. (1999) were included in their diagram the orbital period modulation would have become clear.  \\begin{figure} \\includegraphics[scale=0.47]{simul.eps} \\caption{Influence of data sampling (upper panel), time coverage (middle panel) and accuracy of the eclipse timings (lower panel) in the detection of orbital period modulations. The dotted line represents (O$-$C) values constructed from a period modulation of 20~yr and amplitude 50~s (20~s in the lower panel). Gaussian noise of amplitudes 5~s (upper and middle panels) and 20~s (lower panel) were added to the (annually sampled) timings to simulate the uncertainties of a real data set. The solid line depicts the best-fit ephemeris obtained for the synthetic data used (solid circles) to illustrate each case.} \\label{fig:simul} \\end{figure}  The middle panel of Fig.~\\ref{fig:simul} illustrates the effect of the time coverage on the detection of a period modulation. In this case the observations cover only about half of the cycle period (solid circles), leading to an incorrect inference of a long-term orbital period decrease (solid curve). A longer baseline is needed to allow identification of the cyclical nature of the period changes. Z~Cha is an illustrative example of this case. Robinson et~al. (1995) inferred a significant period increase from an (O$-$C) diagram covering 18~yr of observations. Only when the time coverage was increased to 30~yr the cyclical behavior of the period changes became clear (Baptista et~al. 2002).  The lower panel of Fig.~\\ref{fig:simul} shows how the accuracy of the eclipse timings affect the ability to detect period modulations. In this case the amplitude of the modulation was reduced to match the larger uncertainty of the eclipse timings (20~s). The period modulation is lost in the noise, despite the fact that the (O$-$C) diagram has good sampling and time coverage (solid circles), and the best-fit ephemeris is the linear one (solid line). FO~Aqr, with inclination $i \\sim 70^{\\circ}$ and grazing eclipses (Hellier et~al. 1989), may be an example of this case. The large uncertainty of its eclipse timings ($\\sim 180$~s) is enough to mask cyclical period modulations of amplitude similar to those seen in other eclipsing CVs.  In summary, in order to be able to detect a period modulation one needs a well-sampled (one data point every 1$-$3~yr, no big gaps) (O$-$C) diagram covering at least a decade of observations, constructed from precise eclipse timings (uncertainty $\\simlt 20$~s).  In order to construct a sample according to these requirements, we searched the CVcat database\\footnote{CVcat/TPP is a web-based interactive database on cataclysmic variable stars ({\\tt http://cvcat.net/}).} for all eclipsing CVs with inclination $i \\geq 70^{\\circ}$. The accuracy of eclipse timings below this limit is not enough to allow detection of period modulations with amplitudes $\\simlt 200$~s. We find 14 eclipsing CVs satisfying the above criteria, 6 systems below and 8 systems above the period gap. They are listed in Table~\\ref{tab:modcvs}. \\emph{All systems in the sample show cyclical period changes}. With the inclusion of HT~Cas, there is presently no CVs with well sampled and precise (O$-$C) diagram covering more than a decade of observations that do not show cyclical period changes. This underscores the conclusion of Baptista et~al. (2003) that cyclical period changes seem a common phenomenon in CVs, being present equally among systems above and below the period gap.   \\begin{table} \\caption{Observed orbital period modulations in CVs} \\begin{tabular}{lcccl} \\hline Object & $P_{{\\rm orb}}$ & $P_{{\\rm mod}}$ & $\\Delta P/P$       & Ref.\\\\ & (hr)            & (yr)            & ($\\times 10^{-6}$) &     \\\\ \\hline V4140~Sgr & 1.47 &  6.9 & 0.93 & 1 \\\\ V2051~Oph & 1.50 & 22.0 & 0.30 & 1 \\\\ OY~Car    & 1.51 & 35.0 & 0.52 & 2 \\\\ EX~Hya    & 1.64 & 17.5 & 0.55 & 3 \\\\ HT~Cas    & 1.77 & 36.0 & 0.44 & This work.\\\\ Z~Cha     & 1.79 & 28.0 & 0.85 & 4 \\\\ IP~Peg    & 3.80 &  4.7 & 8.00 & 5 \\\\ U~Gem     & 4.17 &  8.0 & 3.00 & 6 \\\\ DQ~Her    & 4.65 & 13.7 & 2.00 & 7 \\\\ UX~UMa    & 4.72 & 7.1,\\,10.7,\\,30.4$^{\\ast}$ & 2.60 & 8 \\\\ T~Aur     & 4.91 & 23.0 & 3.10 & 9 \\\\ EX~Dra    & 5.03 &  5.0 & 7.90 & 10 \\\\ RW~Tri    & 5.57 & 7.6,\\,13.6$^{\\ast}$ &2.10 & 11 \\\\ AC~Cnc    & 7.21 & 16.2 & 5.50 & 12 \\\\ \\hline \\multicolumn{5}{l}{$\\ ^{\\ast}$Multiperiodic (the best determined periodicities are considered).} \\\\ \\multicolumn{5}{l}{\\emph{References}: (1) Baptista et~al. 2003; (2) Greenhill et~al. 2006;} \\\\ \\multicolumn{5}{l}{(3) Hellier \\& Sproats 1992; (4) Baptista et~al. 2002; (5) Wolf et} \\\\ \\multicolumn{5}{l}{al. 1993; (6) Warner 1988; (7) Zhang et~al. 1995; (8) Rubenstein} \\\\ \\multicolumn{5}{l}{ et~al. 1991; (9) Beuermann \\& Pakull 1984; (10) Shafter \\&} \\\\ \\multicolumn{5}{l}{Holland 2003; (11) Robinson et~al. 1991; (12) Qian et al. 2007b.} \\end{tabular} \\label{tab:modcvs} \\end{table}  Apsidal motion is not a viable explanation for such period changes because the orbital eccentricity for close binaries is negligible. The presence of a third body in the system has often been invoked as an alternative explanation. However, a light-time effect implies a strictly periodic modulations in orbital period, which is usually not observed when data covering several cycles of the modulation are available. If one is to seek a common explanation for the orbital period modulation seen in CVs, then all periodic effects (such as a third body in the system) must be discarded, since the observed period changes in several of the systems (e.g., UX~UMa, RW~Tri, V2051~Oph) are cyclical but clearly not strictly periodic.  The best current explanation for the observed cyclical period modulation is that it is the result of a solar-type magnetic activity cycle in the secondary star. Amongst several mechanisms proposed to explain such modulations, the hypothesis of Applegate (1992) seems the most plausible. It relates the orbital period modulation to the operation of a hydromagnetic dynamo in the convective zone of the late-type component of close binaries.  More precisely, Applegate's hypothesis assumes that a small fraction of the internal angular momentum of the active component is cyclically exchanged between an inner and an outer convective shell due to a varying internal magnetic torque. This affects the oblateness and the gravitational quadrupole moment of the active component, which oscillates around its mean value. When the quadrupole moment is maximum, the companion star feels a stronger gravitational force, so that it is forced to move closer and faster around the centre of mass, thus attaining the minimum orbital period.  On the other hand, when the quadrupole is minimum, the orbital period exhibits its maximum. Lanza et~al. (1998) and Lanza \\& Rodon\\`{o} (1999) have elaborated more on this idea. The model was applied to a sample of CVs by Richman et~al. (1994). The fractional period change $\\Delta P/P$ is related to the amplitude $\\Delta(O-C)$ and to the period $P_{{\\rm mod}}$ of the modulation by (Applegate 1992),  \\begin{equation} \\frac{\\Delta P}{P}= 2\\pi\\,\\frac{\\Delta(O-C)}{P_{\\rm mod}}= 4\\pi\\,\\frac{A}{C}\\; . \\label{eq:dpp} \\end{equation}  \\noindent Using the values of $A$ and $C$ in Tables~\\ref{tab:ephem}, we find $\\Delta P/P = 4.4 \\times 10^{-7}$ for HT~Cas. The values of $\\Delta P/P$ for all the CVs in our sample are listed in the fourth column of Table~\\ref{tab:modcvs}.  The critical aspect of Applegate's hypothesis is the connection of his model of gravitational quadrupole changes to a realistic cyclic dynamo model capable of produce such modulations. In this regard, R\\\"{u}diger et~al. (2002) presented an $\\alpha^{2}$~dynamo model for RS~CVn stars, adding a dynamo mechanism to Applegate's model. Also, theoretical improvements and observational constrains appeared recently in the literature in an attempt to overcome the limitations faced when Applegate's model is applied to RS~CVn stars (Lanza \\& Rodon\\`{o} 2002, 2004; Lanza 2005, 2006a, 2006b). All these results can be scaled to CVs above the period gap because their secondaries also have convective envelopes.  On other hand, secondaries of CVs below the gap are thought to be fully convective (i.e, masses $M_{2} < 0.3\\,M_{\\odot}$). Because fully convective stars have no overshoot layer, the usual dynamo mechanisms cannot be at work in these stars (Dobler 2005). However, isolated late-type main sequence stars (spectral type M5 and later) show indications of the presence of strong magnetic fields (Hawley 1993; Baliunas et al.  1995; West et~al. 2004). Moreover, if magnetic activity in the secondary star is used as an explanation for the observed period modulations, the CVs below the period gap represent the first sample of fully convective dwarfs with magnetic cycles known. Alternative small-scale dynamo models have been proposed to sustain magnetic fields and induce magnetic activity cycles in fully convective stars (Durney et~al. 1993; Haugen et~al. 2004; Brandenburg et~al. 2005; Dobler 2005 and references therein). These models indicate that an overshoot layer is not a necessary ingredient for the generation of large scale magnetic fields.  The common occurrence of $P_{{\\rm orb}}$ modulations in CVs is consistent with the results of Ak et~al. (2001). They found cyclical variations in the quiescent magnitude and outburst interval of a sample of CVs above and below the period gap, which they attributed to solar-type magnetic activity cycles in the secondary stars. They also found no correlation of the cycle period with the rotation regime of the secondary star (i.e., orbital period, for the phase-locked secondary stars in CVs). Considering the period modulations of a variety of close binary systems, Lanza \\& Rodon\\`{o} (1999) found similar observational evidence for system above and below the period gap.  \\subsection{A comparison of the observed orbital period modulations above and below the period gap} \\label{sub:diff}  Given the sample of well observed CVs which exhibit cyclical orbital period changes, this section attempts to quantify the common behavior as well as to address the systematic differences between the observed modulations in systems above and below the period gap. We extended the comparison by considering the period modulations observed in other longer-period close binaries (Algols, RS CVn and W UMa stars). Their magnetic activity should resemble, in some sense, the magnetic activity of the CVs above the gap since they also have a late-type (active) component with a convective envelope.  Figure~\\ref{fig:dppo} shows a diagram of the fractional period change $\\Delta P/P$ versus the angular velocity $\\Omega$ $= 2\\pi/P_{{\\rm orb}}$ of the active, late-type component star in close binaries (rotation is a key ingredient of the dynamo action, see Lanza \\& Rodon\\`{o} 1999). It includes data from the 14 eclipsing CVs listed in Table~\\ref{tab:modcvs}. The period gap is indicated by vertical dashed lines. The open triangles depict $\\Delta P/P$ values for other close binaries with cyclical period modulations (Lanza \\& Rodon\\'{o} 1999; Qian et al. 1999, 2000a, 2000b, 2002, 2004, 2005, 2007a; Lanza et al. 2001; Kang et al. 2002; Qian 2002a, 2002b, 2003; Yang \\& Liu 2002, 2003a, 2003b; Zavala et al.  2002; \\c{C}ak{\\i}rl{\\i} et al. 2003; Kim et al. 2003; Qian \\& Boonrucksar 2003; Af\\c{s}ar et al. 2004; Lee et al. 2004; Qian \\& Yang 2004; Yang et al. 2004, 2007; Zhu et al. 2004; Borkovits et al.  2005; Qian \\& He 2005; Erdem et al. 2007; Pilecki et al. 2007; Szalai et al.  2007).  Systems with independent evidence that the observed (O$-$C) modulations can be explained by a third body were excluded (actually, it is not possible to exclude the possibility that the observed modulation of some of the systems plotted may still be caused by a third body. However, for most of the systems the observation of the modulation covers more than one cycle and there is some indication that the variation is non-periodic). There is a clear correlation between the fractional period change and the angular velocity, which can be expressed by a relationship of the type $\\Delta P/P \\propto \\Omega^{\\gamma}$ for all close binaries above the period gap. A linear least-squares fit yields $\\gamma = -0.7 \\pm 0.1$.  The best-fit solution and the 3-$\\sigma$ confidence level are indicated by a solid and dotted lines, respectively, in Fig.~\\ref{fig:dppo}.  \\begin{figure} \\includegraphics[scale=0.36]{dppo.eps} \\caption{Diagram of the fractional period change $\\Delta P/P$ versus the angular velocity $\\Omega$ of the active component star in short-period CVs (filled circles), long-period CVs (filled squares) and other close binaries from the literature (open triangles). The period gap is represented by the vertical dashed lines. The best-fit linear relation and its 3-$\\sigma$ confidence level are shown as solid and dotted lines, respectively. The diagram includes the measurements of the CVs listed in Table~\\ref{tab:modcvs}.} \\label{fig:dppo} \\end{figure}  We find that the average fractional period changes of the short-period CVs ($\\overline{\\Delta P/P} \\simeq 8 \\times 10^{-7}$) are systematically smaller than those of the long-period CVs ($\\overline{\\Delta P/P} \\simeq 4 \\times 10^{-6}$) by a factor $\\simeq 5$. The mean value of $\\Delta P/P$ of the short-period CVs is more than 3-$\\sigma$ below the value expected from the linear relation inferred for the other, longer period close binaries. An attempt to simultaneously explain CVs above and below the gap by vertically displacing the $\\Delta P/P$ versus $\\Omega$ fitted relation results in a poor fit with high $\\chi^{2}_{\\nu}$. Similarly, including the CVs below the period gap in the sample degrades the fit and significantly increases the $\\chi^{2}_{\\nu}$. Despite the small sample (14 objects), there is a statistically significant difference between the $\\Delta P/P$ values of the CVs above and below the period gap.  This difference cannot be eliminated even if we take into account the fitted $\\Delta P/P$ versus $\\Omega$ relation that predicts decreasing $\\Delta P/P$ values for increasing $\\Omega$.  Thus, we are forced to conclude that the CVs below the period gap do not fit the relation valid for the binaries above the period gap.  If the interpretation of cyclical period changes as the consequence of a solar-like magnetic activity cycle is correct, the existence of cyclical period changes in binaries with fully convective active stars is an indication that these stars do have magnetic fields, not only capable of inducing strong chromospheric activity (e.g, Hawley 1993), but also measurable magnetic activity cycles (Ak et~al.  2001; this paper). On the other hand, the fact that the fractional period changes of binaries with fully convective stars is systematically smaller than those of stars with radiative cores is a likely indication that a different mechanism is responsible to generate and sustain their magnetic fields. In this regard, the observed lower $\\Delta P/P$ values yield a useful constrain to any model that might be developed to account for magnetic fields in fully convective stars.  The investigation of cyclical period changes in CVs will greatly benefit from the increase in the (presently small) sample of systems with well-sampled (O$-$C) diagrams covering more than a decade of observations. This demands the patient but systematic collection of precise eclipse timings over several years. It is worth mentioning that there is a good number of well-known eclipsing CVs for which a few years of additional eclipse timings observations would suffice to overcome the ambiguities illustrated in Fig.~\\ref{fig:simul} and to allow statistically significant detection of period modulations. We also remark that none of the systems \\emph{inside} the period gap has been yet observed for long enough time to allow identification of cyclical period changes. It would be interesting to check whether these systems will fit the $\\Delta P/P \\propto \\Omega^{-0.7}$ relation or will display a behavior similar to that of the CVs below the period gap.  In order to explain the CV period gap, the disrupted braking model predicts a significant reduction in magnetic braking efficiency between the systems above and below the period gap (Hameury et~al. 1991). The observed difference in $\\Delta P/P$ between CVs above and below the period gap could be used to test this prediction if a connection between fractional period changes and braking efficiency could be established."
    },
    "0711/0711.4795_arXiv.txt": {
        "abstract": "Traditionally globular clusters and dwarf spheroidal galaxies have been distinguished by using one or more of the following criteria: (1) mass, (2) luminosity, (3) size, (4) mass-to-light ratio and (5) spread in metallicity. However, a few recently discovered objects show some overlap between the domains in parameter space that are occupied by galaxies and clusters. In the present note it is shown that ellipticity can, in some cases, be used to help distinguish between globular clusters and dwarf spheroidal galaxies. ",
        "introduction": "Recent discoveries of exceedingly faint dwarf spheroidal galaxies have shown quite dramatically that luminosity alone cannot be relied on to distinguish between dwarf spheroidal galaxies and globular clusters. Interesting examples of faint dwarfs are Bo\\\"{o}tes II [M$_{v}$ = -3.1] \\citep{      wjw07}, Coma [M$_{v}$ = -3.7] \\citep{bel07} and Willman 1 [M$_{v}$ = -2.5] \\citep{mar07}. All of these galaxies are fainter than the overwhelming majority of globular clusters.  The size of a globular clusters is best described by its half-light radius R$_{h}$, because this parameter remains almost invariant over $\\sim$ 10 cluster relaxation times. The vast majority of globular clusters have R$_{h} <$ 10 pc. However, some well-established globular clusters have quite large radii. Examples are NGC 2419 (R$_{h}$ = 18 pc) and Palomar 14 (R$_{h}$ = 25 pc). On the other hand Willman 1 \\citep{mar07}, which has a half-light radius of $\\sim$ 20 pc, has generally been regarded as a galaxy.  This overlap in size and luminosity raises deep questions about the nature of the distinction between these two classes of objects. It should be noted that the distinction between galactic nuclei and globular clusters is also somewhat artificial. Over time tidal forces will, for example, strip away much of the stellar population of the Sagittarius dwarf spheroidal, leaving behind only its nucleus the luminous globular cluster M54. \\citet{vdbm04} and \\citet{macvan05} have suggested that the position of an object in a plot of M$_{v}$ versus log R$_{h}$ might be used to discriminate between globular clusters and the stripped cores of dwarf galaxies, such as $\\omega$ Centauri. However, this proposal now seems less attractive than it once did . In particular \\citet{vdb07} has recently noted that the brightest red objects in the halo of the elliptical galaxy NGC 5128 appear to form a continuum in the M$_{v}$ versus log R$_{h}$, {\\bf indicating that globular clusters and dwarf spheroidals may not be clearly separate and distinct types of objects.  The same point has recently been made about the brightest objects in NGC5128 by \\citet{barmby07} and by \\citet{rej07}}. Furthermore NGC 2419, which on the basis of its position in the M$_{v}$ versus log R$_{h}$ plane, had been classified as a stripped galaxy core, turns out to have a small metallicity dispersion and lacks a significant population of extra-tidal stars \\citep{bellrip07}. Both of these factors militate against the hypothesis that NGC 2419 is actually a stripped galaxy core. On the other hand the stripped core suspect B514 in the Andromeda galaxy does seem to be embedded in a very low surface density dwarf spheroidal \\citep{fre07}. In summary it appears that the location of a galaxy in the Mv versus log Rh plot is not always a reliable way of separating globular clusters from the stripped cores of dwarf spheroidal galaxies.  The pioneering investigations of \\citet{bab39} and \\citet{oor40} first demonstrated that dark matter provides a significant contribution to the masses of individual galaxies. More recently \\citet{mat98} showed that such dark matter becomes more and more dominant as one proceeds to study ever dimmer galaxies. It has therefore become customary to regard the presence of dark matter as the touchstone that allows one to unambiguously distinguish between galaxies and star clusters. Furthermore, it is often difficult and time consuming to obtain velocity dispersions of faint stars in distant dwarf galaxies. The problems outlined above suggest that it might be useful to have additional criteria to help discriminate between dwarf spheroidal galaxies and globular clusters. ",
        "conclusions": "Recent discoveries have shown that there is some overlap in parameter space between the regions occupied by globular clusters and dwarf spheroidal galaxies. It is shown that cluster flattening may be used as an additional parameter to help to distinguish between these two classes of objects. There are, of course, many cases [such as the Fornax dwarf and its globular clusters] where it is quite obvious which object is the dwarf spheroidal and which ones are its companion globular clusters. However, there are other cases where this distinction is not quite so obvious \\citep{lee07}. For example, Boo I and Boo II appear to have similar distances and are located at a projected separation of only 1.8 kpc \\citet{wal07}. Are these two objects companion galaxies, or is Boo II (M$_{v}$ = -3.1, R$_{h}$ =72 pc) a globular cluster companion to Boo I (M$_{v}$ = -5.7, R$_{h}$ = 227 pc)? Inspection of the outer isophotes of Boo II published by \\citet{wal07} suggests that 0.1 $\\la$ (a-b)/a $\\la$ 0.2, which is consistent with its being either a galaxy or a cluster.  In summary it is concluded that dwarf spheroidal galaxies are, on average, significantly flatter than globular clusters. In some cases this may help to distinguish these two classes of objects. Globular clusters that are probably the remnant cores of dwarf spheroidal galaxies appear to have present day flattening values that are, on average, intermediate between these two classes of objects."
    },
    "0711/0711.0077.txt": {
        "abstract": "Observations provide increasingly strong evidence that the universe is accelerating. This revolutionary advance in cosmological observations confronts theoretical cosmology with a tremendous challenge, which it has so far failed to meet. Explanations of cosmic acceleration within the framework of general relativity are plagued by difficulties. General relativistic models are nearly all based on a dark energy field with fine-tuned, unnatural properties. There is a great variety of models, but all share one feature in common -- an inability to account for the gravitational properties of the vacuum energy. Speculative ideas from string theory may hold some promise, but it is fair to say that no convincing model has yet been proposed. An alternative to dark energy is that gravity itself may behave differently from general relativity on the largest scales, in such a way as to produce acceleration. The alternative approach of modified gravity (or dark gravity) provides a new angle on the problem, but also faces serious difficulties, including in all known cases severe fine-tuning and the problem of explaining why the vacuum energy does not gravitate. The lack of an adequate theoretical framework for the late-time acceleration of the universe represents a deep crisis for theory -- but also an exciting challenge for theorists. It seems likely that an entirely new paradigm is required to resolve this crisis. ",
        "introduction": "The current ``standard model\" of cosmology is the inflationary cold dark matter model with cosmological constant, usually called LCDM, which is based on general relativity and particle physics (i.e., the Standard Model and its minimal supersymmetric extensions). This model provides an excellent fit to the wealth of high-precision observational data, on the basis of a remarkably small number of cosmological parameters~\\cite{data}. In particular, independent data sets from cosmic microwave background (CMB) anisotropies, galaxy surveys and supernova luminosities, lead to a consistent set of best-fit model parameters -- which represents a triumph for LCDM.  The standard model is remarkably successful, but we know that its theoretical foundation, general relativity, breaks down at high enough energies, usually taken to be at the Planck scale, \\begin{equation} E \\gtrsim M_p \\sim 10^{16}\\,\\mbox{TeV}\\,. \\end{equation} The LCDM model can only provide limited insight into the very early universe. Indeed, the crucial role played by inflation belies the fact that inflation remains an effective theory without yet a basis in fundamental theory. A quantum gravity theory will be able to probe higher energies and earlier times, and should provide a consistent basis for inflation, or an alternative that replaces inflation within the standard cosmological model (for recent work in different directions, see e.g. Refs.~\\cite{qg}).  An even bigger theoretical problem than inflation is that of the late-time acceleration in the expansion of the universe~\\cite{de}. In terms of the fundamental energy density parameters, the data indicates that the present cosmic energy budget is given by \\be \\label{olom} \\Omega_\\Lambda &\\equiv& {\\Lambda \\over 3H_0^2}\\approx 0.75\\,,~~ \\Omega_{m} \\equiv {8\\pi G\\rho_{m0} \\over 3H_0^2}\\approx 0.25\\,, ~~ \\Omega_{K} \\equiv {-K\\over H_0^2} \\approx 0 \\,, \\\\ \\Omega_r &\\equiv& {8\\pi G\\rho_{r0} \\over 3H_0^2}\\approx 8\\times 10^{-5}\\,. \\ee Here $H_0$ is the present value of the Hubble parameter, $\\La$ is the cosmological constant, $K$ is spatial curvature, $\\rho_{m0}$ is the present matter density and $\\rho_{r0}$ is the present radiation density. $G$ is Newton's constant. The Friedman equation is \\begin{eqnarray} \\left({\\dot a\\over a}\\right)^2 \\equiv H^2 &=& {8\\pi G\\over 3}(\\rho_m+\\rho_r)+{\\Lambda \\over 3}-{K \\over a^2} \\nonumber\\\\ &=& H_0^2\\left[\\Omega_m(1+z)^3+ \\Omega_r(1+z)^4 +\\Omega_\\Lambda+ \\Omega_K(1+z)^2 \\right], \\label{h} \\end{eqnarray} where $a$ denotes the scale factor which is related to the cosmological redshift by $z=a^{-1}-1$. We normalize the present scale factor to $a_0=1$. Together with the energy conservation equation this implies \\begin{equation} {\\ddot a \\over a} = -{4\\pi G \\over 3} \\left(\\rho_m+2\\rho_r\\right) + {\\Lambda \\over 3} \\,. \\label{acc} \\end{equation} The observations, which together with Eq.~(\\ref{h}) lead to the values given in Eq.~(\\ref{olom}), produce via Eq.~(\\ref{acc}) the dramatic conclusion that the universe is currently accelerating, \\be \\ddot a_0>0\\,. \\ee  This conclusion holds only if the universe is (nearly) homogeneous and isotropic, i.e., a Friedmann-Lema\\^\\i tre model. In this case the distance to a given redshift $z$, and the time elapsed since that redshift, are tightly related via the only free function of this geometry, $a(t)$. If the universe instead is isotropic around us but not homogeneous, i.e., if it resembles a Tolman-Bondi--Lema\\^\\i tre solution with our galaxy cluster at the centre, then this tight relation between distance and time for a given redshift would be lost and present data would not necessarily imply acceleration. This point is discussed in detail in the contribution by Enqvist~\\cite{kari}. Of course isotropy without homogeneity violates the Copernican principle as it puts us in the centre of the Universe. However, it has to be stressed that up to now observations of homogeneity are very limited, unlike isotropy, which is firmly established. Homogeneity is usually inferred from isotropy together with the Copernican principle. With future data, it will in principle be possible to distinguish observationally an isotropic but inhomogeneous universe from an isotropic and homogeneous universe (see e.g.~\\cite{Goodman}). In the following, we disregard this possibility and assume that the Copernican principle applies.  The data also indicate that the universe is currently (nearly) spatially flat, \\be |\\Omega_K|\\ll 1\\,. \\ee It is common to assume that this implies $K=0$ and to use inflation as a motivation. However, inflation does {\\em not} imply $K=0$, but only $\\Omega_K\\to 0$. In the late universe, the distinction may be negligible. But in the very early universe, a nonzero curvature can have significant effects (see e.g.~\\cite{ellmaa}). In fact, if curvature is small but non-vanishing, neglecting it in the analysis of Supernova data can sometimes induce surprisingly large errors, as  discussed in the contribution by Hlozek et al.~\\cite{brucechris}.  These results are illustrated in Fig.~\\ref{sn} (taken from~\\cite{Knop:2003iy,Wood-Vasey:2007jb}). A detailed discussion of the experimental aspects of the late-time acceleration is given in the contributions by Leib\\-und\\-gut~\\cite{Leibu}, Nichol~\\cite{Nichol} and Sarkar~\\cite{Sarkar}.    \\begin{figure*} \\begin{center} \\includegraphics[height=3.25in,width=3.in]{SNCMBGC.eps}\\quad \\includegraphics[height=3.25in,width=3.in]{SNLS_ESS_Gold_jan07.eps} \\end{center} \\caption{Observational constraints in the $(\\Omega_{m},\\Omega_\\Lambda)$ plane: joint constraints (left) (from~\\cite{Knop:2003iy}); recent compilation of supernova constraints (right) (from~\\cite{Wood-Vasey:2007jb}). The line $\\Om_K=\\Om_m+\\Om_\\La-1=0$ is also indicated.} \\label{sn} \\end{figure*}    The simplest option is probably a cosmological constant, i.e., the LCDM model. Even though the cosmological constant can be considered as simply an additional gravitational constant (in addition to Newton's constant), a cosmological constant enters the Einstein equations in exactly the same way as a contribution from the vacuum energy, i.e., via a Lorentz-invariant energy-momentum tensor $T^{\\rm vac}_{\\mu\\nu}= -(\\Lambda/8\\pi G) g_{\\mu\\nu}$. The only observable signature of both a cosmological constant and vacuum energy is their effect on spacetime -- and so a vacuum energy and a classical cosmological constant cannot be distinguished by observation. Therefore the `classical' notion of the cosmological constant is effectively physically indistinguishable from a quantum vacuum energy.  Even though the absolute value of vacuum energy cannot be calculated within quantum field theory, {\\em changes} in the vacuum energy (e.g. during a phase transition) can be calculated, and they do have a physical effect -- for example, on the energy levels of atoms (Lamb shift), which is well known and well measured. Furthermore, differences of vacuum energy in different locations, e.g., between or on one side of two large metallic plates, have been calculated and their effect, the Casimir force, is well measured~\\cite{casi}. Hence, there is no doubt about the reality of vacuum energy. For a field theory with cutoff energy scale $E$, the vacuum energy density scales with the cutoff as $\\rho_{\\rm vac} \\sim E^4$, corresponding to a cosmological constant $\\La_{\\rm vac} =\\rho_{\\rm vac}/(8\\pi G)$.  If $E=M_{p}$, this yields a renormalization of the'cosmological constant' of about $\\La_{\\rm vac} \\sim 10^{38}$GeV$^2$, whereas the measured effective cosmological constant is the sum of the 'bare' cosmological constant and the contribution from renormalization, \\be \\La_\\mathrm{eff}= \\La_{\\rm vac}+\\La \\simeq 10^{-83}\\,\\mathrm{GeV}^2~. \\label{v1} \\ee Hence a cancellation of about 120 orders of magnitude is required. This is called the {\\em fine tuning} or {\\em size } problem of dark energy: a cancellation is needed to arrive at a result which is many orders of magnitude smaller than each of the terms.\\footnote{In quantum field theory we actually have to add to the cut-off term $\\La_{\\rm vac} \\simeq E_c^4/M_{pl}^2$ the unmeasureable 'bare' cosmological constant. In this sense, the cosmological constant problem is a fine tuning between the unobservable 'bare' cosmological constant and the term coming from the cut-off scale.} It is possible that the quantum vacuum energy is much smaller than the Planck scale. But even if we set it to the lowest possible SUSY scale, $E_{\\rm susy}\\sim 1$TeV, arguing that at higher energies vacuum energy exactly cancels due to supersymmetry, the required cancellation is still about 60 orders of magnitude. These issues are discussed in the contributions by Padmanabhan~\\cite{Paddy} and Bousso~\\cite{rafi}.  A reasonable attitude towards this open problem is the hope that quantum gravity will explain this cancellation. But then it is much more likely that we shall obtain directly $\\La_{\\rm vac}+\\La=0$ and not $\\La_{\\rm vac}+\\La \\simeq 3\\rho_m(t_0)/(8\\pi G)$. This unexpected observational result leads to a second problem, {\\em the coincidence problem}: given that \\be \\rho_\\Lambda = {\\Lambda_\\mathrm{eff} \\over 8\\pi G}=\\,\\mbox{constant}\\,,~\\mbox{ while }~\\rho_m \\propto (1+z)^3\\,, \\ee why is  $\\rho_\\Lambda$ of the order of the {\\em present} matter density $\\rho_m(t_0)$? It was completely negligible in most of the past and will entirely dominate in the future.  Instead of a cosmological constant, one may also introduce a scalar field or some other contribution to the energy-momentum tensor which has an equation of state $w<-1/3$. Such a component is called `dark energy'. So far, no consistent model of dark energy has been proposed which can yield a convincing or natural explanation of either of these problems. A variety of such models is discussed in the contribution by Linder~\\cite{linder}.  Alternatively, it is possible that there is no dark energy field, but instead the late-time acceleration is a signal of a {\\em gravitational} effect. Within the framework of general relativity, this requires that the impact of inhomogeneities somehow acts to produce acceleration, or the appearance of acceleration (within a Friedman-Lema\\^{i}tre interpretation). One possibility is the Tolman-Bondi--Lema\\^\\i tre model discussed in this volume~\\cite{kari}. Another possibility is that the `backreaction' of inhomogeneities on the background, treated via nonlinear averaging, produces effective acceleration. This is discussed in the contribution by Buchert~\\cite{thomas}.  A more radical version is the `dark gravity' approach, the idea that gravity itself is weakened on large-scales, i.e., that there is an ``infrared'' modification to general relativity that accounts for the late-time acceleration. Specific classes of models which modify gravity are discussed in the contributions by Capozziello and Francaviglia~\\cite{CapFran} and by Koyama~\\cite{koyama}. Schematically, we are modifying the geometric side of the field equations, \\be G_{\\mu\\nu}+G^{\\rm dark}_{\\mu\\nu} = 8\\pi G T_{\\mu\\nu}\\,, \\label{mod} \\ee rather than the matter side, \\be G_{\\mu\\nu} = 8\\pi G \\left(T_{\\mu\\nu}+ T^{\\rm dark}_{\\mu\\nu} \\right)\\,, \\ee as in the general relativity approach. Modified gravity represents an intriguing possibility for resolving the theoretical crisis posed by late-time acceleration. However, it turns out to be extremely difficult to modify general relativity at low energies in cosmology, without violating the low-energy solar system constraints, or without introducing ghosts and other instabilities into the theory. Up to now, there is no convincing alternative to the general relativity dark energy models -- which themselves are not convincing.  The plan of the remainder of this paper is as follows. In Section~2 we discuss constraints which one may formulate for a dark energy or modified gravity (dark gravity) theory from basic theoretical requirements. In Section~3 we discuss models that address the dark energy problem within general relativity. In Section~4 we present modified gravity models. The ideas outlined in Sections~3 and 4 are discussed in more detail in the specific contributions of this issue which are devoted to them. In Section~5 we conclude.  %We use units such that $c=\\hbar=1$. ",
        "conclusions": "\\label{s:con}    The evidence for a late-time acceleration of the universe continues to mount, as the number of experiments and the quality of data grow -- dark energy or dark gravity appear to be an unavoidable reality of the cosmos. This revolutionary discovery by observational cosmology, confronts theoretical cosmology with a major crisis -- how to explain the origin of the acceleration. The core of this problem may be ``handed over\" to particle physics, since we require at the most fundamental level, an explanation for why the vacuum energy either has an incredibly small and fine-tuned value, or is exactly zero. Both options violently disagree with naive estimates of the vacuum energy.  If one accepts that the vacuum energy is indeed nonzero, then the dark energy is described by $\\Lambda$, and the LCDM model is the best current model. The cosmological model requires completion via developments in particle physics that will explain the value of the vacuum energy. In many ways, this is the best that we can do currently, since the alternatives to LCDM, within and beyond general relativity, do not resolve the vacuum energy crisis, and furthermore have no convincing theoretical motivation. None of the contenders appears any better than LCDM.  Presently, perhaps the simplest and most appealing contender is the DGP brane-world model. However, the simplicity of its Friedman equation is deceptive, and the complexity of its cosmological perturbations includes the problem of its ghost.  In view of all this, it is fair to say that at the theoretical level, there is as yet no serious challenger to LCDM. It remains worthwhile to continue investigating alternative dark energy and dark gravity models, in order better to understand the space of possibilities, the variety of cosmological properties, and the observational strategies needed to distinguish them.  At the same time, it is in principle possible that cosmological observations, having discovered dark energy/ dark gravity, could rule out LCDM, by showing, to some acceptable level of statistical confidence, that $w\\neq -1$.  Finally, the theoretical crisis does not have only negative implications: dark energy/ dark gravity in the cosmos provides exciting challenges for theory and observations."
    },
    "0711/0711.2450_arXiv.txt": {
        "abstract": "The mixing between sterile and active neutrinos is taken into account in the calculation of Big Bang Nucleosynthesis. The abundances of primordial elements, like D, $^3$He, $^4$He and $^7$Li, are calculated by including sterile neutrinos, and by using finite chemical potentials. It is found that the resulting theoretical abundances are consistent with WMAP data on baryonic densities, and with limits of LSND on mixing angles, only if $^7$Li is excluded from the statistical analysis of theoretical and experimental results. ",
        "introduction": "In recently published papers \\cite{kishimoto,smith} the sensitivity of the $^4$He primordial abundance, upon distortions of the light neutrino spectrum induced by couplings with a sterile neutrino, was analyzed. The effects due to the mixing between sterile and active neutrinos \\cite{kishimoto,smith} reflect upon Big Bang Nucleosynthesis (BBN) in a noticeable manner. Previous studies on this matter can be found in \\cite{cirelli1,cirelli2,cirelli3}, where the effects of mixing upon BBN in presence of primordial leptonic asymmetry have been investigated, and stringent limits on the mixing due to BBN have been presented. Similar studies have been presented in \\cite{maxin1,maxin2}. A review on inclusion of sterile neutrino in cosmology was presented in \\cite{maxin3}.  The results of \\cite{kishimoto} may be taken as a solid starting point for a systematic analysis of the sterile-active neutrino mixing upon cosmological observables, like the primordial abundance of light elements. By the other hand, the mixing mechanism between sterile and active neutrinos has been studied in detail (see \\cite{keranen03} and references therein), so that the calculation of neutrino distribution functions can readily be performed. The information about the neutrino distribution function, in the flavor basis and at a given temperature, is an essential element in the calculation of the neutron decay rate, which is a critical quantity entering BBN \\cite{BBF88,Esma91}.  Direct physical consequences upon BBN, due to the mixing between active and sterile neutrinos, have been explored in \\cite{abazajian,bell}. Following the arguments presented in \\cite{abazajian}, and in the framework of the standard cosmological model, sterile neutrinos would produce a faster expansion rate for the Universe and a higher yield of $^4$He. This is, indeed, a severe constraint on neutrino mixing since a higher predicted abundance of $^4$He may be in conflict with observational data \\cite{abazajian}. Another constraint on active-sterile neutrino mixing is the neutrino mass derived from Cosmic Microwave Background Anisotropy (CMB) \\cite{hannestad}.The analysis of constraints presented in \\cite{seljak} focus on the mixing scheme at the level of the neutrino mass hierarchy, and it suggests the adequacy of the non-degenerate mass hierarchy to set limits on the mass difference between active and sterile neutrinos, $\\delta m^2_{a-s}$. The study of \\cite{seljak} confirms the notion about the convenience of the three active + one sterile neutrinos scheme.  In standard BBN calculations, the mixing of sterile and active neutrinos affects the leptonic fractional occupancies, which are essential quantities appearing in the expression of the weak decay rates. Thus, one needs to know, as input of the calculations, the parameters of the proposed mixing scheme, the neutrino mass hierarchy and the leptonic densities \\cite{bell}. With these elements one can calculate neutron-decay-rates and neutron abundances, by assuming the freeze-out of weak interactions \\cite{BBF88,abazajian}. The effective number of neutrino generations, $N_{\\nu}$, is fixed by the analysis of CMB \\cite{steigman,barger1,barger2}. Current limits on the neutrino degeneracy parameter, for light (electron) neutrinos, $\\eta_l$ \\cite{smith}, runs from $-0.1$ to $0.3$ \\cite{barger1,barger2}. For a detailed presentation of the formalism, in the context of relic-neutrino asymmetry evolution see \\cite{bell}.  In this work we focus on the calculation of the abundances of D, $^3$He, $^4$He and $^7$Li, in presence of sterile-active neutrino mixing in the three flavor scenario, and for the normal and inverse neutrino mass hierarchies \\cite{keranen03}. We have compared the calculated values with data \\cite{muchosd,muchosh,muchosl} and determined the compatibility between them by performing a $\\chi^2$ statistical analysis. Since the theoretical expressions depend on the mixing angle $\\sin^2 2\\phi$, the square mass difference $\\delta m^2_{14}$ (normal mass hierarchy) or $\\delta m^2_{34}$ (inverse mass hierarchy), and the baryonic density $\\Omega_B h^2$, we have adopted the LSND limits on the mixing angle \\cite{eitel00,athana95,McGregor03} and the WMAP results on the baryonic density \\cite{wmap06}, as constraints.  The paper is organized as follows. In Section \\ref{formalism} we briefly present the essentials of the formalism. Section \\ref{decay} is devoted to the calculation of the neutron decay rate and BBN abundances. In Section \\ref{results} we present and discuss the results of the calculations. Conclusions are drawn in Section \\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  In this work we have calculated BBN abundances, by including the mixing between active and sterile neutrinos. As pointed out by Kishimoto et al, the BBN abundances are sensitive to active-sterile neutrino mixing, indeed. Kishimoto et al. \\cite{kishimoto} have demonstrated the sensitivity of the $^4$He abundance on the distortion of the light neutrino spectrum produced by the mixing with a sterile neutrino. In our case, the statistical analysis of the compatibility between theoretical and observed nuclear abundances, indicates the existence of a clear sensitivity of the results upon active-sterile neutrino mixing, too. In performing our analysis, we have considered the WMAP baryonic density, together with the LSND constraint on the sterile-active neutrino mixing. The comparison between calculated and observed abundances indicates some sort of anomaly in the abundance of $^7$Li. Similar difficulties, related to the determination of the abundance of $^7$Li, have been reported previously \\cite{richard05}, in the context of the physics of the interior of stars. We found that the consideration of the abundance of $^7$Li, in presence of active-sterile neutrino mixing, excludes the WMAP value of the baryonic density. This exclusion is not observed when the other nuclear abundances are not included in the analysis."
    },
    "0711/0711.2720_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:SNeintro}  Type Ia supernovae (SNe)  are very useful tools to investigate our universe. For example, the distance-redshift relation obtained as a result of the observations of these strongly suggests the present acceleration of the cosmic volume expansion of our universe\\cite{Riess:1998cb,Perlmutter:1998np, Knop,Riess:2004nr}. To confirm this indication, further projects, such as the ESSENCE project at NOAO,\\footnote{http://www.ctio.noao.edu/wproject/} the Large Synoptic Survey Telescope\\footnote{http://www.lsst.org/lsst\\_home.shtml} and Supernova/Acceleration Probe (SNAP),\\footnote{http://snap.lbl.gov} are now active or planned.  Current observational data points on the distance-redshift plane are somewhat scattered, and this is an origin of the error in estimating the cosmological parameters. One of the reasons for this distance dispersion is gravitational lensing effects due to the mass inhomogeneities in our universe. Although the dispersion due to gravitational lensing is relatively smaller than other effects at low redshift, it may become more prominent and comparable to the intrinsic dispersion at high redshift $z\\gtrsim 1.2$ \\cite{Sarkar:2007sp,Holz:2004xx}. In this sense, the mass inhomogeneities are obstacles to the determination of the cosmological parameters through the observation of Type Ia SNe. On the other hand, the mass inhomogeneities contain rich information about the physical process of the evolution of our universe. Thus, the mass inhomogeneities themselves are very significant subjects in cosmology.   We can understand the property of the mass inhomogeneities by comparing the observational data about type Ia SNe with theoretical predictions obtained through the investigations of the gravitational lensing effects on the light rays from them. This is just the purpose of this paper.  There are many works about gravitational lensing related to SNe: effects on SNe observations \\cite{1991ApJ...374...83R, 1997ApJ...475L..81W, Frieman:1996xk, 1998ApJ...506L...1H, 2000ApJ...532..679P, 2000MNRAS.318..195B, 2000A&A...354..767V, Holz:2004xx, 2006PhRvL..96b1301C, Martel:2007fh}, their availability for investigation of our universe \\cite{1999ApJ...519L...1M, 1999MNRAS.305..746M, 1999A&A...351L..10S, 2000ApJ...534...29H, 2001ApJ...559...53M, 2001ApJ...556L..71H, 2007arXiv0711.0743L}, the extraction of the evidence for lensing effects from observational data \\cite{2004MNRAS.351.1387W, 2005MNRAS.358..101M, Wang:2004ax, 2007JCAP...06....2J} and others \\cite{1998PhRvD..58f3501H, 2000PhLB..486..249G, Wang:1999bz, 2002MNRAS.335.1061S, Wang:2002qc, Gunnarsson:2003fy, Wang:2003gz, 2006astro.ph..1683M, 2007PhRvL..98g1302M}.  In this paper, we particularly focus on the effects of small-scale structures. The smaller scale inhomogeneity may affect the apparent magnitude of sources owing to gravitational focusing. The mass scale of the inhomogeneities, which is considered in this paper, is between $10^{-2}M_\\odot$ and $10^{10}M_\\odot$. There are several works related to this subject \\cite{1998PhRvD..58f3501H,1999ApJ...519L...1M,2007PhRvL..98g1302M,2001ApJ...559...53M,1999A&A...351L..10S}. In these works, compact lens objects were mainly studied, whereas, in this paper, we consider the extended lens objects. Even in the case of the extended lens objects, the strong lensing effects that cause multiple images are of considerable importance if the density profiles of the objects are significantly steep at the center.  We  consider  the flat $\\Lambda$CDM model with $\\Omega_{\\Lambda 0}=0.71, \\Omega_{\\rm m0}=0.29$ as the starting point. In the cold dark matter cosmology, the first objects to form are of subgalactic size. Then, larger structures form through tidal interaction and mergers of smaller objects. Therefore, the dark matter may form numerous clumps. For simplicity, we assume that all matter in the universe takes the form of randomly distributed objects, each of which has finite size and is transparent for light rays. We investigate the magnification effect due to the gravitational lensing on light rays from the point sources such as Type Ia supernovae. The light rays from distant supernova events suffer multiple gravitational lensing effects from the clumps in the universe. The propagation of the light rays is thus stochastic owing to the randomness of the distribution of the clumps. In this paper, we focus on magnification probability distribution function (MPDF). We introduce a simple method to simulate multiple gravitational lensing effects in the clumpy universe model, which is based on our previous study \\cite{Yoo:2006gn}. We perform Monte-Carlo simulations for several lens models, and we investigate the dependence of the MPDF on lens models of dark matter clumps.   Throughout this paper, it is assumed that light sources are pointwise. This assumption is valid only when the linear extent of a source is much smaller than the Einstein radius of each lens object. The linear extent of the SNe is of the order of $\\sim10^{15}$ cm, which is the typical radius of SNe at the peak brightness. The Einstein radius of a lens object of mass $M_L$ is of the order $\\sqrt{GM_LD/c^2}$, where $D$ is the distance to supernovae, $G$ and $c$ are the gravitational constant and speed of light, respectively. Therefore, we assume that $M_L$ satisfies $\\sqrt{GM_LD/c^2}\\gg10^{15}{\\rm cm}$. Because in the case of high-redshift supernova events, $D\\sim 1{\\rm Gpc}$, this assumption leads \\begin{equation} M_L\\gg10^{-2}M_\\odot. \\end{equation}  Our method is complementary to the methods using the N-body simulation. In the N-body simulation, the nonlinear time evolution and the spatial correlation of the dark matter distribution can be studied. Very large scale N-body simulations are possible now. Nevertheless, the questions that can be definitively answered with N-body simulations are still limited by finite resolution in mass and in distance \\cite{BinneyTremaine}. Since the small-scale structure of the dark matter is essential for the gravitational lensing effects, the simulations using analytic density profiles are important to confirm the results from the N-body simulation, and to investigate new effects that are not found in the N-body simulations.     This paper is organized as follows. In \\S\\ref{sec:rev}, we briefly review the gravitational lensing. \\S\\ref{sec:cluni} is devoted to describe the clumpy universe model. In \\S\\ref{sec:mag}, we show a definition of magnification and the relation between the distance and the magnification. In \\S\\ref{sec:nucalre}, numerical calculation and our results are illustrated. Finally, \\S\\ref{sec:consum} is devoted to the conclusions and summary. Throughout this paper, we use the unit of $G=c=1$. ",
        "conclusions": "\\label{sec:consum} We have studied lensing effects on observations of point sources such as Type Ia supernovae by the Monte-Carlo simulation. It has been assumed that all matter in the universe takes the form of randomly distributed objects each of which has finite size and transparent to light rays. In addition, we have assumed that each of the lens objects is axially symmetric along the line of sight. We have found analytically that the magnification probability distribution functions (MPDFs) depend on the mass $M_L$ and the size $R$ of a lens object only through the form $R/\\sqrt{M_L}$. We have calculated MPDFs for various values of the source redshift  and $R/\\sqrt{M_L}$. We found that the resultant MPDFs can be categorized into two groups according to whether the gamma distributions fit well the MPDFs or not. In the case of the lens models that have smooth density profiles, the gamma distributions fit well the MPDFs if $R\\sqrt{H_0/M_L}$ is sufficiently large. Furthermore, we have found that this result holds in the lens model that has the same degree of density concentration at the center as the Navarro-Frenk-White (NFW) lens model. In contrast, the gamma distributions do not fit well the MPDFs for any value of $R/\\sqrt{M_L}$ in the case of the singular isothermal sphere (SIS) lens model. These results suggest that we might be able to distinguish the NFW lens model from the SIS lens model on the basis of MPDFs.  We have shown that MPDF is well fitted using the gamma distribution even if the value of $R/\\sqrt{M_L}$ is uniformly distributed within $6\\leq R/\\sqrt{M_L}\\leq 10$ in the case of lens model (b). This result suggests that MPDF is universally well fitted using the gamma distribution if lens models have sufficiently gradual density profiles. To obtain further evidence for this hypothesis, we need to clarify the effect of the distribution of $R/\\sqrt{M_L}$ in more detail. This is one of our future works.    We introduced the power law cusp lens model, and we investigated the dependence of the goodness of the fitting with the gamma distribution on the power $n$, which represents the steepness of the cusp at the center. We have generated 100,000 samples of magnification in the Monte-Carlo simulations. We divided the 100,000 samples into 100 datasets of 1000 samples, and performed the $\\chi^2$-test for each dataset, since the same analysis might be possible for the observational data of SNAP. Then, we found the significant difference in the rate of rejection between $n\\leq 0.7$ and $n\\geq0.8$. This result suggests that we may obtain the information about a slope of the density profiles of the central cusp of dark matter using MPDF of Type Ia supernovae.  A mathematical explanation of the reason why we obtain the gamma distribution is unavailable now. In Refs.\\citen{Wang:2002qc,Bergstrom:1999xh,2007arXiv0711.0743L}, and \\citen{Wang:2004ax}, analytic fitting functions different from ours for MPDF are proposed. The relation between our analytic fitting function and those should be studied. We note that we need to examine the effect of simplifications carried out in this study, before we compare our results with observations. For example, we have assumed that there are no spatial correlations of the distribution of the clumps. The effect of the spatial correlation of the clumps should be clarified. We leave these issues as our future works.  The density profile of dark matter halos is a hot topic in astrophysics both theoretically and observationally (See e.g. Ref.\\citen{BinneyTremaine}). Our results suggest that, although there are some simplifications in comparison with realistic situations of our universe, the statistical gravitational lensing effect may shed a new light on the observational investigation of the density profile of dark matter halos."
    },
    "0711/0711.0725_arXiv.txt": {
        "abstract": "% We present Suzaku X-ray observations of the recurrent nova T CrB in quiescence.  T CrB is the first recurrent nova to be detected in the hard-X-ray band ($E \\sim 40.0$ keV) during quiescence. The X-ray spectrum is consistent with cooling-flow emission emanating from an optically thin region in the boundary layer of an accretion disk around the white dwarf. The detection of strong stochastic flux variations in the light curve supports the interpretation of the hard X-ray emission as emanating from a boundary layer. ",
        "introduction": "}  Symbiotic stars are binary systems in which a compact object accretes from red giant, the wind  of which forms a dense nebula around the binary.  The nebula is ionized by the hot compact source, which is usually a white dwarf \\citep{kenyon}.  Among the known symbiotic stars \\citep{belczinsky00}, four objects have shown recurrent nova (RN) eruptions: T CrB, RS Oph, V3890 Sgr and V745 Sco. These eruptions are triggered by a thermonuclear runaway on the white dwarf (WD) surface after accretion of a critical amount of H rich material from the companion.  \\citet{orio2001} suggested that during quiescence, the X-ray emission from RN could originate in the boundary layer of an accretion disc that was reconstructed around the WD after the outburst. Therefore, X-ray observations should allow one to determine a lower limit for the accretion rate. The low X-ray luminosity measured in 350 novae and recurrent novae in quiescence \\citep{orio2001} imply a low accretion rate, inconsistent with the observed recurrence times. For example, in the case of RS Oph \\citep{orio1993}, the short recurrence time of $\\sim$20 years requires an accretion rate of $\\sim$10$^{-8}$ M$_{\\odot}$/yr, inconsistent with measured X-ray fluxes. This inconsistency is known as the ``missing boundary layer problem\".  T CrB is a well known recurrent nova with recorded outbursts in 1866 and 1946. Its quiescent-state hard X-ray emission (E$_{max}$ $\\sim$ 100 keV) was first revealed during a 9 month $Swift$ survey \\citep{tueller05}. Two additional $Swift$/XRT observations were analyzed by \\citet{jamie}, who reported highly absorbed emission that was well fit with a cooling flow model with $kT_{max}$=23.1 keV.  The absorber consisted of components that both fully covered and partially covered the hard X-ray source.  Here we describe Suzaku X-ray Imaging Spectrometer (XIS) and Hard X-ray Detector (HXD) observations of the recurrent nova T CrB performed on September 6, 2006. ",
        "conclusions": "}  Using the Suzaku X-ray satellite, we observed hard-X-ray emission from the recurrent nova T CrB in quiescence. Assuming that a more detailed timing analysis confirms the lack of periodic variations, the presence of flickering on a timescale of minutes and the cooling-flow type spectrum suggest that the accretion onto the WD proceeds through a disc with an optically thin boundary layer. T CrB is the first recurrent nova to show hard-X-ray emission from a boundary layer during quiescence."
    },
    "0711/0711.0808_arXiv.txt": {
        "abstract": "We describe the Wise Observatory Optical Transient Search (WOOTS), a survey for supernovae (SNe) and other variable and transient objects in the fields of redshift 0.06--0.2 Abell galaxy clusters. We present the survey design and data-analysis procedures, and our object detection and follow-up strategies. We have obtained follow-up spectroscopy for all viable SN candidates, and present the resulting SN sample here. Out of the 12 SNe we have discovered, seven are associated with our target clusters while five are foreground or background field events. All but one of the SNe (a foreground field event) are Type Ia SNe. Our non-cluster SN sample is uniquely complete, since all SN candidates have been either spectroscopically confirmed or ruled out. This allows us to estimate that flux-limited surveys similar to WOOTS would be dominated ($\\sim80\\%$) by SNe~Ia. Our spectroscopic follow-up observations also elucidate the difficulty in distinguishing active galactic nuclei from SNe. In separate papers we use the WOOTS sample to derive the SN rate in clusters for this redshift range, and to measure the fraction of intergalactic cluster SNe. We also briefly report here on some quasars and asteroids discovered by WOOTS. ",
        "introduction": "Searches for extragalactic variable or transient astronomical phenomena are almost as old as modern optical astronomy. Well-known examples are the searches for extragalactic variable stars by Edwin Hubble (e.g., Hubble \\& Humason 1931) and his successors, and the pioneering searches for supernovae (SNe) by Fritz Zwicky (e.g., Zwicky, Humason, \\& Gates 1960). For many decades, these searches were mostly targeted at specific nearby (redshift $z < 0.1$) galaxies or clusters. In the late 1980s, Norgaard-Nielsen et al. (1989) undertook a survey for SNe in more distant galaxy clusters, resulting in the discovery of a Type Ia supernova (SN~Ia) in the cluster AC 118 ($z=0.31$). This breakthrough achievement illustrated the potential of such surveys, though the observational costs of this project were considerable. Perlmutter et al. (1990, 1995) used the emerging new technologies of wide-field CCD imaging and image subtraction to carry out the first efficient surveys for high-redshift SNe in clusters and in the field. These methods, employed by large international collaborations (e.g., Perlmutter et al. 1997, Schmidt et al. 1998), have resulted in the discovery of hundreds of high-$z$ SNe~Ia. Analysis of these objects indicates that the expansion of the Universe is currently accelerating (Riess et al. 1998, 2004, 2007; Perlmutter et al. 1999; Tonry et al. 2003; Knop et al. 2003; Astier et al. 2006; Wood-Vasey et al. 2007); see Filippenko (2005b) for a review.  SNe are not the only variable or transient objects that are luminous enough to be observed in distant galaxies. Other types of objects include variable active galactic nuclei (AGNs; e.g., Hawkins \\& Veron 1993; Gal-Yam et al. 2002; Sarajedini et al. 2006) of various sorts, and optical transients associated with cosmological gamma-ray bursts (GRBs; e.g. van Paradijs et al. 1997). We have undertaken a survey for SNe and other transient or variable sources in the fields of galaxy clusters, the Wise Observatory Optical Transient Search (WOOTS; Gal-Yam \\& Maoz 2000a,b; Gal-Yam et al. 2005). Our main motivation is to study SNe in galaxy clusters. In particular, measurement of the cluster SN Ia rate has interesting implications for SN Ia progenitors and the origin of metals in the intracluster medium (ICM; e.g., Gal-Yam, Maoz, \\& Sharon 2002; Maoz \\& Gal-Yam 2004). Another objective is to discover and quantify the occurrence of intergalactic SNe (Gal-Yam et al. 2003a).  Here, we describe the technical design of this survey, the data-analysis methods (\\S~2), and the SNe we discovered (\\S~3.1), along with a brief report on other types of objects (\\S~3.2). In a companion paper (Sharon et al. 2007; Paper II) we derive from our SN sample a measurement of the rate of SNe in galaxy clusters. The discovery of intergalactic SNe by WOOTS was reported by Gal-Yam et al. (2003a). ",
        "conclusions": "We have conducted a survey for SNe and other transients in the fields of rich Abell galaxy clusters at $0.06 \\le z \\le 0.2$. Using unfiltered imaging with the Wise $1~$m telescope, we have achieved a sensitivity limit of $R \\approx 21.5$ mag for variable sources. An automated pipeline was written to reduce and analyze the data, using image subtraction to detect variability. A dozen SNe were discovered and spectroscopically confirmed. Seven SNe turned out to have occurred in cluster galaxies, while five are field events. Eleven of the SNe (including all cluster events) are apparently SNe~Ia, typically discovered near maximum light. From the statistics of non-cluster SNe, and relying on the fact that the WOOTS SN sample is complete (i.e., all SN candidates have been either confirmed or ruled out), we can estimate that flux-limited SN surveys with similar limiting magnitudes will be dominated ($\\sim80\\%$) by SNe Ia.  Our follow-up spectroscopy shows that AGNs may mimic SNe, and that reliably rejecting all AGNs without spectroscopy is difficult. Two of our cluster SNe had no visible hosts in the WOOTS images, nor in deep Keck images, and were subsequently shown to be the first good candidates for intergalactic SNe, whose progenitor stars are not associated with galaxies (Gal-Yam et al. 2003a). The cluster SN sample is suitable for the calculation of the SN rate in clusters in this redshift range, as reported in Paper II. The survey also detected several quasars behind the cluster sample, as well as numerous asteroids and variable stars."
    },
    "0711/0711.1354.txt": {
        "abstract": "We have conducted a long slit search for low surface brightness Lyman $\\alpha$ emitters at redshift $2.67 < z < 3.75$. A 92 hour long exposure with the ESO VLT FORS2 instrument down to a  $1\\sigma$ surface brightness detection limit of  $8\\times 10^{-20}$ erg cm$^{-2}$ s$^{-1}$ $\\sq\\arcsec^{-1}$ per $\\sq\\arcsec$ aperture, yielded a sample of 27 single line emitters with fluxes of a few $\\times 10^{-18}$ erg s$^{-1}$cm$^{-2}$. We present arguments that most objects are indeed Ly$\\alpha$. The large comoving number density $3\\times10^{-2}$ h$_{70}^3$ Mpc$^{-3}$, the large covering factor $dN/dz \\sim 0.2- 1$, and the often extended Ly$\\alpha$ emission suggest that the emitters be identified with the elusive host population of damped Ly$\\alpha$ systems (DLAS) and high column density Lyman limit systems (LLS). A small inferred star formation rate, perhaps supplanted by cooling radiation, appears to energetically dominate the Ly$\\alpha$ emission, and is consistent with the low metallicity, low dust content, and theoretically inferred low masses of DLAS,  and with the relative lack of success of earlier searches for their optical counterparts. Some of the line profiles show evidence for radiative transfer in galactic outflows. Stacking surface brightness profiles  we find emission out to at least 4\\arcsec . The centrally concentrated emission of most objects appears to  light up the outskirts of the emitters (where LLS arise) down to a column density where the conversion from UV to Ly$\\alpha$ photon becomes inefficient. DLAS, high column density LLS, and the emitter population discovered in this survey  appear to be different observational manifestations  of the same low-mass, protogalactic building blocks of present day $L*$ galaxies. ",
        "introduction": "Perhaps surprisingly, modern astronomy has found ways to study the material structures of the high redshift universe over their entire vast range of densities and sizes, from the underdense large voids, via bright starforming galaxies to the most luminous QSOs.  The dark stages of this sequence, with densities from the universal mean to virial density, too low to produce detectable radiation, have been studied mainly with QSO absorption lines. This approach has taught us where to find  most of the baryons (in the ionized  Lyman $\\alpha$ forest gas; e.g., Rauch et al 1997a) and most of the neutral gas (in damped Lyman $\\alpha$ systems; e.g., Wolfe et al 1995) at high redshift. From the bright end of the cosmic matter distribution, serious inroads into the population of high redshift galaxies have been made by color-selecting stellar continuum emitters, exploiting the Lyman limit continuum decrement (e.g., Steidel \\& Hamilton 1993), or by searching for Ly$\\alpha$ line emission induced by star-formation (e.g., Cowie \\& Hu 1998). The advent of space-based, broad-band galaxy surveys, in particular the Hubble Ultra Deep Field (Bunker et al 2004; Beckwith et al 2006; Bouwens et al 2007) has begun filling in the gap between the essentially dark, lower mass range of sub-galactic or barely star-forming proto-galactic objects probed by absorption lines, and bright high redshift galaxies with large star-formation rates of tens of $M_{\\odot}$yr$^{-1}$ (e.g., Erb et al 2006).  The boundary between dark and bright is delineated by an important astrophysical transition, the one from ionized gas to neutral, self-shielded gas, on a galactic mass scale. The underlying objects in which this transition has happened, known from absorption studies as Lyman limit systems (LLS), or, at higher column densities, damped Ly$\\alpha$ systems (DLAS),  have been rather elusive.  We know quite a bit about their chemistry and ionization state but little in terms of stellar contents, size,  kinematics  or  mass.  Theoretical studies over the past two decades have suggested that a survey of LLS and DLAS for HI Ly$\\alpha$ line emission of sufficient depth may uncover several distinct astrophysical sources of line emission that have the potential to shed considerable light (literally) on the distribution of neutral hydrogen, and thus the bedrock of galaxy formation. If we acknowledge that high redshift DLAS must have something to do with stars (they contain the main reservoir of neutral gas, and have a median metallicity of [Z/H] = -1.5, more than an order of magnitude higher than the coeval abundances in the intergalactic medium) then we may expect the potentially strongest signal to be star formation induced Ly$\\alpha$ and/or a stellar continuum. Attempts at identifying individual  LLS or DLAS with galaxy counterparts have often been frustrated by the difficulty of detecting an extremely faint object (the DLAS host) next to an extremely bright object (a QSO). At low redshift ($z<1$) where we can hope to learn most about the underlying galaxy population, observations have shown that DLA host galaxies represent a range of galaxy types (e.g., LeBrun et al 1997, Chen, Kennicutt \\& Rauch 2005) dominated by faint objects (e.g., Chun et al 2006) with  low star formation rates (Wild, Hewett \\& Pettini 2007). Searches at high redshift (e.g., Warren et al 2001,  Fynbo et al 2003; Kulkarni et al 2000, 2001;  Christensen et al 2007) have so far produced only a handful of confirmed detections of the underlying galaxies (Weatherley et al 2005). Such efforts indicated that DLAS hosts at high redshift are generally drawn from the very faint end (Fynbo, M\\o ller, \\& Warren 1999; Bunker et al 1999)  of the general galaxy population at high redshift, intersected by the QSO line of sight at  small impact parameters (M\\o ller et al 2002). Wolfe \\& Chen (2007) have recently performed a search for spatially extended continuum emission down to very faint levels using the Hubble Ultra Deep Field, and were able to place stringent upper limits on extended star formation in DLAS.  These results agree well with theoretical CDM-based models of galaxy formation (Kauffmann 1996) that envisage DLAS hosts as numerous small, faint, low mass, merging protogalactic clumps (Haehnelt, Steinmetz \\& Rauch 1998, 2000; Ledoux et al 1998;  Johansson \\& Efstathiou 2006; Nagamine et al 2007), rather than the large hypothetical disks (Prochaska \\& Wolfe 1997) once popular.  \\medskip   Cooling radiation from collapsing galaxies (e.g., Haiman \\& Rees 2001; Haiman, Spaans \\& Quataert 2000; Fardal et al 2001, Furlanetto et al 2005) is a second  process able to produce line radiation at fluxes competitive with those from star formation, at least in the halos of relatively massive galaxies (Dijkstra et al 2006a,b). The peculiar spatial distribution of the radiation and certain asymmetries of the spectral line profile, together with the local absence of a stellar continuum, could conceivably help to distinguish this source of radiation from star formation. The few  candidates for such galactic cooling flows reported so far (Keel et al 1999; Steidel et al 2000; Francis et al 2001; Bower et al 2004; Dey et al 2005; Matsuda et al 2006; Nilsson et al 2006; Smith \\& Jarvis 2007), have been atypical for the galaxy population as a whole.    Perhaps the most curious of these emission processes is  Ly$\\alpha$ fluorescence, where HI ionizing photons impinging on optical thick gas are absorbed and converted with high efficiency to Ly$\\alpha$ line radiation, albeit at a very low surface brightness level $ < 10^{-19}$erg cm$^{-2}$s$^{-1}\\sq\\arcsec^{-1}$; Hogan \\& Weymann 1987; Binette et al 1993; Gould \\& Weinberg 1996; Cantalupo, Lily \\& Porciani 2005).  Lyman limit patches should light up the whole gaseous cosmic web to yield a relatively uniform, very faint \"glow\",  perhaps locally enhanced by the proximity of a QSO (e.g., Cantalupo et al 2005, 2007). A number of increasingly deep searches for this effect at high redshift have been performed (Lowenthal et al 1990; Bunker et al 1998, 1999), in pursuit of a  repeatedly shrinking  theoretically predicted signal. A detection, which at the same time would be a measurement of the general UV background, has so far eluded us. Most recently, a number of detections of the enhanced Ly$\\alpha$ fluorescence in the proximity of QSOs have been reported (Fynbo, M\\o ller, \\& Warren 1999; Bunker et al 2003;  Weidinger et al 2004, 2005; Adelberger et al 2006, Cantalupo et al 2007, Hennawi 2007) but the understanding of the effect has been complicated by the large number of degrees of freedom in the properties and behavior of the QSO (e.g., orientation, opening angle of the beam, life time).   Thus there is good reason to conduct a search for extremely low light level Ly$\\alpha$ line emission at high redshift. While most previous Ly$\\alpha$ surveys have been performed as narrow band imaging searches it is clear that the largest contrast with the sky background and the largest redshift (=spatial) depth is obtained by (long) slit spectroscopy. Because of the small volume covered a blind search is not a viable way to find  objects as rare as Lyman break galaxies, which are hard to hit with a single randomly positioned slit. In contrast, the rate of incidence of  LLS with neutral hydrogen column densities exceeding (N(HI) $> 10^{19}$cm$^{-2}$) per unit redshift is approximately unity at redshift $\\sim 3$ (e.g., Peroux et al 2003), i.e., the objects  essentially cover the sky, and there should be numerous hits in a single setting for a typical long slit spectrograph.  We anticipate  seeing about 30 patches of emission on a 2\\arcsec\\ wide and 7\\arcmin\\ long slit at redshift $\\sim 3$ (e.g., Bunker et al 1998). Searching over a  redshift range near $z\\sim 3$ achieves a reasonable comprise between avoiding the ravages of $(1+z)^4$ dimming and succumbing to the relatively poor blue performance of most low resolution spectrographs worldwide. However, in principle it is desirable to go as far to the blue as possible. The UV background as the source of the Ly$\\alpha$ fluorescence is not expected to drop dramatically down to at least redshift 2,  whereas the atmospheric background as the dominant source of noise decreases rapidly toward the blue.   Prior to the current ESO VLT project the most sensitive search for Ly$\\alpha$ fluorescence off optically thick gas, in regions dominated by the general UV background, had been performed by the longslit experiment of Bunker et al (1998, 1999). The first attempt at measuring this effect with the ESO VLT originated in a (shelved)  science verification project for FORS, later submitted unsuccessfully during periods 64-66 as a large project (PI Rauch). The amount of exposure time required even with an 8m class telescope considerably exceeds typical observing time awards so  the strategy was to exploit the newly available service mode observing and insert the observations during bad seeing periods, as the targets were expected to be extended, with radii of several arcseconds. The current incarnation, a combination of the two projects, re-observing  the original target field of Bunker et al (1999), was awarded 120 hours as  ESO Large Project (LP173.A-0440: PI Haehnelt). A parallel attempt to use the Gemini GMOS instruments on the same field chosen for the ESO project led to time awards with both Gemini telescopes (GN 2004-A-Q-91, GN-2005-B-Q-52, GS-2004-B-Q-61, GS-2005-B-Q-36; PI Bunker) and (GN-2004-B-Q-35,GN-2004-B-Q-35, GS-2004A-Q-78, GS-2004-B-Q-8: PI Rauch). The Gemini effort resulted in  $46\\times 3000s$ exposures, of which 30 where taken with Gemini-North, 16 with Gemini-South. The Gemini data were taken at lower resolution (to take advantage of the most commonly used spectroscopic setup) but cover a larger redshift path than the ESO data. The current paper deals with the ESO dataset and we postpone a full analysis of all data to a future paper. However, the Gemini data in a preliminary reduction have  been consulted here for checking the reality of the emitters found in the ESO dataset (see below).   The ESO FORS2 project  resulted in (after overheads) 92 hours of on-source exposure, finally  reaching a $1\\sigma$ surface brightness detection threshold of $8\\times 10^{-20}$ erg cm$^{-2}$ s$^{-1}$ $\\sq\\arcsec^{-1}$ (measured in a one $\\sq\\arcsec$ aperture). This is remarkably close to our expected sensitivity ($6.6\\times 10^{-20}$ erg cm$^{-2}$ s$^{-1}$ $\\sq\\arcsec^{-1}$ in 120 h).  The experiment is described below. Section \\ref{obssec} details the observational setup, and the sensitivity and spectral resolution limits reached. Section \\ref{obsressec} spells out the individual properties of the emitters found, and gives estimates of the sizes, number densities, flux distributions, and derived rates of incidence. Section \\ref{whataretheysec} discusses the possible identification with low redshift galaxies, and Section \\ref{lyasec} ponders the consequences if the emitters are predominantly Ly$\\alpha$. Alternative origins of the Ly$\\alpha$ photons (stellar, QSO-induced, cooling radiation) are discussed in section \\ref{origsec}, including a detailed discussion of some unusual line profiles found. Section \\ref{dlassec} establishes a connection between the emitters and QSO DLAS, followed by a brief discussion of Ly$\\alpha$ fluorescence in section \\ref{fluores} and the conclusions in section \\ref{concsec}. An Appendix discusses the importance of slit losses.    % \\object{DMS B 2139-0405}, %% The \\notetoeditor{TEXT} command allows the author to communicate %% information to the copy editor.  This information will appear as a %% footnote on the printed copy for the manuscript style file.  Nothing will %% appear on the printed copy if the preprint or %% preprint2 style files are used.  %% The eqnarray environment produces multi-line display math. The end of %% each line is marked with a \\\\. Lines will be numbered unless the \\\\ %% is preceded by a \\nonumber command. %% Alignment points are marked by ampersands (&). There should be two %% ampersands (&) per line. ",
        "conclusions": "\\label{concsec}  Our longslit search for Ly$\\alpha$ fluorescence from the intergalactic medium, taking advantage of moderate seeing periods with FORS2 at the VLT, has yielded a sample of 27 faint emitters with line fluxes of a few $\\times 10^{-18}$ erg s$^{-1}$cm$^{-2}$ over a redshift range $2.66 < z< 3.75$.  At least a third of the sample shows emission line profiles or an association with absorption systems in the nearby QSO, strongly suggesting identification with Ly$\\alpha$. Spectroscopic features and the absence of detected continua down to $3-\\sigma$ flux limits of $\\sim 1.5 \\times 10^{-19}$ erg s$^{-1}$cm$^{-2}$ make a direct identification of the other emitters (as HI Ly$\\alpha$, [OII] doublet, or [OIII]/HI Balmer emission lines) difficult, but comparison with known galaxy populations and other statistical arguments indicate that the  majority of emitters is likely to be Ly$\\alpha$ at mean redshift 3.2.  If this identification is correct, the emitters present a steeply rising luminosity function with a total number density more than 20 times larger than the comoving density of Lyman break galaxies ($M_R < 25.5$) at comparable redshifts. About half of the profiles are extended, possibly owing to radiative transfer of Ly$\\alpha$ photons from a central source, and there are candidates for both outflows and infall features. We have investigated several mechanisms for the Ly$\\alpha$ production and find star formation to be the energetically most viable process, with a few objects being candidates for cooling radiation.  The inferred low star formation rates, large line emission cross-sections, high number density, and a fitting total cross-section per unit redshift on the sky seem to provide an excellent match to the low luminosities, low metallicities, low dust content, and rate of incidence of damped Ly$\\alpha$ systems, the main reservoir of neutral gas at high redshift. This suggests that our objects are the long-sought counterparts of DLAS in emission. The properties of the objects paint the DLAS host galaxies as a  population of low mass proto-galactic clumps as suggested by some of us (Haehnelt et al 1998, 2000; Rauch et al 1997b), disfavoring a model dominated by large disk galaxies (Prochaska \\& Wolfe 1997). Recent, apparently contradictory limits on spatially extended star-formation from DLAS (Wolfe \\&  Chen 2007) and on the heating of DLAS (Wolfe, Gawiser \\& Prochaska 2003) are consistent with the amount of star-formation we measure if it is confined to a small unresolved region, irrespective of the fact that both the absorption and the Ly$\\alpha$ emission cross-sections appear much larger, the latter because of the random walk of photons to the edge of optically thick gas. The physical origin and the nature of the extended optically thick gas is uncertain, but it is intriguing that quite a few objects show Ly$\\alpha$ emission line profiles consistent with galactic outflows. Thus the absorption cross section of DLAS could be enhanced by  winds, a possibility raised by Nulsen, Barcons \\& Fabian (1998) and discussed in the context of Lyman break galaxies by Schaye (2001).  Finally, adding up the surface brightness profiles of all objects in the outer 2-6\\arcsec\\ we detect radiation at the  $> 2\\times 10^{-19}$ erg cm$^{-2}$ s$^{-1}$\\sq\\arcsec$^{-1}$ level out to 4\\arcsec . The light level is higher by a factor $2 - 4$ than the Ly$\\alpha$ fluorescence signal expected for a UV background intensity consistent with current estimates for the photoionization rate of the intergalactic medium. The large sizes can be explained if radiative transfer of Ly$\\alpha$ lights up the outskirts of our objects out to a radius where the conversion from UV to Ly$\\alpha$ starts becoming inefficient. A radius of 4\\arcsec\\ combined with the observed comoving density of our sample can explain the rate of incidence of DLAS and LLS with column densities as low as $3\\times10^{18}$cm$^{-2}$, consistent with the possibility  that many  LLS  arise in the outskirts of DLAS which in turn are to be identified with the faint emitters. With our interpretation the gas in DLAS is the reservoir from which typical $L_*$ galaxies must have formed in a CDM based universe.  Some of our conclusions here are speculative. Further study is highly desirable, but in any case it should be obvious that long spectroscopic exposures are a promising way of discovering low-mass galaxies. Performing single field, blank sky searches for Ly$\\alpha$ emitters for longish amounts of observing time (but not longer than regularly used by  radio astronomers and now even in space based UV, optical and X-ray astronomy) can bring a whole new range of astrophysical phenomena within the range of existing ground based, optical telescopes, and obviously, could provide one of the most exciting science projects for a future generation of ultra-large telescopes. \\medskip"
    },
    "0711/0711.4194_arXiv.txt": {
        "abstract": "The Palomar Testbed Interferometer (PTI) archive of observations between 1998 and 2005 is examined for objects appropriate for calibration of optical long-baseline interferometer observations - stars that are predictably point-like and single. Approximately 1,400 nights of data on 1,800 objects were examined for this investigation. We compare those observations to an intensively studied object that is a suitable calibrator, HD217014, and statistically compare each candidate calibrator to that object by computing both a Mahalanobis distance and a Principal Component Analysis.  Our hypothesis is that the frequency distribution of visibility data associated with calibrator stars differs from non-calibrator stars such as binary stars. Spectroscopic binaries resolved by PTI, objects known to be unsuitable for calibrator use, are similarly tested to establish detection limits of this approach. From this investigation, we find more than 350 observed stars suitable for use as calibrators (with an additional $\\approx 140$ being rejected), corresponding to $\\gtrsim 95\\%$ sky coverage for PTI. This approach is noteworthy in that it rigorously establishes calibration sources through a traceable, empirical methodology, leveraging the predictions of spectral energy distribution modeling but also verifying it with the rich body of PTI's on-sky observations. ",
        "introduction": "Visible and near-infrared interferometers are powerful tools for measuring the minute angular sizes of nearby stars. However, establishing absolute system responses in the presence of atmospheric turbulence and instrument imperfections is a challenging proposition that requires careful attention to detail when constructing an observational approach.  Use of objects predicted to be point-like calibration sources is routinely employed in astronomical interferometry in the optical \\citep{moz91,1998SPIE.3350..872B,2002A&A...393..183B, van05}.  However, given the incomplete nature of our knowledge of these sources, it is not enough to merely predict the expected fringe visibility from these objects - calibration sources need to be rigorously evaluated for their actual observed interferometric visibility, and its appropriateness for use in calibrating the instrument. Herein we examine the body of calibration data taken upon the sky from near-infrared, long-baseline interferometric measurements taken with the Palomar Testbed Interferometer (PTI). Similar efforts have begun for the Very Large Telescope Interferometer (VLTI) \\citep{2005A&A...434.1201R}.  PTI is an 85 to 110 m H- and K-band (1.6 $\\mu$m and 2.2 $\\mu$m) interferometer located at Palomar Observatory and is described in detail in Colavita et al. (1999).  It has three 40-cm apertures used in pairwise combination for measurement of stellar fringe visibility on sources that range in angular size from 0.05 to 5.0 milliarcseconds, being able to resolve individual sources $\\theta \\gtrsim 1.0$ mas in size. PTI has been in nightly operation since 1997, with minimum downtime for throughout the intervening years. The data from PTI considered herein covers the range from the beginning of 1998 (when the standardized data pipeline went into place) until the end of 2005 (when the analysis of this manuscript was begun).  Over the 8 years of operation discussed in this study, PTI was on the sky in its single-star visibility mode 1,390 nights\\footnote{The remaining nights were spent on other instrument modes or were lost to maintenance or weather.}, producing over 83,000 125 second stellar `scans' on 1,818 individual objects. PTI has a minimum K-band fringe spacing of $\\approx 4.3$ mas at the sky position of the calibration objects, making many of these stars readily resolvable.  As is standard practice in optical interferometry, a typical set of observations with PTI involves observation of target objects of scientific interest, bracketed by calibration objects.  The calibration objects serve the purpose of characterizing the point-response, or system visibility ($V^2_{sys}$),  of the interferometer in conjunction with the atmosphere, and as such, the interleaved calibrator-target-calibrator observations are done on timescales expected to be shorter than the seeing evolution time of the atmosphere (typically less than 15 minutes). By measuring $V^2_{sys}$, a properly normalized measurement of the target star's visibility may be obtained for meaningful astrophysical interpretation \\citep{van05}. These calibration objects are selected {\\it a priori} to be as close to point-like as is possible, having visibilities measured by PTI to be very nearly unity. The primary limitation on selection of truly point-like calibrators is the sensitivity of the instrument, coupled with considerations of dynamic range. Towards the end of having nearly point-like calibrators, calibrator lists are vetted for binary systems, which are excluded due to the sub-unity system response measured from such systems.  The decrease of the measured visibility from the desired characterization of $V^2_{sys}$ is due to the interferometer's ability to resolve out even milli-arcsecond separation binaries.  Unfortunately, incomplete knowledge of the true physical nature of the calibration objects can result in binary systems not being properly excluded during the vetting process that selects potential calibrators.  Given this reality of calibrator selection, a number of additional steps are taken to ensure rigorously defensible characterization of $V^2_{sys}$. First, given the uncertain nature of possible calibrators, employing multiple calibrators for any given target star observation is also standard operating procedure for PTI observations.  Second, these clusters of calibrators are observed on multiple nights, since chance geometry may make a binary system readily apparent on one night but not on another.  Finally, these objects may be compared to known `good' calibrators for appropriateness as calibration objects, which this study will examine in detail. ",
        "conclusions": ""
    },
    "0711/0711.2317_arXiv.txt": {
        "abstract": "{The distribution of TeV spectral slopes versus redshift for currently known TeV blazars (16 sources with  $z$$\\leq$0.21, and one with $z$$>$0.25) is essentially a scatter plot with hardly any hint of a global trend. We suggest that this is the outcome of two combined effects of intergalactic $\\gamma\\gamma$ absorption, plus an inherent feature of the SSC (synchro-self-Compton) process of blazar emission. First, flux dimming introduces a bias that favors detection of progressively more flaring sources at higher redshifts. According to mainstream SSC models, more flaring source states imply sources with flatter TeV slopes. This results in a structured relation between intrinsic TeV slope and redshift. The second effect, spectral steepening by intergalactic absorption, affects sources progressively with distance and effectively wipes out the intrinsic slope--redshift correlation. ",
        "introduction": "With the increasing number of very high energy (VHE) $\\gamma$-ray detection of blazars, a broad correlation between observed slope and redshift was generally expected, in the sense that higher redshift sources would tend to show systematically steeper spectra.  The reason for this is absorption of the emitted VHE emission by the intergalactic background light (IBL), i.e. the integrated optical/IR emission from the evolving stellar populations in galaxies (e.g., Hauser \\& Dwek 2001). In fact, pair creation due to the interaction of a hard photon (with energy $E$) with a soft photon (with energy $\\epsilon$) is expected provided that $\\epsilon E$$>$$m_ec^2$. The cross section of the $\\gamma\\gamma$$\\rightarrow$$e^\\pm$ interaction, $\\sigma_{\\gamma \\gamma}(E,\\epsilon)$ (Heitler 1960) is maximized when $\\epsilon$$\\sim$$2(m_ec^2)^2/E$. Then, if $E$$\\sim$1 TeV, pair creation will be most likely for $\\epsilon$$\\sim$$0.5\\,(E/{\\rm TeV})^{-1}$ eV, which corresponds to a 2$\\,\\mu$m ($K$-band) photon. The formalism required to evaluate this effect has been largely developed by Stecker and collaborators (e.g., Stecker 1971; Stecker et al. 1992 and 2006).  \\begin{figure} \\vspace{6.0cm} \\special{psfile=slope_z.ps hoffset=20 voffset=-65 hscale=40.0 vscale=40.0 angle=0} \\caption{Observed TeV slope vs redshift for known TeV blazars. Bars denote statistical uncertainties as quoted in the original papers. } \\end{figure}  However, when we plot the currently known observed TeV slopes against redshift (Fig.1), we see a scatter plot with hardly any sign of a global correlation: neglecting slope uncertainties, a global linear correlation is suggested only marginally ($\\sim$2$\\sigma$) by the whole data sample ($r$$=$0.52, N=18). Even so, most of the suggestion's strength comes from only the two points at $z$$>$0: for the bulk of the distribution, which provides a fair sampling of the redshift interval $z$$\\leq$0.2, there is no correlation ($r$$=$0.22, N=16). In this note we propose an interpretation of Fig.1.  \\begin{table*} \\caption[] {TeV blazar data.} \\begin{flushleft} \\begin{tabular}{ l l l l l l l l l l l } \\hline \\hline \\noalign{\\smallskip} Source  & ~~~~~~$z$  & ~~~~~~~~$\\alpha_\\gamma$ & ~~~~~~~~F$_\\gamma$ & CT & ~~~$\\alpha_\\gamma^{\\rm corr}$ & ~~~~~~ F$_\\gamma^{\\rm corr}$ &  F$_{5\\,{\\rm GHz}}$ & $~~~~~~$F$_{\\rm x}$ & ~~F$_{x,0}$  & Notes \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} (1) & ~~~~ (2)  & ~~~~~~~ (3) & ~~~~~~~~(4) & (5)& ~~~~(6) & ~~~~~~~~(7) & ~~(8) & ~~~~~~(9) & ~~(10) & ~(11) \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} Mrk~421       &  ~~~0.031 & $2.33 \\pm 0.08$ & 1.03($\\pm$0.03)E-10 & M & ~~~2.15 & ~~~7.59E-11 & 0.730 & ~~~4.0E-10 & 4.1E-11 & Al+07a \\\\ Mrk~501       &  ~~~0.034 & $2.28 \\pm 0.05$ & 1.71($\\pm$0.11)E-11 & M & ~~~2.08 & ~~~1.32E-11 & 1.400 & ~~~4.0E-10 & 3.0E-11 & Al+07b \\\\ &           & $2.45 \\pm 0.07$ & 3.84($\\pm$1.00)E-12 & M & ~~~2.25 & ~~~4.73E-12 &       & ~~~1.7E-10 &         &        \\\\ PKS~2344+514  &  ~~~0.044 & $2.95 \\pm 0.12$ & 1.21($\\pm$0.10)E-11 & M & ~~~2.67 & ~~~1.67E-11 & 0.231 & $<$7.0E-11 & 1.7E-11 & Al+07c \\\\ Mrk~180       &  ~~~0.045 & $3.30 \\pm 0.70$ & 8.46($\\pm$3.38)E-12 & M & ~~~3.01 & ~~~1.23E-11 & 0.274 & $<$1.0E-10 & 5.0E-12 & Al+06a \\\\ 1ES~1959+650  &  ~~~0.047 & $2.72 \\pm 0.14$ & 3.04($\\pm$0.35)E-11 & M & ~~~2.41 & ~~~4.63E-11 & 0.253 & ~~~1.4E-10 & 1.4E-11 & Al+06b \\\\ BL~Lacertae   &  ~~~0.069 & $3.60 \\pm 0.50$ & 3.28($\\pm$0.26)E-12 & M & ~~~3.12 & ~~~7.15E-12 & 2.490 & ~~~~~~ --- & 5.8E-13 & Al+07d \\\\ PKS~2005-489  &  ~~~0.071 & $4.00 \\pm 0.40$ & 3.32($\\pm$0.48)E-12 & H & ~~~3.50 & ~~~7.47E-12 & 1.190 & ~~~4.0E-11 & 0.9E-11 & Ah+05a \\\\ &           &                 &                     &   &         &             &       &            &         & P+99 \\\\ RGB~J0152+017 &  ~~~0.080 & $2.95 \\pm 0.36$ & 4.43($\\pm$1.24)E-12 & H & ~~~2.38 & ~~~1.15E-11 & 0.050 & ~~~2.7E-12 & 9.1E-13 & Ah+08 \\\\ &           &                 &                     &   &         &             &       &            &         & D+01 \\\\ PKS~2155-304  &  ~~~0.116 & $3.37 \\pm 0.07$ & 2.89($\\pm$0.18)E-11 & H & ~~~2.51 & ~~~1.36E-10 & 0.310 & ~~~2.7E-11 & 2.0E-11 & Ah+05b \\\\ 1ES~1426+428  &  ~~~0.129 & $3.55 \\pm 0.46$ & 2.53($\\pm$0.43)E-11 & W & ~~~2.54 & ~~~1.47E-10 & 0.038 & ~~~2.0E-11 & 1.6E-11 & Hor+02 \\\\ &           &                 &                     &   &         &             &       &            &         & S+97 \\\\ 1ES~0229+200  &  ~~~0.139 & $2.50 \\pm 0.19$ & 4.46($\\pm$0.71)E-12 & H & ~~~1.46 & ~~~3.04E-11 & 0.046 & ~~~1.5E-11 & 8.8E-12 & Ah+07c \\\\ &           &                 &                     &   &         &             &       &            &         &P+96 \\\\ &           &                 &                     &   &         &             &       &            &         & D+05 \\\\ H~2356-309    &  ~~~0.165 & $3.09 \\pm 0.24$ & 2.55($\\pm$0.68)E-12 & H & ~~~1.84 & ~~~2.66E-11 & 0.065 & ~~~1.0E-11 & 2.4E-11 & Ah+06a \\\\ &           &                 &                     &   &         &             &       &            &         & W+84 \\\\ 1ES~1218+304  &  ~~~0.182 & $3.00 \\pm 0.40$ & 1.01($\\pm$0.26)E-11 & M & ~~~1.61 & ~~~1.39E-10 & 0.056 & ~~~~~~ --- & 1.5E-11 & Al+06c \\\\ 1ES~1101-232  &  ~~~0.186 & $2.94 \\pm 0.20$ & 4.35($\\pm$0.69)E-12 & H & ~~~1.46 & ~~~6.08E-11 & 0.066 & ~~~5.1E-11 & 2.5E-11 & Ah+07a \\\\ 1ES~0347-121  &  ~~~0.188 & $3.10 \\pm 0.23$ & 3.86($\\pm$0.73)E-12 & H & ~~~1.67 & ~~~5.77E-11 & 0.008 &~~~2.8E-11  & 6.0E-12 & Ah+07b \\\\ 1ES~1011+496  &  ~~~0.212 & $4.00 \\pm 0.50$ & 6.40($\\pm$0.32)E-12 & M & ~~~2.37 & ~~~1.43E-09 & 0.636 & ~~~~~~ --- & ~~~ --- & Al+07e \\\\ PG~1553+113   & $>$0.25   & $4.20 \\pm 0.30$ & 5.24($\\pm$0.87)E-12 & M & $<$2.28 & $>$2.18E-10 & 0.286 & ~~~5.0E-11 & 1.4E-11 & Al+07f \\\\ \\noalign{\\smallskip} \\hline \\hline \\end{tabular} \\end{flushleft} \\smallskip  \\noindent {\\it Col.1:} source name.  \\noindent {\\it Col.2:} source redshift.  \\noindent {\\it Col.3:} observed 0.2-2 TeV photon spectral index, and associated statistical uncertainty. The corresponding systematic uncertainties are typically $\\sim$0.1 for H.E.S.S. and $\\sim$0.2 for MAGIC.  \\noindent {\\it Col.4:} observed $>$0.2 TeV flux (in erg cm$^{-2}$ s$^{-1}$), and associated statistical uncertainty (from observed spectral normalization only).  \\noindent {\\it Col.5:} Cherenkov telescope (CT) or array with which the data in col.3 have been collected: symbols stand for H$=$H.E.S.S., M$=$MAGIC, W$=$Whipple.  \\noindent {\\it Col.6:} corrected 0.2-2 TeV photon spectral index.  \\noindent {\\it Col.7:} corrected $>$0.2 TeV flux.  \\noindent {\\it Col.8:} 5 GHz flux density (in Jy): data are from the NED, except for 1ES~0347-121 (Fossati et al. 1998) and H~2356-309 (Costamante \\& Ghisellini 2002).  \\noindent {\\it Col.9:} 2-10 keV flux (in erg cm$^{-2}$ s$^{-1}$): the data come from the papers listed in {\\it Col.11} or from references quoted therein; when necessary, fluxes have been converted from {\\it R}XTE/ASM count rates using the conversion 1 ct/s $=$3.33E$-$10 erg cm$^{-2}$ s$^{-1}$.  \\noindent {\\it Col.10:} minimum 2-10 keV flux (in erg cm$^{-2}$ s$^{-1}$).  \\noindent {\\it Col.11:} Reference for VHE $\\gamma$-ray data, and for the baseline X-ray flux when the latter was not taken from the NED -- Al+07a: Albert et al. 2007a; Al+07b: Albert et al. 2007b; Al+07c: Albert et al. 2007c; Al+06a: Albert et al. 2006a; Al+06b: Albert et al. 2006b; Al+07d: Albert et al. 2007d; Ah+05a: Aharonian et al. 2005a; P+99: Perlman et al. 1999; Ah+08:  Aharonian et al. 2008; D+01: Donato et al. 2001; Ah+05b: Aharonian et al. 2005b; Hor+02: Horan et al. 2002; S+97: Sambruna et al. 1997; Ah+07c: Aharonian et al. 2007c; P+96: Perlman et al. 1996; D+05: Donato et al. 2005; Ah+06a: Aharonian et al. 2006a; W+84: Wood et al. 1984; Al+06c: Albert et al. 2006c; Ah+07a: Aharonian et al. 2007a; Ah+07b: Aharonian et al. 2007b; Al+07e: Albert et al. 2007e; Al+07f: Albert et al. 2007f.   \\end{table*} ",
        "conclusions": ""
    },
    "0711/0711.4916_arXiv.txt": {
        "abstract": " ",
        "introduction": "NGC~6939 is a relatively rich open cluster located in the constellation of Cepheus ($\\alpha=20^{\\mathrm{h}}31^{\\mathrm{m}}30^{\\mathrm{s}}$, $\\delta=+60^{\\circ}39'22''$), $12.3^{\\circ}$ above the Galactic plane. The cluster was  subject of numerous investigations. The first publication on NGC~6939 presenting a photographic catalog of 370 cluster's stars is dated on 1923 (Kustner 1923). Cuffey (1944) estimated  distance to the cluster for 1300 pc. Cannon \\& Lloyd (1969) obtained the age of about $10^{9}$ yr. However, the first deep CCD photometry was obtained just at the beginning of current century. Rosvick \\& Balam (2002) collected BVI photometry for stars located in the central part of the cluster. These authors obtained the age of $(1.6\\pm0.3)\\cdot10^{9}$ yr and found premises of variability of interstellar extinction within the cluster. The reddening  was found to range from $E(B-V)$=0.29 mag to $E(B-V)$=0.41 mag. In the following study Andreuzzi et al. (2004) obtained the  age of $(1.0-1.3)\\cdot10^{9}$ yr, variable reddening $E(B-V)=0.34-0.38$ mag, and distance of 1.8 kpc based on UBVI CCD photometry. More recently, Maciejewski \\& Niedzielski (2007), who based on wide-field BV data, obtained the age of 1.25 Gyr, $E(B-V)=0.38^{+0.18}_{-0.10}$ mag, distance modulus $(m-M)=12.15_{-0.72}^{+0.56}$ mag, and distance of $1.56^{+0.64}_{-0.59}$ kpc. Authors also determined cluster's core radius $r_{\\mathrm{core}}=2.2'$ and limiting radius $r_{\\mathrm{lim}}=15.2'$. The cluster occurred to be larger than previous estimations of the apparent diameter had indicated. Moreover, the evidence of the mass segregation within the cluster was found based on mass function investigation.  The metallicity of NGC~6939 was studied in a few papers (Canterna at al. 1986, Geisler et al. 1991, Thogersen et al. 1993, Fiel et al. 2002) and occurred to be below the solar value with [Fe/H] between $-0.19$ and $-0.10$ depending on authors. The cluster was also searched for variable stars. During 16 nights Robb \\& Cardinal (1998) monitored 215 stars located in the central part of the cluster. As a result variability of 6 stars was detected. Three of them were located in the red clump and revealed long term light-curve changes. Remaining next three variables were classified as eclipsing binary systems.  As the apparent diameter of the cluster was found in Maciejewski \\& Niedzielski (2007) much larger than the value previously accepted of 5' (Dias et al. 2002) we performed a dedicated CCD search for variable stars in this cluster. In this paper we present  results of this  survey. The main goal of this study was to constrain the open cluster parameters better through investigation of  the population of variable stars, members of  the cluster. ",
        "conclusions": "\\begin{figure} \\includegraphics[width=65mm]{ngc6939-fig3.eps} \\caption{Relation between maximum brightness of the cluster's eclipsing systems and their distance from the cluster center. A distance error of 1' -- the precision of cluster's center redetermination (Maciejewski \\& Niedzielski 2007) was adopted. The value of $3 \\sigma$ of typical accuracy obtained in our survey in V band was taken as a magnitude error.} \\label{fig3} \\end{figure}  The results of a CCD photometric survey of the field around NGC~6939 were presented. Twenty two variable stars were detected, only four of them were previously known. Only 4 eclipsing systems (3 detached and 1 contact binary) from 13 variables located within a cluster radius occurred to be likely cluster members. Analysis of absolute magnitude of the newly detected contact system gives the distance to the cluster $D = 1.74 \\pm 0.20$ kpc which value is in agreement with previous studies.  The population of contact binary systems was found to be small in NGC~6939. This observation is in agreement with results obtained for other open cluster of similar age (e.g. Mazur at al. 1995 and references therein). Rich populations of contact binaries occur in old clusters of age of a few Gyr. The relatively rich population of short-period ($P<5$ days) detached systems was found. One cannot exclude that these systems are progenitors of W UMa stars which are formed as a result of angular momentum loss in initially detached systems (Vilhu 1982).  It is worth noting that the relation between maximum brightness of the cluster's eclipsing systems (V12, V14, V17, and V20) and their distance from the cluster center was noticed (Fig.~\\ref{fig3}). The relation $V = (0.19 \\pm 0.06) d + (14.60 \\pm 0.44)$ with the correlation coefficient of 0.92 was obtained as a result of the linear least-square best fit. Maximum brightness of eclipsing systems can be interpreted as a rough approximation of system's total mass in a star cluster. As a result of the cluster internal dynamics more massive (i.e. brighter) systems should occupy the central part of the cluster while the less massive (i.e. fainter) ones -- the outer region. The relation presented in Fig.~\\ref{fig3} seems to confirm mass segregation among cluster's binary systems.   \\textit{Acknowledgements}: We thank Svetlana Boeva for the help in a part of the observations. This paper is a result of PAN/BAN exchange and joint research project \\textit{Spectral and photometric studies of variable stars}. This research is also supported by UMK grant 369-A and Polish Ministry of Science and Higher Education grant 1P02D00730."
    },
    "0711/0711.4850_arXiv.txt": {
        "abstract": "The galaxy cluster {\\rxj} is one of the most X-ray luminous and most massive clusters known. Its extreme mass makes it a prime target for studying issues addressing cluster formation and cosmology. Despite the naive expectation that mass estimation for this cluster should be straightforward (high mass and favorable redshift make it an efficient lens, and in addition it is bright in X-rays and appears to be in a fairly relaxed state), some studies have reported very discrepant mass estimates from X-ray, dynamical and gravitational lensing.  In this paper we present new high-resolution HST/ACS and Chandra X-ray data. The high resolution and sensitivity of ACS enabled us to detect and quantify several new multiply imaged sources, we now use a total of eight for the strong lensing analysis. Combining this information with shape measurements of weak lensing sources in the central regions of the cluster, we derive a high-resolution, absolutely-calibrated mass map. This map provides the best available quantification of the total mass of the central part of the cluster to date. We compare the reconstructed mass with that inferred from the new Chandra X-ray data, and conclude that both mass estimates agree extremely well in the observed region, namely within 400$h_{70}^{-1}$kpc of the cluster center. In addition we study the major baryonic components (gas and stars) and hence derive the dark matter distribution in the center of the cluster. We find that the dark matter and baryons are both centered on the BCG within the uncertainties (alignment is better than $<10\\mbox{ kpc}$). We measure the corresponding 1-D profiles and find that dark matter distribution is consistent with both NFW and cored profiles, indicating that a more extended radial analysis is needed to pinpoint the concentration parameter, and hence the inner slope of the dark matter profile. ",
        "introduction": "\\label{sec:intro} Galaxy clusters have been the focus of very intense research over the past decade. They are extremely valuable for studying the empirical properties of dark matter, as well as for exploring the growth of structure at the high mass tail of the mass function. The mass distribution of galaxy clusters is particularly important for cosmological studies because it provides a critical test of the Cold Dark Matter (CDM) paradigm and constrain dark energy models, if their mass and its distribution is reliably determined.  However, in some cases the cluster mass estimates from the weak lensing, strong lensing and X-ray measurements disagree: the origin of these discrepancies is likely to be a combination of erroneous assumption of the cluster potential, projection effects and complicated gas physics not taken into account. In addition, systematic effects like cluster member contamination and unknown redshift distribution of sources, strong lensing image identification, temperature calibration, etc. enter the error budget for X-ray and lensing analyses.   It was first proposed by \\citet{navarro97} that the dark matter halos on a variety of scales should follow a universal profile within the currently accepted $\\Lambda$CDM paradigm. The 3-D density distribution of dark matter should follow $\\rho_{\\rm DM} \\propto r^{-1}$ within a scale radius $r_{\\rm s}$ and falls of steeper at radii beyond that ($\\rho_{\\rm DM} \\propto r^{-3}$). However recent simulations have shown that the profile might be modified and the structural parameters vary with redshift and mass \\citep{neto07,navarro04}. In addition, likely interaction between dark matter and baryons further complicate the picture \\citep{gnedin04, nagai05}. It is therefore crucial to compare simulations and observations keeping the baryonic component in mind. This will allow us to measure the slopes of dark matter profiles, which is a critical test for our currently accepting cosmology as well as understanding the complicated baryonic physics in galaxy clusters. Several works have previously studied mass distribution in number of clusters using combined strong and weak lensing reconstruction \\citep[see e.g.][]{natarajan96,kneib03,smith04, diego05, cacciato05, jee07, limousin07} and combined strong lensing and stellar kinematics data of the dominating central galaxy \\citep{sand07}. These approaches offer valuable extra constraints for determining the mass distributions.   In this paper we study the most luminous X-ray clusters known to date. The cluster {\\rxj} \\citep{schindler95} at a redshift $z=0.451$ has been a subject of intense past research. Rich datasets including X-ray \\citep{schindler95, schindler97,allen02,ettori04,gitti04, gitti07} and optical \\citep{fischer97,sahu98,cohen02,ravindranath02} as well as observations of Sunyaev-Zel'dovich (SZ) effect \\citep{pointecouteau01,komatsu01,kitayama04} have been obtained, yet the mass determinations based on X-ray properties, SZ effect (SZE), velocity dispersion measurement, strong and weak lensing have many times yielded discrepant results (see \\citealp{cohen02} for a summary). In particular the discrepancy between the dynamical mass estimate \\citep{cohen02}, early X-ray mass measurements \\citep{schindler97} and lensing results \\citep{fischer97} yielded a factor of $\\sim 3$ discrepancy in their mass estimates.   \\citet{cohen02} suggested the cluster is likely undergoing a major merger which would reconcile the low velocity dispersion (and hence virial mass) measurement and its high mass as predicted by gravitational lensing and X-ray data. If the velocity dispersion measures predominately one component, and the X-ray gas has virialized to the post-merger cluster mass, this can possibly explain the data. Further evidence for a merger scenario comes also from the optical morphology: the cluster contains two cD galaxies (see e.g. Fig.~\\ref{fig:arcs} - throughout this paper we will refer to the brighter, western one as the BCG, and the other as the second cD). Finally, there is a region of shocked gas in the south-east quadrant discovered and discussed by \\citet{komatsu01} (with SZ effect) and \\citet{allen02} based on X-ray observations with Chandra. We confirm the latter using the newest Chandra X-ray data.   In order to resolve the puzzle of mass estimates in \\citet{bradac04b} we studied this cluster using ground-based optical data, and for the first time determined its mass distribution using a combined strong (multiply-imaged systems) and weak (statistical measurement using shape information of an ensemble of background galaxies) gravitational lensing analysis. The main drawback of that reconstruction was the lack of clearly identified multiply imaged systems and their redshifts. In the current paper we obtained new multi-color HST/ACS data. The high resolution and sensitivity of ACS allowed us to unambiguously identify many new multiply imaged systems. In this analysis we are able to use 9 multiply imaged systems, in contrast to the single one of \\citet{bradac04b}. Furthermore the far greater density of sources that can be used for weak lensing (a factor of 5 improvement) has allowed us to obtain an absolutely-calibrated mass map at high resolution. We also present new X-ray data and compare estimates from both analyses. In addition we describe a new automated search routine for multiply imaged sources we are currently developing; its main advantage is that it combines colors and information on lensing geometry of the whole system to obtain matching pairs in a semi-automated way.   Our final mass reconstruction is again performed using a pixelated model which combines the information from strong and weak gravitational lensing. If spectroscopic redshifts of the multiply imaged systems were available, this would fully eliminate the need of using models that assume a particular shape of the gravitational potential. However, since only one system has a spectroscopically confirmed redshift, we use the predictive power of gravitational lensing to predict the redshifts for the remaining systems. As explained further in \\S\\ref{sec:slensing}, we therefore model strong lensing data using parametrised models at the start to obtain the redshifts. For the final reconstruction we however do not assume a special shape of the gravitational potential, instead exploring a much larger range of cluster mass distributions and profiles (provided the estimated redshifts are reasonably reconstructed). Furthermore, we use cluster to study the interplay between baryons and dark matter, therefore we study separately its stellar, gas and dark matter component.  This allows us to measure the cluster's dark matter profile, one of the strong predictions of $\\Lambda$CDM cosmology.   This paper is structured as follows. In section~\\ref{sec:data} we describe the optical data used in this analysis. The basic image processing and the extraction of the strong gravitational lensing data is described in section~\\ref{sec:slensing} and in section~\\ref{sec:wlensing} we describe the weak gravitational lensing data. We present the X-ray data in section~\\ref{sec:xraydata}.  We infer the mass distribution of the cluster {\\rxj} from lensing data, and compare it with X-ray data, in section~\\ref{sec:resultsw}. The individual contributions of stars, gas, and dark matter to the total mass of the cluster are studied in section~\\ref{sec:dmbar}. The conclusions are presented in section~\\ref{sec:conclusions}. To evaluate the angular diameter distances throughout the paper we assume the $\\Lambda$CDM cosmology with $\\Omega_{\\rm m} = 0.3$, $\\Omega_{\\Lambda} = 0.7$, and Hubble constant $H_0 = 70 {\\rm \\ km \\: s^{-1}\\:\\mbox{Mpc}^{-1}}$. ",
        "conclusions": " \\begin{enumerate} \\item Using the combined strong and weak lensing mass reconstruction we derive a high-resolution, absolutely calibrated mass map; we get projected, enclosed mass $M_{\\rm L}(<350\\:\\mbox{kpc})= (5.9 \\pm 0.5) \\times 10^{14} M_{\\odot}$. Within the same radius the projected mass derived from X-ray data alone is $M_{\\rm X}(<350\\:\\mbox{kpc})= (6.6 ^{+0.6}_{-0.4}) \\times 10^{14} M_{\\odot}$. \\item The mass estimates are still in disagreement with previous dynamical mass estimates by \\citet{cohen02}. It is however possible that \\citet{cohen02} measured the velocity dispersion of the infalling subcluster only \\citep{allen02}, thereby biasing the mass estimate low. If taken as a measurement of the cluster mass, such a low velocity dispersion cluster is in complete disagreement with strong lensing features we observe (regardless of their redshift). We are currently analysing the remainder of spectroscopy data to resolve this issue (Lombardi et al. in preparation). \\item Using the exquisite resolution of ACS data we are able to study the spatial distribution of dark matter with respect to the baryons at unprecedented accuracy. We clearly detect a mass concentration centered on the BCG and a SW extension, which follows the light distribution of the cluster members. Further we fit a generalised NFW model to the dark matter density profile, finding strong degeneracies between the inner slope $\\beta$ and the scale radius $r_{\\rm s}$. Our data does not extend to sufficiently large radii, for more secure determination of $\\beta$, $r_{\\rm s}$, and concentration parameter; we plan to use weak lensing data extending beyond the virial radius in the future to resolve this issue. \\end{enumerate}  The cluster {\\rxj} is an extremely valuable object for understanding the details of cluster formation and evolution. In addition, due to their magnifying power, these most massive clusters are ideal tools to study the high-redshift universe. By searching around the critical line for high redshift sources as predicted by our lens model we were already able to find high redshift multiple imaged systems. We plan to extend this search in the future.   The large wavelength coverage of the data for this cluster has proved to be extremely valuable for a detailed study of the cluster mass distribution. With such data sets most massive clusters of galaxies remain one of the key objects to constrain cosmological parameters and to study the galaxy formation and evolution from the early times until the present."
    },
    "0711/0711.1158_arXiv.txt": {
        "abstract": "We present very long baseline interferometry polarization images of an X-ray selected sample of BL~Lacertae objects belonging to the first High Energy Astronomy Observatory ({\\it HEAO-1}) and the $ROSAT$--Green Bank (RGB) surveys. These are primarily high-energy-peaked BL~Lacs (HBLs) and exhibit core-jet radio morphologies on pc-scales. They show moderately polarized jet components, similar to those of low-energy-peaked BL~Lacs (LBLs). The fractional polarization in the unresolved cores of the HBLs is, on average, lower than in the LBLs, while the fractional polarizations in the pc-scale jets of HBLs and LBLs are comparable. However a difference is observed in the orientation of the inferred jet magnetic fields -- while LBL jets are well-known to preferentially exhibit transverse magnetic fields, the HBL jets tend to display longitudinal magnetic fields. Although a `spine-sheath' jet velocity structure, along with larger viewing angles for HBLs could produce the observed magnetic field configuration, differences in other properties of LBLs and HBLs, such as their total radio power, cannot be fully reconciled with the different-angle scenario alone. Instead it appears that LBLs and HBLs differ intrinsically, perhaps in the spin rates of their central black holes. ",
        "introduction": "BL~Lacertae objects are radio-loud active galatic nuclei (AGNs) characterized by a predominantly non-thermal, highly polarized continuum, that is variable in total intensity and polarization at all observed wavelengths, and exhibits weak or no optical emission lines \\citep{Strittmatter72,SteinOdell76}. These extreme phenomena are understood to be a consequence of relativistic beaming in their nuclei \\citep{BlandfordRees78,BlandfordKonigl79}. This picture has lead to Unification schemes \\citep{UrryPadovani95} wherein the low-luminosity Fanaroff-Riley type-I radio galaxies \\citep[FRI,][]{FanaroffRiley74} are believed to be the parent population of BL Lac objects, based on similarities in orientation-independent properties such as extended radio emission, emission-line luminosity, host-galaxy type and environment \\citep{Browne83,WardleMooreAngel84,Prestage88,UrryPadovani95}.  Most BL~Lacs discovered before the late 1980's were identified based on radio data \\citep[e.g.,][]{LeddenOdell85,Burbidge87}. However, space-based X-ray surveys were highly efficient in discovering new BL~Lac objects \\citep[e.g.,][]{Giommi89,Stocke89}. The observed spectral energy distributions (SEDs) of the X-ray-selected and radio-selected BL~Lacs are systematically different. The majority of X-ray-selected BL Lacs show a synchrotron emission peak in the UV/soft X-ray regime making them `high-energy peaked BL Lacs' or HBLs, while most radio-selected BL~Lacs show synchrotron emission peaks in the near-IR/optical making them `low-energy peaked BL~Lacs' or LBLs \\citep{GiommiPadovani94,PadovaniGiommi95}. The approximate division between LBLs and HBLs has been defined as log~($f_x/f_r$)~$\\sim-5.5$ \\citep{Perlman94,Perlman96,Wurtz96} where $f_x$ and $f_r$ are the 1keV X-ray and 5~GHz radio flux density, respectively. HBLs are defined as BL~Lacs that have log~($f_x/f_r$)~$\\ge-5.5$.  LBLs are typically more core-dominated on kpc-scales \\citep{PerlmanStocke93,Kollgaard96,RectorStocke00,Giroletti04a}, they display higher average optical polarizations and greater variability \\citep{Schwartz89,Jannuzi94}, lower starlight fraction \\citep{Morris91}, and have more powerful radio lobes than HBLs \\citep{Kollgaard92,LaurentMuehleisen93}. Based on nuclear trends, it appears that the LBLs are more ``extreme'' than HBLs, which in turn has lead to the suggestion that LBLs are oriented at smaller angles to the line of sight than HBLs \\citep{Stocke85}. This scenario is supported by population studies \\citep{UrryPadovani90,UrryPadovani91} and has been explained in the framework of an accelerating-jet model, wherein the X-ray photons are emitted in the slower part of the jet and are therefore less beamed than the radio photons, which are emitted in the faster portion of the jet \\citep{Ghisellini89,UrryPadovaniStickel91}. However, the inadequacy of the orientation scenario to explain the systematic differences in the SED peaks of the HBLs and LBLs has given rise to the alternative scenario wherein the X-ray emitting regions of HBLs have higher electron Lorentz factors and/or magnetic field strengths than LBLs \\citep{Sambruna96}. \\citet{Georganopoulos98} have invoked a model that couples jet orientation and electron kinetic luminosity to unify the BL~Lac subclasses.  The apparent ``blazar sequence'' of increasing synchrotron peak frequency with decreasing source luminosity, as observed in the average SEDs of flat-spectrum radio quasars (FSRQs), LBLs and HBLs \\citep[see][]{Sambruna96,Fossati98,Maraschi01} has been challenged by the discovery of high-energy peaked FSRQs \\citep{Sambruna97,Perlman98,Padovani02} and high luminosity HBLs \\citep{Giommi05}. In addition, the gap between LBLs and HBLs has now been filled with the discovery of intermediate-peak BL~Lacs (IBLs) which have SED peaks in the optical/UV wavebands and exhibit intermediate properties between LBLs and HBLs \\citep{NassBade96,Kock96,Perlman98,LaurentMuehleisen99,Stevens99}. Thus the question of HBLs and LBLs being two distinct AGN classes, as opposed to being the extreme ends of a continuous distribution of BL~Lac properties, is an open one.  Owing to their relatively greater radio flux densities, LBLs have been the subject of extensive VLBI observations \\citep[e.g.,][]{Gabuzda00,Kellermann04,Lister05}. These studies have found them to be typically compact, with a core and a relatively small number of fairly discrete jet components. Measured superluminal speeds are typically modest, in the range of 2--4$c$ \\citep{Gabuzda94,Gabuzda00,Kellermann04}. VLBI polarization observations have shown a number of systematic differences between LBLs and FSRQs, initially detected in images obtained at 5--8~GHz \\citep{Gabuzda92,Cawthorne93,Gabuzda00} but generally being retained in the higher-frequency 15~GHz images \\citep{Lister05} $-$ LBLs have appreciably polarized cores (core fractional polarization, $m_c \\sim 2 - 5\\%$) with predominantly transverse magnetic ($B$) fields in their jets, while quasars have weakly polarized cores ($m_c < 2\\%$) and predominantly longitudinal $B$ fields in their pc-scale jets.  HBL polarization properties have, on the other hand, not been explored as much with VLBI experiments due to their relatively lower radio flux densities. 5~GHz total intensity VLBI images of a sample of HBLs belonging to the $Einstein$ ``Slew\" Survey were presented by \\citet{Rector03}. They found that HBLs displayed core-jet morphologies similar to the LBLs. However there was a weak evidence that the HBL jets exhibited smaller misalignments between parsec and kiloparsec scales than LBL jets \\citep[see also][]{Giroletti04a}. This was interpreted as being either due to intrinsically straighter jets in HBLs or their jets being oriented further from the line of sight than LBLs jets. Note however that 8.4~GHz VLA observations of a sample of EMSS and {\\it HEAO-1} BL~Lacs failed to show a systematic trend of smaller misalignments in HBL jets (E. Perlman, private communication).  In this paper, we present VLBI polarization observations of a sample of 18 BL~Lacs belonging to the {\\it HEAO-1} hard X-ray and the $ROSAT$--Green Bank (RGB) surveys. Thirteen of these BL~lacs have been classified as HBLs. Early VLBI polarization results for four of these {\\it HEAO-1} BL~Lacs were presented by \\citet{Kollgaard96}. \\S~2 and 3 describe the BL~Lac sample and the observations. We present our results in \\S~4--6. A discussion follows in \\S~7, and the conclusions in \\S~8. We assume throughout a cosmology where $H_0$ = 71~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_M$ = 0.24, and $\\Omega_{\\Lambda}$ = 0.76. ",
        "conclusions": "\\hspace*{2em}1.~~We have observed eighteen BL~Lacs belonging to the {\\it HEAO-1} and RGB samples with VLBI, and detected parsec-scale polarization in all but three of them. After dividing each of our objects into the HBL or LBL classes and considering the known results on LBLs in the literature, we have carried out a comparison of the relevant properties of the two BL~Lac classes.  2.~~We find that the total intensity pc-scale images reveal a core-jet morphology, similar to that observed in LBLs.  3.~~The VLBP observations of HBLs and LBLs reveal differences in their core fractional polarization, similar to the trend observed on kpc-scales -- the HBLs have a lower core fractional polarization compared to the LBLs.  High degrees of core polarization were observed for the two RGB sources 0749+540 ($\\simeq 10\\%$) and 0925+504 ($\\simeq 14\\%$), but both of these are, in fact, LBLs.  4.~~The jet fractional polarizations in HBLs do not differ systematically from those in LBLs. In both cases, jet degrees of polarization of tens of per cent can sometimes be observed, indicating the presence of well-ordered $B$ fields.  5.~~The relative VLBI core polarization angles do not show any systematic differences between HBLs and LBLs.  6.~~The most intriguing difference between the HBLs and the LBLs is the orientation of the predominant $B$ fields in their pc-scale jets -- the $B$ fields in HBL jets tend to be aligned with the local jet direction, while the $B$ fields of LBLs tend to be perpendicular to the local jet direction. This result needs to re-examined with a larger sample of HBLs.  7.~~Some of the HBLs display evidence for jet magnetic field structures with transverse $B$ field in the inner region of the jet and longitudinal $B$ field at the edges, i.e, a `spine-sheath' structure. Although such $B$-field structures have been taken to indicate a fast spine+slow sheath velocity structure, they can also come about in a natural, simpler way if the jet has a helical $B$ field (without any requirement for a two-layer velocity structure).  8.~~We find tentative evidence that the observed component speeds in HBLs are lower than the typical apparent speeds  observed in LBLs ($\\simeq 1-5c$) and FSRQs ($\\simeq 5-10c$).  This result suggests that either the HBL jets have lower Lorentz factors than LBLs and FSRQs, or that their jets are oriented at larger angles to the line of sight. Since the collected properties of HBLs provide no clear evidence that their jets lie at systematically larger angles to the line of sight than do LBL jets, this suggests that the outflow speeds are intrinsically lower in HBLs, than in LBLs.  9.~~One way to understand our collected results in the context of other previous results in the literature is if HBLs, LBLs and FSRQs form a sequence of increasing average jet Lorentz factor --- the same order as for the sequence of the synchrotron peaks of these objects, which proceed from highest peak frequency for HBLs to lowest peak frequency for FSRQs, although the physical connection between these two sequences is not clear.  10.~~If the jets of all three classes of compact AGN characteristically have helical $B$ fields due to the ``winding up'' of a seed field via the combination of outflow and rotation of the central black-hole--accretion-disk system, then the dominant jet $B$ fields observed in these three classes of compact AGNs $-$ longitudinal in HBLs and FSRQs and transverse in LBLs $-$ could come about due to the ratios of their characteristic rotational and outflow velocities. If HBLs, LBLs and FSRQs do form a sequence of increasing Lorentz factor, the low outflow speeds in HBLs would require relatively slow rotation ($i.e.,$ smaller spin rates) of their central black-hole--accretion-disk systems."
    },
    "0711/0711.2471_arXiv.txt": {
        "abstract": "{Trails of dust inside galactic bars are easily observable in external galaxies. However, information on the dust lanes of the Milky Way bar is harder to obtain due to our position within the Galactic disc.} {By comparing the distribution of dust and gas in the central regions of the Galaxy, we aim to obtain new insights into the properties of the offset dust lanes leading the bar's major axis in the Milky Way.} {On the one hand, the molecular emission of the dust lanes is extracted from the observed CO $l-b-V$ distribution according to the interpretation of a dynamical model. On the other hand, a three dimensional extinction map of the Galactic central region constructed from near-infrared observations is used as a tracer of the dust itself and clearly reveals dust lanes in its face-on projection. Comparison of the position of both independent detections of the dust lanes is performed in the $(l,b)$ plane. } {These two completely independent methods are used to provide a coherent picture of the dust lanes in the Milky Way bar. In both the gas and dust distributions, the dust lanes are found to be out of the Galactic plane, appearing at negative latitudes for $l>0\\degr$ and at positive latitudes for $l<0\\degr$. However, even though there is substantial overlap between the two components, they are offset from one another with the dust appearing to lie closer to the $b=0\\degr$ plane. } { Two scenarios are proposed to explain the observed offset. The first involves grain destruction by the bar shock and reformation downstream. Due to the decrease in velocity caused by the shock, this occurs at lower $z$. The second assumes that the gas and dust remain on a common tilted plane,  but that the molecular gas decouples from the Milky Way's magnetic field, itself strong enough to resist the shear of the bar's shock. The diffuse gas and dust remain coupled to the field and are carried further downstream. This second scenario has recently been suggested in order to explain observations of the barred galaxy NGC 1097. } ",
        "introduction": "Barred galaxies are ubiquitous in the Universe. Many of these are observed to have dust lanes - trails of gas and dust, seen in extinction, leading the bar's major axis. Evidence for the bar nature of the Milky Way's bulge has been around for over 40 years \\citep{deVaucouleurs1964}. However, our position in our Galaxy does not give us a clear view of the inner regions. Photometric studies are hampered by the high source confusion and severe interstellar extinction whereas gas kinematic studies are rendered difficult due to line of sight crowding of the gas emission lines.  Nevertheless, many studies have obtained information on the Milky Way's dust lanes using a variety of techniques. {\\changes \\cite{Calbet1996} detected a longitude assymetry at $10\\degr<|l|<20\\degr$ in the DIRBE surface brightness maps which they attributed to the dust lanes.} \\cite{Babusiaux2005} used red clump stars to determine the distance as a function of longitude to the bar and using their observations concluded that there is no near-infrared extinction due to the dust lanes in our Galaxy. Other studies have based their interpretation of the observations on models in order to study the chemistry in the Galactic centre (GC) and dust lanes \\citep{Rodriguez-Fernandez2006, Nagai2007} as well as the physical processes underway in the dust lane shocks \\citep{Liszt2006}. However, the results may depend on the assumed geometry for the dust lanes as no coherent picture currently exists for them.  In this letter, we combine the results from gas dynamics and stellar reddening studies in order to provide a description of the dust lanes leading the bar's major axis. \\cite{Fux1999} interpreted the HI and CO longitude-velocity diagrams of the Galactic disc using self-consistent N-body and smooth particle hydrodynamics simulations of the stellar and gaseous components. Later, \\cite{Marshall2006} used a galactic stellar population synthesis model to extract the three dimensional extinction information from near infrared stellar observations. Both studies were able to identify the Milky Way's dust lanes. In the following, we present and analyse the dusty and gassy sky distributions of the dust lanes resulting from these studies. ",
        "conclusions": "\\label{sec:conclu}  We have identified the dominant dust lanes of the Galactic bar both by their near-infrared extinction and by their CO emission. The same large scale behaviour is observed in both components, that is a tilt going from $b<0\\degr$ at positive $l$ to $b>0\\degr$ at negative $l$. In the $l,b$ plane projections, the two components are seen to be displaced from one another. This offset may be due to dust destruction at the shock front followed by reformation and/or ice mantle accretion downstream and at lower $z$, or the molecular gas becoming decoupled from the magnetic field which resists the shearing shock of the Milky Way's bar and carries diffuse gas and dust further downstream on a tilted plane. Further observations of our Galaxy and external galaxies, together with theoretical work are necessary in order to unambiguously interpret these observations.  Finally, the coincidence between the two independently recovered gas and dust distribution along the bar reinforces the interpretation from \\cite{Fux1999} regarding the $l-V$ traces of the dust lanes in the observed HI and CO kinematics, and it adds credibility to the Galactic model and the use of it to determine the three dimensional extinction distribution."
    },
    "0711/0711.2192_arXiv.txt": {
        "abstract": "{ The growth of solid particles towards meter sizes in protoplanetary disks has to circumvent at least two hurdles, namely the rapid loss of material due to radial drift and particle fragmentation due to destructive collisions.  In this paper, we present the results of numerical simulations with more and more realistic physics involved. Step by step, we include various effects, such as particle growth, radial/vertical particle motion and dust particle fragmentation in our simulations. We demonstrate that the initial dust-to-gas ratio is essential for the particles to overcome the radial drift barrier. If this value is increased by a factor of 2 compared with the canonical value for the interstellar medium, km-sized bodies can form in the inner disk ($<2$~AU) within $10^4$~yrs. However, we find that solid particles get destroyed through collisional fragmentation. Only with the unrealistically high-threshold velocities needed for fragmentation to occur ($>30$~m/s), particles are able to grow to larger sizes in disks with low $\\alpha$ values. We also find that less than 5\\% of the small dust grains remain in the disk after 1~Myrs due to radial drift, no matter whether fragmentation is included in the simulations or not. In this paper, we also present considerable improvements to existing algorithms for dust-particle coagulation, which speed up the coagulation scheme by a factor of $\\sim 10^4$.} ",
        "introduction": "The coagulation of sub-$\\mu$m dust particles is believed to be the initial step of planetesimal formation in disks around pre-main-sequence stars \\citep{hbook,natta07}. Evidence for dust grain evolution beyond sizes that are found in the interstellar medium is provided by mid-infrared spectroscopy of disks around Herbig Ae/Be stars \\citep{bouwman01,boekelaebe03}, T~Tauri stars \\citep{prz03,kessler07} and also around brown dwarfs \\citep{apai04,apai05,apai2007,sicilia07}. Millimeter interferometry indicates large populations of particles which have grown to even larger sizes, ranging up to several centimeters \\citep{tes03,Wil05,Rod06}.  All these observations give reason to model the evolution of particles in protostellar disks in order to explain the observational data \\citep{DD04,DD05,tan05,nomura06,allesio06,ormel07}. However, these theoretical investigations do not only attempt to model the evolution of the appearance of protostellar disks. They also unveil certain obstacles in the formation of planetesimals by particle coagulation \\citep{you04,domppv07,brauer07}.  One of these obstacles is the radial inward drift of solid bodies towards the central star as first dicussed by \\cite{Whipple72} and \\cite{Wei77}. The gas in a protostellar disk moves slightly sub-Keplerian due to a radial pressure gradient. For this reason, the dust which moves with near Keplerian velocity feels a continuous head wind of gas. Hence, the particle loses angular momentum due to drag forces between gas and dust and spirals inward.  If this radial drift of the particles is not prevented by some mechanism, then the solid particles drift into the inner evaporation zone and are lost for the process of planetesimal formation. To give an example, the radial drift time scale for meter-sized bodies at 1 AU is $\\sim 10^2$ yrs. Within this time scale these boulders at 1 AU drift into the inner regions of the disk and evaporate. A possible way out of this problem is particle growth since the radial drift velocity is fairly dependent on the particle radius. For example, the drift velocity of meter-sized particles at 1 AU in the disk is $\\sim 50$~m/s, but the radial drift velocity of 10~m sized bodies is already 10 times lower. Therefore, swift particle growth could prevent the particles from drifting into the evaporation zone. However, the general disk evolution comprises a considerable particle loss due to evaporation which is hard to prevent. This problem is a major topic of this paper. Other sublimation zones of the disk, e.g. the snow line at $\\sim 2$~AU \\citep{lecar06}, could also play a role for particle drift and coagulation processes. However, we will for now neglect this issue which will be the topic of a forthcoming paper.  Another obstacle is the fragmentation of solid particles. While low-velocity collisions lead to particle growth, high velocity impacts lead to destruction \\citep{bw98,poppe99,BluWur99}. For example, the relative particle velocity of meter-sized bodies in a protostellar disk can be more than $\\sim 30$~m/s \\citep{Wei77,marmizvol91}. \\cite{Benz00} found that meter-sized rocks appear unlikely to survive an impact with a relative low collision velocity of some cm/s. For porous objects, collision velocities higher than 4\\% of the sound speed lead to particle destruction \\citep{sir04}. For this reason, the particle size of roughly a meter seems to pose an upper limit for particle coagulation.  These two obstacles, the radial drift barrier and the fragmentation barrier, are the issue of this paper. We present a disk model including the growth, the radial drift and the fragmentation of the particles. We show how these three effects change the evolution of the disk by including them step by step. \\begin{enumerate} \\item In the first step we only consider particle coagulation due to Brownian motion, vertical settling and turbulent mixing. This step shows to which sizes particles can grow if radial drift and fragmentation are neglected. \\item The second step includes the radial drift and the radial mixing of dust. The particles are now allowed to move inwards and to disappear into the evaporation zone. However, we investigate which disk parameters influence the drift time scales and for which parameters the dust particles overcome the drift barrier. \\item The last step also includes particle fragmentation. We show under which conditions, i.e. in which regions of the disk and for which disk parameters, it is possible for the dust to overcome this barrier. \\end{enumerate} The radial drift barrier is not only of interest for the radial drift itself. Particles close to this barrier are most susceptible to the motions of the gas and the gravitational effects of the dust. For example, particles can be trapped in very elongated gas vortices in magnetorotational turbulence \\citep{balhaw91,bargesomm95}.  These effects can slow down the radial drift by a factor of two \\citep{JohKlaHen06}. Under certain conditions the solid particle layer itself may become gravitationally unstable \\citep{anders07}. In high dust density regions, the particles contract due to their own gravity and may form a planetesimal within a few orbits \\citep{andersnat07}. Moreover, the flow of the gas and the dust can be unstable to the streaming instablity \\citep{yougoo05} which leads to particle clumping, and possibly also to a gravitational collapse of the dust. Apparently, the radial drift barrier is not only connected to the radial motion of the dust particles, but involves various other important effects as well. For this reason, it is vital to answer the question if particles can actually reach the size regime at which non-linear effects become of importance.  In this paper we will implement a 2+1 dimensional model. The first dimension is the radial coordinate of the disk $r$, the second one is the height above the midplane $z$ and the third coordinate is the mass of the dust particles $m$. The dust may move radially due to radial drift and radial mixing. We will numerically solve the continuity equation for this problem for each particle species. The time evolution for the particle size distribution is determined by the coagulation equation. We will numerically solve this equation as well. In the vertical direction, we will always assume that each particle species is in vertical sedimentation/mixing equilibrium \\citep{dms95,food}. Hence, we will not solve the time dependent continuity equation in the vertical direction as done for example by \\cite{DD05}. Nor do we need to solve the coagulation equation at all $z$ explicitly. Instead, we solve the vertically integrated coagulation equation, which significantly saves computational time (cf. Appendix \\ref{vertint}). We also formulated the coagulation equation in an implicit way (cf. Appendix \\ref{idiff}) which saves another factor of $\\sim 100$ of computer simulation time. ",
        "conclusions": "In this paper we study the evolution of the dust in a protoplanetary disk. We expanded on the work done by \\cite{DD05} by adding an improved treatment of dust aggregate fragmentation and by including the self-consistent radial drift and mixing of dust aggregates. In addition to this, our integration method is faster by orders of magnitude than the original work by Dullemond \\& Dominik, by virtue of two new methods. The first of these is the integration of the equations in vertical direction without deleting the vertical structure information, which can be done because the vertical sedimentation-mixing equilibrium is assumed to be known at all times. Secondly, we integrate the coagulation-fragmentation in time using implicit integration, allowing us to overcome the extreme time scale stiffness that prevented Dullemond \\& Dominik from performing full scale models including fragmentation over time scales of millions of years.  Using these models we are able to study the problem of the radial drift and fragmentation barriers in full detail. Our models show that, consistent with current beliefs, the combination of radial drift and fragmentation is a strong limitation to growth of aggregates in disks. Typically, aggregates cannot grow to sizes larger than millimeters throughout the disk if threshold fragmentation velocities up to several m/s are considered, though the precise maximum size is disk model dependent. This statement holds regardless of the amount of turbulence in the disk. For highly turbulent disks, it is the turbulence-induced relative velocities that cause the damage to aggregates, while for nearly non-turbulent disks it is the differential radial drift that limits the growth. Only if we set the fragmentation threshold velocity to the unlikely value of 30~m/s or larger do we find that the fragmentation barrier can be broken and particles grow to larger sizes. Whether this high fragmentation threshold is realistic remains to be verified by high-speed laboratory collision experiments. It has been tentatively shown that high-speed impacts of a small projectile on a large target may result in growth \\citep{wurm05}, but further study in this direction is imperative.  We have also investigated what happens when, for some reason, no fragmentation occurs. We then find that even in this reduced model, the particle radius never exceeded several centimeters at any time at any radius in the disk because of radial particle motion. However, we demonstrated that this radial drift barrier problem is very sensitive to slight changes in the initial dust-to-gas ratio. If slighly higher initial dust-to-gas ratios than the canonical value of 1\\% are adopted in the simulations, then particles can grow to very large sizes in the inner parts of the disk.  We do not include non-linear feedback of the dust back onto the gas in our model. It has been recently shown by \\cite{andersnat07} that such feedback can lead to the rapid formation of gravitationally bound clumps of dust which subsequently form Ceres-size bodies. The `dust' particles, however, must be large (Stokes number near unity) before this scenario can take place. We find that for low-turbulent disks Stokes numbers larger than 0.1 can be reached, but we need further investigation if the amount of dust present in these large grains is sufficient to trigger such a gravitational collapse in locally overdense regions in the midplane of the disk.  In the full model, most of the solid material has drifted into the evaporation zone after 1~Myrs and the remnant disk contains less than 5\\% of the initial dust mass. However, there are reasons to believe that the strong radial drift of particles in such disks may be reduced by non-linear hydrodynamic effects.  Tentative results from \\cite{JohKlaHen06} find a reduction by a factor of 3 in MRI turbulence. Moreover, \\cite{brauer07} show that observations of millimeter grains in the outer regions of protoplanetary disks indicate that the standard radial drift formulae are inconsistent with the observations. Hence, it may be important to investigate what happens to our model if radial drift velocities are reduced by some factor. This will be the topic of a future paper.  The dynamic calculations of \\cite{bargesomm95} and \\cite{KlaHen97} also suggest that particle trapping in pressure maxima could be a solution also to the relative velocity fragmentation barrier. We intend, in future work, to study how the dust evolves if such pressure maxima are present.  By being able to model the dust evolution in a self-consistent way, our code may act as a basis upon wich further modeling of disks is done.  For example, models of chemistry in disks are very dependent on the total surface area of dust grains per unit volume, in particular when grain surface chemistry is taken into account \\citep{aikawa99,sem06}. Models of MRI turbulence in disks depend strongly on the abundance of free electrons. This abundance is also very dependent on the total surface area of dust \\citep{ilg106} and our model could -- possibly with an ad-hoc reduction factor for the drift speed -- provide new insight in this kind of modeling.  Our disk model involves a constant threshold fragmentation velocity which can be put into question. Laboratory experiments show that this threshold velocity is dependent on particle size \\citep{bm93,para07} and we indeed work on that topic to investigate if more realistic aggregate collision models predict a different disk evolution (Brauer et al. in prep.).  Finally, current state-of-the-art models of disk structure and their spectra and images relies on ad-hoc prescriptions of the dust spatial and size distribution. Models of the kind described here will be linked to radiative transfer calculations to investigate the effect of grain evolution on disk structure in the near future."
    },
    "0711/0711.2647_arXiv.txt": {
        "abstract": "A comparison of the hot and cool boundaries of the classical instability strip with observations has been an important test for stellar structure and evolution models of post- and main sequence stars. Over the last few years, the number of pulsating pre-main sequence (PMS) stars has increased significantly: 36 PMS pulsators and candidates are known as of June 2007. This number allows to investigate the location of the empirical PMS instability region and to compare its boundaries to those of the classical (post- and main sequence) instability strip. Due to the structural differences of PMS and (post-)main sequence stars, the frequency spacings for nonradial modes will be measurably different, thus challenging asteroseismology as a diagnostic tool. ",
        "introduction": "The pre-main sequence (PMS) evolutionary phase is the short time span between the birth of a star from interstellar clouds and the onset of hydrogen burning upon its arrival on the main sequence. Gravitational contraction is the main energy source during this stage.  Low-mass PMS stars (i.e., with masses lower than $\\sim$1.5 $M_{\\odot}$) that have recently become visible in the optical range and contract along their Hayashi tracks \\citep{hay61}, presumably evolve to T Tauri stars \\citep{joy42}. They primarily have spectral types ranging from late F to M. Intermediate-mass young stars (i.e., with masses between $\\sim$1.5 and 10$M_{\\odot}$) evolve more rapidly to Herbig Ae/Be (HAEBE) objects \\citep[e.g.,][]{her60,fin84} with spectral types B to early F. The evolutionary stage of a HAEBE star is ambiguous: stars with masses below $\\sim$4 $M_{\\odot}$ are still in their PMS phase, while hotter stars have already started nuclear hydrogen burning before they become visible in the optical range. Members of both groups, T Tauri and HAEBE stars, show regular and irregular photometric and spectroscopic variability on very different timescales, indicating that stellar activity begins in the earliest phases of stellar evolution.  A and early F type PMS stars with masses between 1.5 and 4$M_{\\odot}$ have the right combination of effective temperature, luminosity and mass to become vibrationally unstable. With the now considerably increased number of known PMS pulsators and candidates it is possible to compare the instability regions for pre- and (post-) main sequence stars. ",
        "conclusions": "For 36 PMS stars ($\\delta$ Scuti-like) pulsations were discovered by different authors within the last years. The PMS evolutionary phase of these stars has been assessed from observational evidences: either they are members of young open clusters or Herbig Ae stars. Although only limited information is available in the literature for pulsating PMS stars, it is possible to explore their positions in the observational HR diagram. These positions and the observed period range support predictions of driving mechanisms for PMS pulsation that are similar to $\\delta$ Scuti stars. Furthermore, the hot and cool boundaries of the $\\delta$ Scuti instability strip seem to coincide with the borders for the pulsating PMS stars.  It has to be noted that in general the determination of fundamental parameters for pulsating PMS stars is still unsatisfactory. For only two of the 18 pulsating PMS field stars the parallaxes are accurate enough to be useful for a comparison to the $\\delta$ Scuti stars. Young clusters frequently exhibit strong differential reddening and field Herbig Ae stars are often surrounded by circumstellar gas and dust which severely limits the determination of $(B-V)_0$. Temperature calibrations using $(B-V)$ colors exist for stars on the ZAMS \\citep[e.g., ][]{ree98}, but their applicability to PMS stars that have luminosity classes III to V is questionable. Consequently, no $T_{\\rm eff}$ and $L/L_{\\odot}$ data are provided in this paper.  With a dense enough observed pulsation frequency spectrum it will be possible to discriminate directly between evolutionary stages using asteroseismology, as recently has been successfully demonstrated by \\citet{gue07}. This avoids hazards with poorly determined distances and photometric indices  Whether the lack of PMS pulsators in the ``cool\" corner of the classical instability strip has an astrophysical background or is due to poor number statistics, calls for dedicated observations."
    },
    "0711/0711.0368_arXiv.txt": {
        "abstract": "We present a new method for numerical propagation through Lyot-style coronagraphs using finite occulting masks. Standard methods for coronagraphic simulations involve Fast Fourier Transforms (FFT) of very large arrays, and computing power is an issue for the design and tolerancing of coronagraphs on segmented Extremely Large Telescopes (ELT) in order to handle both the speed and memory requirements. Our method combines a semi-analytical approach with non-FFT based Fourier transform algorithms. It enables both fast and memory-efficient computations without introducing any additional approximations. Typical speed improvements based on computation costs are of about ten to fifty for propagations from pupil to Lyot plane, with thirty to sixty times less memory needed. Our method makes it possible to perform numerical coronagraphic studies even in the case of ELTs using a contemporary commercial laptop computer, or any standard commercial workstation computer.\\\\ ",
        "introduction": "\\label{sect:intro} The field of high contrast imaging has expanded rapidly in the past few years, driven by the prospect of science enabled by the direct detection and characterization of extrasolar planets and faint circumstellar disks. Several observatories have launched studies and development of such projects, e.g. Gemini Planet Imager (GPI) \\cite{MGP06}, SPHERE \\cite{BMM06}, or HiCIAO \\cite{THT06}. In the future, Extremely Large Telescopes (ELT) \\cite{MTD06,Clo07,VKB07} will provide the higher angular resolution necessary to study more distant planet forming regions, and also faint old giant planets. The study of Earth-like planets will probably have to wait for space-based instruments \\cite{GAB06,TT07}. Ground-based imaging requires the association of coronagraphy and extreme adaptive optics (ExAO), which can be studied both theoretically and numerically \\cite{SKM01,AS04,CBB06,GPK06,SFA07}.  In this paper we study the numerical modeling of Lyot type coronagraphs \\cite{L32}, including Apodized Pupil Lyot Coronagraphs (APLC) \\cite{ASF02,SAF03,S05,SPF07}, or phase masks \\cite{RR97,SDA03}. These coronagraphs consist of the succession of binary filters (pupil, occulting mask, Lyot stop). The method we describe does not provide any significant improvement in the case of infinite-size focal plane masks \\cite{RRB00,AVB01,KT02,FPS05}, which we do not discuss further here. Under the classical approximations of Fourier Optics \\cite{Goo96}, a Fourier Transform (FT) relationship exists between the field amplitude at two successive planes of the coronagraph. The pupil, with its finite support, is not band-limited and so cannot be perfectly sampled according to the Shanon-Nyquist sampling theorem \\cite{Bra99}. The same problem pertains to the occulting mask in the focal plane. A remedy is to impose very fine sampling both in pupil and focal planes, in order to represent small features (e.g. segmentation or secondary mirror support structures, occulting spot). Note that Nyquist sampling of the focal plane with two pixels per resolution element $(\\lambda/D)$, is not sufficient here. Indeed, the small occulting mask has a typical size of about $5\\lambda/D$, and would not be well represented by a disk ten pixels in diameter. This is even more problematic for phase masks of size $\\sim\\lambda/D$ \\cite{RR97,SDA03,FSA07}. Significant oversampling of the PSF is therefore required for sufficient fidelity in coronagraphic modeling.  This two-fold sampling requirement in two Fourier-conjugate planes is not only a severe practical problem, but raises a fundamental problem associated with the uncertainty principle. In order to preserve complete information in the FT, a fine sampling in one domain means that the corresponding FT must be calculated to a very high frequency in the reciprocal domain. This is achieved optically when forming successive pupil and focal plane images. Numerically, this can be reproduced using Fast Fourier Transforms (FFT), but this can lead to overwhelming demands on memory and computing power. In Sec.\\ref{Sec:classicalMethod} we show that this fundamental limitation makes classical FFTs poorly suited for Lyot-type coronagraphic computations. In Sec.\\ref{sect:semianal} we analyze the analytical formalism of coronagraphic propagation, and conclude that a partial FT algorithm providing arbitrary sampling in a limited area is more appropriate for coronagraphy. Several methods can be used for this purpose \\cite{BS01,RSR69}, and we detail a matrix-based FT \\cite{Smi07}, which can be implemented easily in any high-level language.  Comparisons between classical propagation with FFTs and our semi-analytical method show a considerable gain both in speed and memory requirements. This therefore opens new possibilities for coronagraphic studies hitherto impossible without large computer cluster facilities. Among them, the study of extremely fine sampling of the pupil plane in the case of ELTs, of extremely fine sampling of the occulting plane mask, or of end-to-end simulations of an ExAO coronagraph to simulate a long exposure, are achievable with standard desktop equipment. ",
        "conclusions": "\\label{sect:concl}  In this paper we introduced a semi-analytical method to calculate the numerical propagation through a Lyot-style coronagraph, without any additional approximation. We showed that FFTs are not appropriate for Lyot coronagraphy because of fundamental sampling limitations, and should be avoided for these calculations. The semi-analytical method is derived straightforwardly from the analysis of the wave propagation, and requires the use of a Fourier transform method producing arbitrary sampling in a limited area. Several methods exist to perform such transforms, we suggest a matrix-based Fourier propagator that can be implemented efficiently in any language.  For coronagraphic studies for ExAO on ELTs and eight meter class telescopes, the typical speed improvement based on computation costs is of at least a factor twenty to fifty compared to the classical FFT coronagraphic propagation method. In addition to speed, our semi-analytical method is particularly efficient in terms of memory, as it does not involve any zero padding, and typical memory reequirements are reduced by a factor of about 50 in our usage of the algorithm. In the case of a full propagation from pupil to final focal plane, if the final FT is performed using a FFT, the semi-analytical method is approximately three times faster than a FFT-based propagation, since the propagation time is entirely dominated by the final FFT. If a reduced field of view is calculated in the final image, which is often the case in extreme adaptive optics coronagraphy, an MFT can be used in lieu of an FFT. For example a factor of 5 can be gained on the last FT for $N_A=400$ and $\\gamma=10$, and the overall gain with the semi-analytical method is a factor of 15.  With the semi-analytical method, the pupil and focal plane sampling are completely independent. This flexibility enables calculations that were hitherto impossible, or only possible on large and expensive computer clusters. For example, the study of fine occulting mask features and alignment, or fine pupil structures like segmentation, is now achievable on a well-equipped but otherwise standard laptop, or on most commercial workstation sold today.  The MFT may also find interesting applications where a FT has to be calculated in a limited area. This was used by Give'on \\textit{et al.} \\cite{GKV05,GKV06} and new applications may be found in adaptive optics for anti-aliasing spatial filters \\cite{PM04}, pyramid wavefront sensing \\cite{Rag96}, etc.  Our method also enables fast computation of optimal apodizers for APLCs, which involves an iterative algorithm and makes computing efficiency even more critical \\cite{SPF07}. APLCs are one of the most promising concepts for high contrast imaging instruments on ELTs, and their study requires the optimization of the mask size according to various criteria, over broad spectral bandpasses. A complete exploration of the effects of chromaticity, segmentation and spiders on the coronagraph is now made possible with acceptable computing times. Our method could be utilized in several current efforts focussed on ground and spaced based coronagraphy."
    },
    "0711/0711.0018_arXiv.txt": {
        "abstract": "{ Flux-dominated solar dynamo models have demonstrated to reproduce the main features of the large scale solar magnetic cycle, however the use of a solar like differential rotation profile implies in the the formation of strong toroidal magnetic fields at high latitudes where they are not observed. In this work, we invoke the hypothesis of a thin-width tachocline in order to confine the high-latitude toroidal magnetic fields to a small area below the overshoot layer, thus avoiding its influence on a Babcock-Leighton type dynamo process. Our results favor a dynamo operating inside the convection zone with a tachocline that essentially works as a storage region when it coincides with the overshoot layer. } ",
        "introduction": "Flux-dominated solar dynamo models, which use a solar like velocity field, together with an estimated diffusivity \\linebreak profile and an $\\alpha$ effect resembling the Babcock-Leighton mechanism for producing poloidal magnetic fields, have \\linebreak  demonstrated to provide solutions that resemble quite well the observations of the large scale solar magnetic cycle \\linebreak (Dikpati et al. 2004, Chatterjee, Nandy \\& Choudhuri 2004, Guerrero \\& de Gouveia Dal Pino 2007, hereafter GDP).  In this kind of models, the usually accepted scenario for the dynamo operation can be described as follows: starting with a dipolar field, the $\\Omega$ effect takes place both at the convection zone and the tachocline, and a toroidal magnetic field is formed and amplified by the radial and latitudinal components of the differential rotation. Once this field reaches values between $10^4$ - $10^5$ G, it becomes buoyant unstable and arises through the convection zone being twisted by the Coriolis force until it emerges at the surface and forms bipolar magnetic regions. This part of the process is known as Babcock - Leighton $\\alpha$ effect. At the surface, the leading part of a bipolar region points to the equator, while the rear part is drifted polarward by the meridional circulation. This migration contributes to the formation of big loops in each hemisphere, which later will reconnect in order to form a new poloidal field of opposite polarity to the initial one.  In view of the present status of the observations, the model proposed above presents several problems, as \\linebreak remarked by Brandenburg (2005), the most important of which is related to the location, the sign and the amplitude of the alpha effect. In the scenario discussed above, the $\\alpha$ term should  reproduce the emergence of magnetic flux tubes of strong toroidal magnetic fields and then the migration of these fields toward the poles. In this way, the most probable location for the dominance of this term would be the upper layers of the convection zone, with a latitudinal  distribution peaking around $30 ^{\\circ}$, which is the latitude at which the strongest magnetic activity is observed. If the solar dynamo is operating at the tachocline, where the shear has positive values, the natural sign of the alpha effect should be negative (positive) in the north (south) hemisphere, in order to get equatorward branches in agreement with the mean field dynamo theory. However, the presence of an equartorward meridional flow at deeper layers could bring the possibility of a positive (negative) $\\alpha$ effect at the north (south) hemisphere and still give the correct migration. Recent works have suggested that the amplitude of the $\\alpha$ term could be important in the determination of the parity of the dynamo (Dikpati \\& Gilman 2001, Bonnano et al. 2002), this effect is not considered in this work since our domain corresponds to one hemisphere only, but it will be addressed in a forthcoming paper (Guerrero \\& de Gouveia Dal Pino 2007, in preparation).  Another important question regards the real importance of the tachocline in the dynamo process. As a layer of strong radial shear, the tachocline could be the ideal site for the amplification of the magnetic field, though the maximum amplitude of  the shear is located at high latitudes, which suggests an intense activity  close to the poles. In this case, meridional circulation could be again an appropriate solution to this problem, since a velocity of about $3$ m s$^{-1}$ could advect down and equatorward the magnetic field across a stable layer until reaching the latitudes where the strong activity is observed (Nandy \\& Choudhuri 2002, Chatterjee et al. 2004). However, the use of a deep meridional  flow in a flux dominated dynamo has been found to be very sensitive to the adopted $\\alpha$ effect and  the diffusion and velocity profiles (Guerrero \\& Mu\\~{n}oz 2004) and can also generate undesirable abundance variations in the convection zone. In addition, several numerical simulations (Gilman \\& Miesh 2004, R\\\"{u}diger, Kitchatinov \\& Arlt 2005) have demonstrated the impossibility of penetration of the meridional flow below the tachocline more than a few kilometers. In order to allow some shallower penetration, other possibilities, such as overshooting or turbulent magnetic pumping have been invoked in the literature (Rogers, Glatzmaier \\& Jones 2006, Ziegler \\& R\\\"{u}diger 2003, respectively), where magnetic flux tubes are drifted to the overshoot layer and can be amplified to the desired magnitudes.  Indeed, if a flow penetrates inside the tachocline, there should be a mechanism that could prevent either the formation of strong toroidal magnetic fields at higher latitudes or the participation of these fields in the cycle which is responsible for the observed phenomena. In recent work (GDP), we have shown that a possible solution to the problem above could be the choice of an appropriate set of physical parameters, since the model is sensitive to the location of the transition between the sub-adiabatic and the super-adiabatic layers, to the diffusivity value at the bulk of the convection zone and to the thickness of the solar tachocline (see Figs. 8a and 8b of GDP). We have also found that the latitudinal shear term is dominant over the radial shear term throughout the cycle, and that the radial shear peaks only in the early years of the cycle and decays afterward  to values about two orders of magnitude smaller than the latitudinal component.  In this paper, we present a detailed description on how the suppression of the magnetic fields at high latitudes is obtained. We also present the results when two different prescriptions for the meridional circulation profile are employed. They seem to be insensitive to these variations. ",
        "conclusions": "We have described a physical scenario in which a simple flux-dominated solar dynamo model with an appropriate set of parameters produces results which are in good agreement with the observations. The main assumptions and results are summarized below:  \\begin{enumerate} \\item Since an analysis of the evolution of the maximum of the shear terms along the cycle shows that the latitudinal shear is always dominant over the radial one (which attains appreciable  values only during the first years and quickly decays afterward), a possible mechanism to prevent the formation of toroidal magnetic fields  by the radial shear at high latitudes could be the  consideration of a thin tachocline. In this case, we find that  the dynamo will form a thin layer of toroidal field below the overshoot layer  with insufficient amplitude to produce buoyant flux tubes. Several numerical tests have demonstrated the robustness of this assumption in order to avoid the appearance of sunspots near the poles. \\item The bulk of the toroidal magnetic field is formed at the interior of the convection zone thanks to the latitudinal shear and is drifted by convection and/or magnetic pumping towards the sub-adiabatic overshoot layer \\linebreak where we assume that it undergoes magnetic buoyancy and rises to the surface at values above of $5$$\\times$$10^4$ G. Note that the turbulent diamagnetism (which was mentioned in the section 2) may also play an important role in order to pile-up the toroidal magnetic field in the overshoot layer, since it behaves like a downward velocity term in the induction equation, as demonstrated by \\linebreak Kitchatinov \\& R\\\"{u}diger (1992). The importance of this term has not yet been studied in detail, but will be considered in forthcoming studies. \\item The above mentioned results suggest  that the dynamo which is responsible for the observed activity does not operate at the tachocline, as usually assumed, but inside the convection zone and is due to the strong latitudinal shear. In this sense, the real role of the tachocline, if it coincides with the overshoot layer, is to store the magnetic flux tubes until they reach the necessary amplification to become buoyantly unstable. \\end{enumerate}"
    },
    "0711/0711.2537_arXiv.txt": {
        "abstract": "We study the strong gravitational lens statistics of triaxial cold dark matter (CDM) halos occupied by central early-type galaxies. We calculate the image separation distribution for double, cusp and quad configurations. The ratios of image multiplicities at large separations are consistent with the triaxial NFW model,  and at small separations are consistent with the singular isothermal ellipsoid (SIE) model. At all separations, the total lensing probability is enhanced by adiabatic contraction. If no adiabatic contraction is assumed, naked cusp configurations become dominant at $\\approx 2.5''$, which is inconsistent with the data.  We also show that at small-to-moderate separations ($\\lesssim$ 5'') the image multiplicities depend sensitively on the alignment of the shapes of the luminous and dark matter projected density profiles. In constrast to other properties  that affect these ratios, the degree of alignment does not have a significant effect on the total lensing probability. These correlations may therefore be  constrained by comparing the theoretical image separation distribution to a sufficiently large lens sample from future wide and deep sky surveys such as Pan-Starrs, LSST and JDEM. Understanding the correlations in the shapes of galaxies and their dark matter halo is important for future weak lensing surveys. ",
        "introduction": "\\label{sec:intro}  Statistics of strongly lensed images promise a wealth of information about luminous and dark matter distributions and the correlations between them. For example, these statistics have yielded general constraints on the radial mass profile of early-type galaxies (\\cite{maoz93}, \\cite{keeton01a},  \\cite{wyithe01}, \\cite{keeton01}, \\cite{rusin-ma01}, \\cite{takahashi01},  \\cite{sarbu01}, \\cite{oguri02}, \\cite{oguri02a}, \\cite{huterer04},  \\cite{kuhlen04}). Lensing by non-spherical mass distributions shows rich phenomenology that may be used to probe the distribution of dark matter in galaxies and clusters. The shape and concentration of the lenses Mass distributions that are approximately spherical produce a higher fraction of double-image systems, whereas more ellipitical distributions produce more quadruple-image systems. In rare cases, extreme ellipticities  can produce strongly magnified three-image configurations known as ``naked  cusp'' systems (or simply ``cusps''). Studying the statistics of different  image multiplicities (quad, double, and cusp lenses) could help us constrain the shapes of dark matter halos, the environment of and substructure in these halos and as we will show the correlations between the shapes of the luminous and dark matter in galaxies.  The correlations in the shapes of the galaxy and the dark matter halo it resides in is important for four reasons. (1) As we show, these correlations are important to include to make accurate predictions for the strong lensing probabilities. (2) A large sample of strongly lensed images will provide a rich data set that can help us understand galaxy formation better. One of the most interesting aspects of galaxy formation, and perhaps the least well studied, is the correlation in the shapes of the luminous and dark matter components. Lensed image statistics can provide constraints on this aspect. (3) An important systematic effect for weak lensing surveys is the spatial correlation in the shapes of the galaxy shapes, which arise is if galaxy shapes are correlated with their dark matter halo and halo environment  (\\cite{croft00}, \\cite{heavens00}, \\cite{catelan01}, \\cite{crittenden01}, \\cite{jing02a}, \\cite{mandelbaum06}, \\cite{hirata07}). All future weak lensing surveys will cover a large fraction of the sky at great depth and also reveal a great number of strongly lensed quasar and galaxy images. With an accurate framework in place, one may use this strong lensing data set to understand systematic effects stemming from  luminous and dark matter correlations for weak lensing better. (4) Strong lensing systems are also important for understanding fundamental physics. As shown by \\cite{bolton06a}, these systems may be used to constrain the post-newtonian gravity parameter $\\gamma$. Constraining this parameter also directly contrains theories of modified gravity since in this context $\\gamma$ is the ratio of the two scalar potentials in the perturbed FRW metric (\\cite{bertschinger06,zhang07}).  The distribution of image multiplicities for elliptical lenses has been modeled in two regimes: small image separations ($<1''$), in which galaxies become the dominant lensing mass, and large (cluster-sized) separations ($>10''$), in which dark matter becomes dominant.  \\cite{rusin-tegmark01} modeled galaxy lenses with a singular isothermal ellipsoid (SIE) mass distribution, comparing to lenses from the Cosmic Lens All-Sky Survey (CLASS) and Jodrell-VLA Astrometric Survey (JVAS).  They showed that the observed number ratio of quads-to-doubles is greatly underestimated by the predictions of the SIE model (dubbed the ``quad problem''). While this problem has not been solved rigorously, it has been suggested that the effect of massive substructures could sufficiently increase the expected quad fraction to be consistent with observations (\\cite{cohn04}).  Other effects which boost the lens probability for quads include misalignment of the centers of galaxies with their host halos (\\cite{quadri03}) and external shear from lens galaxy environments (\\cite{huterer05}).  Large separation lenses, on the other hand, were modeled by \\cite{oguri-keeton04} using the triaxial NFW model based on simulations by \\cite{jing02}. With an inner slope $\\alpha = 1.0$ (pure NFW), the image multiplicities are shown to be remarkably different from those of galaxy lensing: cusp systems have the highest probability, followed by quads, with doubles being least common.  However, because the overall lensing probabilities are orders of magnitude smaller than that of galaxy lensing, such dark-matter dominated lenses are relatively rare---the only observed lens in this regime is a quad at $14.62''$, found in the Sloan Digital Sky Survey (\\cite{inada03}).  In contrast to the paucity of confirmed large-separation lenses, a significant fraction of observed gravitational lenses lie in the intermediate region between these two extremes, with image separations roughly between $1''$ and $10''$.  The total lensing probability in this region has been studied for the case of spherical mass distributions (\\cite{oguri06}, \\cite{ma03}, \\cite{kochanek01}), but the distribution of image multiplicities for nonspherical lenses has not yet been derived.  We note that none of the single-component elliptical models  mentioned thus far can be applied in this regime, since both the dark matter  and the central galaxy contribute significantly at these separations.   Furthermore, any lensing model including baryons and dark matter must take into account the modification of dark matter halos due to baryon cooling, which increases the concentration of the inner parts of halos (\\cite{kochanek01}, \\cite{keeton01}).  Although a multi-component model was employed by \\cite{moeller07} using galaxies with ellipticity, the dark matter was modeled by spherical NFW halos, ignoring adiabatic contraction and the triaxiality of dark matter halos.  In this paper we derive a distribution of image multiplicities that will be approximately valid for all observed image separations, including the intermediate region. However, to accurately model this region, we must combine the triaxial NFW model, adiabatic contraction, and a central lens galaxy. The lensing properties then become sensitive to the ellipticities, concentrations, and orientations of both galaxy and halo, in addition to the fraction of the total mass contained in the galaxy. This requires a more complicated calculation of lensing probabilities than those in the literature. Despite the complexity of the models considered in this work, several important factors (considered prevoiously in isolation) have to be neglected. In particular, we highlight the absence of dark matter substructure in our model and the effect of the lens galaxy environment. We will later discuss the expected impact of these factors.  We find several interesting features that suggest new methods of constraining lens mass distributions. For example, the ratio of image multiplicities is found to be sensitive to the alignment of the major axes of galaxies with that of their host halos.  Further, the transition from small- to large-separation lenses is marked by a ``crossing'' of lens probabilities when the expected lensing rate of quads overtakes that of doubles, and likewise for cusps.  Understanding the precise location of these turnover points may provide further means of constraining galaxy-halo mass distributions once the sample of observed lens systems becomes sufficiently large.  The paper is organized as follows. In \\S \\ref{sec:imgdist} we will present a general form for the image separation distribution, and discuss the relevant model parameters. Section \\ref{sec:lensmodel} will describe the lens model we used and outline   our approach toward applying the adiabatic contraction model to triaxial halos. Section \\ref{sec:calc} covers the details of the calculation, and results are  given in \\S \\ref{sec:results}.  Finally we summarize and discuss our results in  \\S \\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  We have shown that the distribution of strong lensing image multiplicities offers new possibilities for constraining correlations between galaxies and their host dark matter halos---in particular, adiabatic contraction and the degree of axial alignment. Since varying the amount of alignment makes a considerable difference in the quad/double ratio, and given that the observed ratio is so high, one might expect galaxies and halos to be quite closely aligned.  Simulations seem to support this (\\cite{bailin05}, \\cite{kazantzidis04}), and indeed there is some observational evidence from lensing: \\cite{kochanek02} used mass profile modeling of 20 lenses from the CfA-Arizona Space Telescope Lens Survey (CASTLES) to show that the mass and light distributions were aligned within $\\langle \\Delta \\theta^2\\rangle^{1/2} < 10^\\circ$ where $\\theta$ is the projected alignment angle.  The distribution of image multiplicities provides a simple method, far less time-consuming than individual lens modeling, that can be applied to a large statistical sample of lenses. This method also has the advantange that the data for each individual lens does not have to be very high quality. The disadvantage is that we have to include all correlations and enviromental effects that are either constrained by other data or directly obtained from theory.  One important caveat to keep in mind is that alignment of the major/minor axes of galaxy and  halo does not necessarily imply that the \\emph{projected} axes will be aligned. The projected axes can also be misaligned if the galaxy and halo have differing triaxialities (\\cite{keeton97}), analogous to the phenomenon of ``isophote twist''. Our method constrains the alignment of the projected axes. If the projected axes are shown to be closely aligned, then we can draw the strong conclusion that the major/minor axes of galaxy and halo are aligned, and that the triaxial galaxy and halo shapes are similar on average.  There are several ways in which our model can be improved. First, a fully triaxial model for adiabatic contraction from simulations is required. We used analytic fitting functions from the \\cite{gnedin04} model of adiabatic contraction that was derived from halos with spherically averaged density profiles, extending it to triaxial halos by making the replacement $ r \\rightarrow R$ where $R$ is the triaxial radius. But this is clearly only an approximation that may break down at high ellipticity.  The predicted lensing probability for cusps may differ significantly when a more accurate model of adiabatic contraction is used.  In addition, we have neglected the contribution from spiral galaxies.  This is partly justified by the fact that at least 75\\% of galaxy lensing is thought to be due to ellipticals (\\cite{moeller07}). However, since disk galaxies viewed edge-on can have extremely high projected ellipticities, their effect on image multiplicities may be substantial. We have also not included scatter in the mass-luminosity relation and in halo concentration, the effects of which were discussed in the previous section. Although the image multiplicities are not drastically effected by the scatter, it is nevertheless an important effect and has an appreciable impact on the total lensing probability.  We have also ignored the effect of external shear from lens galaxy environments, which can increase lensing probabilities at intermediate image separations, $> 3''$ (\\cite{oguri05}).  Another effect is the misalignment of the center-of-mass of galaxies with that of their host halos, which will slightly enhance the cross section for quads (\\cite{quadri03}).  Perhaps most importantly, we have not included substructure in our calculations.  Work by \\cite{cohn04} suggest that satellites can roughly double the expected quads-to-doubles ratio, thus dramatically altering the predicted ratio from our model. They also give rise to cross sections for more exotic image configurations, with five or more visible images, whose lens probabilities may be significant.  Including all these effects will be essential for deducing rigorous constraints on the profile, shape and correlations of lensing masses from future data."
    },
    "0711/0711.2715_arXiv.txt": {
        "abstract": "The hard X-ray detector (HXD) on board the X-ray satellite Suzaku is designed to have a good timing capability with a 61 $\\mu$s time resolution. In addition to detailed descriptions of the HXD timing system, results of in-orbit timing calibration and performance of the HXD are summarized. The relative accuracy of time measurements of the HXD event was confirmed to have an accuracy of $1.9\\times 10^{-9}$ s s$^{-1}$ per day, and the absolute timing was confirmed to be accurate to 360 $\\mu$s or better. The results were achieved mainly through observations of the Crab pulsar, including simultaneous ones with RXTE, INTEGRAL, and Swift. ",
        "introduction": "\\label{section:hxd_intro} The Hard X-ray Detector (HXD; \\cite{hxd2007a,hxd2007b}) on board Suzaku \\citep{suzaku2007} is a novel cosmic hard X-ray instrument working  in the 10 -- 600 keV range. In addition to the very low background level, the wide energy band, and the tightly collimated field of view, yet another important feature of the HXD is its timing capability. This is because time variability, periodic or aperiodic, is an important characteristic of compact cosmic X-ray sources, including neutron stars, white dwarfs, galactic black holes, and active galactic nuclei. Considering the wide range of time scales involved in these variations, we have required the HXD to have a time resolution of 61 $\\mu$s, and a timing stability of the order of $10^{-9}$ typically for a day. To fulfill these goals, we carefully designed the timing system of the HXD, as summarized in section \\ref{section:hxd_timing}, and in more detail in \\citet{hxd2007a} and \\citet{hxd2007b}.  The HXD sensor (HXD-S) consists of 16 identical ``Well units'' working as main detectors \\citep{IEEE_uchiyama2001}, and 20 surrounding shield counters made of Bi$_4$Ge$_3$O$_12$ (BGO) scintillators \\citep{IEEE_yamaoka2006}. Each Well unit detects incoming X-rays using 2 mm-thick silicon PIN diodes \\citep{IEEE_sugiho2001} and Gd$_2$SiO$_5$:Ce (hereafter GSO) scintillators, both surrounded by a common active BGO shield. The HXD signals from the Well Units (called WEL type data) are processed first by an analog electronics package (HXD-AE), and then by a digital electronics (HXD-DE).  \\begin{figure}[bht] \\begin{center} \\FigureFile(0.95\\columnwidth,){figure1.eps} \\end{center} \\caption{Timing counters in the HXD electronics.} \\label{fig:hxd_timing_system} \\end{figure}  Before launch, we repeatedly tested and confirmed the HXD timing capability, first at each component level, then incorporating HXD-S, HXD-AE, and HXD-DE \\citep{IEEE_kokubun2003,IEEE_tashiro2003,SPIE_kawaharada2004,IEEE_terada2006,IEEE_ohno2006}. The tests included confirming time assignments of single WEL-type events with specified time intervals (1, 10, 100, and 1000 s) from a time origin, detection of periodic signals (with periods of 1, 10, and 100 ms), and measuring time-interval distribution using random events. After the HXD was mounted on the spacecraft, another series of end-to-end tests were repeated incorporating spacecraft central clock, radio transmitters, and ground receivers.  Through the pre-launch verification, we confirmed that the HXD has correct timing function as designed. However, these tests were limited in many aspects. For example, the absolute timing measurement was performed only before  the HXD was mounted on the spacecraft. The ground equipment used in these end-to-end tests was not identical to those actually used at the Suzaku tracking center, the Uchinoura Space Center (USC) in southern Japan. Furthermore, once the spacecraft is put into orbit, we must consider additional complications, such as signal delays due to air propagation to the ground station, and due to cable transmission from the antenna to signal receivers.  \\begin{figure}[bht] \\begin{center} \\FigureFile(0.95\\columnwidth,){figure2.eps} \\end{center} \\caption{A timing chart between HXD-AE and HXD-DE at the TIME LATCH operation. DP\\_TIME and DP\\_Clock are the signals from DP to HXD-DE to acknowledge time information of the spacecraft, and AE\\_Control\\_Clock is provided into HXD-AE from HXD-DE. The signal named ACU\\_Time\\_Latch is produced at a time-latch operation in HXD-DE, and is supplied into HXD-AE to verify timings between them.} \\label{fig:hxd_time_latch} \\end{figure}  The present paper deals with in-orbit timing calibration of the HXD. After  systematic errors affecting the time assignment are estimated  in section \\ref{section:errors}, we analyze  in section \\ref{section:results} highly stable periodic signals from four fast-rotating pulsars, including the Crab pulsar in particular. These tell us the relative accuracy and stability of the time measurements with the HXD. Finally, we verify the  absolute timing accuracy of the HXD, by comparing the HXD measurements of the Crab pulsar against those with other X-ray missions and with radio telescopes.  \\begin{table*}[htb] \\caption{Summary of Timing Counters of the Hard X-ray Detector.} \\label{tbl:hxd_counter} \\begin{center} \\tabcolsep 2pt \\begin{tabular}{llcccl} \\hline % {Name}&{Component}&{Length}   &{Base}  &{Coverage in Telem.}&{Telemetry Type/Packet}\\\\ {}    &{}         &{(bit)}    &{(s)}      &{(s)}        &\\multicolumn{1}{c}{(HK/Event)}\\\\ \\hline % {Base-TI} &{DP / Spacecraft} &{39}    &{$1/2^{19}$(1.9$\\mu$)}&{$1/2^{12}$(244$\\mu$) -- $2^{20}$(1,048,576)}&{HK /time packet}\\\\ {}        &{HXD-DE}          &{39}    &{$1/2^{19}$(1.9$\\mu$)}&{$1/2^{12}$(244$\\mu$) -- $2^{20}$(1,048,576)}&{HK / HXD-HK}\\\\ \\hline % {AcuTime}&{HXD-AE ACU}&{24} &{$1/2^{17}$(7.63$\\mu$)}&{$1/2^{15}$(30.5$\\mu$) -- $2^{9}$(512) }&{HK / HXD-HK}\\\\ \\hline % {TimeStampTime}&{HXD-AE WPU}&{24} &{$1/2^{15}$(30.5$\\mu$)}&{$1/2^{15}$(30.5$\\mu$) -- $2^{9}$(512)}&{HK / HXD-HK}\\\\ {EventTime$^\\dagger$}&{HXD-AE WPU}&{19} &{$1/2^{14}$(61$\\mu$)}&{$1/2^{14}$(61$\\mu$) -- $2^{5}$(32)}&{Event / HXD\\_WEL}\\\\ \\hline % {TransientTime$^\\ddagger$}&{HXD-AE TPU}&{24} &{$1/2^{15}$(30.5$\\mu$)}&{$1/2^{15}$(30.5$\\mu$) -- $2^{9}$(512)}&{Event / HXD\\_WAM}\\\\ {BurstSendTime}&{HXD-AE TPU}&{23} &{$1/2^{6}$(15.6m)}&{$1/2^{6}$(15.6m) -- $2^{18}$(262144)} &{Event / HXD\\_WAM}\\\\ {BurstFrzTime} &{HXD-AE TPU}&{32} &{$1/2^{15}$(30.5$\\mu$)}&{$1/2^{5}$(30.5m) -- $2^{18}$(262144)} &{Event / HXD\\_WAM}\\\\ \\hline % \\end{tabular}  \\begin{minipage}{1.9\\columnwidth} {\\scriptsize $\\dagger$ Time value tagged to each event which detected by a main unit of the HXD sensor.\\\\ $\\ddagger$ Time value used for the HXD-WAM (Wide-band All-sky Monitor \\citep{IEEE_yamaoka2006}). } \\end{minipage} \\end{center} \\end{table*} ",
        "conclusions": "Systematic errors on the HXD timing due to spacecraft location errors, the temperature-dependence of the on-board clock, and between ground contacts (if it is less than 5 days), are $\\sim 30$, $<40$, and $<150$ $\\mu$s, respectively. Using the Crab pulsar, we confirmed the stability of the relative timing to be $1.9 \\times 10^{-9}$ s s$^{-1}$ per day. The systematic errors on the HXD absolute time have been calibrated as about 360 $\\mu$s or higher ($\\sim$ 270 $\\mu$s under daily ground contacts with the 30m antenna system)."
    },
    "0711/0711.0526_arXiv.txt": {
        "abstract": "We report new results obtained from multi-frequency observations of PSR B0826$-$34 with the Giant Metrewave Radio Telescope (GMRT). (1) We find no evidence of weak emission during the typical long null state of this pulsar, simultaneously at 303 and 610 MHz, as well as individually at 157, 325, 610 and 1060 MHz at separate epochs. Our limit of non-detection is at $\\sim$ 1\\% or better of the peak of the active state profile, and corresponds to $\\sim$ 2 mJy at 610 MHz. (2) Significant correlation in the total intensity of the individual pulses between 303 and 610 MHz is reported from the simultaneous dual frequency observations, which is indicative of the broadband nature of the emission. We also report correlation between total energy in the main pulse and inter-pulse region from the high sensitivity single frequency observations at 610 and 1060 MHz. (3) Though we find the drift pattern to be very similar in the simultaneous 303 and 610 MHz data, we observe that the drift band separation ($P_2$) evolves significantly between these two frequencies, and in a manner opposite to the average profile evolution. In addition, we confirm the dependence of $P_2$ on pulse longitude at 303 MHz and find indications for the same at 610 MHz. We also present results for subpulse width ($\\Delta\\Phi_s$) at different frequencies, and as well as a function of pulse longitude. (4) As a natural out-come of the simultaneous dual frequency observations, we obtain an accurate DM value, equal to  52.2(6) $pc/cm^3$, for this pulsar. ",
        "introduction": "\\label{sec:intro} The phenomenon of subpulse drifting (first reported by \\cite{Drake_etal}), is manifested as an organised subpulse behaviour $-$ subpulses appear at progressively changing longitude in the pulse window following some particular path. The path followed by the subpulses is specific to the individual pulsar concerned and is known as drift band. It has been recently shown that subpulse drifting may be fairly common among pulsars \\citep{Weltevrede_etal_06, Weltevrede_etal_07}. Wide profile pulsars showing this phenomenon can be especially useful tools to understand the distribution of the radio emission regions above the neutron star's polar cap and the physics of the pulsar emission mechanism. Since the wide profiles are usually interpreted as emission from one magnetic pole of a highly aligned pulsar (i.e. magnetic axis almost parallel to rotation axis), our line of sight samples a large part of the polar cap. This leads to enhanced possibility of seeing multiple drift bands in such pulsars, which can provide valuable insights into the pulsar magnetospheres, as exemplified by some of the new and substantive results for PSR B0826$-$34 \\citep{Gupta_etal,Esamdin_etal}, PSR B0815$+$09 \\citep{Qiao_etal} and PSR B0818$-$41 \\citep{Bhattacharyya_etal}.  PSR B0826$-$34 is a pulsar with one of the widest known pulse profiles. The earlier studies of this pulsar \\citep{Durdin_etal,Biggs_etal,Gupta_etal} have brought out some unique properties $:$ strong evolution of the average profile with frequency, apparent nulling for 70\\% of time and a remarkable subpulse drift property $-$ multiple, curved drift bands with frequent changes and sign reversals of drift rate.  Multi-frequency observations, especially when done simultaneously at the different frequencies, are very useful for the study of various aspects of the frequency dependence of pulsar radiation \\citep{Taylor_etal_75,Bartel_etal_78a,Bartel_etal_81a,Bhat_etal}. Dispersion measure (DM) of pulsars can also be measured from single epoch simultaneous dual frequency observations \\citep{Ahuja_etal}. Simultaneous, multi-frequency observations of PSR B0826$-$34 can be useful on a variety of counts: (i) Determination of accurate DM, as the values reported and used in the literature span a fairly wide range, from 47 to 65 $pc/cm^3$. (ii) Study of the pulsar in the null state: \\cite{Esamdin_etal} claim the presence of weak emission, at 1374 MHz, during the long duration nulls of this pulsar. They identify these null states as some weak mode of emission where the pulse intensity is 2\\% of that of the regular, strong mode, with a pulse profile that is similar to the strong mode at lower frequencies. Sensitive, multi-frequency observations at the lower frequencies should provide useful information about the behaviour in the null state. (iii) Study of simultaneous subpulses: Correlated behaviour of subpulses at different frequencies \\citep{Taylor_etal_75,Bartel_etal_78a,Bartel_etal_81a,karastergiou_etal,kramer_etal} is useful to probe the broadband nature of the emission mechanism. For pulsars with drifting subpulses, evolution of $P_2$ (longitude separation between adjacent drift bands) with frequency can also be significant \\citep{Izvekova_etal,Rankin_etal_05,Smits_etal} and a comparison of this with the frequency evolution of other profile parameters can be useful. Earlier study of drift properties of PSR B0826$-$34 \\citep{Biggs_etal,Gupta_etal,Esamdin_etal} report different $P_2$ values at various individual frequencies which do not follow the trend predicted by profile evolution with frequency. In addition, for PSR B0826$-$34, the intricate nature of the multiple drift bands itself deserves a broadband study.  We observed PSR B0826$-$34 using the GMRT, simultaneously at 303 and 610 MHz, and individually at 157, 325, 610 and 1060 MHz in total intensity mode. In this paper we concentrate on the results from the simultaneous dual frequency observations, and supplement these, where necessary, with the results from the single frequencies observations. In Sect.\\ref{sec:observations} we cover the observations, where a new technique for simultaneous dual frequency pulsar observations with the GMRT is presented. Different sub-sections of Sect.\\ref{sec:results} describe the data analysis and results for (1) determination of DM value from simultaneous dual frequency observations, (2) study of the pulsar in its null state and (3) subpulse studies. In Sect.\\ref{sec:Discussions} we discuss the implications of our findings. ",
        "conclusions": "\\begin{tabular}{l|c|c|c|c|c|c|c|c|c|c|c|c|c} \\hline Frequency &State         &No of pulses &No of     &\\multicolumn{2}{|c|}{rms$^\\bullet$}  &Flux of peak  &Non-detection   &\\multicolumn{3}{|c|}{Mean Flux}\\\\ (MHz)     &(Active/Null) &analysed     &antennas  &\\multicolumn{2}{|c|} {}      &              &limit$^\\diamond$&\\multicolumn{3}{|c|} {}       \\\\ &              &             &used      &\\multicolumn{2}{|c|}{(mJy)}&  (mJy)       &   (mJy)        &  \\multicolumn{3}{|c|}{(mJy)}                  \\\\ &              &             &          &   MP   &   OP                     &              &                & FP          & MP$^\\odot$  &  IP$^\\odot$ \\\\\\hline 303$^\\dagger$  & Active  & 750         &    5     &  371   &   28                     & 1687$^\\ast$  &     -          &  443$^\\ast$ & 754 $^\\ast$& 99$^\\ast$ \\\\ &         &             &          &        &                          &              &                &             &            &             \\\\ & Null    & 750         &    5     &   28   &   28                     &     -        &  31$^\\ast$     &             &            &            \\\\ &         &             &          &        &                          &              &                &             &            &            \\\\\\hline 610$^\\dagger$  & Active  & 750         &    10    &   57   &    5                     & 250$^\\ast$   &     -          &  119$^\\ast$ & 178$^\\ast$ & 123$^\\ast$  \\\\ &         &             &          &        &                          &              &                &             &            &             \\\\ & Null    & 750         &    10    &   4    &    4                     &      -       &   5$^\\ast$     &             &            &            \\\\ &         &             &          &        &                          &              &                &             &            &             \\\\\\hline 157$^\\ddagger$ & Active  & 970         &    16    &  427   &   13                     & 1780$^\\ast$  &     -          &  370$^\\ast$ &  656$^\\ast$&  42$^\\ast$ \\\\ &         &             &          &        &                          &              &                &             &            &             \\\\ & Null    & 970         &    16    &  13    &   13                     &      -       &   23$^\\ast$    &             &            &             \\\\ &         &             &          &        &                          &              &                &             &            &            \\\\\\hline 325$^\\S$       & Active  & 940         &    21    &  304   &   5                      & 1330$^\\ast$  &      -         &  393$^\\ast$ &  763$^\\ast$& 62$^\\ast$  \\\\ &         &             &          &        &                          &              &                &             &            &            \\\\ & Null    & 487         &    17    &  5     &   5                      &      -       &   7$^{\\ast\\circ}$ &          &            &             \\\\ &         &             &          &        &                          &              &                &             &            &            \\\\\\hline 610$^\\ddagger$ & Active  & 705         &    17    &  85    &   2                      & 410$^{\\ast\\circ}$ &  -        & 156$^{\\ast\\circ}$ & 248$^{\\ast\\circ}$  & 166$^{\\ast\\circ}$ \\\\ &         &             &          &        &                          &              &                &             &            &              \\\\ & Null    & 705         &    17    &  2     &   2                      &      -       &    2$^{\\ast\\circ}$ &         &            &             \\\\ &         &             &          &        &                          &              &                &             &            &             \\\\\\hline 1060$^\\ddagger$& Active  & 2900        &    21    &  120    &   10                      & 1310$^\\ast$  &       -        & 479$^\\ast$  & 485$^\\ast$ & 989$^\\ast$\\\\ &         &             &          &        &                          &              &                &             &            &                 \\\\ & Null    & 2900        &    21    &   10    &   10                      &      -       &   15$^\\ast$    &             &            &             \\\\ &         &             &          &        &                          &              &                &             &            &              \\\\\\hline \\end{tabular}  $\\dagger$ : Active and null state data recorded from same simultaneous dual frequency observing session \\\\ $\\ddagger$ : Active and null state data recorded from same single frequency observing session \\\\ $\\S$ : Active and null data recorded from single frequency observations from different days \\\\ $\\bullet$ : Value quoted for profile data with time resolution of 4.096 ms \\\\ $\\diamond$ : 3-sigma Non-detection limit for the null state, for profile data with time resolution of 20.048 ms\\\\ $\\ast$ : Flux estimated using standard GMRT system parameters using relative calibration procedure\\\\ $\\circ$ : Flux estimated using absolute calibration with a flux calibrator source\\\\ $\\odot$ : Mean flux is calculated up to 50\\% intensity lavel for MP and IP\\\\ \\end{minipage} \\end{table*} \\begin{figure} \\begin{center} \\includegraphics[angle=0, width=0.45\\textwidth]{bhaswati_fig1.ps} \\caption{Upper two panels plot the average profiles (of the active state) for PSR B0826$-$34 at 303 and 610 MHz. The third panel shows the 303 MHz profile, aligned to the 610 MHz profile after correction for the dispersion delay. MP, IP and OP refer to main pulse, inter-pulse and off pulse, respectively (as described in the text); MP$_{1}$ and MP$_{2}$ refer to the two peaks of the MP.} \\label{fig:dm_delay} \\end{center} \\end{figure} \\label {sec:Discussions} \\subsection {\\bf Determination of accurate DM value}                   \\label{sec:DM_Discussions} Table. \\ref{table:dm_list} compares our result for the DM value of PSR B0826$-$34 with other available values from the literature $-$ these cover a wide range, with the two most recent papers \\citep{Gupta_etal,Esamdin_etal} using values of 47 and 65.6 $pc/cm^{3}$. It is clear that, unlike most normal pulsars, DM determination for this pulsar is a difficult and trick exercise, mainly because the profile is quite complex, very wide and strongly evolving with frequency. Given our more accurate method of determining the DM, we believe that the either of the extreme values above are unlikely to be correct. We note that our DM result is closest to value obtained by \\cite{Hobbs_etal}, using multi epoch timing. The accuracy of both these methods depends on the choice of fiducial point. The advantage of our method of DM determination is that, the observations at a single epoch are self sufficient for obtaining the DM value at that epoch, without requiring any absolute timing information. This method can facilitate the DM determination of other complex and wide profile pulsars. \\subsection {\\bf Study of the pulsar in null state}                    \\label{null_Discussion} We see short duration nulls of a few pulses, which appear to be simultaneous over the frequency range of 303 to 610 MHz. In addition, this pulsar exhibits long durations of what is apparently a null state, which also appears to be broad band over this frequency range. However, we find no evidence for any weak emission in these long, apparently null states of the pulsar. Non-detection of any correlated emission structure in the null state profiles from the simultaneous dual frequency observations (at 303 and 610 MHz) are a clear proof of this.  For frequencies above (up to 1060 MHz) and below (down to 157 MHz) also, we see no evidence for any weak emission. Our best limits on this non-detection are comparable to, or a few times better, than the level of detection (2\\%) that \\cite{Esamdin_etal} report for the null state of this pulsar, at 1374 MHz. These comparisons are all with respect to the peak of the emission profile in the active state. We have also obtained absolute flux limits for the non-detection at various frequencies, which should be a useful comparison standard for any more sensitive studies in the future.  What could be the reason for the difference between our results and those of \\cite{Esamdin_etal}? It could be a case of frequency dependent nulling, as reported by \\cite{Bhat_etal} for PSR B1133$+$16 where the pulsar exhibits cases of short duration nulls at lower frequencies like 325, 610 and 1400 MHz, while there is visible emission at 4850 MHz. However we believe that this explanation is difficult, given that we see no evidence for emission till the frequency of 1060 MHz, coupled with the fact that the emission properties of this pulsar appear to be significantly correlated over an octave range in frequency. Hence, we are unable to reconcile our results with the results of \\cite{Esamdin_etal}. \\subsection {\\bf Study of subpulse emission}                           \\label{subpulse_Discussion} \\begin{table*} \\begin{minipage}{136mm} \\caption{\\bf Frequency dependence of profile width and $P_2$ for different pulsars.} \\label{table:avp_p2} \\begin{tabular}{l|c|c|c|c|c|c} \\hline Pulsar    & Frequency & Profile width          & $P_2$     & Comments            &  Reference  \\\\ &           & $\\Delta\\Phi ^\\diamond$ &           &                     &  \\\\ &  MHz      &  (\\degr)               & (\\degr)   &                     &  \\\\\\hline B0031$-$07& 60        & 75           &  20.6         & $P_2\\propto \\nu^{-0.05}$ &\\cite{Izvekova_etal}\\\\ & 102       & 65           &  19.8         & $\\Delta\\Phi \\propto \\nu^{-0.4}$ &             \\\\ & 406       & 40           & 18.7         &                                 &             \\\\\\hline B0031$-$07& 157       & 44$\\pm$5     & 18.9$\\pm$1.2 & $P_2$ evolution  &\\cite{Smits_etal}\\\\ A-mode    & 243       & 39.5$\\pm$1.5 & 24$\\pm$4     & does not follow    &         \\\\ & 325       & 39.4$\\pm$0.3 & 19.8$\\pm$1.0 & definite trend  &         \\\\ & 607       & 32$\\pm$4     & 21$\\pm$3     &                 &         \\\\ & 840       & 32.2$\\pm$0.8 & 10$\\pm$3     &                 &         \\\\ & 1167      & 25$\\pm$6     & 14$\\pm$5     &                 &         \\\\ & 4850      & 23.6$\\pm$0.6 & 14$\\pm$2     &                 &         \\\\\\hline B0031$-$07& 157       & 46.3$\\pm$1.4 & 19.2$\\pm$0.7 & $P_2$ decreases &\\cite{Smits_etal}\\\\ B-mode    & 243       & 42.3$\\pm$0.6 & 18.7$\\pm$1.6 & or remains unchanged  &         \\\\ & 325       & 40.3$\\pm$0.12& 18.7$\\pm$0.6 & with increasing      &         \\\\ & 607       & 39.3$\\pm$0.9 & 18.7$\\pm$0.6 & frequency            &         \\\\ & 840       & 43$\\pm$0.9   & 14$\\pm$1.6   &                      &         \\\\ & 1167      & 38$\\pm$5     & 12$\\pm$3     &                      &         \\\\\\hline B0320$+$39& 102       & 11           & 3.8          & $P_2 \\propto \\nu^{-0.15}$      &\\cite{Izvekova_etal}\\\\ & 406       & 8.5          & 3.1          & $\\Delta\\Phi \\propto \\nu^{-0.4}$&             \\\\\\hline B0809$+$74& 102       & 34.8         & 14.7         & see Sect.\\ref{subpulse_Discussion} &\\cite{Davies_etal} \\\\ & 406       & 26.0         & 10.8         &                 & \\\\ & 1412      & 29.8         & 8.6          &                 & \\\\ & 1720      & 27.2$^b$     & 8.1          &                 & \\cite{Bartel_etal_81a}\\\\\\hline B0826$-$34& 157       & 140$\\pm$4$^\\dagger$  &  23 $\\pm$0.8  &                                   &   This Paper        \\\\ & 303       & 132$\\pm$4$^\\dagger$  &  22.6$\\pm$0.8 & $\\Delta\\Phi \\propto \\nu^{-0.2}$   & This paper \\\\ & 318       & 134$\\pm$4$^\\dagger$  &  24.9$\\pm$0.8 & $P_2$ increases with               & \\cite{Gupta_etal} \\\\ & 610       & 116$\\pm$4$^\\dagger$  &  27.9$\\pm$0.8 & increasing frequency from         & This paper \\\\ & 645       & 112$\\pm$4$^\\dagger$  &  29$\\pm$2.0    &303 to 645 MHz                   &\\cite{Biggs_etal}\\\\ & 1060      & 117$\\pm$4$^\\dagger$  &  28.7$\\pm$0.8 &                                 &  This Paper         \\\\ & 1374      &  -                   &  27.5         &                                 &\\cite{Esamdin_etal} \\\\\\hline B1237$+$25& 430       & 16$^a$       & 10.6$\\pm$2   & $P_2 \\propto \\nu^{-0.16}$  &\\cite{Wolszczan_etal}\\\\ & 1700      & 14.5$^b$     & 9.1$\\pm$1.5  & $\\Delta\\Phi \\propto \\nu^{-0.07}$&          \\\\\\hline B2016$+$28& 102       & 20           & 8.5          & $P_2 \\propto \\nu^{-0.1}$       &\\cite{Izvekova_etal}\\\\ & 406       & 15           & 6.2          & $\\Delta\\Phi \\propto \\nu^{-0.2}$&                    \\\\ & 1412      & 15           & 5.5          &                                &               \\\\\\hline B2020$+$28& 430       & 17$^a$       & 28.3$\\pm$4   & $P_2$ increases with   &\\cite{Wolszczan_etal} \\\\ & 1700      & 16.2$^b$     & 51.4$\\pm$7   & frequency             &                     \\\\\\hline \\end{tabular} \\\\ $\\diamond$ : Profile width at 10\\% intensity level\\\\ a : average profile width at 400 MHz from \\cite{Gould_etal}\\\\ b : average profile width at 1600 MHz from \\cite{Gould_etal}\\\\ I,II : left and right part of MP respectively, see Fig.1 of \\cite{Biggs_etal}\\\\ $\\dagger$: separation between two components of main pulse\\\\ \\end{minipage} \\end{table*} (i) Pulse energy correlations : Our result that the fluctuations of the main pulse energies at 303 and 610 MHz are correlated up to 47\\% level, indicates that the emission process generating the single pulses has a high degree of correlation over a this octave bandwidth. What is equally noteworthy is that this level of correlation is pretty much the same as that at each of these individual frequencies, indicating that there is practically no loss of correlation over this range of frequency. Table.\\ref{table:cor_table} lists the correlation coefficient of pulse energies for different pulsars between two different observing frequencies, reported in the literature. All the pulsars listed there show significantly high correlation, supporting the case for the broad band nature of pulsar radio emission. At the same time, the significantly lower correlation between the main pulse and inter-pulse energies at 610 and 1060 MHz, implies that the physical processes responsible for the  MP and IP emission may be much more unrelated to each other.  The result that the intensity correlations are positive for large lags implies that there is some kind of memory in the underlying subpulse structure. That this memory is the longest for PSR B0826$-$34 (amongst all known cases), is quantified in column 4 of Table.\\ref{table:cor_table}. This is likely to be related to the large number of subpulses (or drift bands) that each main pulse energy estimate encompasses. The memory could then correspond to the time taken for the emission from a given subpulse spark column to drift through the main pulse window.  The other memory that the main pulse intensity variations exhibit is the intriguing feature of the peak at lag 13 seen in some of the correlation results. This could be an indication of the total time for rotation, around the magnetic axis, of the pattern of drifting subpulses $-$ the carousel rotation time. For example, in the 14 spark model presented by \\cite{Gupta_etal}, the drift rate is such that one subpulse drifts to a location close to that of the adjacent subpulse in one pulsar rotation, leading to a carousel rotation time of $P_4=14 P_1$. However, if we believe that there are 13 sparks, as indicated in the model given by \\cite{Esamdin_etal}, with the same drift rate condition as proposed by \\cite{Gupta_etal}, then we would have $P_4=13 P_1$, matching with the secondary peak feature seen in our data.  (ii) Study of subpulse drifting: PSR B0826$-$34 exhibits substantial variations of $P_2$ at different frequencies, usually reported at different observing epochs (Table.\\ref{table:avp_p2}). The different values of $P_2$ determined from the simultaneous data at 303 and 610 MHz strongly supports the claim that these are genuine variations with frequency, and not due to some time dependent phenomenon. The collected results in Table.\\ref{table:avp_p2} are consistent with a scenario where $P_2$ increases with frequency from 157 to 645 MHz, but then appears to saturate or even turn over at the higher frequencies like 1060 and 1374 MHz. For the low frequency range, we find that this frequency evolution ($P_2 \\propto \\nu^{0.3}$) is of opposite sense to that for the profile width, as measured by the separation between two components of the MP ($\\Delta\\Phi \\propto \\nu^{-0.2}$). Only two other pulsars in Table.\\ref{table:avp_p2} show such an abnormal behaviour : PSR B2020$+$28 shows a clear opposite sense (though only 2 frequency points are available) and PSR B0031$-$07 in the A-mode shows erratic increases and decreases of $P_2$ with increasing frequency.  For most pulsars, though both $P_2$ and the average profile width decrease with increasing frequency, the evolution of these two quantities with frequency does not exactly match (see Table.\\ref{table:avp_p2}). \\cite{Izvekova_etal} showed in detail that $P_2$ is, in general, less dependent on frequency, as compared to the profile width $-$ this is reflected in our Table.\\ref{table:avp_p2} as well. This has been explained \\citep{Gil_etal_96,Gil_etal} to be a natural consequence of the conal model of pulsar beams \\citep{Rankin_etal_93} $:$ the subpulse enhancement follows a narrow bundle of dipole field lines, while the mean enhancement corresponding to the peaks of mean profile components, is distributed over a cone (or a number of nested cones) of dipole field lines. Thus, the frequency dependence of subpulses and pulse widths should in general be different. However, these models do not appear to explain the opposite sense that we see in PSR B0826$-$34.  Another pulsar with anomalous frequency variations of profile width and $P_2$ is PSR B0809$+$74. $P_2$ being 1:8 times greater at 102.5 MHz than at 1720 MHz, for this pulsar, is explained by \\cite{Edwards_etal} as an artifact of the phase step, which is only present at high frequencies. Whereas \\cite{Rankin_etal_05} analysed the $P_2$ variations between 40 to 1400 MHz for this pulsar, and showed that they arise from incoherent superposition of two orthogonal modes of polarisation. It will be interesting to see if detailed polarisation studies of PSR B0826$-$34 can throw some light on the $P_2$ variations with frequency.  The other interesting result is the variation of $P_2$ with pulse longitude, where we confirm the behaviour first reported by \\cite{Gupta_etal} at 318 MHz (a variation from 21.5\\degr to 27\\degr within the MP), and also find some signs for a similar effect at 610 MHz.  Similarly, \\cite{Esamdin_etal} reported that at 1374 MHz, $P_2$ varies from 26.8\\degr to 28\\degr within the MP. \\cite{Gupta_etal} have provided an explanation of this in terms of the geometrical effect of the way our line-of-sight traverses the polar cap.  $\\Delta\\Phi_s$ determined in this work is compared to the result quoted by \\citep{Biggs_etal,Esamdin_etal} at 645 and 1374 MHz respectively in Table. \\ref{table:su_width}. $\\Delta\\Phi_s$ is observed to be frequency dependent, average $\\Delta\\Phi_s$ at 303 MHz is more than 610 MHz value (see Sect.\\ref{subpulse_study}). In addition, $\\Delta\\Phi_s$ is observed to be longitude dependent: $\\Delta\\Phi_s$ in region I is more than region II (Table. \\ref{table:su_width}). Similar longitude evolution of $\\Delta\\Phi_s$ within the MP is reported by \\cite{Biggs_etal} and \\cite{Esamdin_etal}.  To summarise, using results from high sensitivity single pulse observations simultaneously at 303 and 610 MHz, and individually at 157, 325, 610 and 1060 MHz, we have investigated different aspects of PSR B0826$-$34: (1) precise value of DM  (2) emission properties of the long duration null states and (3) simultaneous subpulses and details of the drifting behaviour. \\begin{table} \\caption{Frequency dependence of subpulse width ($\\Delta\\Phi_s$) of PSR B0826$-$34.} \\label{table:su_width} \\begin{tabular}{l|c|c|c|c|c|c|c} \\hline Frequency & \\multicolumn{4}{|c|}{$\\Delta\\Phi_s$} & Reference     \\\\ (MHz)   & \\multicolumn{4}{|c|}{(\\degr)}   &                \\\\ &  MP          & I         &  II       &   IP           &                \\\\ \\hline 157$^\\ddagger$&  15.8$^\\S$   & 22.1$^\\S$ & 10.2$^\\S$ &  -             & This paper               \\\\ 303$^\\dagger$ &  17.4$^\\S$   & 28.4$^\\S$ & 12.6$^\\S$ &  -             & This paper               \\\\ 610$^\\dagger$ &  14.2$^\\S$   & 22$^\\S$   & 12.6$^\\S$ & 12.6$^\\S$      & This paper               \\\\ 610$^\\ddagger$&  14.2$^\\S$   & 22.9$^\\S$ & 12.6$^\\S$ & 12.6$^\\S$      & This paper               \\\\ 645           &  -           & 14$^\\ast$ & 5.9$^\\ast$& 4  $^\\ast$     & \\cite{Biggs_etal}        \\\\ 1060$^\\ddagger$& 14.2$^\\S$   & 20.5$^\\S$ & 12.6$^\\S$ & 17.4$^\\S$      & This paper               \\\\ 1374           & 13.1$\\pm$1.9&    -      &    -      & 19$\\pm$3.5     & \\cite{Esamdin_etal}      \\\\\\hline \\end{tabular} \\\\ $\\dagger$ : From simultaneous dual frequency observing session \\\\ $\\ddagger$ : From single frequency observing session \\\\ I : 118\\degr to 197\\degr  pulse longitude\\\\ II : 197\\degr to 276\\degr  pulse longitude\\\\ $\\ast$ : Lag corresponding to half width at half maximum (HWHM) values\\\\ $\\S$ : Error associated in determination of $\\Delta\\Phi_s$ is $\\pm$1.6\\degr\\\\ \\end{table}"
    },
    "0711/0711.2379_arXiv.txt": {
        "abstract": "The prompt $\\gamma$-ray emission in $\\gamma$-ray bursts is believed to be produced by internal shocks within a relativistic unsteady outflow. The recent detection of prompt optical emission accompanying the prompt $\\gamma$-ray emission appears to be inconsistent with this model since the out flowing plasma is expected to be highly optically thick to optical photons. We show here that fluctuations in flow properties on short, $\\sim1$~ms, time scale, which drive the $\\gamma$-ray producing collisions at small radii, are expected to lead to \"residual\" collisions at much larger radii, where the optical depth to optical photons is low. The late residual collisions naturally account for the relatively bright optical emission. The apparent simultaneity of $\\gamma$-ray and optical emission is due to the highly relativistic speed with which the plasma expands. Residual collisions may also account for the X-ray emission during the early \"steep decline\" phase, where the radius is inferred to be larger than the $\\gamma$-ray emission radius. Finally, we point out that inverse-Compton emission from residual collisions at large radii is expected to contribute significantly to the emission at high energy, and may therefore \"smear\" the pair production spectral cut-off. ",
        "introduction": "In the leading model for $\\gamma$-ray bursts (GRBs), the energy source is a compact object that drives a relativistic unsteady outflow with fluctuating Lorentz factor. Internal shocks within the outflow dissipate the bulk kinetic energy and produce $\\gamma$-rays \\citep{RM94,PX94}. A substantial fraction of the outflow kinetic energy may be dissipated in this model outside the photosphere, allowing one to account for the non-thermal spectra and for the complicated light curves of GRBs \\citep{Kobayashi97}. The internal shocks are expected to generate/amplify magnetic fields and to accelerate electrons, which produce MeV $\\gamma$-rays by synchrotron emission \\citep[For review of internal shock models, see][]{Waxman rev}.   Due to the relatively short duration of the prompt $\\gamma$-ray emission, $T\\sim1$~to~$10^2$~s, the observation of long-wavelength (optical) prompt emission is a difficult task. GRB 990123 was the first event for which optical emission was detected during the burst \\citep{Akerlof99}. Today, thanks to the rapid localization of GRBs by the {\\it Swift} satellite \\citep[see][for recent review]{Zhang07}, a larger number of optical (and longer wavelengths) observations are carried out during the bursting phase \\citep[e.g.][]{Blake05,Vestrand05,Vestrand06,Yost07}.  As shown in \\S~\\ref{sec:absorption} \\citep[see also][]{LZ}, during the emission of prompt $\\gamma$-rays the plasma is expected to be optically thick to optical photons due to strong synchrotron self-absorption. This appears to be inconsistent with the detection of bright optical emission accompanying $\\gamma$-ray emission. We point out here that the internal collisions at small radii, which produce the $\\gamma$-ray emission, are expected to lead to \"residual\" collisions at much larger radii, where the optical depth to optical photons is low. These late residual collisions may naturally account for the relatively bright optical emission. The apparent simultaneity of $\\gamma$-ray and optical emission is due to the highly relativistic speed with which the plasma expands. The time delay between $\\gamma$-ray and optical emission is expected to be shorter than a second, too short to be identified by current optical observations which usually have lower temporal resolution. We discuss in \\S \\ref{sec:large radii} residual collisions in unsteady outflows, and derive the long-wavelength emission they are expected to produce at large radii. The implications of our results are discussed in \\S~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  We have shown that late residual collisions, that occur at radii much larger than those where $\\gamma$-ray producing collisions take place, may naturally account for the observed strong optical emission accompanying the prompt GRB. Internal collisions at small radii reduce the variance of colliding shell velocities. As a result, the energy available for radiation at large radii and the characteristic frequency of radiated photons decrease with radius. We find that one may expect optical to $\\gamma$-ray energy ratio $\\sim10^{-4}$, with large burst-to-burst scatter (see fig.~\\ref{fig}, eq.~\\ref{eq:F_ratio} and discussion at the end of \\S~\\ref{sec:rad}). This is consistent with the results of \\cite{Yost07}, who find that during the prompt emission of GRBs the spectral indices between the optical and the $\\gamma$-ray bands are in the range of $0<\\beta_{{\\rm op}-\\gamma}<0.5$, corresponding to $F_{\\nu_{\\rm op}}/F_{\\nu_\\gamma}\\sim1-10^3$ and implying that the optical emission is only a small fraction, $\\sim10^{-6}-10^{-3}$, of the total emitted energy.  Although the optical emission is produced at large radii, where synchrotron self-absorption is avoided, the expected time delay between $\\gamma$-ray and optical emission, $\\sim0.1$~s (see eq.~\\ref{eq:delay time}), is shorter than the characteristic temporal resolution of the optical observations, which is a few seconds \\citep[e.g.][]{TZ06}. Thus, optical and $\\gamma$-ray emission may appear to be simultaneous. However, better temporal resolution may allow one to detect a systematic time delay between the two wave bands. In addition, one would expect larger observed variability timescales at longer wavelengths, $t_{\\rm var, op}\\sim\\tau_{\\rm delay}$.  \\citet{Wei07} has suggested that optical emission may be generated by strong internal shocks at radii $R/c>10^6$~s, driven by shells emitted with a large time delay, $\\sim10$~s, following those producing the main $\\gamma$-ray emission \\citep[see also][]{Fan05}. Our model is quite different. We show that optical emission is naturally expected to arise, without postulating the existence of delayed shells, by residual collisions at $R/c\\sim10^4$~s, in which the characteristic emitted photon frequency is low, $h\\nu\\sim1$~eV, due to the reduction of the Lorentz factor variance in the flow (rather than by the large radius $R/c>10^6$~s). Moreover, we have shown that the optical luminosity can be estimated from the burst $\\gamma$-ray properties, and that it is consistent with the observations.  The energy released in residual collisions of the relativistic outflow is large. In fact, it would overproduce the optical emission if all the energy is released in the optical band. In the models discussed here, only a small fraction of the energy, $\\sim10^{-2}$, is released as synchrotron radiation, since electrons accelerated in residual collisions cool mainly by IC scattering of the prompt GRB $\\gamma$-rays. This has some important implications to observations at high energy, $> 100$~MeV. Such observations are expected to be useful in determining the bulk Lorentz factor of GRB outflows and the size of the emitting region by detecting the high energy cutoff due to $\\gamma\\gamma$ absorption \\citep[e.g.,][]{Baring00,Lithwick01,LZ03}. The identification of this cutoff may be complicated in the the presence of strong high energy emission from residual collisions, that take place at large radii where the $\\gamma\\gamma$ optical depth is reduced.  A comment is in place here regarding some recent constraints on the size of the GRB emission region, which were inferred from the early X-ray steep decay. Assuming that the steep decay arises from emission by plasma lying away from our line of sight, large radii were inferred, $R_{\\rm em}>2t_{\\rm decay}c/\\theta_j^2\\ga6\\times10^{13}$cm for $t_{\\rm decay}\\ga10^2$s and $\\theta_j\\la0.3$ \\citep{Lazz06,Lyutikov06,Kumar07}. Such radii are larger than typically predicted in internal shock models. It should be realized, however, that if the off-the-line-of-sight emission explanation is adopted, the emission during this phase should peak below the X-ray band, and should therefore arise in a region different than that where $\\gamma$-rays are produced. This is due to the fact that the flat X-ray spectrum ($F_\\nu\\propto \\nu^{0}$) seen in GRB spectra would imply a light curve decay $\\propto t^{-2}$ \\citep[e.g.,][]{Kumar00}, much shallower than observed \\citep[e.g.,][]{Tagli05}. In fact, the spectra during the steep decay are soft, suggesting that indeed the energy peak is below the X-ray band. Thus, if the steep decay is due to off-the-line-of-sight emission, it should originate from a region lying at a larger radius than that where $\\gamma$-rays are produced, producing emission that peaks below the X-ray band. Such emission may be produced, e.g, by residual collisions."
    },
    "0711/0711.3978_arXiv.txt": {
        "abstract": "{ We present results of the photometric campaign for planetary and low-luminosity object transits conducted by the OGLE survey in 2005 season (Campaign \\#5). About twenty most promising candidates discovered in these data were subsequently verified spectroscopically with the VLT/FLAMES spectrograph.  One of the candidates, OGLE-TR-211, reveals clear changes of radial velocity with small amplitude of 82 m/sec, varying in phase with photometric transit ephemeris. Thus, we confirm the planetary nature of the OGLE-TR-211 system. Follow-up precise photometry of OGLE-TR-211 with VLT/FORS together with radial velocity spectroscopy supplemented with high resolution, high S/N VLT/UVES spectra allowed us to derive parameters of the planet and host star. OGLE-TR-211b is a hot Jupiter orbiting a F7-8 spectral type dwarf star with the period of 3.68 days. The mass of the planet is equal to $1.03\\pm0.20$~M$_{\\rm Jup}$ while its radius 1.36 $^{+0.18}_{-0.09}$~R$_{\\rm Jup}$. The radius is about 20\\% larger than the typical radius of hot Jupiters of similar mass. OGLE-TR-211b is, then, another example of inflated hot Jupiters -- a small group of seven exoplanets with large radii and unusually small densities -- objects being a challenge to the current models of exoplanets.} ",
        "introduction": "The discovery of the first transit of an extrasolar planet in HD209458 \\citep{ch00,he00} opened a new era in the extrasolar planet field -- the epoch of searches for extrasolar planets with large scale photometric surveys. While this channel of finding exoplanets developed relatively slowly in the first couple of years after that discovery, the last two years brought a fast acceleration due to maturing of the search methods. At present about 25 transiting planets are known (cf. {\\it http://obswww.unige.ch/~pont/TRANSITS.htm)} constituting about 10\\% of all known exoplanets.  Transiting planets play especially important role among the known exoplanets. These are the only objects for which the most important parameters like radius, mass and density can be precisely derived from observations providing direct comparison with models and, thus, allowing better understanding of the exoplanet structure and evolution. Also a large variety of follow-up observations can be performed on transiting system allowing, for example, studies of planetary atmospheres, search for other planetary companions etc.  Transiting planets are discovered via two channels: photometric follow-up of spectroscopically found exoplanets and classical transit method approach -- large scale photometric surveys providing transiting candidates that have to be then verified and confirmed spectroscopically. Both approaches have been successful in providing examples of very large diversity among transiting planets -- regular hot Jupiters, very hot Jupiters on extremely short 1--2 day period orbits, inflated objects with radii much larger than expected from modeling or small Neptune-sized objects. Large number of new extrasolar transiting planets are expected to be discovered in the next couple of years from the space missions like Corot or Kepler, in particular small size planets undetectable in ground-based searches.  The Optical Gravitational Lensing Experiment (OGLE) was the first successful photometric survey that discovered large number of transiting candidates that were subsequently verified spectroscopically \\citep{ud02a,ud02b,ud02c,ud03}. Five of these candidates, OGLE-TR-56, OGLE-TR-113, OGLE-TR-132, OGLE-TR-111 and OGLE-TR-10 turned out to be extrasolar planetary systems \\citep{ko03,bo04,ko04,po04,bo05,ko05}, the first extrasolar planets discovered with the transit method. Moreover, two planet-sized stars of lowest known masses were also found \\citep{po05b,po06}. Beside that the OGLE photometric transit campaigns provided a huge observational material allowing better understanding the problems of photometric transit searches, data systematics etc. \\citep{go06,pz06}.  OGLE transit campaigns have been conducted every year since 2001. Results of the spectroscopic follow-up of candidates discovered during the Campaign \\#1 (2001) and \\#2 (2002) were published by \\citet{bo05} and \\citet{po05a}, respectively. Recently, analysis of the near threshold candidates from the Campaign \\#2 led to the discovery of the sixth OGLE transiting exoplanet, OGLR-TR-182 \\citep{po07}.  In this paper we present results of the OGLE Campaign \\#5 conducted in 2005 and spectroscopic follow-up of the best transiting candidates detected in these data. A new transiting extrasolar planetary system OGLE-TR-211 was found among them. In the following Sections we provide the details of the photometric and spectroscopic observations of OGLE new candidates and derive parameters of a planetary companion in OGLE-TR-211 system. ",
        "conclusions": "Parameters of the transiting companion of OGLE-TR-211 host star indicate that it is a planetary object belonging to the hot Jupiter class of extrasolar planets. It brings the total number of extrasolar planets discovered by the OGLE survey to seven -- significantly lowering the discrepancy between the number of typical hot Jupiters and shortest period very hot Jupiters in the OGLE sample of extrasolar transiting planets.  OGLE-TR-211b orbits an F7-8 main sequence star. The system parameters closely resemble that of the recently discovered system HAT-P-6 \\citep{no07} with the only difference that the host star in OGLE-TR-211 is more metal abundant by $\\approx0.2$~dex. The position of the planet on the mass-orbital period diagram is similar to the position of other transiting exoplanets of similar mass and period supporting the relation noted first by \\citet{ma05}.  While the mass of a new hot Jupiter OGLE-TR-211b is very similar to the mass of another hot Jupiter OGLE-TR-182b recently found by \\citet{po07} its radius is about 20\\% larger. The mean density of OGLE-TR-211b is about 0.5 g/cm$^3$, placing it among a group of seven objects with the lowest densities: Tres-4, HD209458, HAT-P-1, WASP-1, HAT-P-4 and HAT-P-6 \\citep{kv07}. All these objects have inflated radii compared to other transiting hot Jupiters of similar mass. Such planets pose a problem to the planetary models and are crucial objects for testing the proposed scenarios explaining ``bloated'' hot Jupiters \\citep{bu07,gu06}. It is worth noting that the metallicity of the OGLE-TR-211 host star falls in the middle of the range of metallicities of ``bloated'' Jupiters host stars (recently increased significantly by the HAT-P-6 system) suggesting that metallicity cannot be the only parameter responsible for inflating the radii of hot Jupiters.  \\begin{figure} \\includegraphics[width=9cm, bb=75 392  505 660]{fig4.ps}  \\caption{Mass-radius diagram for transiting hot Jupiters. Closed symbols indicate host stars with $T_{\\rm eq}>1500$ K. OGLE-TR-211 is   indicated by the circled dot. The lines show the models of \\citet{gu05} in two extreme cases:  cold gas planet with a $15 M_\\oplus$ core  and   $T_{\\rm eq}=2000$~K planet without core. The five large planets with   radius error bars above the upper models are all strongly irradiated.} \\end{figure}  On the other hand, the characteristics of the OGLE-TR-211 system reinforce the association of ``bloated'' close-in giant planets with strong irradiation from the host star. Figure~4 shows the mass-radius diagram for the presently known transiting planets, classified according to their equilibrium temperature. \\footnote{$T_{\\rm eq} = T_*((1-A)/4\\,F)^{1/4}(R_*/a)^{1/2}$ using for all planets a Bond albedo $A=0.3$ and a heat redistribution factor $F=1/2$.}  All anomalously large planets have equilibrium temperatures hotter than 1500~K.  This would require an increasingly unlikely coincidence if the anomalously large radii were explained by dynamical effects such as obliquity tide or eccentricity pumping by an unseen companion \\citep{wi05}. It rather tends to favor  explanations directly linking the excess of radius with the incident flux like such as  proposed by \\citet{gu06} (that a fraction of  the incident flux energy is converted to mechanical energy by a yet undetermined physical process) -- over explanations with a more  indirect causal relation, such as that put forward by \\citet{bu07} and \\citet{ch07} invoking increased opacity or internal density gradients.  As we already noted in Section~4.2 the mean level of radial velocity variation of the OGLE-TR-211 host star was shifted by about 180 m/sec during the last April 2007 run. Although the real nature of this shift is unknown at this stage the tempting explanation is that it is caused by the presence of additional companion, possibly of planetary or stellar nature, in the OGLE-TR-211 system. If confirmed by  additional follow-up observations, OGLE-TR-211 would be the first multi-planetary system with a transiting planet or the first moderately-wide binary system hosting a transiting hot Jupiter.  Also, further additional high accuracy photometric follow-up observations of the OGLE-TR-211 transits are necessary for better constraining the inclination and radius of the planet. Although OGLE-TR-211, similarly to other OGLE transiting planets, is too faint for the currently available facilities to carry out some follow-up observations like IR photometry, the precise long term transit timing allowing search for another planetary objects in the OGLE-TR-211 system is feasible."
    },
    "0711/0711.4844_arXiv.txt": {
        "abstract": "Repeat scans by the Sloan Digital Sky Survey (SDSS) of a 278-deg$^{2}$ stripe along the Celestial equator have yielded an average of over 10 observations each for nearly 8,000 spectroscopically confirmed quasars.  Over 2500 of these quasars are in the redshift range such that the C~{\\sc iv} $\\lambda$ 1549 emission line is visible in the SDSS spectrum.  Utilising the width of these C~{\\sc iv} lines and the luminosity of the nearby continuum, we estimate black hole masses for these objects.  In an effort to isolate the effects of black hole mass and luminosity on the photometric variability of our dataset, we create several subsamples by binning in these two physical parameters.  By comparing the ensemble structure functions of the quasars in these bins, we are able to reproduce the well-known anticorrelation between luminosity and variability, now showing that this anticorrelation is independent of the black hole mass.  In addition, we find a correlation between variability and the mass of the central black hole.  By combining these two relations, we identify the Eddington ratio as a possible driver of quasar variability, most likely due to differences in accretion efficiency. ",
        "introduction": "}   The luminosities of quasars and other active galactic nuclei (AGN) have been observed to vary on time-scales from hours to decades, and from X-ray to radio wavelengths. The majority of quasars exhibit continuum variability on the order of 20 per cent on time-scales of months to years \\citep[e.g.,][]{hook94,vandenberk04}.  In fact, variability has long been used as a selection criterion in creating quasar samples from photometric data \\citep[e.g.,][]{ koo86, ivezic04, rengstorf04}. Many simple correlations between photometric variability and various physical parameters have been known for decades.  These relationships are summarised by \\citet{helfand01} and \\citet{giveon99}.  Numerous studies \\citep[e.g.,][]{hawkins02,devries03} have shown variability to correlate with time lag.  Anti-correlations have been found between variability and luminosity \\citep[e.g.,][]{uomoto76,cristiani97} and wavelength \\citep[e.g.,][] {giveon99, trevese01}.   \\citet[hereafter VB04]{vandenberk04}, using a sample of $\\sim25,000$ quasars, confirmed these known correlations, and parametrized relationships between variability and time lag, luminosity, rest-frame wavelength and redshift.  Our understanding of the physics of the central black hole in quasars and active galactic nuclei (AGN) has long been tied to the variability of the quasar's luminosity. Intra-day variability in X-ray and optical light \\citep[see, e.g.,][]{terrell67, kinman68, boller97} point towards a compact object, specifically a supermassive black hole, at the centre of an AGN.  More recently, reverberation mapping techniques have been used to determine the radius of the broad line region and, indirectly, to measure the mass of the central black hole \\citep{blandford82, peterson93,kaspi00}.  Recently, \\citet{wold07}, presented evidence suggesting that the photometric variability of quasars is linked to the mass of the central black hole.  Strength of variability for $\\sim$ 100 quasars was approximated by finding the greatest single-epoch $R$-band deviation from the mean, by using optical light curves from the QUEST1 survey \\citep{rengstorf04}.  Black hole masses were estimated from Sloan Digital Sky Survey \\citep[SDSS;][]{york00} spectra using the 5100-\\AA\\ continuum luminosity and H$_{\\beta}$ line width, as calibrated by \\citet{vestergaard06}.  However, while a clear correlation between black hole mass and variability was found for time-scales greater than 100 days, \\citet{wold07} were unable to reproduce the well-known inverse relationship between luminosity and variability for their sample as a whole.  This is likely due to the redshifting of the blue (and more variable) portion of the spectrum into the R-band at higher redshifts; this causes the high-luminosity quasars visible at higher redshifts to appear more variable than one might expect, based on previous published results, such as \\citet{vandenberk04}.  With so few objects, it is difficult to truly isolate the dependence of  variability upon black hole mass, given the correlation of mass with luminosity, which is in turn tied to redshift and wavelength.   \\citet{wold07} do report a correlation between black hole mass and variability at constant luminosity; these intriguing results might be more convincing, however, if the variability-luminosity dependence were in line with expectations, or if the sample were larger.  With this in mind, we have examined the variability properties of a much larger sample of quasars from the Equatorial Stripe (see \\S\\,\\ref{south}) of the SDSS.  With this sample, we are able to reproduce all of the well-established dependences of photometric variability, including the inverse correlation with luminosity, as well as the recently measured correlation with black hole mass.  We briefly describe the quasar sample in \\S\\,\\ref{dataset} and the statistics used to measure variability in \\S\\,\\ref{variability}.  In \\S\\,\\ref{mbhlumvar}, we describe the methods used to estimate black hole masses, including the continuum- and line-fitting techniques, as well as the dependence of variability on black hole mass and luminosity.  Finally, we interpret our results in terms of physical models for quasars in \\S\\,\\ref{discussion}, and we conclude in \\S\\,\\ref{conclusions}.  Throughout the paper we assume the standard concordance cosmology with parameter values $\\Omega_\\Lambda = 0.7, \\Omega_{M} = 0.3,$ and $H_{0}=70\\ $km s$^{-1}$ Mpc$^{-1}$, consistent with \\citet{spergel07}. ",
        "conclusions": "In this paper, we have studied the ensemble variability properties of almost 8,000 spectroscopically identified quasars from the Sloan Digital Sky Survey Equatorial Stripe.  These objects have been observed an average of over ten times each.  By using their  C~{\\sc iv}  line dispersions and nearby continuum luminosities, we have estimated black hole masses for approximately 2,500 of these quasars.  We have binned these quasars in luminosity and black hole mass and examined the variability properties of the quasars in each bin.  We have been able to:  (1) Reproduce the well-known anticorrelation between luminosity and variability, and  (2) Detect a correlation between variability and black hole mass.  By combining (1) and (2), it appears that variability is inversely related to the Eddington ratio in quasars.  This points to variability being related to the quasar's accretion efficiency.  Given that the relation with black hole mass is more evident at longer time lags, we believe future studies involving longer time baselines will shed more light on this new result.  B. C. W. and R. J. B. would like to acknowledge support from Microsoft Research, the University of Illinois, and NASA through grants NNG06GH156 and NB 2006-02049. The authors made extensive use of the storage and computing facilities at the National Center for Supercomputing Applications and thank the technical staff for their assistance in enabling this work.  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 {\\tt 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."
    },
    "0711/0711.3463_arXiv.txt": {
        "abstract": "{The cosmic microwave background provides an image of the Universe 0.4 million years after the Big Bang, when atomic hydrogen formed out of free electrons and protons.  One of the primary goals of observational cosmology is to obtain follow-up images of the Universe during the epoch of reionization, hundreds of millions of years later, when cosmic hydrogen was ionized once again by the UV photons emitted from the first galaxies.  To achieve this goal, new observatories are being constructed, including low-frequency radio arrays capable of mapping cosmic hydrogen through its redshifted 21cm emission, as well as imagers of the first galaxies such as the {\\it James Webb Space Telescope (JWST)} and large aperture ground-based telescopes.  The construction of these observatories is being motivated by a rapidly growing body of theoretical work.  Numerical simulations of reionization are starting to achieve the dynamical range required to resolve galactic sources across the scale of hundreds of comoving Mpc, larger than the biggest ionized regions.  } ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.3180_arXiv.txt": {
        "abstract": "We test the present expansion of the universe using supernova type Ia data without making any assumptions about the matter and energy content of the universe or about the parameterization of the deceleration parameter. We assume the cosmological principle to apply in a strict sense. The result strongly depends on the data set, the light-curve fitting method and the calibration of the absolute magnitude used for the test, indicating strong systematic errors. Nevertheless, in a spatially flat universe there is at least a 5$\\sigma$ evidence for acceleration which drops to 1.8$\\sigma$ in an open universe. ",
        "introduction": "The first evidence for an accelerated expansion of the universe at the present epoch was presented in the late 1990s by Riess et al.~\\cite{riess98} and Perlmutter et al.~\\cite{perlmutter}. Fitting a cosmological model including a cosmological constant matched the supernovae type Ia of the considered data sets significantly better than a model without such a constant. In the following years, a lot of effort was put into getting larger supernova (SN) sets of increasing quality \\cite{jha,snls,gold07,essence}. On the other hand, a variety of cosmological models have been developed that could be fit to the newest data sets. Those models are characterized by certain parameters which are used to derive a redshift-luminosity relation that can be compared to the observed values. The problem with this so-called dynamical approach is that it is impossible to test without assumptions on the matter and energy content of the universe if there really is a phase of acceleration in the expansion history of the universe.  Some authors tried to avoid this problem by taking a kinematical approach, i.e. they only considered the scale factor $a$ and its derivatives, such as the deceleration parameter $q(z)$, without using any model specific density parameters or a dark energy equation of state. The first kinematical analysis of SN data was done by Turner and Riess \\cite{turner} who considered averaged values $q_1$ for redshift $z<z_1$ and $q_2$ for $z>z_1$ concluding that a present acceleration and a past deceleration is favoured by the data. Other authors tested a variety of special parameterizations of $q(z)$ \\cite{gold04,elgaroy} or used  $a(t)$ \\cite{wang} or the Hubble rate $H(z)$ \\cite{john}. Instead of considering a special parameterization, a more general approach has been done by Shapiro and Turner \\cite{shapiro} who expanded the deceleration parameter $q$ into principal components. Rapetti et al.~\\cite{rapetti} expanded the jerk parameter $j$ into a series of orthonormal functions. However, one has to be careful when doing a series expansion in $z$ as SNe with a redshift larger than 1 are not within the radius of convergence. This problem can be solved by reparameterizing redshift \\cite{cattoen}.  In the present work we will neither make any assumptions about the content of the universe, nor about the parameterization of $q(z)$ or another kinematical quantity. Moreover, we do not need to assume the validity of Einstein's equations. Instead, we will only ask the question if the hypothesis holds that the universe never expanded accelerated between the time when the light was emitted from a SN and today. Our assumption for the test is that the universe is isotropic and homogeneous. This is certainly not true for the real universe in a strict sense. However, we follow the standard approach in assuming that the cosmic structure does not modify the observed SN magnitudes and redshifts apart from random peculiar motion. The basic idea of this analysis has already been presented by Visser \\cite{visser} and has been applied to SN data by different groups \\cite{santos, gong}. But while Santos et al.~\\cite{santos} made some mistakes in their analysis (that we will discuss later in section \\ref{concl}) and Gong et al.~\\cite{gong} only state that accelerated expansion is ``evident'', we are able to give a quantitative value for this evidence. Additionally, we study the size of systematic effects.  For the fit one usually marginalizes over a function of the absolute magnitude $M$ and the Hubble constant $H_0$ because these two values cannot be determined independently by only considering SNe. Marginalization is not suitable for our analysis because in order to do so a special cosmological model or at least a parametization of $q(z)$ has to be inserted, which is what we want to avoid. But the SNe can be calibrated by cepheid meassurements and thus the values of $M$ and $H_0$ are determined. Using this additional information, no marginalization is needed. As there is still a controversy \\cite{jackson} about the appropriate calibration method, we will take two quite different results of calibration into account \\cite{riess, sandage}.  In order to test the robustness of our analysis, we consider two different SN data sets (the 2007 Gold sample \\cite{gold07} and the ESSENCE set \\cite{essence}), where the data of the ESSENCE set are once obtained by the multicolour light-curve shape (MLCS2k2) \\cite{jha} fitting method and once by the spectral adaptive light-curve template (SALT) \\cite{guy}. Both calibration methods are applied to each set. The analysis shows that in all cases the data indicate an accelerated expansion. But the confidence level at which acceleration can be stated strongly depends on the data set, the fitting method and the calibration. We begin our analysis by assuming a flat universe, but later also consider the cases of an open and a closed universe. ",
        "conclusions": "We tested the cosmic expansion without specifying any density parameters or parameterizing kinematical quantities. This was not possible without assuming a certain calibration of $M$ and $H_0$. Therefore, we considered two very different calibrations. We also considered two different data sets and two different light curve fitters and varied the peculiar velocity dispersion. Using single SNe for the test has already given a clear indication of acceleration in a flat universe although a few SNe strongly favour deceleration. We find large systematic effects already at the level of individual SNe.  Similar analyses have already been done by other groups \\cite{santos, gong}. But note that the work of Santos et al.~\\cite{santos} exhibits two major shortcomings. First, they did not calibrate the SN data consistently. They took a certain value for $H_0$ but kept the arbitrary value of the absolute magnitude $M$ given in the set which led to distance moduli that are too high. Thus, they concluded erroneously that there was a recent phase of super-acceleration. The second problem is that they did not realize that due to the integration over redshift this method is not suitable to determine the transition redshift $z_{\\mbox{\\scriptsize t}}$ between deceleration and acceleration. Thus, their value of $z_{\\mbox{\\scriptsize t}}$ does not hold. The shortcoming of the work of Gong et al.~\\cite{gong} is that they have not applied any statistics at all.  Instead of only considering single SNe, we obtained a more significant result by averaging over the SN data of each set. Here it is justified to discard nearby SNe in order to decrease systematics from using different telescopes or from local effects. We argue that we reject the null hypothesis of no acceleration for a spatially flat, homogeneous and isotropic universe at high confidence ($>5\\sigma$). However, the results show large differences for each data set, calibration and light-curve fitter. E.g. changing the calibration from Sandage to Riess calibration in a flat universe using ESSENCE(MLCS2k2) increases the evidence for acceleration from 5.2$\\sigma$ to 11.9$\\sigma$. For the two different light-curve fitters applied to the ESSENCE SNe using the Riess calibration, we get values of 11.9$\\sigma$ and 17.0$\\sigma$.  Wood-Vasey et al.~argue that the differences due to different light curve fitters are not important as they disappear when marginalization is applied \\cite{essence}. This is true if the data are only used to determine the parameter values of a certain cosmological model as in this case no calibration of $M$ and $H_0$ is needed. But nevertheless, the fitters should give the same result for a given absolute magnitude $M$. If this is not the case, a systematic error in the determination of the apparent magnitude $m$ is introduced by at least one of the fitters. Systematics due to different fitters have also been described in \\cite{conley} and those within the Gold sample in \\cite{jassal1,jassal2,nesseris}.  The different results obtained by the two calibrations are very surprising. Our test only depends on the reduced distance modulus $\\mathcal{M}=M-5\\log H_0+25$ and not on the absolute magnitude or the Hubble constant individually. (Note that this is the quantity which is marginalized when fitting cosmological parameters.) Starting from different calibrations $M_1$ and $M_2$ leads to different values of $H_0$, determined by observation of SNe, in such a way that $\\mathcal{M}$ should in principle be the same for both values of $M$. Thus our test should not depend on the calibration. However, the fact that the calibration changes our result significantly can be explained as follows: $H_0$ is determined by observing SNe. Riess et al.~and Sandage et al.~use different sets of SNe and different fitting methods for this determination. The systematic errors due to different sets and fitters then of course also influence the result for the Hubble constant leading to different values $\\mathcal{M}_1$ and $\\mathcal{M}_2$.  For a conservative conclusion, we need to take the set that gives the weakest evidence for acceleration, namely ESSENCE (MLCS2k2) in the Sandage calibration. Using this set, accelerated expansion can be stated at 5$\\sigma$ if we assume a spatially flat or closed universe. In an open universe the evidence is much weaker, namely 1.8$\\sigma$. These results only hold if the correct analysis of SNe is somewhere in the range of the cases we considered here. But as we observe enormous systematic effects in our test, we cannot be sure of this assumption. Thus, it is a major issue to better understand how SNe have to be observed and analyzed.  Remember that results of this work are only valid if the universe is homogeneous and isotropic. The situation changes dramatically if we give up those assumptions. There exist enormously large structures in the universe, an example is the Sloan Great Wall with an extension of $\\sim$400 Mpc \\cite{gott}. Besides, there are also big voids and superclusters on a 100 Mpc scale. This roughness of the universe could spoil the validity of equation (1). Indeed it is possible to construct inhomogeneous models that can describe the observational data without the need of an average acceleration \\cite{celerier,havard,enqvist,ishak}.  Some observational evidence for inhomogeneity and anisotropy in SN Hubble diagrams has recently been presented by \\cite{dyer,weinhorst}, probably due to large scale bulk motion and perhaps systematic effects. A non-trivial scale dependence of the Hubble rate due to the so-called cosmological backreaction \\cite{buchert,schwarz,rasanen,kolb,vanderveld,wiltshire} has been shown in \\cite{li}. This effect of averaging can also mimic curvature effects \\cite{li1} and thus strongly influence the reconstruction of the dark energy equation of state, as has been shown in \\cite{bassett}.  Thus we think a major task must be to establish the acceleration of the universe independently of the assumption of strict homogeneity and isotropy.  \\ack We thank Dragan Huterer, Thanu Padmanabhan, Aleksandar Raki\\'{c}, Adam Riess and Bastian Weinhorst for discussions, comments and references to the literature. This work is supported by the DFG under grant GRK 881."
    },
    "0711/0711.4305_arXiv.txt": {
        "abstract": "The 2 Micron All-Sky Survey (2MASS) Tully-Fisher Survey (2MTF) aims to measure Tully-Fisher (TF) distances to all bright inclined spirals in the 2MASS Redshift Survey (2MRS). Essential to this project is a universal calibration of the TF relation in the 2MASS J (1.2 $\\mu$m), H (1.6 $\\mu$m) and K (2.2 $\\mu$m) bands. We present the first bias corrected or universal TF template in these bands. We find that the slope of the TF relation becomes steeper as the wavelength increases being close to $L\\propto v^4$ in K-band and $L\\propto v^{3.6}$ in J and H-bands. We also investigate the dependence on galaxy morphology showing that in all three bands the relation is steeper for later type spirals which also have a dimmer TF zeropoint than earlier type spirals. We correct the final relation to that for Sc galaxies. Finally we study the scatter from the TF relation fitting for a width dependent intrinsic scatter which is not found to vary significantly with wavelength. ",
        "introduction": "The 2 Micron All-Sky Survey (2MASS) Tully-Fisher Survey (2MTF) is making use of existing high quality rotations widths, new HI widths and 2MASS photometry to measure Tully-Fisher (TF) distances for all bright inclined spirals in the 2MASS Redshift Survey (2MRS; Huchra \\etal ~2005). Previous peculiar velocity surveys have struggled to meet their potential because of large errors on individual measurements, poor statistics and uneven sky coverage. A survey based on 2MRS will provide both better statistics and greatly reduce the impact of the ZoA providing a {\\it qualitatively} better sample for velocity--density field reconstructions. A significant amount of new HI observing (both at the Green Bank Telescope, and the Parkes Radio Telescope) has already been completed for the 2MTF project (Masters \\etal ~in prep. see Masters 2007 for more details).  The first step in constructing the 2MTF catalog is to have a universal calibration of the  TF relation in the 2MASS J, H and K-bands. As discussed in Masters \\etal ~(2006; hereafter M06) the estimation of proper bias corrections for deriving TF templates is essential, not only if the template is to be used as a basis for unbiased distance estimates to galaxies with a wide range of rotational widths, but also if the template is to be used as a constraint in theories of galaxy formation. A bias corrected, or universal TF relation in the 2MASS J, H and K-bands does not currently exist. In this paper we provide such a TF template relation for use with 2MTF and in other applications. We also investigate the dependence of the NIR TF relation with wavelength and galaxy morphology and study the extent of the TF relation. Finally we measure the intrinsic scatter of galaxies from the TF template relations as a function of galaxy rotation width. We finish by providing a recipe for deriving TF distances from data consisting of galaxy magnitudes, inclinations and rotational widths. We parameterize quantities dependent on the Hubble constant with $h=H_\\circ/100$ \\kmsMpc ~where possible, and use $H_\\circ = 74$ \\kmsMpc ~ when we need to compare with $H_\\circ$ independent measurements (\\eg ~calibrations of the TF zeropoint from primary distance measures). ",
        "conclusions": "The 2MASS Tully-Fisher Survey (2MTF) aims to measure Tully-Fisher (TF) distances to all bright inclined spirals in the 2MASS Redshift Survey (2MRS). Essential to this project is a universal calibration of the TF relation in the 2MASS J (1.2 $\\mu$m), H (1.6 $\\mu$m) and K (2.2 $\\mu$m) bands. We present the first bias corrected (or universal) TF template in these bands. Previously published NIR TF relations have either ignored the issue of incompleteness bias affecting the slope and zeropoint derived, or have attempted to construct a complete sample from a cluster, or by using only part of their data.  As in earlier work, we find that the slope of the TF relation becomes steeper as the wavelength increases. We find that the slope is close to the simple theoretical expectation of $L\\propto v^4$ in K-bands, while is is slightly shallower at H- and J-band at $L\\propto v^{3.6}$. The change in zeropoint between bands in consistent with the mean colors of the galaxies in the template sample.  We investigate the dependence of the TF relation on galaxy morphology showing that in all three bands the relation is steeper for later type spirals which also have a dimmer TF zeropoint. We correct for this dependence in deriving a global TF relation for Sc galaxies in the NIR. We also investigate the luminosity dependence of the slope of the TF relation, finding hints that at both the highest and lowest masses the slope of the TF relation is shallower than the mean relation -- however this effect is not large.  We fit the TF relation using direct, inverse and bivariate fits. We favor the bivariate fit which takes into account errors on both the width and magnitude measurements (we include the measured intrinsic scatter in the weighting of the fits) finding a TF template for Sc galaxies (but which can be applied to all spirals with a morphological correction) of: \\begin{eqnarray} M_K - 5\\log h & = & -22.030 - 10.017 (\\log W - 2.5), \\nonumber \\\\ M_H - 5\\log h & = & -21.833 - ~9.016 (\\log W - 2.5),  \\\\ M_J - 5\\log h & = & -20.999 - ~9.070 (\\log W - 2.5). \\nonumber \\end{eqnarray} These relations are bases on 2MASS extrapolated magnitudes and rotation widths measured from the HI line of galaxies at the 50\\% level of the (peak$-rms$) on a fit to the edges of the profile ($W_{\\rm F50}$; supplemental widths from optical rotation curves are corrected to the $W_{\\rm F50}$ scale using the prescription described in M06 based on work of Catinella \\etal ~2007). To measure a TF distance an observation of the inclincation, $i$, of the galaxy is also required to correct the observed Doppler broadened width to the full rotational width of the galaxy, and to correct the observed total magnitude for extinction from dust internal to the galaxy. One would measure the distance to a galaxy using the deviation of a galaxy from this TF relation ($\\Delta M$) and its measured recessional velocity using \\be d = \\frac{v_{\\rm CMB}}{H_\\circ}10^{\\frac{\\Delta M}{5}}, \\ee or (equivalently, but now independent of $H_\\circ$) the peculiar velocity in the CMB frame is \\be v_{\\rm pec, CMB} = v_{\\rm CMB}(1 - 10^{\\frac{\\Delta M}{5}}). \\ee Here $\\Delta M = M_{\\rm obs} - M(W)$ is the difference between $M_X(W) - 5 \\log h$ from the TF relations above (Equations 4, where $X=J, H$ or $K$ describes the band the magnitude is measured in) and the corrected absolute magnitude of the galaxy, calculated from the observed apparent magnitude using \\be M_{\\rm corr} - 5 \\log h = m_{\\rm obs} - A_X - I_X + k_X - T_X - 5 \\log v_{\\rm CMB} - 15. \\ee The observed magnitude of a galaxy above, has a correction for extinction due to dust in our Galaxy ($A_X$), and internal to itself ($I_X$), as well as a cosmological K-correction ($k_X$). We use the prescription described in detail in \\citet{M03} to apply these corrections to the 2MASS extrapolated magnitudes. Also needed is a morphological type correction, $T_X(W)$ to correct for change in slope and zeropoint we see between different types of spirals. This correction depends on both the type, $T$ and the rotational width of the galaxy and is described in Section 3 above. The observed width, $W_{\\rm F50}$ must also be corrected for instrumental effects ($\\Delta_s$), inclination, cosmological broadening and the effect of turbulent motions ($\\Delta_t$)  in the disk: \\be W_{\\rm corr} =  \\left[ \\frac{W_{\\rm F50} - \\Delta_s}{1+z} - \\Delta_t \\right]  \\frac{1}{\\sin i}, \\ee The instrumental correction will depend on the details of the telescope used see \\citet{S05} for more information on this.  Based on arguments in Springob \\etal ~(2005) we use $\\Delta_t = 6.5$\\kms.   The scatter in the TF relation is important to understand, both to properly derive distance/peculiar velocity errors derived from TF, and as a constraint to models of galaxy formation. We fit for an intrinsic TF scatter which varies with galaxy luminosity (following the methodology of G97b and M06). We find no evidence for significant variation of the scatter with the wavelength of the magnitudes used. We do observe that later type spirals on the mean have a larger intrinsic scatter than earlier types, however disentangling this effect from the luminosity dependence of the scatter is tricky (since there are more later types at the low mass end of the relation where the scatter is observed to be larger) and beyond the scope of this paper.  The TF relations derived in this paper will be used in future work to measure distances to galaxies in the 2MTF survey. The derivation of an unbiased or ``universal'' TF relation in the NIR J, H and K-bands should also be of significant use to constrain models of galaxy formation."
    },
    "0711/0711.0306_arXiv.txt": {
        "abstract": "The soft gamma-ray repeater \\src\\ has been attracting a lot of attention owing to the fact that in December 2004 it emitted the most powerful giant flare ever observed. Here we present the results of the first \\suz\\ observation of \\src, that seems to have reached a state characterized by a flux close to the pre-flare level and by a relatively soft spectrum. Despite this, the source remained quite active, as testified by several short bursts observed by \\suz. We discuss the broadband spectral properties of \\src\\ in the context  of the magnetar model, considering its recent theoretical developments. ",
        "introduction": "The Soft Gamma-ray Repeater (SGR) \\src\\ was discovered in 1979 as a transient source of high-energy photons \\citep{laros86}. SGRs emit sporadic and short (\\mbox{$\\sim$0.1 s}) bursts of (relatively) soft gamma-rays with luminosity of $10^{40}$--$10^{41}$ \\lum during periods of activity, that are often broken by long intervals of quiescence. SGRs are thought to be magnetars: highly magnetized neutron stars with field strength of \\mbox{$10^{14}$--$10^{15}$ G}, considerably larger than those of the majority of radio pulsars \\citep[][]{thompson95,thompson96}. In the magnetar framework, the ultimate source of energy for the bursts and the quiescent emission is the ultra-strong magnetic field. The persistent X-ray counterpart of \\src\\ was observed for the first time with the \\asca\\ satellite in 1993 \\citep{murakami94} in soft X-rays (\\mbox{$<$10 keV}) and more recently also in the hard X-ray range with \\int\\ \\citep{mgm05}. A \\xte\\ observation led to the discovery of pulsations in the persistent emission with period $P\\simeq7.47$ s and period derivative \\mbox{$\\dot{P}\\simeq2.6\\times10^{-3}$ s  yr$^{-1}$} \\citep{kouveliotou98}. Under the assumption of pure magnetic dipole braking, these values imply a surface magnetic field strength of \\mbox{$8\\times10^{14}$ G}, strongly supporting the magnetar model.\\\\ \\indent Both the burst rate and the X-ray persistent emission of \\src\\ started increasing during 2003 and throughout 2004 \\citep{mte05}, culminating with the giant flare of 2004 December 27, during which \\mbox{$\\sim$$10^{47}$ erg} were released (assuming isotropic luminosity and a distance of 15 kpc), $\\sim$100 times more than in the flares observed from SGR\\,0526--66 in 1979 and from SGR\\,1900+14 in 1998 \\citep{hurley05}. After the giant flare, the persistent X-ray flux of \\src\\ started to decrease, eventually recovering the ``historical'' pre-flare level, and its X-ray spectrum to soften \\citep{met07}. A similar flux decrease have been observed from its hard X-ray and infrared counterparts \\citep{met07}. Here we present the results of the first \\suz\\ observation of \\src\\ and we discuss the spectral properties of its persistent emission. ",
        "conclusions": "Although the  \\suz/HXD does not have imaging capabilities, we know thanks to  \\int\\ observations that no other bright hard X-ray point sources are present in its field of view. The uncertainties in the instrumental background (currently at the $\\sim$5\\% level, \\citep{kokubun07}) and in the modelling of the GR emission are a more concern. Future improvements in the knowledge of these components may eventually allow us to obtain more robust conclusions. With all these caveats in mind we proceed now to discuss the broadband spectrum of \\src.\\\\ \\indent With respect to the XIS, the HXD data show an ``excess'' (see Fig.\\,\\ref{suz-broad}) that cannot be completely ascribed to calibration uncertainties between the instruments. Given the lack of a direct measure of the GR emission around  \\src, we cannot exclude that this excess is due an underestimation of such contribution to the background. If instead it is a real feature of the source spectrum, its broadband spectrum could be empirically modeled adopting a power-law with the photon-index changing from $\\sim$1 to $\\sim$2 at \\mbox{$\\sim$16 keV}, and a blackbody component with \\mbox{$k_BT\\sim0.8$ keV}. This would agree with the results reported by \\citep{gotz06}, who point out that the hard tails of the SGRs are softer than the power-law components measured below \\mbox{10 keV}. If confirmed, the presence of down-break in the 10--20 keV spectrum of \\src\\ has remarkable physical implications. The HXD data suggests that the soft power-law can extend well above the \\xmm\\ and \\cxo\\ spectral band, potentially up to $\\sim$15--20~keV, before it turns downwards to match the steeper power-law spectrum measured at higher energy. Up to now, it is still unclear wether a single physical mechanism is responsible for the observed spectrum in the \\mbox{1--200 keV} band, or the soft ($\\sim$1--10 keV) and the hard ($>$20 keV) emission are produced in different ways. In the latter case, the overall spectral shape changes in time and the difference between the sources may be produced by the different relative strength of the soft and hard components.\\\\ \\indent Recently, theoretical computations of the magnetars broadband spectra have been presented, in the attempt to identify a common physical mechanism which could explain both the hard and soft X-ray emission. As suggested by \\citep{tlk02}, within the twisted magnetosphere model a substantial non-thermal (power-law) component appears as the result of resonant cyclotron up-scattering of soft surface-emitted photons onto charges flowing into the magnetosphere. A simple one-dimensional treatment of resonant compton scattering (RCS) has been presented by \\citep{lyutikov06} and more detailed, 3-dimensional computations have been performed by \\citep{fernandez07}. Since all these treatments are based on the use of the non-relativistic (Thompson) magnetic cross section, spectral predictions are supposedly accurate only up to $\\sim$30--50~keV. This is, however, well above the spectral break detected by \\suz. Interestingly, some of the model spectra presented by \\citep{fernandez07} exhibit a downward break in the tens of keV range and have an overall shape quite reminiscent of the XIS/HXD spectrum of \\src. In particular (see their Fig.\\,6 and Fig.\\,11), when assuming a (non-thermal) top-hat velocity distribution or a broadband velocity distribution for the magnetospheric charges, multiple peaks can appear in the spectrum. Due to the complexity of the model, it is difficult to predict the relative position of the peaks, but it may well be that the downturn detected in our data witnesses the presence of a second ``hump'' (in addition to the main thermal one) in the range \\mbox{10--20~keV}. Nobili, Turolla, \\& Zane (in preparation), assuming a \\mbox{1-dimensional} thermal electron distribution superimposed to a (constant) bulk velocity, found also double humped spectra. In this case the second (and only) hump occurs when resonant scattering is efficient enough to fill the Wien peak at the temperature of the comptonizing particles. A spectral break at $\\sim$15~keV would translate then in a temperature of $\\sim$5 keV for the magnetospheric electrons. Therefore, if  confirmed, the spectral break in the X-ray data of \\src\\ may lend further support to the RCS model and prove useful in constraining the physical parameters of the model.   \\begin{theacknowledgments} PE thanks the McGill University and in particular V.~M.~Kaspi, A.~Cumming, M.~Zamfir, R.~P.~Breton, and J.~Gardner, for the warm hospitality. \\end{theacknowledgments}"
    },
    "0711/0711.2303_arXiv.txt": {
        "abstract": "We present results of {\\it Hubble Space Telescope} and {\\it Chandra X-ray Observatory} observations of globular clusters (GCs) and low-mass X-ray binaries (LMXBs) in the central regions of Centaurus~A. Out of 440 GC candidates we find that 41 host X-ray point sources that are most likely LMXBs. We fit King models to our GC candidates in order to measure their structural parameters. We find that GCs that host LMXBs are denser and more compact, and have higher encounter rates and concentrations than the GC population as a whole. We show that the higher concentrations and masses are a consequence of the dependence of LMXB incidence on central density and size plus the general trend for denser GCs to have higher masses and concentrations. We conclude that neither concentration nor mass are fundamental variables in determining the presence of LMXBs in GCs, and that the more fundamental parameters relate to central density and size. ",
        "introduction": "\\label{sec:intro}  Since the discovery of X-ray sources associated with globular clusters (GCs) in the Milky Way  (MW; e.g., Giacconi et~al.\\ 1974) it has been known that low-mass X-ray binaries (LMXBs) are formed more efficiently in GCs by a factor of ~100 compared to the field (Katz 1975; Clark 1975). This is thought to be a direct consequence of the high central densities in GCs. At high density, dynamical formation mechanisms are greatly enhanced relative to the field (Clark 1975; Fabian et~al.\\ 1975; Hills 1976).  While the dynamical origin of most GC LMXBs was proposed three decades ago, subsequent observational studies of the core properties of GCs that host LMXBs have been largely restricted to the MW. The {\\it Chandra X-ray Observatory} has made it possible to study the LMXB populations of nearby ($\\lesssim 30$ Mpc) galaxies (e.g., Fabbiano 2006 and references therein), but the core properties of the GCs, which are thought to be the most relevant in the dynamical processes that create LMXB progenitors, are not robustly determined even with {\\it HST} for most of those galaxies.  In the MW, Bellazzini et~al.\\ (1995) showed that bright LMXBs appear in denser and more metal-rich GCs,  while Pooley et~al.\\ (2003) and Heinke et~al.\\ (2003) showed that the number of GC X-ray sources correlates with $\\Gamma$, a basic indicator of dynamical encounter rates defined as $\\Gamma\\equiv \\rho_0^{1.5} r_c^2$ (e.g., Verbunt et~al.\\ 2007), where $\\rho_0$ is the central mass density and $r_c$ is the core radius. Bregman et~al.\\ (2006) showed additionally that bright GC LMXBs in the MW appear preferentially in GCs with smaller core and half-light radii and shorter half-mass relaxation times. Sivakoff et~al.\\ (2007; S07) used 11 early-type galaxies to show that encounter rates are a good predictor for the appearance of LMXBs in GCs, deriving $\\Gamma$ from measurements at the half-light radii. This is consistent with earlier results in M87 that used more uncertain derived core properties (Jord\\'an et~al.\\ 2004; J04).  To make further progress in understanding the dynamical processes that create LMXB progenitors in the cores of GCs it is necessary to directly observe LMXBs and the core structure of the GCs that host them for as large a sample as possible. In this work, we use {\\it HST/ACS} and {\\it Chandra} observations of the central regions of Centaurus~A  (or NGC~5128; hereafter Cen~A) in order to extend such studies beyond the Local Group. At the distance of Cen~A, $D\\approx 3.7$ Mpc (average of 5 distance indicators, see \\S6 in Ferrarese et al 2007), it is possible to study the central properties of GCs with {\\it HST}, an ability we exploit in what follows in order to study which structural properties of GCs are most relevant in determining the presence of LMXBs in GCs. ",
        "conclusions": "\\label{sec:conc}  We have shown in this {\\it Letter} that GCs in Cen~A that host LMXBs have significantly higher central densities, smaller sizes, and higher concentrations than the GC population as a whole. The higher concentrations are shown to result from the dependence of LMXB incidence on central density and size, plus the fact that denser (or more massive) GCs are more centrally concentrated on average.  More importantly, we have used our Cen~A sample to compare the properties of LMXB-hosting GCs to those that do not but have the same underlying mass or $\\Gamma$ distribution. We find that GCs that have LMXBs have significantly higher central densities, encounter rates and smaller sizes than GCs of the same mass that do not have LMXBs. Interestingly, we find that the mean concentrations are indistinguishable once the trend for more massive GCs to have higher concentrations is removed. As $r_h/r_t$ depends only on $c$ for King models, our conclusion agrees with the MW data presented in Bregman et~al.\\ (2006; their Figure~1).  We further show that potential GC mass-dependent processes are not fundamental in the formation of GC LMXBs (see also S07; Verbunt et~al.\\ 2007). It is rather smaller sizes and denser cores, or equivalently higher values of the core encounter rate $\\Gamma$, that are the main drivers.  Finally, our finding that concentration is not important in determining the presence of LMXBs in GCs validates the use of half-mass densities and radii to probe the dynamical properties of GCs and their connection to the presence of LMXBs. This is helpful as core properties are uncertain for GCs in galaxies much more distant than Cen~A. The detailed dependence of LMXB incidence on GC core properties will aid in further understanding of the interplay of various dynamical processes in creating X-ray binaries (e.g., Banerjee \\& Ghosh 2006;  Ivanova et~al.\\ 2007). In particular, we find in agreement with previous work (J04; S07) that the dependence of the expected number of LMXBs $\\lambda_t$ on encounter rate, $\\lambda_t \\propto \\Gamma^{0.85}$, is shallower than the naively expected linear behavior."
    },
    "0711/0711.2629_arXiv.txt": {
        "abstract": "Although full interoperativity between theoretical and observational data in the framework of the Virtual Observatory would be a very desirable achievement, the current status of VO offers few approaches to handle theoretical models.  TSAP (Theoretical Spectra Access Protocol) has been proposed as a tool to fill this void, providing a simple scheme to easily operate with this kind of data.  TSAP is useful not only for synthetic spectra but also for other types of theoretical data. As an example we show an Isochrone and Evolutionary Tracks server using TSAP.  Finally, we pay special attention to the correct treatment of the credits an important issue in the field of theoretical models. ",
        "introduction": "Theoretical models are widely used in Astronomy. Synthetic spectra, for instance, can be used to infer the physical properties of an object by comparing its observed spectrum to a theoretical collection of spectra.  A number of libraries of theoretical models are presently available in the Internet. They  can be downloaded as a collection of data files with, in some cases, the help of a web form allowing a previous selection of the files of interest. The results are usually presented in different formats as ASCII or FITS files.  This scenario forces the user to perform a previous work in order to be able to compare theoretical and observational data. The situation is even worse if different sets of theoretical models, developed by different groups, are used. This lack of homogeneity makes it difficult to design automatic tools to simultaneously work with different models xand almost impossible to develop applications able to use the models on the fly.  One of the aims of the Virtual Observatory is to guarantee a full interoperativity not only between observational data but also between them and theoretical data. However, there are a number of issues that make it difficult to achieve this goal. One of them is the clear VO bias towards observational data. This can be seen, for instance, in the main protocols and standards already developed: SSAP, SIAP, ConeSearch,..., all of them require the object position and name as mandatory, parameters that are meaningless when dealing with theoretical models.  Actually, depending on the physical problem to be tackled, the particular approach taken by the author of the model or even his/her own personal preferences, the set of parameters used to characterize the model usually differs from one to another. This makes it difficult to define a general data model for theoretical data, a major topic for the IVOA and Euro-VO Theory Groups.  TSAP is being developed as a collaboration between ESAVO\\footnote{http://esavo.esa.int} and SVO\\footnote{http://svo.laeff.inta.es} as a way to make theoretical data easily available through the VO. It makes use of the already accepted standards although with some necessary modifications. ",
        "conclusions": ""
    },
    "0711/0711.2904_arXiv.txt": {
        "abstract": "We discuss Cosmic Microwave Background constraints on the causal set theory of quantum gravity, which has made testable predictions about the nature of dark energy.  We flesh out previously discussed heuristic constraints by showing how the power spectrum of causal set dark energy fluctuations can be found from the overlap volumes of past light cones of points in the universe.  Using a modified Boltzmann code we put constraints on the single parameter of the theory that are somewhat stronger than previous ones.  We conclude that causal set theory cannot explain late-time acceleration without radical alterations to General Relativity. ",
        "introduction": "Causal sets or causets are an approach to quantum gravity in which spacetime is discretized with causality as the fundamental relation between the discrete elements.  The requirement that the as-yet unknown theory of causet dynamics should give rise to a smooth classical spacetime imposes strong restrictions on the character of the theory.  In particular, the requirement that physics should appear to be Lorentz-covariant implies a Poisson relation between the classical volume of a region of spacetime and the number of discrete elements in it.  The mean number of elements per Planck volume, $\\alpha^2$, in the Poisson process is the only free parameter in the theory on macroscopic scales.  It was noted in \\citet{forksintheroad} that the fluctuations in the Poisson process naturally give an appropriately sized dark energy if the volume of the universe is taken as conjugate to $\\Lambda$.  In \\citet{evl} it was shown that the same principle predicts a randomly fluctuating $\\Omega_\\Lambda$ of roughly constant amplitude $\\alpha$ throughout the history of the universe, with the name everpresent $\\Lambda$ given to the phenomenon.  A fundamental objection to this explanation of dark energy was raised in \\citet{barrow}, which noted that since everpresent $\\Lambda$ is sourced in past light cones (PLCs) of points, regions of the last-scattering surface separated by more than a causal horizon should have uncorrelated $\\Lambda$, and should produce $o(\\alpha)$ CMB anisotropies.  A limit of $\\sim 10^{-5}$ was therefore placed on $\\alpha$, too small for it to explain late-time acceleration.  In this paper we flesh out the objection in \\citet{barrow} with a numerical analysis based on the volumes of overlapping PLCs and place more constraints on $\\alpha$.  To emphasize that we are not considering a cosmological constant we will break from the notation used in \\citet{evl} and use $\\lambda$ to represent the causet dark energy, rather than $\\Lambda$. ",
        "conclusions": "\\label{sec-conclude}  Under both our assumptions we can place strong limits on the value of the fundamental causet parameter $\\alpha$, which controls the amplitude of all the causet fluctuations and thus (quadratically) controls the power spectra scales.  In both cases, and probably in all simple models, the causet $C_\\ell$ are largest at low $\\ell$.  This makes comparing the predictions to data straightforward; we can simply see at what $\\alpha$ the low-$\\ell$ causet \\cl\\ is equal to the measured value.  A higher $\\alpha$ would lead to higher \\cl\\  and is therefore ruled out.  In both cases, as predicted by Barrow, the large scale CMB values limit $\\alpha$ to well below $1$.  The exact constraint depends on the detail of the conversion from the $\\lambda$ perturbation to $\\psi$.  In the case of the zero velocity assumption described above the limit is is approximately $\\alpha < 10^{-6}$.  In the case of the zero shear assumption the limit is approximately $\\alpha < 10^{-7}$.  In both these cases, again as Barrow notes, the given $\\alpha$ is too small to generate late time acceleration.  We would require $\\alpha$ to be at least  $o(10^{-2})$ to achieve that.  We therefore concur with Barrow and find no loophole in the mechanics of causal set structure formation that  would invalidate his conclusion that \\evl\\ is not dark energy.  From a positive perspective, we can think of these results as providing direction to the theory.  A successful version of it must resolve the problem of excess structure formation, but retain acceleration at late times.  This might come about by finding a very unusual stress-energy tensor, by suppressing $\\lambda$ at early times, or by modifying the Einstein equations more dramatically.  The last seems the most promising possibility."
    },
    "0711/0711.2590_arXiv.txt": {
        "abstract": "Granada oscillation code (GraCo) is a software constructed to compute adiabatic and non-adiabatic oscillation eigenfunctions and eigenvalues. The adiabatic version gives the standard numerical resolution, and also the Richardson extrapolation, different sets of eigenfunctions, different outer mechanical boundary conditions or different integration variables. The non-adiabatic version can include the at\\-mosphere-pulsation interaction. The code has been used for intensive studies of $\\delta$ Scuti, $\\gamma$ Doradus, $\\beta$ Ceph., SdO and, SdB stars. The non adiabatic observables ``phase-lag'' (the phase between the effective temperature variations and the radial displacement) and ${\\delta T_{eff}\\over T_{eff}}$ (relative surface temperature variation) can help to the modal identification. These quantities together with the energy balance (``growth rate'') provide useful additional information to the adiabatic resolution (eigenfrequencies and eigenfunctions). ",
        "introduction": "GraCo \\citep{moya} is a software developed to solve non-radial adiabatic and non-adiabatic oscillation equations. It is written in fortran95 language. It can be used for models all over the HR diagram. GraCo is able to work with three different sources of equilibrium models: CESAM \\citep{morel}, Granada Code \\citep{claret} and JMSTAR \\citep{mac}. The numerical technique used is the so called Henyey relaxation method as it is described in \\citet{unno} (Section 18.2). The simple representation of the system of differential equations in terms of second-order centred differences is adopted for the numerical resolution. ",
        "conclusions": "GraCo is a complete software for the resolution of systems of differential equations related with the stellar oscillations. The general integration scheme used in the code is the Henyey relaxation method as it is explained in \\citet{unno}. In the adiabatic frame, different integration schemes can be used to obtain the eigenfrequencies and eigenfunctions: 1) Two outer mechanical boundary conditions, 2) two choices for the set of eigenfunctions, 3) two choices of the integration variables, 4) the use or not of the Richardson extrapolation and, 5) for radial modes the use of the LAWE second order differential equation or set $\\ell=0$ in the standard non-radial differential equations.  On the other hand, the main characteristic of the code is that the non-adiabatic set of differential equations can also be solved. Two different treatments of the photosphere can be used. The first considers the atmosphere as a single boundary layer, and the second describes the atmosphere-pulsation interaction. This makes it possible to obtain more accurate non-adiabatic observables (phase-lag and ${\\delta T_{eff}\\over T_{eff}}$) crucial for modal identification through the multicolor photometry. On the other hand, non-adiabatic studies also allow to study the modal energy balance, giving a theoretical range of overstable modes.  Finally we want to remark that the analytical expressions of the differential equations, eigenfunctions and boundary conditions, and the numerical procedure followed in GraCo can be found in \\citet{unno}."
    },
    "0711/0711.2073_arXiv.txt": {
        "abstract": "We present the first results from the AGN Multiwavelength Survey of Early-type galaxies in the Virgo cluster (AMUSE-Virgo). This large program targets 100 early-type galaxies with the Advanced CCD Imaging Spectrometer on board the {\\it Chandra X-ray Observatory} and the Multi-band Imaging Photometer on board the {\\it Spitzer Space Telescope}, with the aim of providing an unbiased census of low-level super-massive black hole activity in the local universe. Here we report on the \\cxo observations of the first 16 targets, and combine them with results from archival data of another, typically more massive, 16 targets.  Point-like X-ray emission from a position coincident with the optical nucleus is detected in 50 per cent of the galaxies (down to our completeness limit of $\\sim 4\\times 10^{38}$ erg \\se). Two of the X-ray nuclei are hosted by galaxies (VCC1178=N4464 and VCC1297=N4486B) with absolute B magnitudes fainter than $-$18, where nuclear star clusters are known to become increasingly common. After carefully accounting for possible contamination from low mass X-ray binaries, we argue that the detected nuclear X-ray sources are most likely powered by low-level accretion on to a super-massive black hole, with a $\\simlt$11 per cent chance contamination in VCC1178, where a star cluster is barely resolvable in archival {\\it Hubble Space Telescope} images. Based on black hole mass estimates from the global properties of the host galaxies, all the detected nuclei are highly sub-Eddington, with luminosities in the range $-8.4<$log(L$_{\\rm 0.3-10~keV}/\\rm L_{\\rm Edd}$)$<-5.9$.  The incidence of nuclear X-ray activity increases with the stellar mass \\mstar\\ of the host galaxy: only between 3--44 per cent of the galaxies with \\mstar$<10^{10}$ \\msun\\ harbor an X-ray active super-massive black hole. The fraction rises to between 49--87 per cent in galaxies with stellar mass above $10^{10}$ \\msun\\ (at the 95 per cent confidence level). ",
        "introduction": "One of the main recent developments in the study of galaxy formation and evolution has been the realization of the key role played by nuclear activity due to accretion onto super-massive black holes (SMBHs).  Low-level accretion-powered activity has been suggested to be relevant for a variety of phenomena, including regulating star formation at galaxy scales via energy feedback to solve the `downsizing' problem and providing extra energy to solve the `cooling flow' problem (Cowie et al. 1996; Dalla Vecchia et al. 2004, Springel \\etal 2004, Treu et al. 2005a,b; Bundy et al. 2005, 2007; Juneau et al. 2005; De Lucia et al. 2006, Abraham \\etal 2007, Sijacki \\etal 2007, McNamara \\& Nulsen 2007).  The most compelling pieces of evidence supporting a strong connection between galaxy formation and nuclear activity are the tight empirical scaling relations connecting the mass of the central SMBH with global properties of the host galaxy, such as bulge luminosity and mass (Kormendy \\& Richstone 1995; McLure \\& Dunlop 2002; Marconi \\& Hunt 2003; H\\\"aring \\& Rix 2004), galaxy-light concentration (Graham et al. 2001; Erwin 2004), and stellar velocity dispersion (Gebhardt et al. 2000; Ferrarese \\& Merritt 2000).  An issue of fundamental importance in understanding the galaxy-black hole connection is the `duty cycle' of nuclear activity, and its dependence upon, e.g. black hole mass.  If SMBHs are indeed ubiquitous in galactic bulges, little is known about the frequency and intensity of nuclear activity, the more so at the low-mass end (see Greene \\& Ho 2007a).  Even though a minimum level of accretion should be present, fueled by mass loss during stellar evolution (e.g. Ciotti et al. 1991; Ciotti \\& Ostriker 2007), the inferred accretion-powered luminosities are often lower than what expected from standard Bondi-Hoyle accretion (as found e.g. for the Galactic Center SMBH Sgr A$^{\\star}$; Baganoff \\etal 2003).  {From an empirical point of view, optical studies are mostly limited to samples of known active nuclei (e.g. Woo \\& Urry 2002; Heckman et al. 2004; Kollmeier et al. 2006; Greene \\& Ho 2007b) with limited coverage of the black hole mass-Eddington luminosity (\\bhm-L$_{\\rm Edd}$) plane.} Prior to the launch of the {\\it Chandra X-ray Observatory}, searches for low-level accretion powered X-ray emission from apparently inactive galaxies were effectively limited to X-ray luminosities $\\simgt 10^{40}$ erg \\se (e.g. Fabbiano \\& Juda 1997; Allen, Di Matteo \\& Fabian 2000; Sulkanen \\& Bregman 2001). The greatly improved \\cxo sensitivity, together with its fine spatial resolution, has made it possible to investigate nuclear emission associated with SMBH activity down to 3 orders of magnitude deeper, effectively bridging the gap between active galactic nuclei (AGN) and inactive galaxies. Perhaps surprisingly, only very low levels of nuclear X-ray luminosity (\\lx/\\ledd$<10^{-6}$, 2-10 keV) have been observed in nearby {\\it massive} ellipticals (Di Matteo \\etal 2000; Ho \\etal 2001; Loewenstein \\etal 2001; Pellegrini 2005; Soria \\etal 2006a,b; Santra \\etal 2007), despite their containing vast fuel reservoirs in the form of hot X-ray emitting interstellar gas.  While these results rule out radiatively efficient solutions for the accretion flow, the detected X-ray luminosities can vary by orders of magnitude when plotted against the Bondi accretion rate (Pellegrini 2005), with a large fraction of systems being even fainter than predicted by advection-dominated accretion flow models (Narayan \\& Yi 1994).  These observations are closely related to the role of SMBH feedback in inhibiting star formation at a galaxy scale level. Semi-analytical models applied to state of the art cold-dark-matter simulations have recently highlighted the importance of highly sub-Eddington SMBH activity. In the formulation by Croton \\etal (2006), a low level of prolonged activity (the so-called `radio mode') is essential to prevent the reservoir of gas surrounding the most massive galaxies from cooling and producing young stars, and thus reproduce their red colors.  In this scenario, low-level SMBH feedback halts the gas supply to the disk from the surrounding hot halo, truncating star formation and allowing the existing stellar population to redden.  So far however, these studies have been sparse and necessarily focused on the small number of galaxies at the high-mass end of the local population.  At the same time, while the paucity of AGN in the local universe is a known phenomenon, recent studies point towards an actual decline in the spatial density of local active black holes with mass below $10^{6.5-7}$\\msun\\ (Greene \\& Ho 2007b), possibly due to low black hole occupation fraction and/or low bulge fraction in dwarf galaxies.  In turn, this can place constraints on the very mechanism by which SMBHs formed in the early universe, since different models for the formation of black hole seeds predict different black hole occupation fractions at redshift zero. This effect becomes more prominent down the mass function. In particular, models where the black hole seeds are formed in the nuclei of gravitationally unstable pre-galactic discs that form through the collapse of haloes at redshift $\\sim$10 (e.g., Madau \\& Rees 2001, Begelman et al. 2006, Lodato \\& Natarajan 2006) predict the existence of a population of faint low-mass galaxies with no black hole at their center (Volonteri \\etal 2007a,b).  As a matter of fact, `light' SMBHs, presumably harbored by faint dwarf galaxies, remain elusive: the strong limits placed by dynamical studies on the masses of the nuclear objects in M33 (Gebhardt \\etal 2001; Merritt \\etal 2001) and NGC205 (Valluri \\etal 2005) suggest that neither galaxy hosts a SMBH of the mass expected from extrapolation of the known scaling relations in massive bright galaxies.  Ferrarese \\etal (2006a) suggest that, while SMBHs are common in bright (absolute B magnitude M$_{\\rm B} <-$ 20) massive galaxies, they would be progressively replaced by compact stellar nuclei moving down the mass function, and may disappear entirely at the faint end. On the other hand, compelling evidence exists for a $3.7\\times 10^6$ \\msun\\ black hole at the center of our own Milky Way (Ghez \\etal 2005), providing us with the best example of highly radiatively inefficient black hole accretion (the measured X-ray luminosity between 2--10 keV can be as low as $10^{33.3}$ erg \\se; Baganoff \\etal 2003).  Although no { direct dynamical} black hole mass determination exists below $10^6$ \\msun, indirect evidence points towards the existence of such objects in active galaxies (Filippenko \\& Ho 2003; Peterson \\etal 2005; Barth \\etal 2004; Greene \\& Ho 2004, 2007a), globular clusters (Gebhardt \\etal 2002, Gerssen \\etal 2002) as well as (some) ultra-luminous X-ray sources (Miller 2005).  {In this paper we present the first results from an extensive multi-wavelength survey of 100 spheroids --elliptical, lenticular and dwarf spheroidal galaxies-- in the Virgo cluster, conducted with the \\cxo {\\it X-ray Observatory} and the {\\it Spitzer Space Telescope}: AMUSE-Virgo (AGN Multiwavelength Survey of Early-types in Virgo\\footnote{http://tartufo.physics.ucsb.edu/$^{\\sim}$amuse}).  As described in ~\\S~\\ref{sec:amuse} the survey is designed to provide the first unbiased census of low levels of nuclear activity in the local universe as a function of host galaxy mass for early-type galaxies. \\S~\\ref{sec:data} describes our analysis of the new and archival \\cxo data obtained so far (32/100 galaxies), as well as of archival {\\it Hubble Space Telescope} (HST) images used to connect the X-ray detections with their optical counterparts and with the host galaxy properties. In \\S~\\ref{sec:bhm} we use the known correlations with host galaxy properties (stellar velocity dispersion $\\sigma$ and spheroid luminosity \\lb) to estimate masses for the central black holes. \\S~\\ref{sec:res} presents our main results, which are summarized in~\\S~\\ref{sec:sum}. } ",
        "conclusions": "\\label{sec:sum}  This paper presents the first \\cxo results of AMUSE-Virgo, a multi-wavelength survey of early type galaxies in the Virgo cluster, aimed at investigating the incidence and activity of super-massive black holes in the nuclei of 100 local nearby spheroids. The ACSVCS sample (C{\\^o}t\\'e\\ etal 2004) is selected based on the properties of the host galaxies, and therefore it is an unbiased census of nuclear activity as a function of host galaxy stellar mass (and hence presumably black hole mass). Since the stellar mass distribution of the galaxies peaks below $10^{10}$ \\msun, AMUSE-Virgo will provide us the deepest census of low-level accretion-powered activity in early type galaxies over an unprecedented range of masses.  In this paper, we combine \\cxo results from the first 16 targets with the analysis of archival data of 16, typically more massive, targets. The absolute B magnitudes of this sample of 32 objects range from $M_B=-22.5$ to $M_B=-15.0$.  The main results of this study can be summarized as follows:  \\begin{itemize}  \\item We detect point-like X-ray emission from a position consistent with the optical nucleus in 50 per cent of the targets. 12 detection out of 16 belong to the archival observations, but only 9 out of those 12 were previously reported in the literature. The remaining 4 detections (VCC2092, VCC1692, VCC1833, VCC1178) were made using new snapshot observations of low-mass targets.  \\item Two of the detected nuclei (VCC1178 and VCC1297, having \\lx=$4.7\\times 10^{38}$ and $2.6\\times 10^{38}$ erg \\se, respectively) are hosted in galaxies with absolute B magnitude fainter than $-18$ ($M_B = -17.68$ and $-$16.91), or host galaxy stellar mass lower than $10^{10}$ \\msun\\ (M$_{\\star} = 8.1\\times 10^{9}$ \\msun\\ and $5.1\\times 10^9$ \\msun). At these luminosities, massive stellar clusters are known to become increasingly common, and have been suggested to possibly even replace SMBHs (of the kind that define empirical scaling relations at the bright end; Ferrarese \\etal 2006a).  {\\item Analysis of archival ACS HST images reveals a slight excess in the surface brightness profile of VCC1178, with respect to a Sersic model. We conservatively interpret this as due to a nuclear star cluster. The inferred stellar mass does not exceed a few $10^7$ \\msun, implying a less than 11 per cent probability that the nuclear X-ray source is a solar mass compact object based on results by Sivakoff et al. (2007).}  \\item After carefully addressing possible contamination from low mass X-ray binaries in the remaining objects (based on the luminosity function by Gilfanov 2004), we conclude that {\\it the nuclear X-ray sources are most likely due to low-level accretion-powered activity from a super-massive black hole}.  \\item Between 3 and 44 per cent (95 per cent confidence level) of the galaxies with stellar masses lower than $10^{10}$ \\msun\\  harbor an X-ray active SMBH -- down to our completeness limit of $\\sim 4\\times 10^{38}$ erg \\se. The fraction of galaxies hosting an active SMBH increases to between 49 and 87 per cent for host masses above $10^{10}$. Even with only a third of the sample (32/100 galaxies), this study shows that there is {\\it statistically significant increase in the incidence of nuclear activity towards the high mass end}, consistent with what found in late-type Virgo galaxies (Decarli \\etal 2007). This should be folded with the actual black hole mass function in order to properly constrain the distribution of nuclear activity.   \\item The upper limit to the {\\it average} X-ray luminosity in the stacked image of the 12 undetected-nuclei with snapshot (5.4 ks) \\cxo observations amounts to $3.8 \\times 10^{37}$ erg \\se, or $\\langle$L$_{\\rm (0.3-10~keV)}$/\\ledd$\\rangle <3\\times 10^{-8}$ for an average black hole mass of $9.3\\times 10^6$ \\msun.  \\item Based on `fiducial' values for the central black hole mass (based on 'secure' measurements of the dispersion velocity for 24 targets, and on B magnitude otherwise) the ratio L$_{\\rm (0.3-10~keV)}$/\\ledd varies between $10^{-8.4}$ and $10^{-5.9}$ for the detected nuclei. The detected nuclei are under-luminous with respect to Bondi accretion from the interstellar gas. In agreement with earlier works (e.g. Pellegrini 2005; Soria \\etal 2006a,b; Santra \\etal 2007), this argues for an inefficient accretion mechanism, albeit our results can not break the degeneracy between intrinsically low radiative efficiency and/or drastically reduced mass accretion rate onto the black hole (owing to outflows/winds).  \\end{itemize}   A crucial question still to be addressed is that of the amount of power released in the form of kinetic energy. According to a recent study by Merloni \\& Heinz (2007), a (non-linear) correlation exists between the Eddington-scaled kinetic power and the bolometric luminosity. The non-linearity implies that the relative amount of power dissipated by these nuclei in the form of mechanical power decreases toward low X-ray luminosities. If the same scaling is applied to our sub-sample of 16 detected X-ray nuclei, the inferred kinetic power are between $\\sim 10^{-5}-10^{-3}$\\ledd, indicating that low-level SMBH feedback can be effective in faint spheroids as well as in bright massive elliptical galaxies.  Most of the galaxies yet to be observed as part of AMUSE-Virgo have stellar masses around $10^{9.5}$ \\msun, a mass range that remains largely unexplored as far as low-level nuclear SMBH activity is concerned. \\cxo observations of 68 additional faint galaxies are under way, and will further constrain the fraction of galaxies that harbor a nuclear X-ray source. As shown by Ferrarese \\etal (2006a), massive nuclear star clusters become increasingly common down the galaxy mass function, thereby increasing the chance of bright LMXB contamination. Detection of a high brightness temperature compact radio counterpart to the detected X-ray nuclei would provide definitive evidence for an accreting SMBH, as no Galactic X-ray binary can be possibly detected in the radio band at the Virgo cluster distance with current instrumentation. In fact, with the exception of the known radio sources VCC1226(M49), VCC1316(M87), VCC1978(M60), VCC763(M84), VCC1535(N4526), VCC1632(M89), none of the sample galaxies have a detected radio core brighter than the limiting flux density of 1.8 mJy at 1.4 GHz.\\footnote{Data based on GOLDMine: http://goldmine.mib.infn.it/} At the average distance of 16.5 Mpc, this corresponds to an upper limit to the radio luminosity of $\\nu \\rm L_{\\nu}< 8.3 \\times 10^{38}$ erg \\se. Deep radio observations of the targeted nuclei, with the Very Large Array, are in progress, and will put further constraints on the spectral energy distribution at low frequencies. At the same time, while LMXBs are known to emit the bulk of the dissipated accretion power in the X-ray band, SMBHs typically emit at longer wavelengths, yielding bolometric corrections as high as 80 (Marconi \\etal 2004); upcoming mid-IR observations, with the \\spi {\\it Space Telescope}, will hopefully enable us to estimate the bolometric luminosity of the detected nuclei, and to uncover obscured SMBH activity."
    },
    "0711/0711.1888_arXiv.txt": {
        "abstract": "We present new photometry of HD~149026 spanning five transits of its ``super-Neptune'' planet. In combination with previous data, we improve upon the determination of the planet-to-star radius ratio: $R_p/R_\\star = 0.0491^{+0.0018}_{-0.0005}$. We find the planetary radius to be $0.71\\pm 0.05$~$R_{\\rm Jup}$, in accordance with previous theoretical models invoking a high metal abundance for the planet. The limiting error is the uncertainty in the stellar radius. Although we find agreement among four different ways of estimating the stellar radius, the uncertainty remains at 7\\%. We also present a refined transit ephemeris and a constraint on the orbital eccentricity and argument of pericenter, $e\\cos\\omega = -0.0014\\pm 0.0012$, based on the measured interval between primary and secondary transits. ",
        "introduction": "Many clues about the processes of planet formation and evolution have been discovered by studying the ensemble properties of exoplanets, such as the ``brown dwarf desert'' (Halbwachs et al.~2000, Marcy \\& Butler 2000) and the tendency for metal-rich stars to have more detectable planets (Santos et al.~2003, Fischer \\& Valenti 2005). However, there are also individual exoplanets whose properties bear directly on theories of planet formation and evolution. One of the best examples is the transiting planet HD~149026b (Sato et al.~2005).  Compared to Saturn, HD~149026b has a similar mass but its radius is 15\\% smaller, despite the intense irradiation from its parent star that should {\\it enlarge}\\, the radius. Sato et al.~(2005) modeled HD~149026b as a dense heavy-element core surrounded by a fluid envelope of solar composition. They found a core mass of 70-80~$M_\\oplus$, which is 65-75\\% of the total mass of the planet. This is larger than the canonical core mass of 10-20~$M_\\oplus$ that is expected from the core-accretion theory of planet formation (Mizuno 1980, Pollack et al.~1996). The finding of a highly metal-enriched composition was confirmed in models by Fortney et al.~(2006), Ikoma et al.~(2006), Broeg \\& Wuchterl~(2007), and Burrows et al.~(2007). The latter authors dubbed HD~149026b a ``super-Neptune'' because the inferred mass fraction of heavy elements is similar to that of an ice giant rather than a gas giant.  Interestingly, the parent star has a rather high metallicity ([Fe/H]~$=+0.36$; Sato et al.~2005). The observation of a large core in such a metal-rich system would seem to support the core-accretion theory as opposed to coreless alternatives such as gravitational instability (Boss 1997). However, the larger-than-expected core mass raises some questions. Why did the growing protoplanet not accrete gas efficiently?  Or if it did, what happened to its envelope of light elements?  Many scenarios have been proposed: a collision of two massive protoplanets (Sato et al.~2005, Ikoma et al.~2006), {\\it in situ}\\, formation in a low-pressure nebula (Broeg \\& Wuchterl~2007), a viscous and evaporating gas disk (Ikoma et al.~2006), and a separation of gas from planetesimals at the magnetospheric ``X point'' (Sato et al.~2005).  More recently, Harrington et al.~(2007) found that the 8~$\\mu$m brightness temperature of HD~149026b exceeds its expected blackbody temperature, even if the planet is assumed to absorb all of the incident stellar radiation. In this sense the planet is anomalously hot. The high temperature may result from novel atmospheric or structural properties. Most recently, Torres et al.~(2007) announced the discovery of a transiting planet, HAT-P-3b, whose measured mass and radius indicate that it too is highly enriched in heavy elements.  In short, HD~149026b seems to be the harbinger of an entirely new kind of planet that current models of planet formation, evolution, and structure cannot accommodate without interesting and possibly exotic modifications. Because of this situation, it is desirable to improve the reliability and the precision of estimates of the system parameters, and especially a key parameter that makes this planet unusual: its small radius.  One can measure the planetary radius by gathering photometry during transits, modeling the light curve, and supplementing the model with external information about the stellar radius.  Previously, Sato et al.~(2005) analyzed 3 light curves, and Charbonneau et al.~(2006) added 3 light curves.  In this paper we present another 5 light curves of comparable or higher quality to the previously published data, and we simultaneously model all of the data to derive the most precise planetary, stellar, and orbital parameters that are currently available. We present our observations and data reduction procedure in \\S~2 and the light-curve modeling procedure in \\S~3. We provide the results in \\S~4, along with an extended discussion about the limiting error: the uncertainty in the stellar radius. The final section summarizes the results and speculates on future prospects for improvement. ",
        "conclusions": "We have presented 5 new transit light curves of the exoplanet HD~149026b and analyzed them along with 5 previously published light curves. The joint analysis has resulted in much more precise determinations of the orbital period and transit ephemerides, and also in a more precise value of the planet-to-star radius ratio. In some cases, this ratio is of primary interest, such as inferring the brightness temperature of the planet from the depth of a secondary eclipse (Harrington et al.~2007), or testing for any wavelength dependence in the radius ratio as a means of identifying planetary atmospheric features (see, e.g., Charbonneau et al.~2002).  However, when it comes to understanding the interior structure of the planet, the quantity of primary interest is $R_p$ itself, and here we can offer no significant improvement. The limiting error is the 7\\% uncertainty in the stellar radius. This error was not reduced by acquiring more light curves, although we did find agreement between the results of 4 different (and intertwined) methods for estimating the stellar radius using all of the available data. Thus, we leave unchanged the interpretation of this planet as a being unexpectedly small for its mass, and likely to be highly enriched in heavy elements (Sato et al.~2005, Fortney et al.~2006, Ikoma et al.~2006, Burrows et al.~2006).  Further improvement will depend upon progress in measuring the stellar radius. Baines et al.~(2007) recently used optical interferometry to measure the angular diameter of the planet-hosting star HD~189733, and combined it with the {\\it Hipparcos}\\, parallax to measure the stellar radius. For HD~149026, similar observations are not likely to result in a more precise value of the stellar radius, at least not in the near future. This is not only because of the 6\\% uncertainty in the parallax, but also because the expected angular diameter is only $\\approx$180~$\\mu$as, which is only 7--8 times larger than the measurement error that was achieved for HD~189733.  Supposing the parallax were known with 10~$\\mu$as precision (as one might hope from a space-based interferometric mission), the error in the Stefan-Boltzmann method for determining $R_\\star$ would be reduced to 2.7\\%. The limiting errors in that case would arise from the effective temperature and bolometric correction. In the nearer term, a possible path forward is the continued acquistion of high-quality transit photometry, in order to improve upon our measurement of $a/R_\\star$ and thereby establish the stellar mean density with greater precision.  At fixed mean density, $R_\\star$ varies as $M_\\star^{1/3}$, and our application of the Yonsei-Yale models to HD~149026 suggests that the stellar mass has already been pinned down to within 4.6\\%.  If $a/R_\\star$ were known exactly, the fractional error in the stellar radius would be approximately 1.5\\% (i.e., one-third as large as the fractional error in the stellar mass). In effect, transit photometry measures $M_\\star/R_\\star^3$, and the stellar models generally constrain a different combination of $M_\\star$ and $R_\\star$ (see, e.g., Cody \\& Sasselov 2002).  We encourage observers to be persistent in gathering additional seasons of ground-based photometry and look forward to the results of space-based photometry for this system."
    },
    "0711/0711.3844_arXiv.txt": {
        "abstract": "We present 2.0-2.4$\\mu$m integral field spectroscopy at adaptive optics spatial resolution ($\\sim$0.''1) obtained with the Near-infrared Integral Field Spectrograph (NIFS) at Gemini North Observatory of six Classical T Tauri stars: T Tau, DG Tau, XZ Tau, HL Tau, RW Aur and HV Tau C.  In all cases, the {\\it v}=1-0 S(1) (2.12$\\mu$m) emission is detected at spatially extended distances from the central stars.  HL Tau, T Tau and HV Tau C have H$_2$ emission that extends to projected distances of more than $\\sim$200 AU from the stars.  The bulk of the H$_2$ emission is typically not coincident with the location of continuum flux.  The observed morphologies vary between emission that is spatially continuous but decreasing away from the star (HV Tau C, DG Tau), and H$_2$ that shows discrete knots and arcs with 0.''2-0.''3 ($\\sim$28-42 AU) spatial extents (T Tau, XZ Tau).  Multiple transitions detected in the K-band spectra show that H$_2$ level populations are typical of gas in thermal equilibrium with excitation temperatures in the 1800K-2300 K range.  Two dimensional maps of the extinction and H$_2$ excitation temperature were estimated from the high signal-to-noise (S/N) observations of T Tau using maps of the spatially resolved {\\it v}=1-0 Q(3)/{\\it v}=1-0 S(1) and {\\it v}=2-1 S(1)/{\\it v}=1-0 S(1) line ratios.  These maps show that A$_v$ and T$_{ex}$ values in the vicinity of T Tau vary on a significant level on spatial scales of less than 100 AU.  Three of the stars have H$_2$ velocity profiles that are centered at the stellar radial velocity, and three show velocity shifts with respect to the system.  RW Aur exhibits high and low velocity H$_2$ emission, and the observed morphology and kinematics of the high velocity material are consistent with atomic emission from its red-shifted micro-jet.  The morphologies detected in H$_2$ from DG Tau and HV Tau C resemble emission seen in scattered light from stars with inclined or edge-on circumstellar disks.  The locations and relative brightnesses of H$_2$ knots detected within $\\sim$1$''$ of T Tau has varied on a timescale of $\\sim$3 years.  Each of the stars studied here show observed excitation temperatures, spatial extents, and kinematics of the H$_2$ that are most consistent with shock excited emission from the inner regions of the known Herbig-Haro energy flows or from wide-angle winds encompassing the outflows rather than predominantly from UV or X-ray stimulated emission from the central stars.  The data presented in this study highlights the sensitivity of adaptive optics-fed integral field spectroscopy for spatially resolving emission line structures in the environments of bright young stars. ",
        "introduction": "The {\\it v}=1-0 S(1) ro-vibrational transition of molecular hydrogen at 2.12$\\mu$m was discovered in the spectra of a range of astronomical objects during the rapid expansion of infrared (IR) astronomy capabilities in the 1970s.  In star formation science, this emission was first detected from shock excited outflows in the Orion nebula (Gautier et al. 1976).  Since this time, molecular hydrogen from electronic transitions in the ultraviolet (UV), ro-vibrational transitions in the near-IR, and pure rotational emission in the mid and far IR have been studied in the spectra of young stars and their environs (Ardila et al. 2002; Saucedo et al. 2003; Walter et al. 2003; Herczeg et al. 2004; 2006; Davis et al. 2001; 2002; Bary et al. 2003; Takami et al. 2004; Duchene et al. 2005; Thi et al. 1999; 2001; Richter et al. 2002; Sheret et al. 2003; Sako et al. 2005).  Today, detection of {\\it v}=1-0 S(1) H$_2$ (2.12$\\mu$m) emission in the spectra of Classical T Tauri Stars (CTTSs), embedded proto-stars, and Herbig-Haro (HH) energy sources is commonplace.  Molecular hydrogen is a primary constituent of cool gas in the circumstellar environments of young stars.  A better understanding of the evolution of the gas, and of the H$_2$, is needed to ascertain how this material in circumstellar environments will ultimately evolve into planets.  Yet, many studies of molecular hydrogen emission in the inner $\\sim$200 AU regions of CTTSs are for systems that are known to drive HH outflows (Saucedo et al. 2003; Walter et al. 2003; Herczeg et al. 2006; Takami et al. 2004; Kasper et al. 2002; Duchene et al. 2005).  In these inner environments of CTTSs, the main H$_2$ stimulation mechanisms, kinematics and morphological character from gas in outflows versus in quiescent circumstellar disks is not well established.  However, understanding the nature of the H$_2$ is especially of interest in the inner $\\sim$50 AU where thermal emission from gas in the planet-forming regions of circumstellar disks is also believed arise (Nomura \\& Millar 2005; Nomura et al. 2007; Bary et al. 2003; Ramsay Howat \\& Greaves 2007).  The ro-vibrational transitions of molecular hydrogen observed in the near IR spectra of young stars can be excited by UV fluorescence pumped by FUV or Ly $\\alpha$ photons, stellar UV or X-ray heating, and by shocks in outflows.  Shock-excited H$_2$ typically shows thermal level populations in ro-vibrational features detected in K-band spectra (2.0-2.4$\\mu$m), sometimes with high and low H$_2$ velocity features arising from different environments in the outflow or wind (Burton et al. 1989; Curiel et al. 1995; Eisloffel et al. 2000; Davis et al. 2001; 2002; Giannini et al. 2002; Takami et al. 2004; Takami et al. 2007).  In the case of UV fluorescence, molecular hydrogen is pumped into an electronically-excited state by Ly$\\alpha$ or FUV flux (Ardila et al. 2002; Walter et al. 2003; Herczeg et al. 2004; 2006), and lines in the near IR arise from the corresponding cascade through the vibration-rotation transitions (Black \\& van Dishoeck 1987; Sternberg \\& Dalgarno 1989; Shull 1978).  Electronically excited states (and subsequently, high-v states) can be stimulated by UV fluorescence pumped by Ly$\\alpha$ arising either from the star or in shocks, and in regions of low gas density the resulting level populations from this excitation can be quite different from thermal excitation (Black \\& van Dishoeck 1987; Hasegawa et al. 1987; Ardila 2002).  Stellar UV and X-ray flux from the central star also heat the ambient circumstellar gas and dust.  The H$_2$ molecules are correspondingly excited into thermal equilibrium and level populations of low-v states can be identical to those observed in shocked regions (Maloney et al. 1996; Tine et al. 1997; Nomura \\& Millar 2005; Nomura et al. 2007).  However, H$_2$ emission in the disk excited by UV pumping or heated by UV or X-ray photons from flux specifically arising from the young star is not expected to extend beyond a 30-50AU distance, and it should not be shifted in velocity by more than $\\sim$10-15km/s (Maloney et al. 1996; Tine et al. 1997; Nomura \\& Millar 2005; Nomura et al. 2007).  In CTTSs, {\\it v}=1-0 S(1) ro-vibrational emission at 2.12$\\mu$m and UV electronic transitions of molecular hydrogen were discovered in the spectra of the T Tau multiple system (Beckwith et al. 1978; Brown et al. 1981).  Since this time, T Tau has been a laboratory for studying the processes which govern molecular hydrogen in the inner environments of young stars.  IR and UV studies of T Tau have shown that extended knots and arcs of H$_2$ emission exist out to $\\sim$20$''$ distances from the central stars (van Langevelde et al. 1994; Herbst et al. 1996; 1997; Walter et al. 2003).  These emission structures are best explained by shocks in extended regions of the known HH 155 and 255 outflows.  UV studies of the inner regions of the system have also found strong levels of fluorescent H$_2$ pumped by Ly$\\alpha$ flux (Saucedo et al. 2003; Walter et al. 2003; Herczeg et al. 2006).  High spatial resolution imaging spectroscopy presented by Kasper et al. (2002) did not detect any near IR {\\it v}=1-0 H$_2$ emission in the inner $\\sim$70 AU region of T Tau, but the adaptive optics spectra of Duchene et al. (2005) did reveal significant emission near the position of the southern component, T Tau S.  Beck et al. (2004) also detected near-IR H$_2$ in emission toward T Tau S and found from multiple spectral measurements that that the line flux level could be variable on year-long timescales.  Herbst et al. (2007) has presented high spatial resolution imaging spectroscopy of H$_2$ in the region and found many spatially extended knots and arcs encompassing the position of T Tau S.  T Tau is arguably the best studied young star for understanding the processes that characterize molecular hydrogen in the inner circumstellar environments.  Yet, a clear definition of the location, structure and excitation of the H$_2$ has eluded astronomers because of different wavelength regions and sensitivities sampled by the observations and possible intrinsic variation in the emission.  To date, little is known about the spatial structure, excitation, and variability of near IR molecular hydrogen emission in the inner $\\sim$200 AU environments of most T Tauri stars.  In these regions, emission stimulated by UV pumped flux, UV and X-ray heating, and by shocks from known outflows are all expected to contribute to the observed H$_2$.  Investigation of the morphology, emission line ratios, and velocity structure in the H$_2$ transitions detectable in the K-band could distinguish between excitation mechanisms and clarify the nature of the emission in the inner regions of YSOs.  In just one pointing of a telescope, imaging spectroscopy with integral field units (IFUs) can provide three-dimensional x, y, $\\lambda$ datacubes at high spatial resolution with simultaneous coverage of many molecular hydrogen emission lines of interest.  During the last few years, there has been a marked increase in the capabilities for adaptive optics (AO) fed near IR integral field spectroscopy at 8-10 meter class observatories (Eisenhauer et al. 2000; McGregor et al. 2003; Larkin et al. 2006).  IFUs optimized for AO spectroscopy have the power to spatially resolve molecular hydrogen emission structures with less than 0.$''$1 extents over the full wavelength ranges sampled by typical IR spectrographs.  As such, the new generation of IFUs provide the means to study the molecular hydrogen morphologies, excitations, and kinematics in the environments of CTTSs.  In this paper, we present the first star formation science results from the Near IR Integral Field Spectrograph at the Gemini North Observatory.  We report on the H$_2$ emission characteristics in the environments of six classic Classical T Tauri Stars:  T Tau, RW Aur, XZ Tau, DG Tau, HL Tau and HV Tau C. ",
        "conclusions": "We have presented adaptive optics integral field spectroscopy of molecular hydrogen in the inner 200 AU environments of Classical T Tauri Stars using the Near-infrared Integral Field Spectrograph at Gemini North Observatory.  The key findings from our study are:  1) Spatially extended {\\it v}=1-0 S(1) (2.12$\\mu$m) molecular hydrogen emission is detected in six Classical T Tauri Stars that drive Herbig-Haro energy flows.  For HL Tau, HV Tau C, RW Aur and XZ Tau these observations present the first spatially resolved detections of near IR H$_2$ in the vicinity of the young stars (Takami et al. 2004; Duchene et al. 2005).   The majority of the H$_2$ emission detected in our study is not coincident with the continuum flux from the stars.  2) In several cases, the morphologies of the H$_2$ emission are consistent with the emisison from the known HH energy flows.  Particularly, RW Aur, HL Tau and T Tau roughly show H$_2$ that follows the location and knots+arcs morphologies seen in atomic emission species from the outfows.  3) In the K-band region, extraction of spectra using small apertures from integral field spectroscopic data increases the detection sensitivity to weak H$_2$ emission lines over the continuum flux.  In addition to the {\\it v}=1-0 S(1) (2.12$\\mu$m) feature, we detect up to 8 other H$_2$ transitions in the spectra of the stars.  4)  Line ratios of {\\it v}=1-0 Q(3)/{\\it v}=1-0 S(1) and {\\it v}=1-0 Q(2)/{\\it v}=1-0 S(0) are used to estimate extinctions toward each of the stars, assuming attenuation by material which follows an ISM extinction law.  2-D maps of line ratios are used to spatially resolve the extinction and excitation temperature in the vicinity of the T Tau triple star.  We find large variations in line ratios, and hence extinction and excitation temperature, on spatial scales of less than 100 AU in the environment of T Tau.  5) The level population diagrams derived from the multiple H$_2$ features show that the emission is consistent with gas in thermal equilibrium with temperatures in the 1800K to 2300K range.  6)  The centroid velocities and velocity dispersions of the {\\it v}=1-0 S(1) emission line are determined for all stars in spectral pixels that had S/N greater than 12.0  (Figures 9 and 10). The resulting maps of the spatially resolved velocity and dispersion structure were compared to the calculations carried out on a single sky emission line (used to estimate the instrumental performance).  Three of the six stars show shifts with respect to their systemic radial velocities, and an additional two show velocity structure in the velocity dispersion or standard deviations.  RW Aur has a H$_2$ profile that is very broad because of the contribution of high velocity emission, and the H$_2$ kinematic structure is consistent with the known atomic emission from its HH energy flow.  XZ Tau is the only star which shows velocity and dispersion kinematics consistent with spectrally unresolved H$_2$ emission centered at the stellar radial velocity.  7)  For all six stars, the observed spatial extents and excitation temperatures of the H$_2$ are factors of $\\sim$2-3 or more than predicted by models of UV or X-ray heating by the parent star of quiescent gas in circumstellar disks.  Additionally, the H$_2$ spatial extent and the ro-vibrational level populations are not consistent with emission exclusively from FUV or Ly$\\alpha$ stimulated fluorescence arising from the central star.  8) The observed H$_2$ emission line ratios are compared to theoretical values calculated for UV excited fluorescent emission (Black \\& van Dishoeck 1987) and for J and C type shock models (Smith 1995).  The H$_2$ line ratios are found to be most consistent with shock excitation for all cases but HL Tau, which appears to have line fluxes that are adversely affected by strong levels of scattered light.  However, these models for UV fluorescent emission and shock excited H$_2$ emission were generated for ISM environments that have different physical conditions than typically found within 200 AU of Classical T Tauri Stars.  9) Taken in combination, the H$_2$ emission morphologies, detection of molecular hydrogen at distances of greater than $\\sim$50 AU from the star, excitation temperatures of more than 1800K, kinematics of the features, and the consistency with existing shock models combine to indicate that the bulk of the H$_2$ in these stars arises from shock excited emission associated with the known HH energy outflows rather than from quiescent H$_2$ disk gas stimulated into emission by flux from the central star."
    },
    "0711/0711.3133_arXiv.txt": {
        "abstract": "% We analyze the evolution of nitrogen resulting from a set of spiral and irregular galaxy models computed for a large number of input mass radial distributions and with various star formation efficiencies. We show that our models produce a nitrogen abundance evolution in good agreement with the observational data. Differences in the star formation histories of the regions and galaxies modeled are essential to reproduce the observational data in the N/O-O/H plane and the corresponding dispersion. ",
        "introduction": "When observations, specially the metal-poor galaxy data, are ploted in the plane N/O {\\sl vs} O/H, it is evident that a primary N contribution must exist. Applying the classical {\\sl Closed Box Model}, the points may be limited by three lines as those plotted in Fig.~\\ref{no_hii}a: 1) This one called NS, which shows the evolutionary tracks of N/O when N need a seed of O to be created: $\\frac{N}{O}=\\frac{p_{N}(O)}{p_{O}}\\propto O $. 2) This one called NP that appear when N is created directly from H or He: $\\frac{N}{O}=\\frac{p_{N}}{p_{O}}= constant$. 3) NS$+$NP when both contributions there exist, which seems to be the apparent behavior reproduced by the data.  It is, however, necessary to take into account three factors not included in that simple scenario: 1) the mean lifetimes of stars, 2) the star formation efficiency or the existence of different star formation histories or star formation rates in different galaxies or objects, and 3) the metallicity dependent of stellar yields.  If the NP is ejected by low and intermediate mass (LIM) stars, it will appear in the interstellar medium very abruptely after a time delay, which means when O/H has already reached a certain value.  Obviously, this level will be higher or lower depending on the mean-lifetime of these stars. But it also depends on the efficiency to form stars: for a high efficiency the NP will appear later in the time, or at higher O/H, than for a low one. Both facts produce a certain dispersion in the resulting abundances. Besides that the metallicity dependent stellar yields produce tracks elongated in comparison with the one produced when only a value of yield is used \\cite[for details see][]{mol06}. ",
        "conclusions": ""
    },
    "0711/0711.3954_arXiv.txt": {
        "abstract": "{} {We report on observations of the eclipsing and interacting binary $\\beta$ Lyrae from the {{\\it Suzaku}} X-ray telescope. This system involves an early B star embedded in an optically and geometrically thick disk that is siphoning atmospheric gases from a less massive late B~II companion.} {Motivated by an unpublished X-ray spectrum from the {{\\it Einstein}} X-ray telescope suggesting unusually hard emission, we obtained time with Suzaku for pointings at three different phases within a single orbit. } {From the XIS detectors, the softer X-ray emission appears typical of an early-type star.  What is surprising is the remarkably unchanging character of this emission, both in luminosity and in spectral shape, despite the highly asymmetric geometry of the system.  We see no eclipse effect below 10~keV. The constancy of the soft emission is plausibly related to the wind of the embedded B star and Thomson scattering of X-rays in the system, although it might be due to extended shock structures arising near the accretion disk as a result of the unusually high mass-transfer rate.  There is some evidence from the PIN instrument for hard emission in the 10-60~keV range.  Follow-up observations with the RXTE satellite will confirm this preliminary detection.} {} ",
        "introduction": "The binary $\\beta$ Lyr is a nearly edge-on, semi-detached interacting system that has undergone mass reversal and remains in a phase of large-scale mass transfer.  The primary, mass-losing star (the ``Loser'') is a B6-B8 IIp star. The mass-gaining star (the ``Gainer'') is embedded in an optically thick accretion disk and is not directly visible. Although the embedded source had been considered as a possible compact object (Devinney 1971; Wilson 1971), it is probably a main sequence B0 star (Hubeny \\& Plavec 1991).  The system is very complex, having bipolar jet-like structures (Harmanec \\etal\\ 1996; Hoffman \\etal\\ 1998), a circumbinary envelope (Batten \\& Sahade 1973; Hack \\etal\\ 1975), and a substantial kilo-Gauss magnetic field (Leone \\etal\\ 2003).  \\begin{table} \\begin{center} \\caption{Properties of $\\beta$ Lyrae \\label{tab1}} \\begin{tabular}{lcc} \\hline\\hline Component$^a$ & Gainer & Loser \\\\ \\hline \\rule[0mm]{0mm}{3.5mm}$T_{\\rm eff}$[K] & 32,000 & 13,300 \\\\ \\rule[0mm]{0mm}{4mm}Spectral Type & $\\approx$B0 V & B6-8 IIp \\\\ \\rule[0mm]{0mm}{4mm}$M/M_\\odot$ & $\\approx$13 & $\\approx$3 \\\\ \\rule[0mm]{0mm}{4mm}$\\dot{M}^b[M_{\\odot}~{\\rm yr}^{-1}]$ & $4.7 \\times 10^{-8}$ & $7.2 \\times 10^{-7}$ \\\\ \\rule[0mm]{0mm}{4mm}$ \\varv_\\infty^b$[km s$^-1$] &  1470  & 390 \\\\ \\\\\\hline\\hline System Property$^a$ & \\multicolumn{2}{c}{Value} \\\\ \\hline \\rule[0mm]{0mm}{3.5mm}Orbital Period & \\multicolumn{2}{c}{12.9 days}\\\\ \\rule[0mm]{0mm}{3.5mm}Viewing Inclination & \\multicolumn{2}{c}{$86^\\circ$}\\\\ \\rule[0mm]{0mm}{4mm}Binary Separation & \\multicolumn{2}{c}{$55-60 R_\\odot$}\\\\ \\rule[0mm]{0mm}{4mm}Distance & \\multicolumn{2}{c}{270 pc} \\\\ \\rule[0mm]{0mm}{4mm}$\\log N_H$ (cm$^{-2}$)  & \\multicolumn{2}{c}{20.76}  \\\\ \\rule[0mm]{0mm}{4mm}{\\it ROSAT} PSPC$^c$: & \\multicolumn{2}{c}{0.07 cps} \\\\ \\rule[0mm]{0mm}{4mm}{\\it Einstein} SSS$^d$: & \\multicolumn{2}{c}{0.11 cps}\\\\ \\hline \\end{tabular}  \\parbox{2.5in}{\\small $^a$ Component and system properties taken from Harmanec (2002) unless otherwise noted. \\\\ $^b$ Mass-loss rate and terminal speed for the stellar winds of the respective binary components (Mazzali 1987).\\\\ $^c$ Bergh\\\"{o}fer \\& Schmitt 1994)\\\\ $^d$ Waldron (private comm.)} \\end{center} \\end{table}  The optical light curve of the system features a primary minimum that is $\\approx 1$ magnitude deep and a secondary minimum $\\approx 0.4$ magnitudes deep (see Fig.~\\ref{fig1}); however, the secondary minimum is deeper than the primary minimum at shorter wavelengths, and below Ly$\\alpha$ the eclipses no longer appear (Kondo \\etal\\ 1994).  A summary of the system properties is given in Table~\\ref{tab1}.  The orbital period is 12.9~days, and the mean light curves appear stable with epoch. The UV spectrum of $\\beta$ Lyr is dominated by an anomalous continuum and emission lines with unusually strong P~Cygni profiles typical of hot star winds (Hack \\etal\\ 1975; Aydin \\etal\\ 1988; Mazzali 1987).  Despite numerous and ongoing modeling attempts (e.g., Wilson 1974; Linnell \\& Hubeny 1996; Bisikalo \\etal\\ 2000; Linnell 2002; Nazarenko \\& Glazunova 2003, 2006ab), no model is yet capable of matching the observed light curves of $\\beta$ Lyr from the IR through the UV.  Strangely, $\\beta$~Lyr has been largely unstudied in the X-ray regime, despite the strong X-ray flux detected by the {\\it ROSAT} HRI (Bergh\\\"{o}fer \\& Schmitt 1994).  An unpublished spectrum taken with the {\\it EINSTEIN}/SSS in 1979 reveals X-ray emission at relatively high energies, suggesting that phase-dependent observations may provide new clues to resolving the puzzle of the $\\beta$ Lyr geometry and interactions.  Exploiting {\\it Suzaku's} excellent sensitivity to hard X-rays, we conducted three pointed observations of $\\beta$~Lyr within the same orbit.  In the following section the observations and reduction of data are detailed. Analyses of the spectra with phase are described in \\S 3, and a discussion of the results is presented in \\S 4.  \\begin{figure}[t] \\resizebox{\\hsize}{!}{\\includegraphics{f1.eps}} \\caption{ Illustration of the Suzaku pointings with respect to the binary phase.  The solid line represents the normalized V-band light curve (Fourier fit by Harmanec et al. 1996). Primary eclipse (phases 0.0 and 1.0) occurs when the mass-losing star is occulted by the disk. The three intervals marked on the curve represent our Suzaku pointings, which occurred during a single orbit in May 2006.  Phases were calculated from the quadratic ephemeris of Harmanec \\& Scholz (1993). \\label{fig1}}  \\end{figure} ",
        "conclusions": "The results of our {\\it Suzaku} study are both revealing and perplexing. We had expected to see an eclipse of soft X-rays and essentially no hard emission.  Instead, we found that the soft X-rays were nearly constant, and there is an indication of a quite hard emission component in the system.  The near constancy of the soft spectrum suggests that the X-ray source must be axially symmetric.  This is quite surprising, since the system is intrinsically non-axisymmetric.  It is worth recalling that the UV lightcurve below Ly$\\alpha$ is notable for lacking any eclipses as well (Kondo \\etal\\ 1994).  A somewhat analogous source to $\\beta$~Lyr is the near edge-on interacting binary W~Ser.  Weiland \\etal\\ (1995) interpret UV spectra of W~Ser in terms of an extended boundary layer between the star and the accretion disk. Thus, it may be possible that the largely steady FUV and soft X-ray emission could arise from a similar region in the $\\beta$~Lyr system. However, we do not currently favor this explanation since the most recent modeling of the disk and star components of \\blyr\\ (Linnell 2002) indicate that gainer star is entirely obscured at all phases.  An alternative view is suggested by hydrodynamic simulations.  Bisikalo \\etal\\ (2000) modeled the mass transfer of $\\beta$~Lyr and found that they could reproduce a disk-like structure but without a well-identified ``hot spot''.  Instead, their simulations yielded a complex distribution of extended shocks.  They also were able to generate a bipolar flow similar to the jet detected by Harmanec \\etal\\ (1996) and Hoffman \\etal\\ (1998).  However, radiative cooling was not treated self-consistently in those simulations, but approximated through the equation of state. It is unclear whether this model can generate gas at sufficiently high temperatures to produce the observed X-rays.  Hydrodynamic modeling of the $\\beta$~Lyr system was also pursued by Nazarenko \\& Glazunova (2003; 2006ab), first in 2D simulations and then in 3D simulations. These researchers included cooling curves in their simulations and allowed for a stellar wind by the gainer.  Their models also lead to a disk, a bipolar flow, and an environment permeated with elongated shocks at different azimuthal orientations about the gainer.  ``Warm'' gas temperatures up to about 200,000~K are achieved, but such gas can not contribute significant emission to the X-ray band.  The gas dynamical simulations are qualitatively promising in terms of predicting shocked structures that are spatially distributed in radius and azimuth throughout the disk.  In such a model, the hot gas could potentially that hot gas could be viewable at every phase.  However, X-ray emissions were not the focus of those studies, and it would be useful to have new simulations that emphasize the hot plasma structures.  We prefer an interpretation in which the X-rays arise from distributed shocks in the wind of the gainer star.  At $kT \\approx 0.6$~keV, the spectral characteristics of the {\\it Suzaku} spectra are compatible with a typical early main sequence B star, and the XIS luminosity $L_X({\\rm XIS}) \\approx 6.6\\times 10^{30}$ erg s$^{-1}$ is commensurate with the soft emission expected from an early B0-B1 star (Cohen, Cassinelli, \\& MacFarlane 1997).  Although the favored model for the binary geometry may preclude a direct view of the gainer star, it does permit a reasonably deep view into its wind, even at secondary eclipse (with the giant star at front).  Moreover, we know from the polarimetric study of Hoffman \\etal\\ (1998) that there is substantial scattering opacity above and below the disk plane.  The near constancy of the X-rays seen with the XIS may result from an extended ``halo'' of scattered X-ray light, similar in spirit to the scattered optical light observed in images of some Herbig-Haro objects seen edge-on to their disks, such as HH~30 (Burrows \\etal\\ 1996). In this picture the soft X-rays ultimately originate in the shocked wind of the early gainer star.  Part of the emission is observed directly from hot plasma at large radii in the wind, and part is scattered into our line-of-sight from above and below the disk.  Thus the 0.6~keV temperature of the hot gas would represent an upper limit value owing to the fact that photoabsorption of X-rays is more severe toward the soft end of the spectrum.  On the other hand, the observed X-ray luminosity must be a lower limit. \\\\  We have considered three distinct origins for the observed X-ray emissions. At this point we favor the gainer wind as the source of the soft X-rays, but acknowledge that future hydrodynamic modeling may change this picture. In addition, there seems to be an unusually hard component of emission in the system, in excess of 10~keV.  Such hard emission is more typically associated with compact objects, and not the winds of hot stars (e.g., Bergh\\\"{o}fer \\etal\\ 1997).  However, the nature of the gainer is still ambiguous; perhaps the X-ray properties of \\blyr\\ could be understood in relation to a central compact object (Devinney 1971; Wilson 1971). Although this seems unlikely, the X-ray observations certainly do not conform to our original expectations, and so at least a reconsideration of the possibility seems in order.  The X-ray pointings by {\\it Suzaku} offer the tantalizing promise of providing new and valuable information about the \\blyr\\ system, but it is clear that a far more rigorous sampling of the X-ray lightcurve and source spectrum and new models for the system will be needed to test the suppositions that we have put forth. We have obtained time with the RXTE satellite to create a more complete X-ray light curve of $\\beta$~Lyr, which should provide new insight into this complex system in the near future."
    },
    "0711/0711.3486_arXiv.txt": {
        "abstract": "We present an analysis of the optical spectra of a volume-limited sample of 375 radio galaxies at redshift {$0.4 < z < 0.7$} from the 2dF-SDSS Luminous Red Galaxy and QSO (2SLAQ) redshift survey. We investigate the evolution of the stellar populations and emission-line properties of these galaxies. By constructing composite spectra and comparing with a matched sample of radio-quiet sources from the same survey, we also investigate the effect on the galaxy of the presence of an active nucleus.  The composite spectra, binned by redshift and radio luminosity, all require two components to describe them, which we interpret as an old and a younger population. We found no evolution with redshift of the age of the younger population in radio galaxies, nor were they different from the radio-quiet comparison sample. Similarly, there is no correlation with radio power, with the exception that the most powerful radio sources ($P_{1.4}>10^{26} \\mathrm{W/Hz}$) have younger stars and stronger emission lines than the less powerful sources.  This suggests that we have located the threshold in radio power where strong emission lines ``switch on'', at radio powers of around $10^{26}\\;\\mathrm{W/Hz}$. Except for the very powerful radio galaxies, the presence of a currently-active radio AGN does not appear to be correlated with any change in the observed stellar population of a luminous red galaxy at $z\\sim 0.5$. ",
        "introduction": "\\label{sec:Introduction}  It is now accepted that most, if not all, massive galaxies host super-massive black holes at their centre. The observed tight correlation between the mass of the black hole and the global properties of the host galaxy \\citep{mtr+98,gkh+00} suggests that the growth of the black holes at their centres is intimately related to the formation and assembly of these galaxies.  However, at any one time only $\\sim$5\\% of galaxies are seen as powerful radio sources or show other evidence of an active galactic nucleus (AGN).  The question arises: are the host galaxies of such AGN different from the hosts of inactive sources?  It has been known since the early 1960s that the hosts for powerful radio galaxies are massive ellipticals \\citep[e.g.][]{mms64}.  The trigger for the transition from quiescence to an active state has been suggested to be gravitational interaction between galaxies. Signatures of tidal interactions are common amongst radio galaxies \\citep[e.g.][]{hsb+86}, which presumably result in a large increase in the amount of material feeding the black hole, thus triggering the radio emission.  Some individual radio galaxies also show significant association with young stellar populations \\citep[e.g.][]{att+01,tdm+02,wtrm02,jhcs05}; this suggests that the merger process which triggers the radio emission can also lead to star formation in the host galaxy. However, numerical studies of galaxy interactions and mergers show that star formation is not necessarily enhanced, due to the amount of gas swept out \\citep{mcms07}.  Several questions remain to be answered in this scenario: What is the relationship among mergers, the triggering of the active nucleus, and star formation: are all three components necessarily present? Do all early-type galaxies go through radio galaxy phases as they assemble?  In order to investigate these questions, we need to study a large and uniform sample of radio galaxies.  Samples of this kind have recently been studied in the local universe ($z<0.3$) by \\citet{bkh+05} and \\citet{ms07}, but until now there has been no comparable sample available at higher redshift.  \\citet{bkh+05} investigated the properties of a sample of 2215 nearby radio galaxies (most with 1.4\\,GHz radio luminosity below $10^{25}\\,\\mathrm{W/Hz}$) from the Sloan Digital Sky Survey.  They showed that the fraction of galaxies hosting a radio-loud AGN was a strong function of stellar mass, and that there was no correlation between radio luminosity and optical emission-line luminosity for the galaxies in their sample.  \\citet{bkh+05} concluded that optical AGN and low-luminosity radio-loud AGN are independent phenomena which are triggered by different physical mechanisms.  \\citet{ms07} studied a sample of 2661 radio-loud AGN selected from the 6dF Galaxy Survey. They confirmed the findings of \\citet{bkh+05} that radio-loud AGN preferentially inhabit the brightest and most massive host galaxies, and showed that the fraction of all galaxies which host a radio-loud AGN scales with the infrared K-band luminosity as $L_\\mathrm{K}^2$.  We have constructed a volume-limited sample of radio sources by combining optical data from the 2SLAQ LRG survey \\citep{cde+06} with the radio data from the VLA FIRST survey \\citep{bwh95}. The 2SLAQ sample is focused on luminous red galaxies, thus pre-selecting for the large elliptical galaxies most likely to host powerful radio sources. \\citet{scm+07} used this sample to study the evolution of the radio sources; in the current work, we investigate the properties of the host galaxies.  The signal-to-noise ratio of individual 2SLAQ spectra is adequate for redshift determinations, but is generally too low for accurate measurement of weak spectral features.  We overcome this limitation by combining spectra to form composites. This approach allows us to investigate the variation of composites with both redshift and with radio power, and to address the key question: How does the stellar population of a galaxy correlate with the presence or absence of an active radio-loud AGN in its centre?  Throughout this paper, we use $H_0$ = 71 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m$ = 0.27 and $\\Omega_\\Lambda$ = 0.73. ",
        "conclusions": "\\label{sec:conclusions}  We have investigated the optical properties of a volume-limited sample of radio galaxies in the 2SLAQ Luminous Red Galaxy survey. We created composite spectra in order to investigate how the stellar populations of these galaxies depend on redshift or radio power. By constructing a comparison sample of galaxies matched in redshift and luminosity, we can eliminate the possibility that any effects we find are a result of the fact that radio-loud galaxies are on the average more luminous than radio-quiet galaxies.  We modelled the underlying starlight using single-age stellar population models. None of the spectra were adequately described by a single model; they all demanded an additional stellar component. We achieved satisfactory fits assuming that each spectrum consisted of an ``old'' (7~Gyr) population plus a fraction of light contributed by a ``young'' population, whose age was allowed to vary between 10~Myr and 5~Gyr. Neither a power-law component nor a model with continuous star formation described the spectra as well as the two-stellar population model.  The stellar populations of all the composites were almost identical, independent of redshift, radio power, and whether or not the galaxy contains an active radio source.  The fraction of light contributed by the young population ($\\sim$ 40\\% at 4050\\AA) increased marginally with redshift, but is not correlated with radio power. The age of the ``young'' component was consistent with an age of $1000^{+1000}_{-500}$~Myr for all composites. The sole exception was the composite made from the eleven most powerful radio sources, with $\\log P > 26$, which was fitted with a young population of age 100~Myr. This was not seen in the corresponding composite of radio-quiet sources, which means it is unlikely to be due to a systematic effect, which would affect both radio-loud and radio-quiet galaxies equally. In addition, \\citet{scm+07} showed that the median redshift of 2SLAQ radio galaxies is independent of radio power, so there are unlikely to be subtle aperture effects affecting the sample.  The two-population model is probably not a very realistic description of the underlying stellar light. However, our study is a \\emph{differential} test, comparing the properties of the radio-loud galaxies with their radio-quiet counterparts. In a future paper, we will investigate other models to represent the stellar populations of these galaxies.  The emission line properties of the population show more correlation with radio activity. Approximately 15\\% of the radio-loud sources show significant \\oii\\ emission, compared to 7.5\\% of the radio-quiet sample; similarly, the radio-loud sources were roughly twice as likely to show \\neiii\\ and/or \\nev\\ emission than their radio-quiet counterparts (12\\% and 6\\%, respectively). Still, it is notable that most of the radio galaxies in the 2SLAQ would not have been recognised as AGN on the basis of their optical spectra alone. The different \\oii/Ne line ratios seen in different objects suggest there is more than one emission mechanism operating to ionise the gas: these are likely to include hot stars and post-AGB stars as well as ionisation by the central AGN.  These observations support the model that all luminous red galaxies contain a central black hole, only some of which are active at any one time; the current activity of the central black hole, as measured by the radio emission, does not appear to correlate with the stellar population in the galaxy. The sole exception appears to be that the most powerful radio galaxies show more activity and younger stars than the rest of the population. Even so, with young ($\\sim 100~\\mathrm{Myr}$) stars contributing 30\\% of the light, the fraction of mass in young stars is only 0.5\\% of the mass of the galaxy.  We need to be careful to distinguish between differences related to the activity of the central black hole and luminosity effects, since more massive galaxies are more likely to harbour AGN. We controlled for this bias by constructing a matched sample of radio-quiet galaxies.  The comparison sample is matched in redshift and apparent magnitude, so the radio-quiet galaxies match the radio-loud sources closely in absolute luminosity. Since we saw differences in the most powerful radio galaxies, but not in their radio-quiet counterparts, the difference most likely relates to the presence of an active black hole. In Paper~I, the 2SLAQ radio galaxies were found to cluster more strongly than the 2SLAQ LRG population as a whole \\citep{scm+07}. It was speculated in that paper that this might be a luminosity effect, since the 2SLAQ radio galaxies are more luminous than the rest of the sample. However, recent work using samples matched in luminosity suggest this is not the explanation (Wake et~al., in prep.). This suggests the radio galaxies really are preferentially located in denser environments, with more gas available to feed the black hole. \\citet{blk+07} have recently suggested that a similar effect is seen in lower-redshift radio galaxies from SDSS.  Our results support the findings of \\citet{bkh+05}, who found that the probability that a galaxy is radio-loud is independent of whether it is optically classified as an AGN, and that the host galaxies of radio-loud AGN match their radio-quiet counterparts. If optical emission lines trace current feeding of black holes, while radio emission points to mostly dormant black holes, accreting at low rates, then the current rate of growth of black holes in these galaxies is very low.  One notable point is that the age of the ``young'' population remains unchanged at 700~Myr (Table~\\ref{tab:composites}) despite the different ages of the galaxies: the highest redshift galaxies in our sample ($z=0.76$) are 3.1~Gyr younger than the lowest redshift galaxies ($z=0.31$).  If the stellar populations in the lowest redshift galaxies were born 1~Gyr before $z=0.31$, we should see the era of star birth, since it is covered in our sample (at a redshift $z=0.39$).  We see no evidence, however, for an epoch of star formation at that redshift.  So where are the galaxies which are forming stars as we observe them? There are two possibilities. If the single-age stellar population is a bad assumption, we may have low-level ongoing star-formation instead of star formation in discrete bursts. The second possibility, however, is that the initial colour selection of the 2SLAQ sample is biased towards accepting galaxies with these ages; in other words, the counterparts of our $z=0.3$ galaxies are indeed forming stars at $z=0.4$, but they don't appear as red galaxies while forming stars, so are not included in the 2SLAQ sample.  In fact, our modelling suggests that we may not notice these galaxies using the 2SLAQ \\emph{gri} colour selection criteria. A model galaxy with 5\\% of the mass in young stars formed in a single burst (which contributes 40--50\\% of the light around 1~Gyr after the burst) is only 1 mag bluer than the bulk of old galaxies. Furthermore, since most of the colour evolution occurs very early on, in the first 50~Myr, the chances of finding such an object in our sample is not high. In order to locate these galaxies, we would need to assemble a large sample of luminous galaxies using a different colour selection."
    },
    "0711/0711.1356_arXiv.txt": {
        "abstract": "{ We will briefly discuss the importance of sensitive X--ray observations above 10 keV for a better understanding of the physical mechanisms associated to the Supermassive Black Hole primary emission and to the cosmological evolution of the most obscured Active Galactic Nuclei. ",
        "introduction": "The fraction of the hard X-ray background (XRB) resolved into discrete sources by deep {\\it Chandra} and XMM-{\\it Newton} observations smoothly decreases from practically  100\\% below 2--3 keV to about 50\\% in the 6--10 keV energy range (Worsley et al. 2005), becoming negligible at energies above 10 keV, where the bulk of the energy density is produced. \\par The shape and intensity of the unresolved XRB calculated by Worsley et al. (2005) is well described by a ``peaky\" spectrum which is similar to that computed by folding the average spectrum of heavily obscured ($N_H > 10^{23}$ cm$^{-2}$) and Compton Thick ($N_H > 10^{24}$ cm$^{-2}$; hereinafter CT) AGN over the redshift range $z \\sim$ 0.5--1.5 (e.g., Comastri 2004a). \\par The search for and the characterization of this population has become increasingly important in the last years and represents the last obstacle towards a complete census of accreting Supermassive Black Holes (SMBHs). According to the most recent version of XRB synthesis models (Gilli et al. 2007), heavily obscured and CT AGN are thought to be numerically as relevant as less obscured Compton Thin AGN. As a consequence, their impact on the astrophysics and evolution of AGN population as a whole cannot be neglected. \\par It is convenient to consider absorption in the CT regime in two classes: mildly and heavily CT (see Comastri 2004b). The column density of the absorbing gas of the former class is just above the unity optical depth for Compton scattering, while it is much larger than unity for the latter. The primary radiation in mildly CT AGN is able to penetrate the obscuring gas and is visible above $\\sim$ 10 keV, while Compton down-scattering mimics absorption over the entire energy range for heavily CT AGN. As far as the XRB is concerned, the most relevant contribution is expected to come from mildly CT AGN which are expected to be as bright as unobscured AGN above 10--15 keV. Though the flux of the primary emission is strongly depressed, also heavily CT AGN may provide some contribution to the XRB which depend on their scattering/reflection efficiency, most likely correlated with the geometry of the obscuring gas. \\par At present, the best estimates of the basic properties of mildly CT AGN rely on the observations obtained with the {\\tt PDS} instrument onboard {\\it BeppoSAX} (Risaliti et al. 1999; Guainazzi et al. 2005). CT absorption, at least among nearby AGN, appears to be rather common. The relative fraction depends on the adopted selection criterium and can be as high as 50--60\\% for optically selected Seyfert 2 galaxies. More recently, {\\tt BAT}/{\\it Swift} (Ajello et al. 2007) and {\\tt IBIS}/{\\it INTEGRAL} (i.e., Bird et al. 2007) have surveyed the hard X--ray sky. Although these surveys have provided relatively large numbers of hard X--ray sources (e.g., Krivonos et al. 2007), they are limited to bright fluxes ($S_{10-100 keV} > 10^{-11}$ erg cm$^{-2}$ s$^{-1}$), thus sampling only the very local Universe. As a consequence, the resolved fraction of the hard ($>$ 10 keV) XRB is of the order of a few percent (Sazonov et al. 2007). \\par Assuming the absorption distribution determined in the local Universe and folding it with cosmological evolution in AGN synthesis models, it is possible to predict the relative number of CT AGN in a purely hard X--ray selected sample. Though the error bars suffer from small number statistics, there is a fairly good agreement with the model predictions (Figure~1), implying that CT absorption should be common also at high redshift and luminosities. However, only a few CT AGN are known beyond the local Universe (e.g. Comastri 2004b; Della Ceca et al. 2007). Even the deepest X--ray surveys in the 2--10 keV band (e.g., Tozzi et al. 2006; Georgantopoulos et al. 2007) uncovered only a dozen of candidate heavily obscured CT AGN and await for better spectroscopic X--ray data for a confirmation. \\par It has been recently suggested that selection via mid--IR and optical colors is promising to pick up high-$z$ heavily obscured AGN (Fiore et al. 2007a; Daddi et al. 2007). Stacking the {\\it Chandra} counts of sources selected on the basis of mid infrared (24$\\mu$m) excess emission wrt the near--infrared optical flux, a strong signal is revealed in the hard band (up to 4--6 keV), which implies, at the average source redshift ($z$$\\sim$2), $N_H$$\\sim10^{24}$~cm$^{-2}$. Despite the similarities in the source selection and stacking techniques, the estimated space density of the candidate high-$z$ CT AGN differs by up a factor 3--5. The difference can be ascribed at least in part to the slightly different luminosity ranges sampled. According to Fiore et al. (2007a) the ``missing\" population of luminous ($L_X > 10^{44}$ erg s$^{-1}$) candidate CT AGN at $z \\sim$ 2 would have the right size to explain the {\\it unresolved} XRB, while following Daddi et al. (2007) it may be significantly more numerous than predicted by synthesis models especially at low ($L_X < 10^{43}$ erg s$^{-1}$) luminosities. The above estimates are affected by large uncertainties, however they strongly point towards the existence of a sizable population of accreting heavily obscured SMBHs at high redshift. The direct detection of the CT AGN population over the redshift range most densely populated by less obscured sources contributing to the XRB ($z \\sim$ 0.7--1) bears important implications for our understanding of SMBH evolution. Moreover, a reliable determination of their luminosity function is a key parameter to reconcile the relic SMBH mass function with the local one estimated through the local $M_{BH} - \\sigma / M_{BH} - M_{bulge}$ relationships and galaxies luminosity function (Marconi et al. 2004; Merloni et al. 2004).  \\begin{figure}[] \\resizebox{\\hsize}{!}{\\includegraphics[clip=true]{cfr.pico.thick.ps}} \\caption{ \\footnotesize The predicted fraction of CT AGN (adapted from Gilli et al. 2007) as a function of X-ray flux in two different energy ranges, as labeled. The points with error bars correspond to the estimates in the {\\it Chandra} Deep Field South (CDFS; Tozzi et al. 2006) at faint fluxes, and from Swift/BAT survey (Markwardt et al. 2005). The upper point is computed assuming that the unidentified sources are all CT AGN. } \\label{pico} \\end{figure}    \\section {Current hard X--ray observations}   The expected fraction of CT AGN in the 15--200 keV band as a function of flux is reported in Figure~1. The advantage of the hard X--ray selection wrt the 2--10 keV band is evident. The {\\tt INTEGRAL} and {\\tt Swift} all sky hard X--ray surveys have provided the first flux limited samples which can be compared with model predictions. There is a fairly good agreement, but the statistics is still dominated by fluctuations associated to small numbers. Moreover, the CT nature of {\\tt Swift} and {\\tt INTEGRAL} sources is estimated combining low signal to noise ratio hard X--ray ($>$ 10 keV) spectra with  observations at lower energies, mainly with {\\it Chandra} and XMM--{\\it Newton}. Follow--up observations of candidate CT AGN from the surveys described above were performed with {\\it Suzaku}. The hard X--ray detector and the XIS CCD camera onboard the Japanese satellite are sensitive over a broad energy range ($\\sim$ 0.5--60 keV) and return good quality spectra for relatively bright hard X--ray sources. Interestingly enough, {\\it Suzaku} follow--up observations of hard X--ray selected {\\tt INTEGRAL} and {\\tt Swift} sources (Ueda et al. 2007; Comastri et al. 2007), only barely detected below 10 keV, suggest that a population of extremely hard sources, appearing only above 5--6~keV, may have escaped detection by surveys at lower energies. Their high energy spectrum ($>$ 5--10 keV) is dominated by a strong reflection component from optically thick material, while there is no evidence for the presence of a soft X--ray component, presumably due to scattering of the nuclear radiation, and common among Seyfert 2 galaxies. A possible explanation is that the geometrical distribution of the obscuring/reflecting gas is different from known Type~2 AGN.   By requiring that their integrated emissivity does not exceed the XRB level and associated uncertainties at 30 keV (e.g. Churazov et al. 2007;  Frontera et al. 2007), and assuming the same evolution of X--ray selected AGN (e.g. Hasinger et al. 2005; La Franca et al. 2005) it is, in principle, possible to constrain their number density. The hypothetic population of reflection dominated AGN could be up to a factor 4 more numerous than CT AGN (Figure ~2). Such an estimate is highly uncertain and strongly depends upon the assumption of a reflection dominated spectrum. The bottom line of such an exercise is that there might be room for a previously unknown population of obscured AGN emerging only above 5--10~keV.   \\begin{figure}[] \\resizebox{\\hsize}{!}{\\includegraphics[clip=true]{xrb_simbolx.ps}} \\caption{ \\footnotesize The contribution of reflection dominated AGN in the Gilli et al. (2007) model (lower black line) can be increased up to a factor 4 (upper black line) and still be consistent with present uncertainties on the XRB level. } \\label{frac} \\end{figure} ",
        "conclusions": "Deep {\\it SimbolX} surveys will be able to resolve about half of the XRB in the 10--40 keV energy range (Figure~5). This achievement will represents a major step forward. In fact,  at present, less than a few percent of the XRB in that band is resolved. A comparison with an extrapolation of an AGN synthesis model (Comastri 2004a) tuned to reproduce the resolved background below 10 keV is reported in Fig~5. Although the level of the XRB predicted to be resolved by {\\tt SimboX} is close to the model extrapolation, it will be possible, with the present configuration, to probe the entire range of absorption of the XRB sources. Several ``new\"  CT AGN (see Fiore et al. 2007b), which are undetected even in the deepest  XMM--{\\it Newton} and {\\it Chandra} surveys (Fig.~1), will be revealed. Moreover, given that the fraction of CT AGN is expected to steeply increase just below $\\sim 10^{-14}$ erg cm$^{-2}$ s$^{-1}$, it would be highly rewarding to push the limiting flux to somewhat lower fluxes. It is worth remarking that the resolution of a significant fraction of the XRB  has always brought significant advances in our understanding of AGN evolution. {\\it ROSAT} deep surveys showed that luminous unobscured quasars at high redshift ($z \\sim$ 1.5--2) were responsible for most of the soft (around 1 keV) XRB (Lehmann et al. 2001). Later on, thanks to {\\it Chandra} and XMM--{\\it Newton} deep surveys, it was demonstrated that the bulk of the XRB at least up to 5--6 keV is originating in relatively low luminosity sources, most of them obscured, at $z \\sim$ 1 (Brandt \\& Hasinger 2005). While we expect to uncover the so far elusive population of CT AGN  up to relatively high redshift, it may well be possible that the content of the X--ray sky above $\\sim$ 10 keV is different from what predicted. We look forward to new unexpected findings which could be obtained by pushing imaging observations in the so far unexplored hard X--ray band.  \\begin{figure*}[t!] \\resizebox{0.98\\hsize}{!}{\\includegraphics[clip=true]{sx_res_xrb.ps}} \\caption{\\footnotesize The fraction of the 10--40 keV XRB which will be resolved by a deep 1 Ms {\\it SimbolX} pointing (point with error bars). The XRB measurements  below 10 keV are from {\\it Chandra} and XMM--{\\it Newton} as well as the resolved fraction in deep fields (gold points and red squares). while the points extending up to $\\sim$ 15 keV are from RXTE (Revnivtsev et al. 2003). At high energy, the HEAO1--A4 spectrum (Gruber et al. 1999; green points) is reported along with the recent \\sax\\ PDS observations (Frontera et al. 2007; red points). The red solid line represents an AGN synthesis model (see Comastri 2004a) which reproduces the observed level of the resolved XRB below 10 keV. } \\label{eta} \\end{figure*}"
    },
    "0711/0711.0166_arXiv.txt": {
        "abstract": "We present long-slit, optical spectra of the merging system Arp 220, taken with the William Herschel Telescope (WHT) on La Palma. These data were taken as a part of a large study of ultraluminous infrared galaxies (ULIRGs) with the aim of investigating the evolution and star formation histories of such objects. Spectral synthesis modelling has been used to estimate the ages of the stellar populations found in the diffuse light sampled by the spectra. As the closest ULIRG, it proved possible to perform a detailed study of the stellar populations over the entire body of the object. The data show a remarkable uniformity in the stellar populations across the full 65 arcsec covered by our slit positions, sampling the measurable extent of the galaxy. The results are consistent with a dominant intermediate-age stellar population (ISP) with age 0.5 $<$ t$_{ISP}$ $\\leq$ 0.9 Gyr that is present at all locations, with varying contributions from a young ($\\leq$ 0.1 Gyr) stellar population (YSP) component. However, it is notable that while the flux contribution of the YSP component in the extended regions is relatively small ($\\leq$ 40\\%), adequate fits in the nuclear region are only found for combinations with a significant contribution of a YSP component (22 - 63\\% ). Moreover, while a low intrinsic reddening ($E(B - V)$$\\lsim$ 0.3) is found for the ISPs in the extended regions, intrinsic reddening values as high as $E(B - V) \\sim $ 1.0 are required in the galactic center. This clearly reflects the presence of a reddening gradient, with higher concentrations of gas and dust towards the nuclear regions, coinciding with dust lanes in the HST images. Overall, our results are consistent with models that predict an epoch of enhanced star formation coinciding with the first pass of the merging nuclei (represented by the ISP), with a further episode of star formation occurring as the nuclei finally merge together (represented by the YSP and ULIRG). ",
        "introduction": "Since the early 1980s, several studies have revealed the presence of galaxies with the spectral energy distribution dominated by infrared emission \\citep{Houck84,Houck85,Soifer84a,Soifer84b,Soifer87,Lefloch05}. These galaxies are classified as Luminous (L$_{ir} > 10^{11}$ L$_{\\odot}$) or Ultraluminous ($L_{ir} > 10^{12}$ L$_{\\odot}$) infrared galaxies (LIRGs/ULIRGs). Their prodigious infrared emission is generally attributed to the optical/UV light of luminous central sources reprocessed by dust. Starburst activity is ongoing in most, if not all, of these sources, coexisting with AGN activity in some cases \\citep{Genzel98,Surace99,Veilleux99}. Therefore these objects provide us with an excellent oportunity to study both AGN and starburst phenomena. Moreover ULIRGs are almost unvariably associated with galaxy mergers and interactions (see \\cite{Sanders96} for a revision). Thus, they also represent ideal objects to test models of galaxy evolution via major galaxy mergers.  (e.g. Mihos \\& Hernquist, 1996; Barnes \\& Hernquist, 1996; Springel et al., 2005.)  In this context Arp220, as the closest ULIRG (z = 0.018), is a key object for our understanding of star-formation in galaxy mergers. \\begin{figure*} \\begin{minipage}{170mm} \\begin{tabular}{cc} \\hspace*{-1cm}\\psfig{file=arp220_new_extend_stepF.ps,width=10cm,angle=0.}& \\psfig{file=arp220_new_nuc_stepD.ps,width=8cm,angle=0.} \\end{tabular} \\caption{Hubble space Telescope (HST) images taken with the F814W filter, showing the locations of the slits and the extraction apertures. Left: WFPC2 image showing both the nuclear region and the extended regions. Right: ACS HRC Image showing the region within the dashed-line box in the left pannel. } \\end{minipage} \\end{figure*} At infrared wavelengths, it appears as a double nucleus system \\citep{Graham90,Scoville98} separated 0.95 arcsec, with two tidal tails visible at optical wavelengths \\citep{Joseph85}. Although the mid-IR spectra suggest starburst activity as the main source of power for the infrared continuum \\citep{Genzel98,Lutz99}, the galaxy is classified as a LINER in the optical \\citep{Veilleux99}. Arp 220 was also observed by \\cite{Scoville97} at 1.3 and 2.6 $\\mu$m, sampling the CO (2-1), CO (1-0) and dust continuum emission. They determine a total H$_2$ mass of 3.2 $\\times$ 10$^{10}$$M_{\\odot}$, of which 2/3 is contained within 250 pc and confined to a thin, centrally located, dense disk, embedded in a CO disk extended by $\\sim$1 Kpc . The dynamics of the nuclear region are consistent with a collision between prograde and retrograde disks overlapping with the main gas disk ( Mundel et al. 2001, see figure 7 in their paper). Recently, \\cite{Wilson06} performed a detailed photometric study of the star clusters in Arp220. They found evidence that the cluster population can be divided into two groups: a centrally located young population of age t$_{YSP}$ $<$ $\\lsim$ 10 Myr, and an intermediate-age population of age $\\sim$300 Myr spread towards the north of the galaxy. However, given the photometric uncertainties, it is not straightforward to break the age/reddening degeneracy problem unambiguously in such studies, based only on colour-colour diagrams with relatively few filters.  To date, there have been few spectroscopic studies of stellar populations in ULIRGs. To remedy this situation we present in this paper a detailed study of the stellar populations in Arp220, based on spectroscopic observations of the extended diffuse light along three slit positions. The results are discussed in the context of evolutionary models of star formation in mergers.  Throughout the paper we assume $H_0= 71$ km s$^{-1}$, $\\Omega_{0} = 0.27$ and $\\Omega_{\\Lambda} = 0.73$, resulting in a scale of 0.363 kpc arcsec$^{-1}$ and a distance of 77.6 Mpc at z = 0.018. ",
        "conclusions": "The results presented in this paper demonstrate the utility of spectroscopic studies of the diffuse light for investigating the evolution of the stellar populations in ULIRGs.  We have found evidence for a complex star forming history in the ULIRG Arp 220, with at least two distinct episodes of starburst activity. The results are consistent with an intermediate-age stellar population of age 0.1 $<$ t$_{ISP}$ $\\leq$ 0.9 Gyr, dominating the optical emission throughout the body of the galaxy. We have also detected young ($\\leq$ 100 Myr) stellar populations (YSPs) located in the central $\\sim$ 2.5 kpc that contribute significantly to the optical emission. This latter population is likely to be related to the last enhancement of the activity towards the end of the merger event.  On the other hand, we have estimated that the bolometric luminosity accounted by the optically visible stellar populations and compared it with the IR-luminosity of the galaxy. We conclude that we are not directly detecting the bulk of the starburst activity responsible for the ULIRG phenomenon in Arp220, which is likely to be related to the radio supernovae found by \\cite{Smith98}, representing another ``extra'' episode of star formation activity apart from the two suggested by our modelling results. It is clear that our results for Arp220 are consistent with the emerging evidence for a complex and multimodal star formation activity in merging systems ( The Antennae: \\cite{Whitmore99}, NGC7252: \\cite{Maraston01}, PKS1345+12: \\cite{Rodriguez-Zaurin07}).  In terms of reddening, while the presence of highly reddened stellar component is unlikely in the extended regions of the galaxy, high reddening values are required to match the data for the apertures sampling the nuclear regions, coinciding with dust lanes in the HST images. Our results show the importance of accounting for reddening when trying to model stellar populations in star forming galaxies.  Finally, we note that the modelling results for the three large apertures, representing the integrated light of the galaxy, are consistent with those of the detailed study based on smaller apertures. This gives confidence to studies of ULIRGs at medium and high redshifts, suggesting that age measurements based on relatively large aperture studies of the diffuse light at such redshifts, are representative of the stellar populations of the galaxies as a whole."
    },
    "0711/0711.2952_arXiv.txt": {
        "abstract": "The period changes of six eclipsing binaries have been studied with focus on the Light-time effect. With the least squares method we also calculated parameters of such effect and properties of the unresolved body in these systems. With these results we discussed probability of presence of such bodies in the systems with respect to possible confirmation by another method. In two systems we also suggested a hypothesis of fourth body or magnetic activity for explanation of the \"second-order variability\" after subtraction of the light-time effect of the third body. ",
        "introduction": "Eclipsing binary systems provide a good opportunity for studying the presence of an unresolved third body by observing their minima times because of the light-time effect (hereafter LITE). It was explained by Irwin \\shortcite{Irwin} and its necessary criteria have been mentioned by Frieboes-Conde \\& Hertzeg \\shortcite{Frieboes} and also by Mayer \\shortcite{Mayer}. Presence of the third body in the system is possible only if the times of minima behave in agreement with a theoretical LITE curve, the resultant mass function has reasonable value and corresponding variations in radial velocities are measured. In the last decade also a confirmation by astrometry seems to be plausible.  In each case we have calculated new light elements of the eclipsing pair and also the parameters of the predicted third body orbit. The tables I. and II. present results for each system, where $M_i$ are masses of components, $p_i$ computed period of the unresolved body, $A_i$ semiamplitude of LITE, $e_i$ eccentricity, $\\omega_i$ length of periastron, $f(M_i)$ mass function and $M_{i,min}$ minimum mass (for $i$ = 90$^\\circ$) of predicted body, respectively. The subscript 3 and 4 denotes the parameter of the third and fourth body, respectively.  \\begin{figure} \\hspace{-2mm} \\scalebox{0.72}{\\includegraphics*[26mm,40mm][196mm,270mm]{figures.ps}} \\caption[]{The O-C diagrams of selected eclipsing binaries. The curves represent proposed light-time effect caused by the third and fourth bodies with periods and semiamplitudes given in the tables I and II. The individual primary and secondary minima are denoted by dots and circles, respectively. Larger symbols correspond to the photoelectric or CCD measurements which were given higher weights in our computations. In OO Aql and T LMi cases the LITE caused by the third body is represented by the dashdotted line together with the LITE caused by the fourth body by the solid line. The residuals after subtraction of LITE caused by the third body are shown in figures on the right.} \\label{fig} \\end{figure} ",
        "conclusions": "We have derived new LITE parameters for six eclipsing binaries by means of an $O\\!-\\!C$ diagram analysis. In two cases, OO Aql and T~LMi, another variation was found, so there is a possibility of a presence of the fourth body in the system, or magnetic activity in them. But we have not enough data to make a final decision. So the consequence is, that for the confirmation of presence of the LITE in these systems, we need detailed photometric, spectroscopic or astrometric data of these binaries."
    },
    "0711/0711.2163_arXiv.txt": {
        "abstract": "We explore the prospects of predicting emission line features present in galaxy spectra given broad-band photometry alone.  There is a general consent that colours, and spectral features, most notably the 4000 \\AA\\ break, can predict many properties of galaxies, including star formation rates and hence they could infer some of the line properties.  We argue that these techniques have great prospects in helping us understand line emission in extragalactic objects and might speed up future galaxy redshift surveys if they are to target emission line objects only.  We use two independent methods, Artifical Neural Neworks (based on the ANNz code) and Locally Weighted Regression (LWR), to retrieve correlations present in the colour N-dimensional space and to predict the equivalent widths present in the corresponding spectra. We also investigate how well it is possible to separate galaxies with and without lines from broad band photometry only. We find, unsurprisingly, that recombination lines can be well predicted by galaxy colours.  However, among collisional lines some can and some cannot be predicted well from galaxy colours alone, without any further redshift information. We also use our techniques to estimate how much information contained in spectral diagnostic diagrams can be recovered from broad-band photometry alone.  We find that it is possible to classify AGN and star formation objects relatively well using colours only. We suggest that this technique could be used to considerably improve redshift surveys such as the upcoming FMOS survey and the planned WFMOS survey. ",
        "introduction": "\\begin{figure*} \\begin{center} \\includegraphics[width=8.5cm,angle=0]{./figs/sdss_z.eps} \\includegraphics[width=8.5cm,angle=0]{./figs/deep_z.eps} \\caption{The redshift distribution for both of the samples considered in this work. The median redshifts of SDSS and DEEP2 are  $z \\sim 0.1$ and $1.0$ respectively. \\label{fig:nz}} \\end{center} \\end{figure*}  Current galaxy formation models can be very successful in predicting the general spectral energy distribution (SED) of galaxies at moderate redshifts by evolving stellar population models and co-adding a certain number of stellar populations with different ages \\citep[e.g.][]{2004A&A...425..881L}. These models have become mainstream in obtaining stellar population properties (like ages) given current data. With them, constraints can be imposed on the formation times of old elliptical galaxies, the epoch of reionization and cosmology \\citep[e.g.][]{2002ApJ...573...37J,2004MNRAS.350.1322F}.  However most stellar population models do not include the modelling of emission lines in galaxies. This is mainly due to the need for modelling the different phases present in the interstellar medium. Some spectrophotometric codes such as Pegase \\citep{2004A&A...425..881L} do include simple prescriptions for emission lines. Other models such as Starburst99 \\citep{1999ApJS..123....3L} actually provide equivalent width and flux estimates for the most important emission lines usually observed in galaxy spectra.  The broad band shape of the SED distribution can be related empirically to the presence of an emission line. Roughly speaking, we know that the emission line strength depends on the condition of the gas inside a galaxy; this is strongly related to the amount of star formation ongoing in this galaxy; the star formation can be inferred to a certain extent from the colour, i.e. we know that the SED of a star burst galaxy is much flatter than the corresponding SED of a more passively evolving red galaxy. Therefore the shape of the SED must be correlated in some way to the presence of a strong emission feature in the same galaxy.  We propose to find empirical correlations between the equivalent widths of different lines and the broad band colours from imaging data by utilising novel statistical methods.  One method, Artificial Neural Networks, is based on the ANNz code \\citep[e.g.][]{2004PASP..116..345C,2007ApJ...666..747B,2007arXiv0705.1437A}, previously used to predict photometric redshifts from galaxy colours. The second statistical method we explore here is Locally Weighted Regression (LWR).  We can then use such correlations to understand better the process of galaxy formation and the conditions under which strong lines are produced.  We suggest that with these statistical results, any model attempting to predict emission features in galaxy spectra could be considerably improved by being calibrated to have the same statistical properties as found in nature.  Another strong motivation for this work is the predictability of an emission feature present in high redshift galaxies. The Sloan Digital Sky Survey (SDSS) and the 2dF Galaxy Redshift Survey (2dFGRS) have probed the nearby Universe. More current (VVDS and DEEP2) as well as future efforts will attempt to probe a large part of the high redshift Universe. More specifically if one is only interested in redshift surveys and not on the more specific SED of a galaxy then it is important to minimize the amount of time spent looking for a galaxy redshift. The most efficient way of doing this is to perform redshift surveys targeted at emission line objects. If there is a possibility of predicting statistically which galaxies have strongest emission lines then the potential time spent locating all these galaxies in redshift space can be reduced and the scope for larger surveys increased in the same way. These approaches would be useful in particular for the forthcoming FMOS survey \\citep{2006SPIE.6269E.136D} and the planned WFMOS survey.  The outline of this paper is as follows. In Sec.\\ref{sec:data} we present the data used to provide correlations found, we examine a low redshift sample (SDSS) as well as a high redshift sample (DEEP2). Section 3 presents the methodology used to extract the correlations from the data and describes the results found, we attempt to use LWR and ANNz for this purpose. In Section 4 we present methodology for classification of AGNs and starforming galaxies. We comment on the prospects that this technique has in speeding up future redshift surveys in Section 5 and conclude in Section 6.  We also provide two short appendices describing the mathematics of the LWR and ANNz methods used. ",
        "conclusions": "\\label{sec:conclusion}  We have looked, in this paper, for empirical relations between the equivalent width for several different lines and the broad band colours for the same objects. We used an automated way to explore the correlations found in the data via the use of LWR and ANN. The two methods give very similar results.  We have performed the analysis in two samples, a low redshift sample obtained from the Sloan Digital Sky Survey covering the nearby Universe and a high redshift sample from the first data release of the DEEP survey covering a deeper sample off the Universe. We found that in the DEEP data one could predict the equivalent width of the \\oii\\ line relatively well by looking at broad band colours only. With the SDSS data six lines were predicted from a training set sample. In general collisional lines presented little correlation although some of them were reasonably predicted by the colours. There was a stronger correlation found in recombination lines.  We have compared the power of prediction of both methods used. We have concluded that both Artificial Neural Networks and the Locally Weighted Regression methods are capable of recovering most of the information encoded in the training set with little statistical difference between both methods. We have used both methods for classification purposes and found that the Neural Networks are capable of classifying the objects into line-emitting and non line-emitting but that the Locally Weighted Regression method was unable to do so. We have shown that it is possible to classify galaxies into AGN, passive galaxies and line emitting galaxies well without galaxy spectra, using solely colours and a training set of the order of 15000 galaxies.  We have discussed the prospects of speeding up redshift surveys with this and we conclude that with a reasonable training set of the order of 10000 galaxies one would be able to considerably speed up future surveys done with instruments such as FMOS and WFMOS. Furthermore, in this paper we have only taken advantage of the colour data; however the methods described can accommodate very easily other information such as morphology or size."
    },
    "0711/0711.2480_arXiv.txt": {
        "abstract": "We investigate the shapes and mutual alignment of voids in the large scale matter distribution of a $\\Lambda$CDM cosmology simulation.  The voids are identified using the novel WVF void finder technique.  The identified voids are quite nonspherical and slightly prolate, with axis ratios in the order of $c:b:a\\approx 0.5:0.7:1$. Their orientations are strongly correlated with significant alignments spanning scales $>30h^{-1}\\hbox{Mpc}$.  We also find an intimate link between the cosmic tidal field and the void orientations. Over a very wide range of scales we find a coherent and strong alignment of the voids with the tidal field computed from the smoothed density distribution. This orientation-tide alignment remains significant on scales exceeding twice the typical void size, which shows that the long range external field is responsible for the alignment of the voids.  This confirms the view that the large scale tidal force field is the main agent for the large scale spatial organization of the Cosmic Web. ",
        "introduction": "Galaxy redshift surveys and computer simulations of cosmic structure formation have shown that matter traces out a frothy pattern, the Cosmic Web.  \\cite{bondweb1996} established how the tidal forces induced by the inhomogeneous cosmic matter distribution are the main agent shaping the Cosmic Web.  The resulting cosmic tidal force field is the source of the large scale, coherent, anisotropic forces which shape the cosmic matter distribution into characteristic filamentary and sheetlike patterns.  \\begin{figure*} \\includegraphics[width=0.45\\textwidth]{density_field2.ps}\\hspace{0.5cm} \\includegraphics[width=0.45\\textwidth]{density_bar_zoom.ps} \\vspace{0.5cm}\\\\ \\includegraphics[width=0.45\\textwidth]{skel_ellipses_zoom.ps}\\hspace{0.5cm} \\includegraphics[width=0.45\\textwidth]{skel_bars_tidal_zoom.ps}  \\caption{The four panels show a slice through the density field with the void-segmentation superposed. The first panel (a) a slice through the entire simulation. The image shows the density field of the simulation. The density field is determined by DTFE. Superimposed are the black solid lines indicating the boundaries of the WVF voids. The white box indicates the section which forms the subject of the next three frames. Second panel: density field in the subbox with the WVF void boundaries superimposed (solid lines). The blue bars indicate the direction of the projection of the longest void axis. Third panel: the WVF void boundary contours (solid lines) with the fitted void shape ellipsoids sperimposed. Fourth panel: on the landscape with WVF void boundaries the tidal field compressional component is represented by tidal bars (red), representing the direction and strength of the tidal field. Also depicted are the void shape bars (blue). } \\label{fig:slices} \\end{figure*}  Voids are a dominant component of the Cosmic Web \\citep[see e.g.][]{tully2007, weyrom2007}, occupying most of the volume of space. In this paper we present evidence for significant alignments between neigbouring voids and establish the intimate dynamic link between voids and the cosmic tidal force field.  The Watershed Void Finder (WVF) procedure \\citep{platen2007} that we use to identify voids is crucial to the succes of this analysis.  The WVF technique is based on the topological characteristics of the spatial density field and thereby provides objectively defined measures for the size, shape and orientation of void patches.  Large voids form around deep density troughs in the primordial density field.  The main aspects of the evolution of a large void may be understood on the basis of the expansion of simple spherical and isolated under-densities (e.g. \\cite{edbert1985}).  Under-dense regions expand with respect to the background Universe and in general will have the tendency to grow more spherical with time \\citep{icke1984}.  However, \\cite{shandarin2006} and \\cite{parklee2007a} have emphasized that, in realistic cosmological circumstances, voids will be nonspherical.  This is quite apparent in images of, for example, the Millennium simulation \\citep{springmillen2005}.  Moreover, substructures within voids display a manifest alignment along a preferred direction.  This is, in part, a consequence of the relatively strong influence of the surrounding inhomogeneous matter distribution on the void's structure and evolution.  Voids evolve hierarchically: as they expand with respect to the background Universe they merge with their surrounding peers, building voids of ever larger sizes \\citep{regoes1991, weykamp1993, gottloeber2003, colberg2005}.  Small voids can survive as substructure within the larger voids or may disappear under the gravitational impact of surrounding overdense structures \\citep{sahni1994, shethwey2004}.  The result is a void population whose scale is distributed around a characteristic void size: we shall show this using our WVF sample.  \\begin{figure*} \\mbox{\\hskip -0.6truecm\\includegraphics[width=0.36\\textwidth]{void_distribution_LCDM.ps}} \\mbox{\\hskip -0.5truecm\\includegraphics[width=0.36\\textwidth]{axis_ratio.ps}} \\mbox{\\hskip -0.5truecm\\includegraphics[width=0.36\\textwidth]{epsilon_distribution.ps}} \\caption{Void characteristics and statistics. Left: the (peaked) void size distribution as a function of void radius $R_v$. Centre: scatter diagram of the two void axis ratios $\\eta_{32}=a_3/a_2$ versus $\\eta_{21}=a_2/a_1$. Right: distribution of void ellipticity $\\epsilon=1-a_3/a_1$.} \\label{fig:shape} \\end{figure*} A major manifestation of large scale tidal influences is that of the alignment of shape and angular momentum of objects \\citep[see][]{bondweb1996, desjacques2007}.  The alignment of the orientations of galaxy haloes, galaxy spins and clusters with larger scale structures such as clusters, filaments and superclusters have been the subject of numerous studies \\citep[see e.g.][]{binggeli1982, bond1987, rhee1991, plionis2002, basilakos2006, trujillo2006, aragon2007, leevrard2007, parklee2007b,leespringel2007}.  Recent analytical and numerical work by \\cite{parklee2007a} discussed the magnitude of the tidal contribution to the shape of voids \\citep[see also][]{leepark2006}.  Describing the evolution of voids under the influence of tidal fields by means of the Lagrangian Zel'dovich approximation, they found that the ellipticity distribution of voids is a sensitive function of various cosmological parameters and remarked that the shape evolution of voids provides a remarkably robust constraint on the dark energy equation of state \\citep{leepark2007}.  In section 2 we introduce the N-body simulation which we used to identify the voids, along with a brief description of the WVF method we use to identify the voids in our sample. The results for individual void characteristics, in particularly the size and shape of voids, are presented in section 3.  Then in section 4 we present the analysis of the alignment of voids in the sample.  The possible origins for the void alignment are discussed in section 5, and in section 6 we focus on the role of tidal fields in determining the evolutionary history of voids.  Finally, section 7 concludes with a short discussion and prospects. ",
        "conclusions": "We have investigated the shapes and alignments of voids found in the large scale matter distribution of the VLS GIF N-body simulation of structure formation in a $\\Lambda$CDM universe.  The void sample has been identified using the Watershed Void Finder (WVF) technique.  \\subsection{Overview of Results} Although voids tend to be less flattened or elongated than the halos in the dark matter distribution, they are nevertheless quite nonspherical: they are slightly prolate with axis ratios on the order of $c:b:a\\approx 0.5:0.7:1$.  Two important factors contribute to this flattening.  Internally, voids tend to become more spherical as they expand \\citep{icke1984}. However, they will never be able to reach perfect sphericity before meeting up with surrounding structures. Furthermore, external tidal forces will contribute substantially to the anisotropic development of the voids. In this study we have established the influence of the tidal fields as being an important agent in the dynamical evolution of voids.  Indeed, our study provides strong confirmation that tidal forces dominate the shaping of the Cosmic Web, as emphasized long ago by \\cite{bondweb1996}.  We have also investigated the relative orientation of the voids. The orientation of voids appears to be strongly correlated with alignments spanning distances $>30h^{-1}\\hbox{Mpc}$.  This is true of both the shortest and longest axes of the void shape.  This coherence is quite conspicuous in plots of the density distribution. To test for such void alignments in observational data a follow-up study will apply the WVF formalism to galaxy redshift surveys.  We find an intimate link between the cosmic tidal field and the void orientations. Over a range of scales we find a strong alignment of the voids with the tidal field arising from the smoothed density distribution. Given that the alignment correlations remain significant on scales considerably exceeding twice the typical void size, our results show that the long range external field is responsible for the alignment of the voids. This confirms the view that the large scale tidal force field is the main agent for the spatial organization of the Cosmic Web.  \\subsection{Insights} We have argued that the large scale tidal field is not only the main agent responsible for the shapes of the voids but also for the resulting alignments of the voids. Locally, the orientation of a void turns out to be strongly aligned with the tidal force field generated by structures on scales up to at least $20-30h^{-1}\\hbox{\\rm Mpc}$. On scales comparable to the average void size we also found indications of nonlinear effects, with the void's orientation reacting to the small-scale nonlinear influences.  This conclusion agrees with that of a similar study of halo alignments by \\cite{leespringel2007}. However, the alignment of halos tends to be strongly attenuated by local nonlinear effects \\citep[e.g.][]{haarwey1993, aragon2007} rendering the final halo alignment signal weaker than that for voids.  One aspect we have not yet adressed in detail is the entanglement of initial conditions and the influence of the tidal field. On large scales, where the density fluctuations are in the linear regime, these are different aspects of the inhomogeneous matter distribution: if one is known the other can be determined. In this respect the indication of nonlinear effects in the tidal-void alignment ${\\mathcal A}_{TS}$ (fig.~\\ref{fig:cor2}) at small scales is interesting. They may very well be an expression of the strong influence of coherent overdense filaments and dense compact clusters on the evolution of small voids in the outer regions of large voids.  \\cite{weysheth2004} argued that the collapse of small voids, an essential aspect of the hierarchical evolution of voids \\citep{shethwey2004}, manifests itself in a tidally induced anisotropic contraction of small underdensities at the void boundaries. The dynamics and the fate of these small underdensities is often decisively influenced by external anisotropic forces \\citep{weybabul1996}. While such forces may affect the nonlinear evolution of halos and clusters to a considerable extent \\citep{bondmyers1996,sheth2001}, they do not go as far as deciding their fate.  This work is part of a study of the evolution of voids under the influence of external influences. A proper framework for following the gradual hierarchical buildup of void regions under the influence of external tidal forces is provided by the void-patch description \\citep{platen2008, bondweb1996}. To some extent the void-patch formalism is a more accurate approximation of reality than for the equivalent peak-patch: because of their approximate uniformity void evolution is accurately described by the ellipsoidal model.  A perhaps even more complex issue is that of the identity of the outer regions of voids, the regions where the void merges into the surrounding matter distribution. We have argued that this will play an important role in determining the shape of voids. However, a proper definition of a void boundary does not (yet) exist. Arguably the WVF has proven to represent a major advance along these lines and we should soon be able to be more definite on the issue.  Having established the external influences on a void's evolution, we may try to understand the sensitivity of the void population to the underlying cosmology. This depends to a considerable degree on the differences in scale dependence of the tidal force fields. \\cite{leepark2007} did find a considerable influence. This will open the door to the use of voids in determining the global cosmology."
    },
    "0711/0711.2449_arXiv.txt": {
        "abstract": "It has been proposed that RX J1914.4+2456 is a stellar binary system with an orbital period of 9.5 mins. As such it shares many similar properties with RX J0806.3+1527 (5.4 mins). However, while the X-ray spectrum of RX J0806.3+1527 can be modelled using a simple absorbed blackbody, the X-ray spectrum of RX J1914.4+2456 has proved difficult to fit using a physically plausible model. In this paper we re-examine the available X-ray spectra of RX J1914.4+2456 taken using {\\xmm}. We find that the X-ray spectra can be fitted using a simple blackbody and an absorption component which has a significant enhancement of neon compared to the solar value.  We propose that the material in the inter-binary system is significantly enhanced with neon. This makes its intrinsic X-ray spectrum virtually identical to RX J0806.3+1527. We re-access the X-ray luminosity of RX J1914.4+2456 and the implications of these results. ",
        "introduction": "The X-ray source RX J1914.4+2456 (hereafter RX J1914+24) has been the subject of much debate since its discovery during the {\\ros} all-sky survey. A number of competing models have been put forward to account for its observed properties, but all of them involve a stellar binary system. The models can be split into accretion and non-accretion models. The non-accretion model is the unipolar inductor (UI) model where dissipation of large electrical currents heat the magnetic white dwarf (Wu et al 2002).  It shares many similar observational characteristics to the X-ray source RX J0806.3+1527 (hereafter RX J0806+15, see Cropper et al 2004a for a review).  One of the main uncertainties to understanding the nature of RX J1914+24 is accurately determining its X-ray luminosity, $L_\\mathrm{X}$. In the UI model, $L_\\mathrm{X}$ is proportional to the rate of change of the orbital period and the degree of asynchronisation between the binary orbital period and the primary star (Wu et al 2002, Dall'Osso et al 2007).  \\begin{table*} \\begin{center} \\begin{tabular}{lrrrrrrrrr} \\hline Revolution & Date of & \\multicolumn{2}{c}{EPIC PN} & \\multicolumn{2}{c}{EPIC MOS1} & \\multicolumn{2}{c}{EPIC MOS2} & RGS & Flux\\\\ & Observation & Mode &  Duration & Mode & Duration & Mode & Duration & Duration & \\ergscm\\\\ \\hline 0718 & 2003-11-09 & sw & 6641 & timing & 8933 & sw & 8993 &  9880 & $1.25\\times10^{-12}$ \\\\ 0721 & 2003-11-15 & sw & 2888 & timing & 3865 & sw & 3865 & 4000  & $1.20\\times10^{-12}$ \\\\ 0880 & 2004-09-28 & ff & 11869 & ff & 14954 & ff & 14954 &  16954 & $1.27\\times10^{-12}$ \\\\ 0882 & 2004-10-02 & ff & 15274 & ff & 18419 & ff & 18419 & 18880  & $1.40\\times10^{-12}$ \\\\ \\hline \\end{tabular} \\end{center} \\caption{The log for the observations of RX J1914+24 made using {\\xmm}. We show the mode the detector was in where `sw' refers to `small window', `ff' to `full frame' and the duration of `good' time -- ie excluding time intervals of enhanced background. The RGS detectors were configured in `Spectroscopy' mode. In the last column we show the observed integrated flux in the 0.2--10keV energy band using the EPIC pn detector data and fitting an absorbed blackbody plus broad emission line spectral model.} \\label{log} \\end{table*}  In practise it has been difficult to get an accurate value for $L_\\mathrm{X}$. This is partly due to the fact that RX J1914+24 is highly reddened. The other factor is that its X-ray spectrum is rather unusual and therefore difficult to determine the underlying emission model. X-ray spectra obtained using {\\xmm} are not well fitted using a simple absorbed blackbody, showing large residuals near 0.7 keV (Ramsay et al 2005). Ramsay et al (2005) found that an absorbed blackbody with a broad emission line centered at 0.59keV gave a much improved goodness of fit and for a distance of 1kpc implied $L_\\mathrm{X}\\sim10^{35}$ \\ergss.  Using two further longer series of observations of RX J1914+24 also taken using {\\xmm}, Ramsay et al (2006) found that an absorbed low-temperature thermal plasma model with an edge at 0.83keV gave a significantly improved fit compared to the previous best model fit. For a distance of 1kpc this optically thin emission model gave $L_\\mathrm{X}\\sim10^{33}$ \\ergss. On the other hand if the distance was much lower than 1 kpc (Steeghs et al 2006, Barros et al 2007) then $L_\\mathrm{X}$ could be as low as $\\sim3\\times10^{31}$ \\ergss -- giving a range in $L_\\mathrm{X}$ of 4 orders of magnitude!  RX J1914+24 has been observed using {\\xmm} at 4 separate epochs (Table \\ref{log}). An analysis of the data taken using the EPIC detectors have been presented in Ramsay et al (2005) (from the first two epochs) and in Ramsay et al (2006) (from the last two epochs). In this paper we examine the data obtained using the RGS detectors; re-examine the data obtained using the EPIC detectors using more recent calibration data, and also re-examine the models used to fit the data. ",
        "conclusions": "For the reasons outlined in \\S \\ref{ins}, we rule out RX J1914+24 being an isolated neutron star. Since all the known neutron star UCBs have X-ray emission extending up to many 10's of keV we also rule out an accreting neutron star UCB model. We cannot rule out that a neutron star is in a binary system where a secondary star was not filling its Roche Lobe. In this scenario, an X-ray bright system would have to be powered by UI.  It is highly unlikely that the line of sight absorption to RX J1914+24 has a chance enhancement of neon. It is much more likely that this overabundance is concentrated in the binary system. Juett \\& Chakrabarty (2005) noted that for some neutron star UCBs the Ne/O abundance varied from epoch to epoch and hence the observations could not be used to determine the abundance of the secondary, mass-donating star, in the binary system. In the case of RX J1914+24 there is clear evidence for a significant over-abundance of neon in the absorption component at each epoch. At this stage it is not clear if this over-abundance is due to circumbinary material left over from a previous stage in the binary formation process or can give us a direct insight into the chemical composition of the secondary star (if accretion is occurring).  What are the implications of our findings regarding the X-ray luminosity of RX J1914+24? Steeghs et al (2006) discuss the extinction and distance estimates to RX J1914+24. While the distance is rather uncertain, it is likely that it is greater than $\\sim$1 kpc.  We can rule out the lower estimates ($L_\\mathrm{X}\\sim10^{33}$ \\ergss for a distance of 1 kpc) which were derived using a low temperature thermal plasma model.  Taking the unabsorbed bolometric fluxes derived using the blackbody with absorption component with variable abundances and assuming a distance of 1 kpc we find $L_\\mathrm{X}=2\\times10^{34} - 1.6\\times10^{35}$ \\ergss.  Dall'Osso et al (2007) made a detailed investigation of the UI model in the context of RX J1914+24 and RX J0806+15. They predicted that for low luminosities, $L_\\mathrm{X}\\sim10^{33}$ \\ergss, the asynchronism between the orbit and the magnetic star in RX J1914+24 would have to be $\\alpha\\sim0.9-$0.98, where $\\alpha=\\omega_{1}/\\omega_{o}$, and $\\omega_{1}$ is the rotation frequency of the primary star and $\\omega_{o}$ is the orbital frequency. Unless RX J1914+24 was located at a distance significantly less than 1kpc, we can rule these low values of asynchronism. For high luminosities ($L_\\mathrm{X}=10^{34-35}$ \\ergss), Dall'Osso et al (2007) predicted that an asynchronism of a few was required ($\\alpha\\sim$ 4). For an observed period of 569 sec, $\\alpha=2-10$ gives a predicted period of $\\sim$60-300 sec. There is no evidence for power at these periods in the power spectra of the X-ray light curves (Ramsay et al 2006)."
    },
    "0711/0711.3603_arXiv.txt": {
        "abstract": "We study the effect of primordial non--Gaussianity on the development of large-scale cosmic structure using high-resolution $N$-body simulations. In particular, we focus on the topological properties of the ``cosmic web'', quantitatively characterized by the Minkowski Functionals, for models with quadratic non-linearities with different values of the usual non--Gaussianity parameter $f_{\\rm NL}$. In the weakly non-linear regime (the amplitude of mass density fluctuations $\\sigma_0<0.1$), we find that analytic formulae derived from perturbation theory agree with the numerical results within a few percent of the amplitude of each MF when $|f_{\\rm NL}|<1000$.  In the non-linear regime, the detailed behavior of the MFs as functions of threshold density deviates more strongly from the analytical curves, while the overall amplitude of the primordial non--Gaussian effect remains comparable to the perturbative prediction. When smaller-scale information is included, the influence of primordial non--Gaussianity becomes increasingly significant statistically due to decreasing sample variance. We find that the effect of the primordial non-Gaussianity with $|f_{\\rm NL}|=50$ is comparable to the sample variance of mass density fields with a volume of $0.125(h^{-1}{\\rm Gpc})^3$ when they are smoothed by Gaussian filter at a scale of $5h^{-1}$Mpc. The detectability of this effect in actual galaxy surveys will strongly depend upon residual uncertainties in cosmological parameters and galaxy biasing. ",
        "introduction": "According to the standard scenarios for the formation of large-structure in the Universe, the present-day cosmic density field evolves from small-amplitude initial fluctuations which are described by Gaussian statistics. The hypothesis of primordial Gaussianity is  supported by present observations of the Cosmic Microwave Background (CMB), particularly those from the Wilkinson Microwave Anisotropy Probe (WMAP) \\citep{Komatsu2003,Spergel2007}. These results are consistent with an inflationary origin for the primordial perturbations, since the simplest forms of cosmic inflation produce nearly Gaussian fluctuations.  In order to understand the early Universe in more detail, however, it is necessary to measure (or at least constrain) the departures from non--Gaussianity that inevitably arise at some level during the inflationary epoch. For example, the simplest slowly-rolling single field inflation model predicts very small levels of primordial non--Gaussianity, while multi-field inflation models and models with a non-standard kinetic term for the inflation may yield larger effects which could be detected in ongoing or next-generation observations \\citep[e.g.][]{BMR2002,BU2002,lyth2003,dvali2004,ACMZ2004,AST2004,BKMR2004,Chen2007,BB2007}. Only when such phenomena are detected will it be possible to distinguish between the hundreds of currently viable variations on the theme of inflation by understanding the dynamical behavior of the inflation field.  In order to model the primordial non--Gaussianity that might arise during inflation, the following simple form including quadratic corrections to the curvature perturbation $\\Phi$ \\citep{B1980} during the matter era has been often adopted \\citep{Gangui1994,Verde2000,KS2001}: \\begin{equation} \\label{eq:ngpotential2} \\Phi=\\phi+f_{\\rm NL}(\\phi^2-\\langle\\phi^2\\rangle), \\end{equation} where $\\phi$ represents an auxiliary random-Gaussian field and $f_{\\rm NL}$ characterizes the amplitude of a quadratic correction to the curvature perturbations in a dimensionless way. In principle, $f_{\\rm NL}$ could be scale-dependent, but current observations are not sufficiently sensitive to detect any such variation, so a constant $f_{\\rm NL}$ remains a useful parametrization of the level of non--Gaussianity. Recent analyses of the angular bispectrum for WMAP provides strong constraints on $f_{\\rm NL}$ to lie in the range from $-54$ to $114$ at the 95 percent confidence level \\citep{Komatsu2003,Spergel2007,Creminelli2006}.  The Large-Scale Structure (LSS) of the distribution of galaxies in the Universe provides another potentially powerful probe of primordial non--Gaussianity \\citep{FS1994,CB1996,Verde2000,Scocci2004,HKM2006,SK2007}.  The three-dimensional spatial information arising from LSS is potentially a richer source information about primordial non--Gaussianity than the two-dimensional information arising from the CMB.  For example, constraints from upcoming cluster surveys should be comparable with current CMB limits and those from galaxy surveys, which could be as tight as $|f_{\\rm NL}|\\sim 10$ for the planned surveys and $|f_{\\rm NL}|\\sim 0.2$ for an all-sky survey of galaxies up to redshift $z=5$ \\citep{SK2007,Dalal2007}.  A variety of large-scale projects of LSS observation covering Gpc$^3$ volumes are being proposed, such as an extension of the Sloan Digital Sky Survey; APO-LSS survey; The Hobby-Eberly Dark Energy Experiment (HETDEX) \\citep{Hill2004}; Wide-Field Multi-Object Spectrograph (WFMOS) \\citep{Glazebrook2005}; and the Cosmic Inflation Probe (CIP) mission \\citep{Melnick2004}.  It is consequently important to study the optimal way to extract information about primordial non--Gaussianity from such surveys.  The statistical analysis of non--Gaussianity has been mainly performed through the calculation of the bispectrum \\citep{Verde2000,Scocci2004,SK2007}.  Strong motivation for this is that the bispectrum is the simplest statistical function that can measure quadratic non-linearity \\citep[e.g.][]{WC2003}. Although the quadratic model provides an extremely useful benchmark for statistical analysis techniques, one must always bear in mind that there are many different ways for a random field to be non--Gaussian. In general, there is no one statistic that completely characterizes the statistical nature of a non--Gaussian random field, so a battery of higher-order statistics must be deployed. In particular, when the full nature of non--Gaussianity is virtually unknown, such as is really the case for primordial perturbations, the theoretical model assumed should be validated before its parameters are constrained. Different statistics reflect different aspects of non--Gaussianity so the use of different statistics plays a vital role in this kind of consistency check.  In this paper we use a set of invariant characteristics of the topology of the cosmic web, known as the Minkowski Functionals (MFs). These have already been used to describe the morphological properties of cosmic density fields in a variety of contexts \\citep{MBW1994,SB1997,SG1998,Hikage2003}. Four MFs are defined in three-dimensional density fields such as LSS: the volume fraction ($V_0$); surface area ($V_1$); mean curvature ($V_2$); and Euler characteristic ($V_3$).  Using a perturbative approach, \\citet{HKM2006} derived analytical formulae for the behavior of the MFs for LSS including primordial non--Gaussianity (as a function of $f_{\\rm NL}$ as given in equation [\\ref{eq:ngpotential2}]), in addition to the non--Gaussianity due to non-linear gravity and galaxy biasing. The validity of the perturbative analysis is, however, limited to the weakly non-linear regime.  Smaller-scale modes also contain rich information about the primordial density fields, and this could help place more stringent constraints on primordial non--Gaussianity. In this paper, we use high-resolution $N$-body simulations to study the effect of primordial non--Gaussianity on the MFs from the mildly to strongly non-linear regime. There are two reasons for using the full numerical analysis: one is to see how well the perturbative formulae describe the simulated MFs to check their applicability; the other is to study how the primordial non--Gaussian effect behaves in the strongly non-linear regime and thus to estimate the significance of the effect on the MFs.  The paper is organized as follows. In Section \\ref{sec:pb}, we review the perturbative formulae for the MFs.  The details of the $N$-body simulations and the computing method of the MFs are summarized in Section \\ref{sec:sim}. In Section \\ref{sec:results}, we compare the perturbative formulae of MFs with simulated results to study the primordial non--Gaussian effect in non-linear regime. Section \\ref{sec:summary} is devoted to the summary and conclusions. ",
        "conclusions": "\\label{sec:summary}  We have studied the imprint of primordial non--Gaussianity on the topological properties of LSS using the MFs. Characterizing primordial non--Gaussianity as a quadratic correction to the primordial potential fluctuation with constant amplitude $f_{\\rm NL}$, we compare the MFs with different values of $f_{\\rm NL}$ from the mildly to the strongly non-linear regime using high-resolution $N$-body simulations. Perturbative formulae of the MFs based on the multivariate Edgeworth expansion well reproduce the MFs of simulated mass density fields in the weakly non-linear regime. When the amplitude of the density fluctuation $\\sigma_0<0.1$ and $|f_{\\rm NL}|<1000$, the deviations of the perturbative formulae from simulations are less than a few percent of the amplitude of each MF. They are also $10\\sim 40$ percent in respect to the non-Gaussian contributions alone.  As the fluctuations become more strongly non-linear, the simulated MFs begin to deviate significantly from the perturbative predictions owing to non-linear gravitational evolution. In order to include small-scale information in realistic cosmological data sets, detailed numerical analysis is therefore essential.  When we include information from smaller scale fluctuations, the effects of primordial non--Gaussianity are indeed significant. Using $\\chi^2$ statistics, we find that the primordial non-Gaussianity with $f_{\\rm NL}=50$ has  a significance level corresponding to $1\\sigma$, considering the sample variance of mass density fields at $R=5h^{-1}$Mpc with a volume of $0.125(h^{-1}{\\rm Gpc})^3$. This implies that measuring the MFs in a SDSS-like survey could constrain $f_{\\rm NL}$ at a level comparable with current CMB limits. This is an interesting result, since other observations, like the cluster abundance, that can effectively constrain $f_{\\rm NL}$ at high redshifts, become useless at $z=0$ when non--Gaussian features generated by non-linear dynamics completely obliterate primordial ones \\citep{Grossi2007,Kang2007}.  The actual detectability of the primordial non--Gaussianity is, however, strongly dependent on the degeneracy between the cosmological parameters and the primordial non--Gaussian effect. Understanding the properties of the galaxy biasing is also very important in determining the primordial non--Gaussianity accurately. We will consider this issue in a forthcoming paper."
    },
    "0711/0711.0746_arXiv.txt": {
        "abstract": "We use two very large cosmological simulations to study how the density profiles of relaxed $\\Lambda$CDM dark halos depend on redshift and on halo mass. We confirm that these profiles deviate slightly but systematically from the NFW form and are better approximated by the empirical formula, $d\\log \\rho/d\\log r \\propto r^{\\alpha}$, first used by Einasto to fit star counts in the Milky Way. The best-fit value of the additional shape parameter, $\\alpha$, increases gradually with mass, from $\\alpha\\sim 0.16$ for present-day galaxy halos to $\\alpha\\sim 0.3$ for the rarest and most massive clusters. Halo concentrations depend only weakly on mass at $z=0$, and this dependence weakens further at earlier times. At $z\\sim 3$ the average concentration of relaxed halos does not vary appreciably over the mass range accessible to our simulations ($M \\, \\gsim \\, 3\\times 10^{11} \\, h^{-1} M_\\odot$). Furthermore, in our biggest simulation, the average concentration of the {\\em most massive}, relaxed halos is constant at $\\langle c_{200}\\rangle \\sim 3.5$ to 4 for $0\\leq z\\leq 3$.  These results agree well with those of Zhao et al (2003b) and support the idea that halo densities reflect the density of the universe at the time they formed, as proposed by Navarro, Frenk \\& White (1997). With their original parameters, the NFW prescription overpredicts halo concentrations at high redshift. This shortcoming can be reduced by modifying the definition of halo formation time, although the evolution of the concentrations of Milky Way mass halos is still not reproduced well. In contrast, the much-used revisions of the NFW prescription by Bullock et al. (2001) and Eke, Navarro \\& Steinmetz (2001) predict a steeper drop in concentration at the highest masses and stronger evolution with redshift than are compatible with our numerical data. Modifying the parameters of these models can reduce the discrepancy at high masses, but the overly rapid redshift evolution remains.  These results have important implications for currently planned surveys of distant clusters. ",
        "introduction": "Over the past decade, cosmological N-body simulations have shown consistently that equilibrium dark matter halos have spherically-averaged mass density profiles which are approximately ``universal'' in form; i.e., their shape is independent of mass, of the values of the cosmological parameters, and of the linear power spectrum from which nonlinear structures have grown. As a result, it is useful to parametrize halo profiles by simple empirical formulae, such as that proposed by Navarro, Frenk \\& White (1995, 1996, 1997, hereafter NFW): \\begin{equation} \\frac{\\rho(r)}{\\rho_{\\rm crit}}= \\frac{\\delta_{\\rm c} }{(r/r_{\\rm s})(1+r/r_{\\rm s})^2}, \\label{eq:nfw} \\end{equation} where $\\rho_{\\rm crit}=3H^2/8\\pi G$ is the critical density for closure{\\footnote{We express Hubble's constant as $H(z)$ and its present-day value as $H(z=0)=H_0=100 \\, h$ km s$^{-1}$ Mpc$^{-1}$.\\label{ftn:hz}}}, $\\delta_c$ is a characteristic density contrast, and $r_s$ is a scale radius. Note that this formula contains two scale parameters but no adjustable shape parameter.  As discussed in some detail by NFW and confirmed by subsequent numerical work, the two parameters of the NFW profile do not take arbitrary values, but are instead correlated in a way that reflects the mass-dependence of halo assembly times (e.g., Kravtsov, Klypin \\& Khokhlov 1997; Avila-Reese et al.  1999; Jing 2000; Ghigna et al. 2000; Bullock et al. 2001; Eke, Navarro \\& Steinmetz 2001; Klypin et al. 2001). The basic idea behind this interpretation is that the characteristic density of a halo tracks the mean density of the universe at the time of its formation.  Thus, the later a halo is assembled, the lower its characteristic density, $\\delta_c$, or, equivalently, its ``concentration'' (see Section ~\\ref{ssec:hcat} for a definition).  Although the general validity of these trends is well established, a definitive account of the redshift and mass dependence of halo concentration is still lacking, even for the current concordance cosmology. This is especially true at high masses, where enormous simulation volumes are required in order to collect statistically significant samples of these rare systems. Simulating large cosmological volumes with good mass resolution is a major computational challenge, and until recently our understanding of the mass profile of massive halos has been rather limited, derived largely from small numbers of individual realizations or from extrapolation of models calibrated on different mass scales (NFW; Moore et al. 1998; Ghigna et al. 2000; Klypin et al. 2001; Navarro et al. 2004; Tasitsiomi et al. 2004; Diemand et al. 2004; Reed et al. 2005).  Individual halo simulations may result in biased concentration estimates, depending on the specific selection criteria used to set them up.  In addition, they are unlikely to capture the full scatter resulting from the rich variety of possible halo formation histories.  Extrapolation based on poorly tested models can also produce substantial errors, as recently demonstrated by Neto et al. (2007). These authors analyzed the mass-concentration relation for halos identified at $z=0$ in the {\\it Millennium Simulation} of Springel et al. (2005, hereafter {\\tt MS}) and confirmed the earlier conclusion of Zhao et al. (2003b) that the models of Bullock et al. (2001, hereafter B01) and Eke et al. (2001, hereafter ENS) (which were calibrated to match galaxy-sized halos) severely underestimate the average concentration of massive clusters, by up to a factor of $\\sim 3$.  Estimates of concentrations can also be biased by the inclusion of unrelaxed halos. These often have irregular density profiles caused by major substructures. Smooth density profiles are often poor fits to such halos, and the resulting concentration estimates are ill-defined because they depend on the radial range of the fit and choice of weighting. They can also lead to spurious correlations (see for example figure~9 of Neto et al. 2007). Consequently, in this paper we follow Neto et al. and select only relaxed halos for analysis. This is not without its own problems. Such selection biases against recently formed halos, which may preferentially have lower concentrations. However we believe this is preferable to polluting the sample with meaningless concentration estimates of the kind that arise when smooth spherical models are force-fit to lumpy, multi-modal mass distributions. Hayashi \\& White (2008) stacked {\\it all} halos, regardless of dynamical state, in the {\\tt MS} and studied the resulting mean profiles as a function of halo mass. The relatively small differences between their results and those found below shows that the inclusion of unrelaxed halos has rather little effect in the mean.  A further preoccupation concerns indications that halo profiles deviate slightly but systematically from the NFW model (Moore et al. 1998, Jing \\& Suto 2000, Fukushige \\& Makino 2001, Navarro et al. 2004, Prada et al. 2006, Merritt et al. 2006), raising the possibility that estimating concentrations by force-fitting simple formulae to numerical data may result in subtle biases that could mask the real trends. This is especially important because of hints that such deviations depend systematically on halo mass (Navarro et al. 2004, Merritt et al. 2005). Evaluating and correcting for such deviations is important in order to establish conclusively the mass and redshift dependence of halo concentration.  These uncertainties are unfortunate since observations, especially at high redshift, often focus on exceptional systems. For example, massive galaxy clusters are readily identified in large-scale surveys of the distant universe, and understanding their internal structure will be critical for the correct interpretation of cluster surveys intending to constrain the nature of dark energy. These will make precise measurements of the evolution of cluster abundance in samples detected by gravitational lensing, by their optical or X-ray emission, or through the Sunyaev-Zel'dovich effect (see, e.g., Carlstrom et al. 2002, Hu 2003, Majumdar \\& Mohr 2003, Holder 2006, and references therein).  There is at present no {\\it ab initio} theory that can reliably predict the internal structure of CDM halos. The models of Zhao et al (2003a), Wechsler et al (2002) (see also Lu et al. 2006) are interesting, but as shown in Neto et al (2007) they account for only a small fraction in the measured dispersion in concentration at a fixed mass. There is currently no substitute for direct numerical simulation when detailed predictions are needed for comparison with observation.  We address these issues here by combining results from the {\\tt MS} with results from an additional simulation which followed a substantially smaller volume but with better mass resolution.  This allows us to extend the range of halo masses for which we can measure concentrations and to assess how these measures are affected by numerical resolution. Our analysis procedure follows closely that of Neto et al. (2007). In particular, we concentrate in this paper on the properties of halos which are relaxed according to the criteria defined by these authors; mean density profiles for {\\it all} {\\tt MS} halos of given mass, regardless of dynamical state, are presented by Hayashi \\& White (2008). We begin in Section~\\ref{sec:numexp} by describing briefly the numerical simulations and the halo catalogue on which this study is based. In Section~\\ref{sec:res} we present our main results for the dependence of profile shape and concentration on halo mass and redshift. We conclude with a brief discussion and summary in Section~\\ref{sec:conc}.  \\begin{figure*} \\hspace{-0.5cm} \\resizebox{8cm}{!}{\\includegraphics{figs/f1a.ps}} \\resizebox{8cm}{!}{\\includegraphics{figs/f1b.ps}} \\resizebox{8cm}{!}{\\includegraphics{figs/f1c.ps}} \\resizebox{8cm}{!}{\\includegraphics{figs/f1d.ps}} \\caption{{\\it Left panels:} The stacked density profile of 464 halos of virial mass $ \\sim 2 \\times 10^{14}h^{-1}M_\\odot$, identified at $z=0$ in the {\\tt MS}. The curves in the upper left panel show different NFW fits obtained by varying the radial range fitted: $[0.02,1]r_{\\rm vir}$ (solid red), $[0.05,1]r_{\\rm vir}$ (dashed green) and $[0.1,1]r_{\\rm vir}$ (dot-dashed blue).  Note the small disagreement between the actual profile shape and the NFW model (upper panels). This leads to concentration estimates which depend slightly on the innermost radius of the fit. Fits using eq.~(\\ref{eq:rhoalpha}) are more robust to such variations in fitting range, as shown in the lower left panel. Panels on the right show the residuals corresponding to the fits shown on the left.} \\label{fig:r2rho} \\end{figure*} ",
        "conclusions": "\\label{sec:conc}  We have used data from the {\\it Millennium Simulation} and from a smaller but higher resolution simulation to examine how the density profiles of relaxed $\\Lambda$CDM halos vary with halo mass and redshift. We study profile shape and concentration for halos with mass exceeding $\\sim 3 \\times 10^{11} h^{-1} M_\\odot$ over the redshift range $0\\leq z\\leq 3$.  Our main results may be summarized as follows.  \\begin{itemize}  \\item We confirm the conclusion of previous studies that, in the mean, the shape of spherically averaged $\\Lambda$CDM halo density profiles deviates slightly but systematically from the two-parameter fitting formula proposed by Navarro, Frenk \\& White (1995, 1996, 1997). A more accurate description is provided by the three-parameter Einasto (1965) profile, for which the logarithmic slope is a power-law of radius, $d\\log \\rho/d\\log r \\propto r^{\\alpha}$. We show that this fitting formula gives concentration estimates which are insensitive to the radial range fitted, albeit at the price of an additional shape parameter.  Although Einasto fits avoid small but significant biases that arise if all halos are fit to the NFW model, we emphasise that individual halos typically deviate from either model by more than the difference between them.  \\item Using stacked profiles of halos of similar mass, we show that the shape parameter, $\\alpha$, depends systematically on halo mass and on redshift. These dependences can be collapsed into a dependence on the single ``peak height'' parameter, $\\nu(M,z)=\\delta_{\\rm crit}(z)/\\sigma(M,z)$.  Halos with large $\\nu$ are rare objects in the high-mass tail of the halo mass distribution and have large $\\alpha$ values (see eq.~(\\ref{eq:alphanu})). This provides an important clue for models that attempt to interpret the dependence of halo density profiles on mass and redshift.  \\item The dependence of halo concentration on mass becomes progressively weaker with increasing redshift, as found earlier by Zhao et al (2003b).  At $z = 3$, concentrations are almost independent of mass over the mass range accessible to our simulations.  Relaxed halos with $\\nu > 3$ (the rarest and most massive systems) have similar concentrations, $\\langle c_{200} \\rangle = 3.5$ to 4, at all the redshifts we have studied .  \\item The models of Bullock et al. (2001) and Eke et al. (2001) fail to reproduce our measured concentrations for high-mass and high-redshift objects, predicting a stronger mass dependence and more evolution than is seen in the simulations. Parameters in these models can be changed to reduce the strength of their mass dependence, but the predicted redshift evolution, while fitting galaxy mass halos well up to redshift $z=1$, remains substantially too strong at high mass and at higher redshifts. As a result, the predictions of these models for high-redshift galaxy clusters can be in error by large factors. The original model of Navarro et al. (1997) overpredicts the concentrations of such objects at redshifts beyond 1 (by up to $\\sim 40\\%$ at $z=3$) but a modified NFW model with a different definition of formation redshift reproduces the simulation results substantially better over the redshift and mass ranges we have examined. Both the original and the modified NFW models underestimate the concentration evolution of relaxed $10^{12}h^{-1}M_\\odot$ halos, leading to 30\\% discrepancies at $z=0$ and $z=3$.  \\end{itemize}  We hope that our simulation results will stimulate theoretical work aimed at a deeper understanding of the factors which determine the internal structure of $\\Lambda$CDM halos. Such work may eventually result in simple recipes like those of NFW, ENS and B01. Substantial errors are found, however, when these published prescriptions are extrapolated beyond the regime where their authors tested them, in particular, to the regime relevant to high-redshift galaxy clusters. This demonstrates that careful numerical work is mandatory before any recipe can be applied in a new regime. When making forecasts for surveys of distant massive clusters, our simulations show our modified NFW recipe to give reliable results at least out to $z=3$."
    },
    "0711/0711.0422_arXiv.txt": {
        "abstract": "Transfer equation in a geometrically thin accretion disk is reexamined under the plane-parallel approximation with finite optical depth. Emergent intensity is analytically obtained in the cases with or without internal heating. For large or infinite optical depth, the emergent intensity exhibits a usual limb-darkening effect, where the intensity linearly changes as a function of the direction cosine. For small optical depth, on the other hand, the angle-dependence of the emergent intensity drastically changes. In the case without heating but with uniform incident radiation at the disk equator, the emergent intensity becomes isotropic for small optical depth. In the case with uniform internal heating, the limb brightening takes place for small optical depth. We also emphasize and discuss the limb-darkening effect in an accretion disk for several cases. ",
        "introduction": "Accretion disks are now widely believed to be energy sources in various active phenomena in the universe: in protoplanetary nebulae around young stellar objects, in cataclysmic variables and supersoft X-ray sources, in galactic X-ray binaries and microquasars, and in active galaxies and quasars. Accretion-disk models have been extensively studies during these three decades (see Kato et al. 1998 for a review).  Radiative transfer in the accretion disk has been investigated in relation to the structure of a static disk atmosphere and the spectral energy distribution from the disk surface (e.g., Meyer, Meyer-Hofmeister 1982; Cannizzo, Wheeler 1984; K\\v ri\\v z and Hubeny 1986; Shaviv, Wehrse 1986; Adam et al. 1988; Hubeny 1990; Ross et al. 1992; Artemova et al. 1996; Hubeny, Hubeny 1997, 1998; Hubeny et al. 2000, 2001; Davis et al. 2005; Hui et al. 2005). In many cases the diffusion approximation or Eddington one was employed; it provides a satisfactory description at large optical depth, although the emergent radiation field originates at optical depth of the order of unity. Furthermore, gray and non-gray models of accretion disks were constructed under numerical treatments (K\\v ri\\v z and Hubeny 1986; Shaviv and Wehrse 1986; Adam et al. 1988; Ross et al. 1992; Shimura and Takahara 1993; Hubeny, Hubeny 1997, 1998; Hubeny et al. 2000, 2001; Davis et al. 2005; Hui et al. 2005) and under analytical ones (Hubeny 1990; Artemova et al. 1996).  In these studies, however, the vertical movement and the mass loss were not considered. Hence, recently, in relation to the radiative disk wind, radiative transfer in a moving disk atmosphere was also investigated (e.g., Fukue 2005a, b, 2006a, b). In contrast to the static atmosphere, in the moving atmosphere the boundary condition at the surface of zero optical depth should be modified (Fukue 2005a, b). Moreover, the usual Eddington approximation violates in the highly relativistic flow (Fukue 2005b; see also Turolla, Nobili 1988; Turolla et al. 1995; Dullemond 1999), and the velocity-dependent variable Eddington factor was proposed (Fukue 2006b).   In the usual studies of radiative transfer in the disk, the emergent intensity has not been fully obtained, since the attention were usually focused on the disk internal structure such as a temperature distribution. In addition, the effect of limb darkening has not been well examined except for a few cases relating to cataclysmic variables (e.g., Diaz et al. 1996; Wade, Hubeny 1998; see also Fukue 2000; Hui et al. 2005 for high energy cases).   In this paper, we thus reexamine radiative transfer in the accretion disk with {\\it finite optical depth} under the plane-parallel approximation, and analytically obtain an emergent intensity for the cases with or without internal heating. Besides cataclysmic variables, we also emphasize and discuss the limb-darkening effect in an accretion disk for various cases.   In the next section we describe the basic equations. In section 3, we show analytical solutions. In section 4, we discuss several cases of accretion disk models, and emphasize the importance of the limb-darkening effect. The final section is devoted to concluding remarks. ",
        "conclusions": "In this paper we analytically solve the radiative transfer problem of a geometrically thin accretion disk with finite optical depth, and obtain analytical expressions for emergent intensity from the surface of the disk. For small optical depth, the angle-dependence of the emergent intensity drastically changes from the case with large optical depth. In the case without heating but with uniform incident radiation at the disk equator, the emergent intensity becomes isotropic for small optical depth. In the case with uniform internal heating, the limb brightening takes place for small optical depth. We also emphasize the importance of limb darkening in the accretion disk study, and discuss several cases, including relativistic silhouettes.  In the calculation of spectra from, e.g., cataclysmic variables, the effect of limb darkening was considered (e.g., Diaz et al. 1996; Wade, Hubeny 1998). In the cases of high energy and relativistic regimes, we should also consider limb darkening (cf. Hui et al. 2005).  In addition, if there exist intense radiation sources, such as neutron stars or radiating jets, irradiation takes place and the outer boundary condition changes (e.g., Hubeny 1990). In such cases under strong incident radiation, limb brightening may occur (cf. Stibbs 1971), even for large optical depth.   Finally, if there exists mass loss from the disk surface, the radiative transfer problem becomes much more complicated. In order to examine spectra, fluxes, eclipsing light curves, we must calculate the emergent intensity analytically or numerically.  \\vspace*{1pc}  This work has been supported in part by a Grant-in-Aid for Scientific Research (18540240 J.F.) of the Ministry of Education, Culture, Sports, Science and Technology."
    },
    "0711/0711.1754_arXiv.txt": {
        "abstract": "%  This is the written account of the second of two panel discussions, on {\\it Reconciling observations and modeling of star formation at high redshifts}. The chair of the panel was Pavel Kroupa, and panel members were Marc Balcells, John Beckman, Christopher Conselice, and Joseph Silk. After a short introduction by each of the panelists, panel and audience entered into a lively discussion, centered around the following six themes: the mass function of pre-stellar gas clouds; a possible top-heavy initial mass function at high redshifts versus mini-quasars as the first sources of ionization; the integrated galactic initial mass function; possible differences in specific star formation rates in disks and in massive galaxies; whether merging rates yield a wrong prediction for massive galaxies, and what is the physics behind the onset of the red sequence of galaxies; and the case of dark matter-dominated dwarf galaxies versus tidal dwarf galaxies. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.4989_arXiv.txt": {
        "abstract": "To study the local Hubble flow, we have run constrained dark matter (DM) simulations of the Local Group (LG) in the concordance \\LCDM\\ and OCDM cosmologies, with identical cosmological parameters apart from the $\\Lambda$ term. The simulations were performed within a computational box of $64\\hmpc$ centred on the LG. The initial conditions were constrained by the observed peculiar velocities of galaxies and positions of X-ray nearby clusters of galaxies.  The simulations faithfully reproduce the nearby large scale structure, and in particular the Local Supercluster and the Virgo cluster. LG-like objects have been selected from the DM halos so as to closely resemble the dynamical properties of the LG.  Both the \\LCDM\\ and OCDM simulations show very similar local Hubble flow around the LG-like objects. It follows that, contrary to recent statements, the dark energy (DE) does not manifest itself in the local dynamics. ",
        "introduction": "\\label{sec:intro}     It has been recently stated that the cosmological constant ($\\Lambda$) or its generalisation the dark energy (DE), manifests itself in the dynamics of the local universe (\\citeauthor{bar01} \\citeyear{bar01}, \\citeauthor{che04} \\citeyear{che04}, \\citeauthor{tee05} \\citeyear{tee05}, \\citeauthor{che06} \\citeyear{che06}, \\citeauthor{Chernin_07b} \\citeyear{Chernin_07b}, and \\citeauthor{Chernin_07a} \\citeyear{Chernin_07a}). In these papers, the coldness of the local Hubble flow around the Local Group (LG) has been attributed to the existence DE. This has been supported by \\citet{MGH05} who analysed a set of N-body simulations and concluded that indeed .\\textit{..[their] results provide new, independent evidence for the presence of dark energy on scales of a few   megaparsecs}. These results, if correct, would have provided an independent corroboration  to  the DE component   whose existence is otherwise inferred from observations of distant objects and the early Universe. These authors used the term 'local' as describing the neighborhood of the LG out to a distance of a few Mpc.   Much of the dynamical implications of the cosmological constant for the large scale structure were worked out by \\cite{lah91}. The LG constitutes a quasi-linear object and therefore its dynamics cannot be modeled by the linear theory or the spherical top-hat model. Consequently, we have recently studied the local universe by means of constrained simulations (CSs, \\citeauthor{kra02} \\citeyear{kra02}, \\citeauthor{kly03} \\citeyear{kly03}, \\citeauthor{mvyh07} \\citeyear{mvyh07} and \\citeauthor{HLYD07} \\citeyear{HLYD07}).  The unique feature of the CSs is that their initial conditions are generated as constrained realizations of Gaussian random fields \\citep{HR91}. The initial conditions are constrained by observational data and hence they are designed to reproduce the main gross features of the local large scale structure. As such they provide the optimal tool for studying the dynamics of the LG, being a given individual but not an atypical object. In particular the recent constrained flat $\\Lambda$ dominated ( \\LCDM) and open (OCDM) cold dark matter N-body simulations of \\citet{mvyh07} were designed to study the local dynamics in cold dark matter cosmologies with and without a DE component.  These simulations are to be used here as a laboratory for thesting the hypothesis that the cold local Hubble flow is a signature of dark energy. We are less interested here in the actual coldness of the flow and more in the possibility that the DE affects the local flow. A thorough analysis of the  issue of  the coldness of the local flow is to be given elsewhere (Martinez-Vaquero et al, in preparation). In what follows 'local' is defined as the region contained in a sphere of radius $R=3$ Mpc centred on the LG.   The structure of the paper is as follows. A very brief review of the simulations of \\citet{mvyh07} is presented in Section \\ref{sec:CS}. The selection criteria for LG candidates are summarised in Section \\ref{sec:LG}. The flow fields around the simulated LG candidates in the \\LCDM\\ and OCDM simulations are presented in Section \\ref{sec:Hubble}. In Section \\ref{sec:grav} we compare the gravitational field around the LG candidates. A general discussion concludes the paper (Section \\ref{sec:disc}). ",
        "conclusions": "\\label{sec:disc}   The most striking result of this paper is that the local Hubble flow around LG-like objects in the OCDM model is almost indistinguishable from the \\LCDM\\ flow. To the extent that the models do differ it is the OCDM model that has somewhat colder Hubble flow than the \\LCDM\\ one. It follows that the local flow is not affected by the DE and does not manifest the present epoch dominance of the DE.     One should not be surprised by the departure of the simulated local Hubble flow from the prediction of the simple model proposed by \\cite{Chernin_07a}. First, the actual gravitational field deviates considerably from the predicted one. Second, the gravitational dynamics is not local and the tidal field plays an important role in the quasi-linear regime (\\citeauthor{YH86} \\citeyear{YH86}, \\citeauthor{YH89} \\citeyear{YH89}, \\citeauthor{zar93} \\citeyear{zar93}, \\citeauthor{rien94} \\citeyear{rien94}, \\citeauthor{delpopo1} \\citeyear{delpopo1}). It follows that the dynamics depends not only on  the local field, but is also  affected by the shear, namely the tidal field. The shear breaks the simple linear relation of the density and velocity fields of the linear regime and therefore the local density field cannot account for the local Hubble flow.  The comparison of the \\LCDM\\ and OCDM simulations shows that they yield very similar LG-like objects with virtually identical local Hubble flows. It follows that the dynamical properties of LG-like objects and their environments, in the linear and quasi-linear regime, depend mostly on the cold matter content of the universe, namely $\\Omega_m$, and only weakly on the DE. This is another manifestation of the fact that the properties of the cosmic web, expressed in co-moving coordinates, depend mostly on $\\Omega_m$ and hardly on the DE \\citep{HLYD07}."
    },
    "0711/0711.4171_arXiv.txt": {
        "abstract": "We report on  two {\\sl Chandra} observations of the 3-Myr-old pulsar B1929+10, which reveal a faint compact ($\\sim 9'' \\times 5''$) nebula elongated in the direction perpendicular to the pulsar's proper motion, two patchy wings, and a possible short ($\\sim 3''$) jet emerging from the pulsar. In addition, we detect a tail extending up to at least $4'$ in the direction opposite to the pulsar's proper motion, aligned with the $\\sim15'$-long tail  detected in  {\\sl ROSAT} and {\\sl XMM-Newton} observations. The overall  morphology of the  nebula suggests that the shocked pulsar wind is confined by the ram pressure due to the pulsar's supersonic speed. The shape of the compact nebula in the immediate vicinity of the pulsar seems to be consistent with the current MHD models. However, since these models  do not account yet for the change of the flow velocity at larger distances from the pulsar, they are not able to constrain the  extent of the long pulsar tail. The luminosity of the whole nebula as seen by {\\sl Chandra} is $L_{\\rm PWN} \\sim 10^{30}$ ergs s$^{-1}$ in the 0.3--8 keV band, for the distance of 361 pc. Using the {\\sl Chandra} and {\\sl XMM-Newton} data, we found that the pulsar spectrum is comprised of non-thermal (magnetospheric) and thermal components. The non-thermal component can be described by a power-law model with photon index $\\Gamma\\approx 1.7$ and luminosity $L_{\\rm PSR}^{\\rm nonth}\\approx 1.7\\times 10^{30}$ ergs s$^{-1}$ in the 0.3--10 keV band. The blackbody fit for the thermal component, which presumably emerges from hot polar caps, gives the temperature $kT\\approx 0.3$ keV and projected emitting area $A_{\\perp}\\sim 3\\times 10^3$ m$^2$, corresponding to the bolometric luminosity $L_{\\rm bol}\\sim (1$--$2)\\times 10^{30}$ ergs s$^{-1}$. ",
        "introduction": "\\label{intro}  Pulsar wind nebulae (PWNe), generated when the magnetized relativistic pulsar winds shock in the ambient medium, have been observed around $\\approx 50$ pulsars \\citep[see][for a recent review]{2008AIPC..983..171K}. The shocked pulsar winds produce synchrotron radiation  from radio frequencies through $\\gamma$-rays, revealing properties of the winds and the ambient medium.   The PWN morphology depends on the pulsar's velocity,  pressure, temperature, and magnetic field of the ambient medium, and the geometry of the wind outflow. For instance, the torus-jet PWN structures observed  around young pulsars in supernova remnants (SNRs), such as the Crab and Vela PWNe \\citep[]{2000ApJ...536L..81W,2003ApJ...591.1157P}, are commonly interpreted as anisotropic outflows with equatorial and polar components confined by the pressure of the hot gas in the host SNR interiors, where the speed of sound exceeds the pulsar's speed \\citep[]{2003MNRAS.344L..93K,2004A&A...421.1063D}. When the pulsar is moving with a supersonic speed with respect to the ambient medium, bow-shock PWNe are formed \\citep{2005A&A...434..189B,2005ApJ...630.1020R}, such as the PWNe around PSR J1747$-$2958 \\citep[the Mouse PWN;][]{2004ApJ...616..383G} and PSR B1957--20 \\citep{2003Sci...299.1372S}. The pulsar's supersonic speed causes the ram pressure to exceed  the  ambient pressure, so that  the wind termination shock acquires  a bullet-like shape \\citep[e.g.][]{2005A&A...434..189B} and forms a PWN,   sometimes with a long tail extending behind the pulsar. The pulsar's speed becomes supersonic as it encounters a relatively cold interstellar medium (ISM), either after leaving the SNR in which it was born, or when the host SNR dissolves in the ISM background. These last stages of the SNR evolution occur after $\\sim 10^5$ years.    The majority of middle-aged ($\\sim 10^5$--$10^6$ yr) and old ($\\gtrsim 10^6$ yr) pulsars are expected to  move supersonically through the surrounding ISM, and as long as their  winds are sufficiently strong, they are expected to produce observable bow-shock PWNe. Observations of these objects  provide an opportunity to study  bow-shock morphologies and   probe pulsar winds and their surroundings at later stages of pulsar evolution.     Until now, only a handful of PWNe around middle-aged pulsars have been detected: e.g.,  PSR\\,J0538+2817 \\citep{2003ApJ...585L..41R}, PSR\\,B0355+54 \\citep{2007Ap&SS.308..309M}, and Geminga \\citep{2004MmSAI..75..470C,2006ApJ...643.1146P} in X-rays, PSR\\,B0906--49  at radio frequencies \\citep{1998ApJ...499L..69G}, and PSR\\,B1951+32 in X-rays and H$\\alpha$ \\citep{2004ApJ...610L..33M}. However, one of the first bow-shock PWN candidates was detected serendipitously around a much older pulsar, PSR B1929+10 (hereafter B1929). This is one of the nearest pulsars \\citep[$d=361^{+10}_{-8}$ pc;][]{2004ApJ...604..339C}, with a period $P=226.5$ ms and period derivative $\\dot{P}=1.16\\times 10^{-16}$ s s$^{-1}$. Despite the large spindown age, $\\tau\\equiv P/2\\dot{P} = 3.1$ Myr, its spindown power, $\\dot{E}\\equiv 4\\pi^2 I\\dot{P}/P^3 =3.9\\times 10^{33}$ ergs s$^{-1}$ for the moment of inertia $I=1\\times 10^{45}$ g cm$^2$, is still relatively high. As the pulsar's transverse velocity in the plane of the sky, $V_{\\perp}=177^{+4}_{-5}$\\,km\\,s$^{-1}$ \\citep{2004ApJ...604..339C}, substantially exceeds the typical ISM sound speed, one can expect that the pulsar wind outflow forms a bow-shock PWN. An indication of such a PWN was first noticed by \\citet{1993Natur.364..127W}, who detected a very long ($\\sim1.6$\\,pc) tail-like structure behind this pulsar in  {\\sl ROSAT} data. The tail behind B1929 had been the longest extended structure associated with a compact Galactic object until the recent discovery of a 6-pc-long X-ray tail behind the middle-aged pulsar PSR\\,J1509$-$5850, reported by \\citet{2006HEAD....9.0757K} and described in detail by \\citet{2008arXiv0802.2963K} Similar (albeit shorter) tails have been found behind a few other pulsars \\citep{2008AIPC..983..171K} suggesting that such extended PWN morphologies might be ubiquitous.      The B1929 pulsar  has been extensively observed in X-rays, in an attempt to study the magnetospheric and thermal components of its emission and compare them with the models for pulsar radiation. In particular, the thermal emission from polar caps heated by relativistic particles accelerated in the pulsar magnetosphere is predicted by the current pulsar models \\citep[e.g.,][]{2001ApJ...556..987H,2002ApJ...568..862H}, but a limited sample of the observed old pulsars suggests that the nonthermal magnetospheric emission with a power-law spectrum dominates at higher X-ray energies, $E\\gtrsim 2$ keV, while the nature of radiation at lower energies has been a matter of debate \\citep{2004ApJ...615..908B,2005ApJ...633..367B,2004ApJ...616..452Z,2006ApJ...636..406K}. Observations of B1929 with the {\\sl Einstein} \\citep{1983IAUS..101..471H}, {\\sl ROSAT} \\citep{1994ApJ...429..832Y,1997A&A...326..682B}, and {\\sl ASCA} \\citep{1997ApJ...482L.159W,1998AdSpR..21..213K} observatories were not able to constrain the nature of the X-ray emission from the pulsar because its spectrum  could be equally well described  as thermal or  non-thermal. The same result has been obtained by combining the available {\\sl ROSAT} and {\\sl ASCA} data by \\citet{2005A&A...434.1097S}, who found that the spectrum of B1929 could be fitted  either by a power-law or a  two-temperature blackbody model.      Analyzing the images obtained in the 40 ks {\\sl ROSAT} PSPC observation, \\citet{1993Natur.364..127W} noticed an elongated structure extending approximately 10$'$ behind B1929, almost aligned with the direction of its proper motion. \\citet{1993Natur.364..127W} have suggested that the  detected tail is synchrotron radiation from the  wind of B1929 confined by the ram pressure. The detection of this diffuse emission  has also been reported by \\citet{1994ApJ...429..832Y} in the analysis of the same  {\\sl ROSAT} PSPC data. \\citet{1998AdSpR..21..213K} also claim detection of some diffuse emission in the vicinity of B1929 in the {\\sl ASCA} images, but the poor angular resolution of {\\sl ASCA} prevented a detailed analysis and did not allow to separate the contribution from background sources in the alleged PWN emission.    In a multi-wavelength study of B1929 \\citep{2006ApJ...645.1421B}, the tail-like structure extending up to more than 1\\,pc behind the pulsar has been observed by {\\sl XMM-Newton}, confirming the {\\sl ROSAT} finding. The spectral analysis suggests that this extended emission is nonthermal, likely produced by synchrotron radiation of the shocked pulsar wind. \\citet{2006ApJ...645.1421B} also report on an elongated faint diffuse radio emission found in the Effelsberg radio continuum survey data, whose brightness distribution roughly coincides with the X-ray tail. These authors also conclude that the radiation of the B1929 pulsar is predominantly nonthermal (i.e., magnetospheric), with a rather soft power-law spectrum (photon index $\\Gamma\\approx 2.7$).    To understand the nature of the extended emission behind B1929 and prove that it is indeed a tail of a bow-shock PWN, the nebula in the immediate vicinity of the pulsar should be observed with an angular resolution  much better than those of {\\sl ROSAT} PSPC ($25''$) or {\\sl XMM-Newton} ($15''$ half-energy width). Therefore, we conducted two {\\sl Chandra} observations of the B1929 pulsar and its surroundings. The first results, including the detection of a faint nebula in the immediate vicinity of B1929, have already been reported by \\citet{2006HEAD....9.0758M}. In this paper we present a more detailed analysis of these observations. ",
        "conclusions": "\\label{discussion}   \\subsection{The nebula around PSR\\,B1929+10}  \\subsubsection{PWN morphology}   In the analysis of the two ACIS-S observations of B1929, we detected a faint PWN surrounding this old pulsar. We found that the pulsar is immersed in a very compact, $\\approx 9''\\times 5''$ nebula elongated in the direction perpendicular to the pulsar's proper motion (Fig.\\ 1, bottom). Together with the two patchy ``wings'' (best seen in slightly smoothed images), the overall compact structure looks like a bow shock PWN, with an apex at $\\approx 3''$ ($\\approx 1.6\\times 10^{16}$ cm at $d=361$ pc) ahead of the pulsar. In addition, there is marginal evidence of a short jet-like feature emerging from the pulsar between the wings, in the direction opposite to the pulsar's proper motion (seen up to $3''$--$4''$ from the pulsar in the unsmoothed combined image and in the pulsar-subtracted image in Fig.\\ \\ref{fig-unsmoothed-images}). There is also some very faint emission  in front of the apparent bow-shock apex, seen up to $6''$--$7''$ from the pulsar (see Fig.\\ \\ref{fig-profile}).   In addition to the compact PWN, a faint, inhomogeneous tail-like structure was detected, extending in the direction approximately opposite to the pulsar's proper motion up to at least $\\sim1.5'$--2$\\arcmin$. This structure is well aligned with the much longer tail (up to 10\\arcmin--15\\arcmin) detected in the {\\sl ROSAT} and {\\sl XMM-Newton} data \\citep{1993Natur.364..127W,2006ApJ...645.1421B}.   The surface brightness of the detected PWN is not uniform: e.g., the two wings are brighter than the other PWN structures (see Table~\\ref{table-regions}). Furthermore, the {\\sl Chandra\\/} and {\\sl XMM-Newton\\/} data suggest a possible variability of the detected PWN, although the number of detected photons was too small to establish it unambiguously.    The approximate axial symmetry with respect to the direction of the pulsar's proper motion strongly suggests that the PWN is shaped by the oncoming flow of the ambient matter in the pulsar's reference frame. The corrected pulsar's proper motion and distance correspond to the pulsar's speed $v_{\\rm PSR} = 178\\, (\\sin i)^{-1}$ km s$^{-1}$, where $i$ is the angle between the velocity vector and the line of sight. It exceeds the ISM sound speed, $c_s = (5kT/3\\mu m_H)^{1/2} = 37 \\mu^{-1/2} T_5^{1/2}$ km s$^{-1}$ (where $\\mu$ is the molecular weight, $T_5=T/10^5$ K), for $T< 2.3\\times 10^6 \\mu\\, (\\sin i)^{-2}$ K, i.e., for virtually any plausible ISM temperature, which supports the bow-shock interpretation of the detected nebular emission.    Our current understanding of the pulsar wind dynamics in a bow-shock PWN mostly relies upon numerical models. A recent work in this direction has been done by \\citet{2003A&A...404..939V}, \\citet{2005A&A...434..189B}, and \\citet{2007MNRAS.374..793V}. These models show that at supersonic pulsar speeds the  termination shock (TS) of an {\\em isotropic} pulsar wind acquires a bullet-like shape, with a distance \\begin{equation} R_h \\simeq (\\dot{E}/4\\pi c p_{\\rm ram})^{1/2}\\, \\end{equation} between the pulsar and the bullet head, where \\begin{equation} p_{\\rm ram} = \\rho_{\\rm amb}\\,v_{\\rm PSR}^2 \\end{equation} is the ram pressure, and $\\rho_{\\rm amb}$ is the ambient density. The shocked pulsar wind is confined between the TS and the contact discontinuity (CD) surface, while the forward bow shock (FBS) separates the shocked ambient medium (between the CD and the FBS) from the unshocked one. For large Mach numbers, $\\mathcal{M}\\equiv v_{\\rm \\sc PSR}/c_{s} =(3p_{\\rm ram}/5p_{\\rm amb})^{1/2} \\gg 1$, and small values of the magnetization parameter\\footnote{The magnetization parameter in the pre-shock wind is defined as the ratio of the Poynting flux to the kinetic energy flux.} of the pre-shock pulsar wind, $\\sigma < 0.1$,  the bullet's cylindrical radius is $r_{\\rm TS}\\sim R_{h}$ and the distance of its back surface from the pulsar is $R_{b}\\sim 6R_{h}$ \\citep[see Figs.\\ 1--3 in][]{2005A&A...434..189B}. The shape of the CD surface ahead of the pulsar is similar to that of the TS but with the apex at $R_{\\rm CD}\\approx 1.3 R_h$, while  it acquires a cylindrical shape with a radius $r_{\\rm CD}\\approx 4 R_h$ behind the TS bullet. The shocked wind flows away from the pulsar within this cylinder, forming a PWN tail. The flow velocity in the central part of the tail (the inner channel, $r\\lesssim r_{\\rm TS}$) is $\\sim 0.1c$--$0.2c$, while it is as high as $v_{\\rm tail}=0.8c$--$0.9c$ in the bulk of the tail's volume, at $r_{\\rm TS} \\lesssim r < r_{\\rm CD}$ (where $r$ is the cylindrical radius). The magnetic field, purely toroidal in these models, also depends on $r$, being enhanced toward the tail axis and the CD surface. Its typical value in the tail is \\begin{equation} B_{\\rm tail}\\sim 2.8\\, (c/v_{\\rm tail})\\, (\\sigma p_{\\rm ram})^{1/2} {\\rm G}, \\end{equation} where the $p_{\\rm ram}$ is in units of dyn cm$^{-2}$ \\citep[see eq.\\ (13) of][]{2005A&A...434..189B}. In the numerical simulations, the flow velocity and the magnetic field do not show substantial changes along the tail, but these simulations are limited to relatively short distances from the pulsar \\citep[e.g., $<13 R_h$ in][]{2005A&A...434..189B}.      To interpret the B1929 PWN in the framework of the current MHD models, we attempted to match the observed morphological structures with those predicted by the simulations. As the brightest X-ray emission is expected from the shocked wind between the TS and CD in the front part of the PWN, we assume that the very compact structure in the immediate vicinity of the pulsar (Fig.~\\ref{fig-unsmoothed-images}, bottom) is the PWN head and estimate the distance from the pulsar to the CD apex to be $\\sim 3''$. This corresponds to $R_{\\rm CD}\\sim 1.6\\times 10^{16}$ cm, $R_h \\sim r_{\\rm TS}\\sim 1.2\\times 10^{16}$ cm (or $2.3''$), $R_b\\sim 7.5\\times 10^{16}$ cm (or $14''$), and $r_{\\rm CD}\\sim 4.8\\times 10^{16}$ cm (or $9''$). In this picture, the patchy wings in the PWN image are produced by the shocked wind between the TS and CD,  while the ``blobs'' farther behind the pulsar are in the pulsar tail. However, the faint emission ahead of the pulsar (region 3, or the Front, in Fig.~\\ref{fig-unsmoothed-images}, top, and Table~\\ref{table-regions}) is not explained in this model. We can only speculate that this emission might be produced by a polar outflow along the pulsar's spin axis (not included in the models), possibly scattered by the head wind, or it might suggest that, because of the Rayleigh-Taylor instability, the CD surface acquires an irregular shape ahead of the pulsar, with ``fingers'' of the shocked wind penetrating into the shocked ISM. As the above estimates of sizes of the PWN elements are based on simplified models and shallow images, one should consider them as order-of-magnitude estimates rather than accurately established values. Therefore, in the following estimates, we will retain the explicit dependence on $R_h$, scaling it as $R_h=R_{h,16}\\times 10^{16}$ cm.   Using this estimate for $R_h$ and equation (1), we can calculate the ram pressure: \\begin{equation} p_{\\rm ram}= \\dot{E}f_\\Omega (4\\pi cR_h^2)^{-1} = 1.0\\times 10^{-10} f_\\Omega R_{h,16}^{-2}\\,\\, {\\rm dyn}\\,\\,{\\rm cm}^{-2}, \\end{equation} where we introduced the factor $f_\\Omega$ to account for possible anisotropy of the pulsar outflow (e.g., $f_\\Omega <1$ for a mostly equatorial outflow if the equatorial plane is perpendicular to the pulsar's velocity). Since, on the other hand, \\begin{equation} p_{\\rm ram}=\\rho_{\\rm amb}V_{\\rm PSR}^2 =5.3\\times 10^{-10} n_{\\rm amb}\\, (\\sin i)^{-2}\\,\\,{\\rm dyn}\\, {\\rm cm}^{-2}\\,, \\end{equation} where $n_{\\rm amb}\\equiv \\rho_{\\rm amb}/m_H$, we obtain an estimate for the ambient number density: \\begin{equation} n_{\\rm amb}= 0.19 f_\\Omega R_{h,16}^{-2} \\sin^2i\\,\\,\\,  {\\rm cm}^{-3}\\,. \\end{equation}    Using equation (4), we can estimate the Mach number: \\begin{equation} \\mathcal{M}= 7.7 f_\\Omega^{1/2}R_{h,16}^{-1}p_{\\rm amb,-12}^{-1/2}\\,, \\end{equation} where we scaled $p_{\\rm amb}$ to $10^{-12}$ dyn cm$^{-2}$, a typical value for the ISM. From the corresponding sound speed, $c_s = 23 f_\\Omega^{-1/2} R_{h,16} p_{\\rm amb,-12}^{1/2} (\\sin i)^{-1}$ km s$^{-1}$, we estimate the temperature: \\begin{equation} T=3.9\\times 10^4 \\mu f_\\Omega^{-1} R_{h,16}^2 p_{\\rm amb,-12}(\\sin i)^{-2}\\, {\\rm K}\\,. \\end{equation} The ambient pressure in the unperturbed ISM is in the range $p_{\\rm amb,-12}\\approx 0.2$--2 \\citep{2001ApJ...551L.105H,2001RvMP...73.1031F}. The pressure in the pulsar's neighborhood might be somewhat higher if the pulsar's UV and soft X-ray emission heats and ionizes the surrounding medium, which is supported by the lack of an H$_\\alpha$ bow-shock PWN in the observations reported by \\citet{2006ApJ...645.1421B} \\citep[see][for a detailed discussion on conditions of observability of H$_\\alpha$ bow-shock PWNe]{2001A&A...375.1032B}. As the factor $\\mu f_\\Omega^{-1} R_{h,16}^2 (\\sin i)^{-2}$ is also very likely greater than unity, we expect $T\\sim 10^5$ K to be a realistic estimate for the ambient temperature.   We should note that although the comparison of our data with the numerical models leads to reasonable estimates for the parameters of the ambient medium, these two-dimensional models are based on a number of assumptions that are not necessarily realistic. In particular, the models assume an isotropic pulsar outflow, whereas we know that it is predominantly equatorial, at least in the case of young pulsars. Obviously, a substantial anisotropy would distort the shape of the bow shock in the pulsar vicinity (e.g., decreasing $R_h$ in the case of an equatorial outflow in the plane perpendicular to the pulsar velocity). For B1929, a hint of an equatorial outflow might be suggested by the elongation of the most compact PWN structure perpendicular to the proper motion (Fig.~\\ref{fig-unsmoothed-images}, bottom). Such a structure, together with the ``wings'', might be interpreted as a torus of the shocked equatorial wind, compressed and possibly bent by the oncoming wind. Such an interpretation might also be supported by the marginal detection of the jet-like structure immediately behind the pulsar and the faint emission ahead of the pulsar (a crushed counter-jet?).         Application of the PWN simulations to the observed tail leads to some discrepancies. For instance, although the tail may look like a cylinder with the predicted diameter of $\\sim 18''$ at the relatively small distances, $\\lesssim 2'$, from the pulsar in the ACIS images, its transverse size apparently increases at larger distances, possibly reaching $\\sim 5'$ at the edge of the MOS field of view ($15'$ from the pulsar). The tail's surface brightness looks very patchy, which might suggest some interaction of the shocked wind with the surrounding medium, such as the Kelvin-Helmholtz instability caused by the shear at the CD surface. This instability may produce clumps of shocked ISM embedded in the shocked wind, which would distort the flow pattern and compress the wind material \\citep{2005A&A...434..189B}. The same instability may result in temporal variations in the X-ray emission of the tail, with a time scale of $r_{\\rm CD}/c \\sim 0.5$ months, which might explain the alleged variability of the PWN (\\S~\\ref{variability}).   Using equations (3) and (4), the magnetic field in the tail can be estimated as \\begin{equation} B_{\\rm tail} = 56 \\sigma^{1/2} (0.5c/v_{\\rm tail}) f_\\Omega^{1/2} R_{h,16}^{-1}\\,\\,\\,\\mu{\\rm G}\\,. \\end{equation} We should note that this field becomes rather small, comparable to the interstellar magnetic field, at the  usually assumed small values of the magnetization parameter, $\\sigma\\lesssim 10^{-2}$.    Using this estimate of $B_{\\rm tail}$, we can estimate the synchrotron cooling time for relativistic electrons in the tail, $\\tau_{\\rm syn}=5.1\\times 10^{10} \\gamma_8^{-1}B_{-5}^{-2}$ s, where $\\gamma=10^8 \\gamma_8$ is the electron Lorentz factor and $B_{-5}=B/10^{-5}\\,{\\rm G}$. As the characteristic energy of synchrotron photons is \\begin{equation} E \\sim 0.5 B_{-5} \\gamma_8^2\\,\\,\\, {\\rm keV}, \\end{equation} photons with energy $E$ are emitted by electrons with $\\gamma_8\\sim 2 (E/1\\,{\\rm keV})^{1/2} B_{-5}^{-1/2}$, and the synchrotron cooling time corresponding to $E=8$ keV photons (upper energy of our band) is \\begin{equation} \\tau_{\\rm syn}= 6.8\\times 10^8 (\\sigma f_\\Omega)^{-3/4} (v_{\\rm flow}/0.5c)^{3/2} (E/8\\,{\\rm keV})^{-1/2} R_{h,16}^{3/2}\\,\\, {\\rm s}. \\end{equation} It follows from this equation that the projected tail length, $l_{\\rm tail}\\sim v_{\\rm flow} \\tau_{\\rm syn}\\,\\sin i$, as observed at energy $E$, can be estimated as \\begin{equation} l_{\\rm tail} \\sim 1\\times 10^{19} \\left(\\frac{v_{\\rm flow}}{0.5c}\\right)^{5/2} \\left(\\frac{8\\,{\\rm keV}}{E}\\right)^{1/2} \\left(\\frac{R_{h,16}^2}{\\sigma f_\\Omega}\\right)^{3/4}\\sin i\\,\\,\\, {\\rm cm}. \\end{equation} For instance, for $\\sigma=0.01$, $R_{h,16}f_\\Omega^{-1/2}=2$, $v_{\\rm flow}=0.5c$, $\\sin i=0.5$, and $E<8$ keV we obtain $l_{\\rm tail}\\gtrsim 150$ pc (i.e., $\\sim 24^\\circ$), two orders of magnitude larger than observed by {\\sl XMM-Newton} and {\\sl ROSAT}. The main reason for this discrepancy is the assumption that the flow speed remains very high along the entire tail in the ideal MHD model. \\citet{2005A&A...434..189B} have noticed that a similar discrepancy arises when the same model is applied to the X-ray tail of the Mouse PWN: the observed tail length is much shorter than that estimated from the model at any $\\sigma <1$. These authors suggest that the flow can be slowed down by the interaction with the ambient medium, which could also explain the divergence of the Mouse tail in the radio. As these factors have not been included in the current models, we can only empirically estimate an average flow velocity that would be consistent with the observed length of 1.5 pc: $v_{\\rm flow} \\sim 0.1 c\\, \\sigma_{-2}^{3/10} (f_\\Omega/R_{h,16}^2)^{3/10} (\\sin i)^{-2/5}$, assuming that equation (3) is still applicable (here $\\sigma_{-2}=\\sigma/10^{-2}$). We note that this velocity is still much larger than the pulsar velocity. This means that the equation \\begin{equation} l_{\\rm tail}= v_{\\rm PSR,\\perp} \\tau_{\\rm syn} \\approx 0.3 (v_{\\rm PSR,\\perp}/178\\,\\,{\\rm km}\\,\\,{\\rm s}^{-1}) \\gamma_8^{-1} B_{-5}^{-2}\\,\\,\\, {\\rm pc}, \\end{equation} often used for estimating the magnetic field \\citep[e.g.,][]{2006ApJ...645.1421B}, is inapplicable in this case\\footnote{Equation (13) implies that the shocked wind flow acquires the speed of the ambient ISM matter in the immediate vicinity of the pulsar, so that the observed elongated PWN is actually a ``trail'' of the decelerated wind left behind the moving pulsar. Such an assumption strongly contradicts all the models of bow-shock PWNe \\citep[e.g.,][]{2005ApJ...630.1020R,2005A&A...434..189B}. \\citet{2006ApJ...645.1421B} obtained a reasonable estimate for the magnetic field, $B\\lesssim 12$ $\\mu$G, using eq.\\ (13) because they assumed $\\gamma=10^6$. However, at such $B$ and $\\gamma$, the energy of a synchrotron photon is $E=6$ eV (see eq.\\ [10]), well below the X-ray band.}.    It is worthwhile to mention that the length of the tail of B1929 measured so far might be limited by the field of view of the available {\\sl XMM-Newton} observations, and further deep observations along the tail are needed to determine its full extent. Recent observations with {\\sl Chandra} and {\\sl XMM-Newton} have discovered a number of other PWNe with very long tails \\citep{2008AIPC..983..171K}, including the 6\\,pc tail ($l_{\\rm tail}/R_{h}\\sim 600$) behind PSR J1509$-$5850, the longest pulsar tail known to date \\citep{2008arXiv0802.2963K}. Detection of more such objects would help to facilitate further modeling that would account for cooling in the varying magnetic field  at larger distances from the pulsar.   \\subsubsection{X-ray luminosity and spectrum of the PWN}   From our  estimate of the total unabsorbed flux of the B1929 PWN detected with {\\sl Chandra}, we estimated the X-ray PWN luminosity, $L_{\\rm PWN}=(5.3$--$8.6)\\times 10^{29}$ ergs s$^{-1}$, which corresponds to the PWN efficiency $\\eta_{\\rm PWN} \\equiv L_{\\rm PWN}/\\dot E = (1.4$--$2.2)\\times 10^{-4}$ in the 0.3--8\\,keV band. Similar values of the luminosity and efficiency have been measured in the {\\sl XMM-Newton} data by \\citet{2006ApJ...645.1421B} from a larger area farther southwest from the pulsar. As both measurements were taken for small fractions of the PWN, the luminosity and efficiency of the entire PWN can be higher than these estimates.   The estimated efficiency of the B1929 PWN is within the range $\\eta_{\\rm PWN}\\sim 10^{-4.5}$--$10^{-2}$ found by \\citet{2007ApJ...660.1413K} in a recent study of several energetic middle-aged Vela-like pulsars and their PWNe. Compared to a few known PWNe with long tails, which, on average, show higher efficiencies, $\\eta_{\\rm PWN}\\sim 10^{-3.8}$--$10^{-1.7}$, than more compact PWNe \\citep[see][]{2008arXiv0802.2963K}, the B1929 PWN is among the least efficient ones; however, we stress that a deeper observation of B1929 is needed to measure the efficiency more accurately.  The spectral slope of the PWN spectrum is also consistent with the values measured for other PWNe \\citep[e.g., see][]{2007ApJ...660.1413K,2008AIPC..983..171K}.   \\subsection{Spectral properties of PSR\\,B1929+10} \\label{spectral-discussion}  As we mention in \\S\\ref{intro}, it has been a matter of debate whether the X-ray emission from B1929 is predominantly magnetospheric or it has a significant thermal component emitted from the neutron star surface. Most recently, \\citet{2006ApJ...645.1421B} have concluded, based on the analysis of the {\\sl XMM-Newton} data, that the spectrum of B1929 is best described by a PL model (i.e.\\  the emission is predominantly magnetospheric), while the contribution of the thermal component, modeled as BB radiation, does not exceed $\\sim$7\\%. Our analysis of the {\\sl Chandra} and {\\sl XMM-Newton} spectra has shown, however, that adding the BB component to the PL model significantly improves the fit to the combined {\\sl Chandra} plus {\\sl XMM-Newton} spectrum, with the the best-fit BB component providing $\\sim$45\\% of the luminosity in the 0.3--10 keV band.   In addition to the better fit, there are other serious arguments in favor of the PL+BB model. First of all, the hydrogen column density, $N_{\\rm H}=(2.23\\pm0.27)\\times 10^{21}$ cm$^{-2}$, obtained from the PL fit, looks unreasonably large. In particular, it is much larger than the standard estimate, $N_{\\rm H}\\sim 10 N_e \\approx 1\\times 10^{20}$ cm$^{-2}$, obtained from the pulsar's dispersion measure, DM = 3.178 pc cm$^{-3}$, under the usual assumption of a 10\\% ISM ionization (where $N_e$ is the electron column density). In other words, the PL model requires a very low ISM ionization, $\\sim$0.4\\%, (i.e., a very large ratio $N_{\\rm H}/N_e \\sim 230$, much larger than for any other radio pulsar detected in X-rays). Also, taking into account that B1929 is a nearby pulsar ($d=361$ pc), the $N_{\\rm H}$ obtained from the PL fit is uncomfortably close to the HI column density, $N_{\\rm HI}=3.5\\times 10^{21}$, throughout the entire Galaxy in the direction of B1929 \\citep{2005A&A...440..775K}. On the contrary, $N_{\\rm H}= 0.17^{+0.23}_{-0.17}\\times 10^{21}$ cm$^{-2}$ inferred from the PL+BB fit is consistent with the usual estimate based on the dispersion measure as well as with the $N_{\\rm H}$ estimated from the PWN spectrum (see Fig.\\ \\ref{cont-PWN}), and it is much lower than the total Galactic $N_{\\rm HI}$.   Also, the PL+BB model looks more attractive than the PL model because it gives the slope of the PL component, $\\Gamma\\approx 1.7$, within the range of spectral slopes found for a large sample of young pulsars \\citep[$\\Gamma \\approx 1-2$;][]{2003ApJ...591..361G}, while the PL fit to the B1929 spectrum results in a considerably steeper slope, $\\Gamma \\approx 3$. One could argue that the PL fits to the spectra of other old pulsars \\citep[e.g., B0950+08, B2224+65, B0823+26, B0628--28, B1133+16, B0943+10;][]{2004ApJ...616..452Z,2004ApJ...615..908B,2005ApJ...630L..57T,2006ApJ...636..406K,2005ApJ...624L.109Z} also show rather steep slopes, $\\Gamma\\approx 2-3$, suggesting that pulsar spectra might soften with increasing age or decreasing spin-down power \\citep{2006ApJ...636..406K}. However, a plausible alternative interpretation of the softer spectra of old pulsars is that they are, in fact, comprised of a (soft) thermal component and a PL component with a more gradual slope\\footnote{Note that the observed spectra of these pulsars can be satisfactorily fitted with quite different models because these objects are very faint in the X-ray range.}, similar to those found in young pulsars \\citep{2004ApJ...616..452Z}.   Thus, based on the goodness of fit and the astrophysical arguments, we conclude that the X-ray emission from B1929 most likely includes both thermal and magnetospheric components. The BB model for the thermal component gives the apparent temperature $kT_a\\approx0.3$\\,keV and projected emitting area $A_{\\perp,a}\\sim 3000$ m$^2$, at the distance of 361 pc, suggesting that the thermal emission originates from small heated spots (e.g., polar caps). These temperature and area correspond to the apparent radius $R_a^{\\rm es}=(A_{\\perp,a}/\\pi)^{1/2}\\approx 33$ m and bolometric luminosity $L_{{\\rm bol},a}^{\\rm es}=4 A_{\\perp,a}\\sigma T_a^4 \\approx 1.1\\times 10^{30}$ ergs s$^{-1}$ of an equivalent sphere. The true size and luminosity of the polar caps depend on the geometry and the gravitational redshift factor, $g_r=(1-R_g/R_{\\rm NS})^{1/2}$, where $R_g=2.95 M_{\\rm NS}/M_\\odot$ km, $R_{\\rm NS}$ and $M_{\\rm NS}$ are the mass and radius of the neutron star. If there are two identical hot spots at the poles of a centered magnetic dipole, the polar cap radius and the luminosity of two polar caps, as measured at the neutron star surface, are $R_{\\rm pc}= R_{a}^{\\rm es}f^{-1/2}$ and $L_{\\rm bol,pc}= L_{{\\rm bol},a}^{\\rm es}/(2fg_r^4)$, where $f\\leq 1$ is a geometrical factor depending on the angles $\\zeta$ beween the line of sight and the spin axis and $\\alpha$ between the spin and magnetic axes, as well as on $g_r$ \\citep{2007ApJ...664.1072P}. For instance, assuming $M_{\\rm NS}=1.4 M_\\odot$, $R_{\\rm NS}=10$ km (i.e.\\ $g_r=0.766$), $\\alpha=36.0^\\circ$, and $\\zeta=61.5^\\circ$ \\citep{2001ApJ...553..341E} and using the approach described by \\citet{1995AstL...21..149Z} and \\citet{2002ApJ...566L..85B}, we find $f=0.897$, which gives $R_{\\rm pc}= 1.06 R_{a}^{\\rm es}\\approx 35$ m and $L_{\\rm bol,pc}=1.62 L_{{\\rm bol},a}^{\\rm es} \\approx 1.8\\times 10^{30}\\, {\\rm ergs\\, s}^{-1}=4.7\\times 10^{-4}\\dot{E}$.          The estimated polar cap radius is a factor of $\\sim$9 smaller than the conventional polar cap radius of a radio pulsar, $\\tilde{R}_{\\rm pc} = (2 \\pi R_{\\rm NS}^3/cP)^{1/2}$ \\citep[see e.g.,][]{1991tnsm.book.....M}, which is $\\approx 300$ m for B1929, assuming $R_{\\rm NS} =10$ km. However, the observed polar cap radii can be substantially smaller or larger than the conventional value. In particular, BB fits of several old pulsar spectra show $R_{\\rm pc}\\sim (0.1$--$0.2)\\, \\tilde{R}_{\\rm pc}$ \\citep[e.g.,][]{2004ApJ...616..452Z,2005ApJ...624L.109Z,2006ApJ...636..406K}. For some pulsars, such a discrepancy can be alleviated assuming that the polar cap is covered by a hydrogen or helium atmosphere, in which case the effective temperature would be a factor of 2 lower, and the radius a factor of 3--10 larger while the bolometric luminosity would not change significantly \\citep{1995lns..conf...71P,2004ApJ...616..452Z}. Another explanation for such a discrepancy was proposed by \\citet{2005ApJ...624L.109Z}, who suggested that only a small fraction of the polar cap surface, associated with footprints of sparks produced by intermittent breakdowns of an ``inner gap'' above the polar cap, is hot enough to emit X-rays.    The luminosity of the detected thermal component, which is less dependent on the assumed spectral model than the radius and temperature, can be compared with the models for pulsar polar cap heating \\citep{2002ApJ...568..862H,2001ApJ...556..987H}. For B1929, these models predict the polar cap luminosity $L_{\\rm bol,pc}\\sim 10^{31}$\\,ergs\\,s$^{-1}$ if the polar cap is heated by positrons produced through curvature radiation of electrons accelerated in the neutron star magnetosphere. Although the polar cap thermal luminosity estimated from the PL+BB spectral model is a factor of a few lower than that predicted by \\citet{2001ApJ...556..987H}, this can be considered as a reasonably good agreement, with account of the uncertainties of both the theoretical model and the observational results. We expect that deeper observations of other old pulsars would also detect thermal components from their polar caps and help us better understand the polar cap heating mechanisms.    As Figure~\\ref{fig-combined-spectra-pl}  shows, the thermal component dominates at $0.5\\,{\\rm keV}\\lesssim E \\lesssim 2\\,{\\rm keV}$. This means that we can expect different pulse shapes within and outside this energy range, because the pulsations of thermal radiation should be smoother and shallower than those of the magnetospheric radiation. The study of the energy dependence of pulsations not only can confirm the presence of the thermal component, but it would also allow one to infer the polar cap geometry and emission mechanism (e.g., BB vs.\\ a light-element atmosphere), and even measure the mass-to-radius ratio for the neutron star \\citep{1995AstL...21..149Z,1997ApJ...490L..91P,2004ApJ...616..452Z}. Since the current {\\sl XMM-Newton} data with sufficient time resolution do not have enough counts for such an analysis, a new deep observation of B1929 is needed to perform this important test.   The spectral slope of the PL (magnetospheric) component inferred from the PL+BB fit is $\\Gamma\\approx 1.7$, similar to those of younger pulsars. The luminosity of this component, $L_X= 4\\pi d^2 F_X^{\\rm unabs} \\approx 1.7\\times 10^{30}$ ergs s$^{-1}$ in the 0.3--10 keV band, correponds to the nonthermal X-ray efficiency $\\eta_{\\rm nonth}=L_X/\\dot{E}\\approx 4.4\\times 10^{-4}$, typical for the whole population of radio pulsars detected in the X-rays, and comparable to the thermal (polar cap) efficiency. Notice that the one-component PL fit results not only in a much softer spectrum ($\\Gamma\\approx 3.0$) but also in a higher luminosity, $L_X\\approx 6.6\\times 10^{30}$ ergs s$^{-1}$, and efficiency, $\\eta_{\\rm nonth}\\approx 1.7\\times 10^{-3}$, above the typical value for X-ray emitting radio pulsars. A similar trend has been observed in other old pulsars: PL fits of their spectra yield X-ray efficencies noticeably higher that those of young and middle-aged pulsars, which might suggest that the X-ray efficiency in this energy range grows with increasing age and decreasing spin-down power, perhaps at the expence of the efficiency at higher photon energies. However, this conclusion is based on the assumption that the contribution of the thermal component is negligible in the soft X-ray range, which we believe is not valid at least for the best-studied old pulsars, B1929 and B0950+08. To understand the evolution of X-ray properties of pulsars, deep observations of a larger sample of old pulsars are warranted.        \\subsection{Summary}  Two ACIS-S observations of B1929 revealed a faint PWN surrounding this old, nearby pulsar. The observed morphology includes a compact nebula with two patchy wings in the immediate vicinity of B1929, and a tail extending in the direction opposite to the pulsar's proper motion, aligned with the much longer tail detected in the previous {\\sl ROSAT} and {\\sl XMM-Newton} observations. The shape of the nebula and its spectral properties are  consistent with the proposed bow-shock classification. The properties of the compact nebula suggest that the pulsar wind outflow is anisotropic, possibly concentrated toward the equatorial plane perpendicular to the pulsar velocity. The size of the observed tail implies an average flow velocity of $\\sim 0.1 c$. The whole observed PWN radiates about $2\\times 10^{-4}$ of the pulsar spin-down power in X-rays. By comparing the observed PWN properties with the bow-shock MHD models, we estimated the temperature  of the local ISM to be $\\sim10^5$\\,K, which is consistent with the lack of the H$\\alpha$ bow-shock emission around this pulsar.   In the spectral analysis of the combined {\\sl Chandra} and {\\sl XMM-Newton} data, we detected  a thermal component in the pulsar radiation, whose luminosity, $\\sim (1$--$2) \\times 10^{30}$\\,ergs\\,s$^{-1}$, is comparable with the magnetospheric luminosity in the X-ray band. The thermal radiation is likely emanating from polar caps heated by positrons created by the curvature radiation of ultrarelativistic electrons accelerated in the pulsar magnetosphere. The spectrum and X-ray efficiency of the magnetospheric component are similar to those found in young and middle-aged pulsars. Further X-ray observations of the pulsar will allow one to better characterize its spectral and timing properties and understand the nature of the magnetospheric radiation and the mechanisms of polar cap heating.     We thank Patric Broos for the useful discussion about the MARX simulations, Oleg Kargaltsev and Leisa Townsley for  the help with the high-resolution imaging techniques, and Slava Zavlin for the discussion about the {\\sl XMM-Newton} spectral results. We wouls also like to thank the referee for very helpful comments and suggestions. This work was partially supported by NASA grant NAG5-10865 and {\\sl Chandra} award SV4-74018."
    },
    "0711/0711.2946_arXiv.txt": {
        "abstract": "{} {The knowledge of the properties of the youngest radio sources is very important in order to trace the earliest phase of the evolution of the radio emission. RXJ1459+3337, with its high turnover frequency ($\\sim$ 25 GHz) provides a unique opportunity to study this class of extreme objects.} {High-sensitivity multi-frequency VLA observations have been carried out to measure the flux-density with high accuracy, while multi-frequency VLBA observations were performed, aimed at determining the pc-scale structure. Archival ROSAT data have been used to infer the X-ray luminosity.} { The comparison between our new VLA data and those available in the literature shows a steady increment of the flux-density in the optically-thick part of the spectrum and a decrement of the turnover frequency. In the optically-thin regime, the source flux density has already started to decrease. Such a variability can be explained in terms of an adiabatically-expanding homogeneous radio component. The frequency range spanned by our VLBA observations, together with the resolution achieved, allows us to determine the source size and the turnover frequency, and then to derive the magnetic field directly from these observable quantities. The value obtained in this way is in good agreement with that computed assuming equipartition condition. A similar value is also obtained by comparing the radio and X-ray luminosities.} {} ",
        "introduction": "The radio emission of extragalactic sources is synchrotron radiation produced by relativistic electrons with a power-law energy distribution. The relativistic electrons produce a power-law radio spectrum which reaches its maximum in correspondence to the turnover frequency $\\nu_t$. \\\\ At frequencies below $\\nu_{t}$, the spectrum turns over likely due to Synchrotron Self-Absorption (Snellen et al. \\cite{sn00}), although free-free absorption may also play a role (Mutoh et al. \\cite{mutoh02}, Kameno et al. \\cite{kameno00}).\\\\ The evolutionary models proposed to interpret the various stages of the life-cycle of radio sources relate the typical spectral peak of very small and then compact region with age. In the evolutionary scenario, the peak of the spectrum progressively moves toward lower frequencies as the radio source expands/grows. In this framework, GHz-peaked spectrum (GPS) and compact steep-spectrum (CSS) radio sources, the former with $\\nu_{t} \\sim 1\\; {\\rm GHz}$ and the latter with $\\nu_{t} \\sim 100-500\\; {\\rm MHz} $, are considered to represent early stages in the individual radio source evolution.\\\\ The timescale of the spectral evolution in extremely young radio sources is very short, of the order of a few tens of years (Dallacasa \\cite{dd03}). This implies that the ideal targets to investigate how the radio spectrum evolves, and that are the main mechanisms at work, must be sought among sources with very high turnover frequencies. The existence of ``extreme'' GPS sources with turnover frequency above 10 GHz, termed ``high frequency peakers'' (HFPs) by Dallacasa et al. (\\cite{dd00}), is expected from all the radio source evolutionary models. Their detection is, however, quite difficult since they are short-lived objects evolving into GPS and then CSS sources, and to observational limitations preventing their selection. For a source to be recognized as an extreme GPS object, the characteristic turnover frequency must lie within the frequency range sampled by large area surveys. The frequencies of the surveys currently available are not high enough to allow an efficient selection of high-frequency peaking objects, since there is insufficient information above $\\sim$ 5 GHz.\\\\ \\begin{figure*}[!h] \\begin{center} \\special{psfile=8098fg1a.eps voffset=-180 hoffset=-10 hscale=25 vscale=25} \\special{psfile=8098fg1b.eps voffset=-180 hoffset=150 hscale=25 vscale=25} \\special{psfile=8098fg1c.eps voffset=-360 hoffset=-10 hscale=25 vscale=25} \\special{psfile=8098fg1d.eps voffset=-360 hoffset=150 hscale=25 vscale=25} \\special{psfile=8098fg1e.eps voffset=-540 hoffset=-10 hscale=25 vscale=25} \\special{psfile=8098fg1f.eps voffset=-540 hoffset=150 hscale=25 vscale=25} \\vspace{19cm} \\caption{VLBA images of the source J1459+3337 at 1.7, 5, 8.4, 15, 22 and 43 GHz. On each image we give the following information on the plot itself: a) Peak flux density in mJy/beam; b) First contour intensity (mJy/beam), which is generally equal to 3 r.m.s. of the noise level; contour levels increase by a factor of 2; c) the dynamic range (the ratio of the peak brightness and 1$\\sigma$); d) The restoring beam, plotted on the bottom left corner of each image. } \\label{immagini} \\end{center} \\end{figure*} The radio source RXJ1459+3337 was one of the first objects to be recognized as an ``extreme'' GPS with the turnover frequency at about 30 GHz and a peak flux density of about 800 mJy as reported by Edge et al. (\\cite{edge96}). This object is associated with a quasar at z = 0.6448. It was detected by ROSAT (Brinkmann et al. \\cite{brink00}) and shows an X-ray luminosity $L_{X} \\sim 2.8\\times 10^{45} {\\rm erg/s}$.\\\\ Multi-epoch VLA observations indicate strong variability in the optically-thick part of its radio spectrum. In particular, the flux-density variability at 5 GHz (i.e. in the optically-thick regime) seems to steadily increase. Although flux-density variability is quite rare in GPS sources (O'Dea \\cite{odea98}), an increment of the flux density at frequencies below $\\nu_{t}$ may be a consequence of the evolution of a synchrotron self-absorbed spectrum in an expanding component.\\\\ In this paper we present the results of multi-frequency VLA and VLBA observations conducted in 2003 and 2005, respectively. In order to identify the mechanisms at the basis of the spectral evolution we combine the information on the physical conditions provided by our new VLA data with that from archival data obtained in 1996 and 1999. The frequency range sampled by the new VLBA observations, together with their high-resolution, allows us to compute the magnetic field by means of observable quantities only, such as peak frequency, peak flux density and the source size, as from the synchrotron theory (Kellermann \\& Pauliny-Toth \\cite{kpt81}; see also Section 4.2). The comparison between this value with the field strength obtained assuming equipartition provides important information on the physical conditions of the radio source.   Throughout this paper, we assume the following cosmology: $H_{0} = 71\\; {\\rm km\\,s^{-1}\\,Mpc^{-1}}$, $\\Omega_{\\rm M} = 0.27$ and $\\Omega_{\\lambda} = 0.73$, in a flat Universe.\\\\ ",
        "conclusions": "We have presented new multi-frequency VLBA and VLA observations of the radio source RXJ1459+3337. Given the position of the spectral peak at very high frequency, this object is classified as an {\\it extreme} GPS radio source, also known as HFP.\\\\ By comparing our VLA observations with data from the literature spanning about 17 years, we find that the flux-density in the optically-thick part of the spectrum has been continuously increasing, while the turnover frequency has been moving toward lower frequencies. Such a regular variability, although not common in GPS sources peaking below 5 GHz, can be described in the case the radio emission originates from an adiabatically-expanding homogeneous component.\\\\ If in this context we compare either the flux-densities at the same frequency measured at different times or the turnover frequencies, we constrain the age of the radio emission, which is about $\\sim$50$\\pm$10 years. \\\\ The resolution achieved by the VLBA observations, together with the frequency range sampled, allowed us to ``observe'' the turnover frequency and to {\\it directly} measure parameters, such as the peak flux density and angular size of the source with great accuracy. In this way we estimate the magnetic field to be 0.16 $\\pm$ 0.03 G by means of observable quantities only. The consistency of this value with the field derived for a minimum energy condition strongly supports the idea that such extreme objects, as RXJ1459+3337 may already be in equipartition. The availability of ROSAT X-ray observations allowed us to infer the magnetic field by comparing radio and X-ray luminosities. Even in this case, the value inferred for the magnetic field agrees with the previous results.\\\\"
    },
    "0711/0711.1518_arXiv.txt": {
        "abstract": "We study the effect of the magnetic field geometry on the oscillation spectra of strongly magnetized stars. We construct a configuration of magnetic field where a toroidal component is added to the standard poloidal one. We consider a star with a type I superconductor core so that  both components of the magnetic field are expelled from the core and confined in the crust. Our results show that the  toroidal contribution does not influence significantly the torsional oscillations of the crust. On the contrary, the confinement of the magnetic field in the crust drastically affects on the torsional oscillation spectrum. Comparison with estimations for the magnetic field strength, from observations, exclude the possibility that magnetars will have a magnetic field solely confined in the crust i.e. our results suggest that the magnetic field in whatever geometry has to permeate the whole star. ",
        "introduction": "\\label{sec:Intro}  It is known that magnetic fields  play important role in the astrophysical phenomena, such as supernovae, gamma-ray bursts, galaxy jet, and so on. Recently, there is a growing consensus in explaining Soft Gamma Repeaters (SGRs) via the magnetar model \\citep{DUN1992}. Magnetars are believed to be neutron stars with strong magnetic field which is responsible for the observed flare activity. Three giant flares,  SGR 0526-66, SGR 1900+14, and SGR 1806-20, have been detected so far. The peak luminosities are in the range of $\\simeq 10^{44}-10^{46}$ erg/s. This huge amounts of energy can be explained by the presence of a strong magnetic field whose strength is estimated to be larger than $4\\times 10^{14}$ Gauss for SGR 1900+14 \\citep{SGR1900} and of the order of $ 8\\times 10^{14}$ Gauss for SGR 1806-20 \\citep{SGR1806}. In addition to the initial short and hard peak, these three events show a decaying tail lasting hundreds of seconds. In this late part careful analysis revealed the existence of characteristic Quasi Periodic Oscillations (QPOs); see \\cite{SW2006} for references. The oscillation frequencies in these QPOs are in the range of a few tenths of Hz to kHz and thought to be associated with crust torsional modes of magnetars. These signals might be the first evidence of a direct detection of crust oscillations of neutron stars:  a first attempt to explain these frequencies by the crust torsional oscillations has been done by using a model with dipole magnetic fields and realistic EoS (see \\cite{SA2007} and \\cite{SKS2007}) and this attempt was partially successful. On the other hand \\cite{GSA2006} suggested an alternative scenario to explain the observational frequencies via  Alfv\\'{e}n oscillations while \\cite{Levin2006} was the first to realize that the Alfv\\'{e}n oscillations of magnetars could be a continuum. More recently \\cite{Levin2007} pointed out that the edges or turning point of the continuum can be long-lived QPOs. In a recent paper, \\cite{SKS2007b} showed that the Alfv\\'{e}n oscillations of magnetars could be a continuum which can explain  the lower ones of the observed frequencies.   However, it is quite difficult to construct realistic models for neutron star magnetic fields and/or to understand the details of their dynamics. In fact, the observational data suggest that  the surface temperature of neutron stars with a strong magnetic field is not uniform (e.g., \\cite{SHHM2005,HTV2006}). This observational evidence can be explained, for example,  by the presence of magnetic fields in the crust region with a strong toroidal components or a more complicated structure than the poloidal one \\citep{GKP2006,PG2007}. Such crustal magnetic fields can be generated  if the core is a  type I superconductor. In this  case, the magnetic fields are expelled from the core due to the Meissner effect after about $10^4$ years \\footnote{Note that there is a possibility that the flux expulsion is more local, and will not lead to the entire core being void of flux (e.g., \\cite{SC1998}).}.   Recently, the analysis of the spectrum of timing noise for SGR 1806-20 and SGR 1900+14 has supported the idea that the core region can be a type II superconductor \\citep{ACT2004}. For the type II superconductor, even if an initial uniform magnetic field is present in the core, it will be split in thin magnetic fluxes tubes (\\textit{fluxoids}), which are parallel to the initial magnetic field. These fluxoids interact with the neutron vortices and are driven by these vortices outside the core, towards the crust on a timescale of $10^8-10^9$ years \\citep{Ruderman}. In any case, it is too difficult to discriminate between all these models by using only the past observational data.  This is due to uncertainties associated with the chemical composition of the interior and the crust of neutron stars as well as about the interior magnetic field configuration and its influence on the properties of neutron star.   The aim of this paper is to calculate the axisymmetric crust torsional modes of magnetars with both poloidal and toroidal magnetic field, where  both the components are confined in the crust. Then our numerical results show that this magnetic configuration cannot explain the actual observational data of SGRs. In this paper, we adopt the unit of $c=G=1$, where $c$ and $G$ denote the speed of light and the gravitational constant, respectively, and the metric signature is $(-,+,+,+)$. ",
        "conclusions": ""
    },
    "0711/0711.0773_arXiv.txt": {
        "abstract": "Spectra of high redshift QSOs show deep Gunn-Peterson absorptions on the blue sides of the \\Lya emissions lines. They can be decomposed into components called \\Lya leaks, defined to be emissive regions in complementary to otherwise zero-fluxed absorption gaps. Just like \\Lya absorption forests at low redshifts, \\Lya leaks are both easy to find in observations and containing rich sets of statistical properties that can be used to study the early evolution of the IGM. Among all properties of a leak profile, we investigate its equivalent width in this paper, since it is weakly affected by instrumental resolution and noise. Using 10 Keck QSO spectra at $z\\sim6$, we have measured the number density distribution function $n(W,z)$, defined to be the number of leaks per equivalent width $W$ and per redshift $z$, in the redshift range $5.4 - 6.0$. These new observational statistics, in both the differential and cumulative forms, fit well to hydro numerical simulations of uniform ionizing background in the $\\Lambda$CDM cosmology. In this model, Ly $\\alpha$ leaks are mainly due to low density voids. It supports the early studies that the IGM at $z\\simeq6$ would still be in a highly ionized state with neutral hydrogen fraction $\\simeq 10^{-4}$. Measurements of $n(W,z)$ at $z>6$ would be effective to probe the reionization of the IGM. ",
        "introduction": "The absorption spectra of QSOs at low redshift show \\Lya forests, which have played an important role to understand the physical status of diffused cosmic baryon gas and the ionizing background. At redshift $z > 5$, however, they are no longer to show forests features, but consist of complete absorption troughs separated by spikes of the transmitted flux (e.g. Becker et al. 2001; Fan et al. 2006). That is, although the cosmic hydrogen gas at $z>5$, in average, is opaque for Ly$\\alpha$ photons, there are many tiny regions, which are Gunn-Peterson transparent and lead to Ly$\\alpha$ photon leaking.  The nature of the leaking is crucial to understand the physics of reionization. According to commonly accepted scenario of reionization, at early stage, only isolated patches around ionizing sources are highly ionized. The subsequent growing and overlapping of the ionizing patches lead to a uniform ionizing background and the end of reionization (e.g., Ciardi et al. 2003; Sokasian et al. 2003; Gnedin 2004; Mellema et al. 2006) The ionization fraction of the IGM and the ionizing radiation underwent an evolution from highly non-uniform patches to a quasi-homogeneous field. Before the patch-to-uniform transition, only ionized patches would be transparent to \\Lya photons. After the transition, the low density voids will also be Gunn-Peterson transparent.  Therefore, the origin of Ly$\\alpha$ leaks will constrain the epoch of the patch-to-uniform transition. In this {\\it Letter}, we study the origin of \\Lya leaks in the observed spectra of QSOs at $z\\simeq6$.  Several statistics have been introduced to describe the transmitted flux of Ly$\\alpha$ absorption at high redshifts, including the probability distribution function (PDF) of the flux (Fan et al. 2002; Becker et al. 2007), the distribution of the size of dark gaps (Songaila \\& Cowie 2002; Fan et al. 2006), and the largest peak width distribution (Gallerani et al. 2007). We will focus on the profile of the leaking features in the transmitted flux. We fit these statistical features with samples of hydrodynamic simulation with a uniform ionizing background, and analyze the possibility of explaining the leaks by ionized patches embedded in neural IGM background. ",
        "conclusions": "The transmitted fluxes between Gunn-Peterson troughs of high redshift QSO's absorption spectra are not only one peak of the flux, but contain rich structures, which can be decomposed into \\Lya leaks. The \\Lya leaks have profiles similar to emission lines, and can be measured by equivalent width $W$. The equivalent width functions, $n(W,z)$ are effective statistical measurement of the process of reionization. We show that the equivalent width functions of the observed spectra at redshifts $5.4<z<6.0$ can be well fitted by hydro simulation of the $\\Lambda$CDM cosmology assuming the ionizing background to be uniform. In this model, all the Ly$\\alpha$ leaks are leaking through low density voids.  The mean transmitted flux at $z$ is given by $\\bar{F}\\propto\\int n(W,z)W dW$. Therefore, by adjusting the photoionization rate to match the observed Gunn-Peterson optical depth, the mean of $W$ for simulation samples should be the same with observation. Thus, Figure 1 actually is to show that once we adjusted the mean flux to be the same as observation, the simulation yields the same distribution of the observed $W$. In other word, the scattering of $W$ is caused by the fluctuations of mass density field of HI. Therefore, a small inhomogeneity of the ionizing background would be allowed. That is, the distribution of $W$ would still be able to be fitted with a fluctuated ionizing background if its variance is much less than that of HI.  In addition to the distribution of $W$, the evolution of $W$ is also helpful when differentiating models. For example, in voids scenario, the evolution of $W$ reflects the evolution of low density voids; while for ionized patches, it reflects the evolution of the UV luminosity function of ionizing sources.  We show that the ionized patches of galaxies embedded in a fully neutral IGM at redshift $z\\simeq 6$ are not enough to produce the observed leaks. We also show that leaks can only be produced by patches around strong ionizing UV photon sources, but not weak sources. Especially, big leaks ($W>0.5$ \\AA, \\ or $F>0.1$) have to come from very strong sources. Therefore, at higher redshift, Ly$\\alpha$ leaks only probe strong ionizing sources. Thus, from the existence of many big leaks at $z\\leq 6$ we can conclude that the patch-to-uniform transition of the ionizing background would occur at $z>6$, and most of the IGM at $z\\simeq6$ are still in a highly ionized state of neutral fraction $f_{HI}\\simeq 10^{-4}$. This result is consistent with the analysis of the transmitted flux PDF (Becker et al. 2007), the QSO proximity zones (Lidz et al. 2007) and the luminosity function of \\Lya emitting galaxies (Dijkstra et al. 2007).  It should be pointed out that the resolution of the observed data is low, $\\sim3$ \\AA, which corresponds to a comoving size $\\simeq 0.7$ h$^{-1}$Mpc. In contrast, most of the simulated leaks possess a intrinsic width $<3$\\AA. Thus the low resolution data provide only a test of smoothed leaking features. One cannot see whether the smoothed features is due to individual or clustered leaks. Higher resolution spectra would be able to test not only the width functions, but also the spatial correlations of the leaks. They can also provide other measurements of Ly$\\alpha$ leaks, like the FWHM, which would be effective for a confrontation between real data and models.  The statistics of Ly$\\alpha$ leaks at redshifts $\\leq 6$ actually is the statistics of voids formed in the early universe. The equivalent width functions, $n(W,z)$ of Ly$\\alpha$ leaks are similar to the mass function of galactic clusters. Thus, one can expect that the width functions of voids are sensitively dependent on cosmological parameters, and play a similar role as the mass function of clusters. For instance, the formation of large voids is found to be sensitively dependent on the mass parameter $\\Omega_m$ (Miranda \\& de Araujo 2001). With data of leaks, we may set constraint on cosmological parameters at high redshifts. This approach will be reported separately.  We thank X. Fan for providing the observational spectra and instructive suggestions. J. L. acknowledges the support of the International Center for Relativistic Center Network (ICRAnet). This work is supported in part by the US NSF under the grant AST-0507340."
    },
    "0711/0711.2540_arXiv.txt": {
        "abstract": "We review our work on Galactic open clusters in recent years, and introduce our proposed large program for the LOCS (LAMOST Open Cluster Survey). First, based on the most complete open clusters sample with metallicity, age and distance data as well as kinematic information, some preliminary statistical analysis regarding the spatial and metallicity distributions is presented. In particular, a radial abundance gradient of - 0.058$\\pm$ 0.006 dex kpc$^{-1}$ was derived, and by dividing clusters into age groups we show that the disk abundance gradient was steeper in the past. Secondly, proper motions, membership probabilities, and velocity dispersions of stars in the regions of two very young open clusters are derived. Both clusters show clear evidence of mass segregation, which provides support for the ``primordial'' mass segregation scenarios. Based on the great advantages of the forthcoming LAMOST facility, we have proposed a detailed open cluster survey with LAMOST (the LOCS). The aim, feasibility, and the present development of the LOCS are briefly summarized. ",
        "introduction": "Open clusters (OCs) are considered excellent laboratories for studies of stellar evolution. Studies in the research area dealing with OCs showed a rapid growth in 1990's and this area continues to develop vigorously. There may be many reasons for this recent growth of OC studies. Many new techniques are greatly beneficial for OC observations, including, for example, the application of wide-field, high quality CCD cameras and, more recently, multi-object spectrographs. In addition, OC studies play a very important and unique role in determining the structure and evolution of the Galactic disk.  OCs have long been used to trace the structure and evolution of the Galactic disk. From an observational point of view, there are some important advantages of using OCs as opposed to field stars. In general, in terms of OCs, we deal with groups of stars of nearly the same age, a similar composition and at a similar distance. We can observe OCs to large distances, and the distances to most of the OC sample have already been derived. In particular, OCs will have relatively stable orbital motion, which can be used as a better tracer of the Galactic disk structure. OCs also have a wide range of ages, so that -- combined with their spatial distribution and kinematic information -- we can study dynamical evolution effects, and very young OCs are ideal objects for studies of the stellar initial mass function. Furthermore, when combined with abundance data, we can investigate the chemical evolution history of the Galactic disk (\\cite[Carraro et al. 1998]{Carraro_etal98}, \\cite[Chen et al. 2003]{Chen_etal03}).  At present, the total number of detected clusters and associations is around 1700 (\\cite[Dias et al. 2002]{Dias_etal02}), about $60\\%$ of which have distance and age information and for about half of which proper motions are available. Less than one-fourth of the sample have both proper motion and radial velocity parameters, and only a small portion (about $8\\%$) of OCs have well-determined mean abundance values. We have compiled an updated OC catalog (Chen et al., in prep.), for which data have been extracted from various sources -- mainly from the Dias' catalog (\\cite[Dias et al. 2002]{Dias_etal02}), Kharchenko's data(\\cite[Kharchenko et al. 2005a]{Kharchenko_etal05a}, \\cite[Kharchenko et al. 2005b]{Kharchenko_etal05b}) and the WEBDA database (\\cite[Mermilliod \\& Paunzen 2003]{MermilliodPaunzen03}). Based on this large sample, we can embark on statistical studies of the Galactic OC system. ",
        "conclusions": ""
    },
    "0711/0711.0545_arXiv.txt": {
        "abstract": "We present a second catalog of HI sources detected in the Arecibo Legacy Fast ALFA Survey. We report \\ntot \\ detections over 135 deg$^2$, within the region of the sky having $22{\\rm h}<\\alpha<03{\\rm h}$ and $+26^{\\circ}<\\delta<+28^{\\circ}$. We present here the detections that have either (a) $S/N>6.5$, where the reliability of the catalog is better than 95\\% or (b) $5.0<S/N<6.5$ and a previously measured redshift that corroborates our detection. Of the \\ntot \\ objects presented here, \\nhvc \\ are High Velocity Clouds or clumps thereof with negative heliocentric recession velocities. These clouds are mostly very compact and isolated, while some of them are associated with large features such as Wright's Cloud or the northern extension of the Magellanic Stream. The remaining \\ngal \\ candidate detections are identified as extragalactic objects and have all been matched with optical counterparts. Five of the six galaxies detected with M$_{HI}<10^{7.5}$M$_{\\odot}$ are satellites of either the NGC672/IC1727 nearby galaxy pair or their neighboring dwarf irregular galaxy NGC784. The data of this catalog release include a slice through the Pisces-Perseus foreground void, a large nearby underdensity of galaxies. We report no detections within the void, where our catalog is complete for systems with HI masses of $10^8$M$_{\\odot}$. Gas-rich, optically-dark galaxies do not seem to constitute an important void population, and therefore do not suffice at producing a viable solution to the void phenomenon. ",
        "introduction": "The Arecibo Legacy Fast ALFA Survey (ALFALFA) is an on-going project conducted at the Arecibo Observatory. Using the 305m telescope and the seven-beam ALFA receiver, the survey aims at covering 7000 deg$^2$ of high galactic latitude sky in order to detect some 20,000 extragalactic HI-bearing objects \\citep[]{alfalfa1}. As first illustrated in \\citet{alfalfa2}, the survey makes significant improvement over previous blind HI surveys in terms of sensitivity, resolution, and number of detections. Earlier this year, in a first ALFALFA data release, we presented a catalog of 730 HI detections for the northern Virgo Cluster region \\citep{alfalfa3}. Of these detections, only $5\\%$ are present in the northern HIPASS catalog \\citep{wong} over the same area of the sky, and $69\\%$ of the cataloged sources are new HI detections. Clearly, the contribution of ALFALFA is significant.  The survey was designed to cover a wide variety of cosmic environments. The ``Spring'' ALFALFA survey covers the right ascension range from 7.5 to 16.5 hours, which includes several nearby overdensities such as the Virgo Cluster, the Coma Cluster and the Leo Group, while the ``Fall'' survey maps the region from 22 to 3 hours of right ascension over most of the declination range accessible from Arecibo ($\\delta=0^{\\circ}-36^{\\circ}$). This region includes the Pisces-Perseus Supercluster and also a large nearby cosmic void in its foreground. We present in this paper a first catalog of galaxies from this ``Fall'' survey. During the first season of ALFALFA observing, all observations necessary to obtain full coverage over a $2^{\\circ}$ wide strip were completed. This specific region of sky, centered on $\\delta=+27^{\\circ}$ and stretching from $\\alpha=22$h to $\\alpha=3$h, was chosen to begin the survey observations because it was already targeted during the ALFALFA precursor observations \\citep{alfalfa2}, it goes through the very nearby galaxy group NGC672, it covers the southernmost portion of the Pisces-Perseus Supercluster ridge, and surveys a large nearby void.  The scientific goals and applications of ALFALFA are very broad and diverse \\citep{alfalfa1}; an important one  is the determination of the HI mass function (HIMF), and its dependence on environment. Several blind HI surveys have been used to address this question, both using single dish telescopes \\citep[e.g.][]{briggs,rosenberg,hipass,hijass} and synthesis arrays \\citep[e.g.][]{weinberg,deblok}. While there is no consensus yet on the slope of the faint end of the HIMF \\citep[for a recent compilation of the different results, see][]{springob05}, the different surveys measure an abundance of low-mass HI systems that is far less than the number of low-mass dark matter halos produced by numerical simulations \\citep[e.g.][]{klypin99}. With its sensitivity and large sky coverage, ALFALFA should detect hundreds of galaxies with HI masses less than $10^{7.5}$M$_{\\sun}$, compared to the handful contained in other surveys, and should therefore help determine with high confidence the abundance of low-mass HI systems in the local Universe and put constraints on the ``missing satellites'' problem.  There is another related problem when it comes to comparing observations of galaxies with the predicted abundances from simulations, which can also be addressed by ALFALFA. This problem concerns the abundance of low-mass systems in underdense regions of space and has been coined as the ``void phenomenon'' \\citep{peebles01}. Using numerical simulations, \\citet{gott03} find that voids are filled with dark matter filaments, just like denser environments, except that the amplitudes of the perturbations are much smaller. They predict that a void with a diameter of 20 $h^{-1}$ Mpc should contain almost 1000 dark matter haloes with masses of $10^9 h^{-1}$ M$_{\\sun}$ and about 50 with masses of $10^{10} h^{-1}$ M$_{\\sun}$. Unless there are mechanisms at play that would inhibit the formation of galaxies in underdense regions of space, there should therefore be a significant population of dwarf galaxies in voids. Such a population has so far not been observed. \\citet{hoyle05} have built a luminosity function of void galaxies using data from the Sloan Digital Sky Survey. They do not find any significant dependence of the slope of the faint end of the luminosity function on environment. A steeper slope in the voids would have revealed the existence of a more important population of dwarfs.  If we accept the predictions of the numerical simulations, there are two main avenues to explain the discrepancy with the results of the optical surveys: (1) the processes that are required to trigger the star formation process in galaxies are not available in voids, or (2) there is no material out of which these stars could be formed.  While the second scenario implies that some physical processes have suppressed star formation in dwarf galaxies by removing their reservoirs of cool gas, according to the first scenario the voids should contain a large number of halos that are filled with HI but devoid of stars, the so-called ``dark galaxies''. So far, such a population has not been observed, though a few candidates have been reported \\citep{schneider83,gh89,kilborn00,ryder01,minchin}. After further investigation, these have however all been found to be either associated with or the companions of a faint optical counterpart, or part of the tidal features of neighboring galaxies \\citep[][Haynes et al. 2007 in preparation]{schneider89,chengalur95,doyle05}. The largest blind HI survey to date, the HI Parkes All Sky Survey \\citep[HIPASS,][]{hipass} has failed to find any truly dark galaxies. With a sensitivity greater than that of HIPASS by a factor of eight, ALFALFA promises to either find a population of dark galaxies if they exist, or at least to put a much stronger constraint on their abundance. In this paper, we provide a new catalog of HI detections from ALFALFA which provides first insight on a slice through the Pisces-Perseus foreground void and allows us to investigate in a preliminary manner what ALFALFA may be telling us about the void phenomenon.  The rest of the paper is organized as follows. In \\S 2, we describe the large-scale structure present in the sky area studied, give an overview of the ALFALFA observing and data reduction processes and present a catalog of \\ntot \\ HI detections. In \\S 3, we present the properties of the sample and a description of the High Velocity Cloud population. Finally, in \\S 4 we discuss the implications of the data for the void problem. A value for the Hubble constant of 70 \\kms \\ Mpc$^{-1}$ has been assumed throughout the paper. ",
        "conclusions": "We present a catalog of \\ntot \\ HI detections from the ALFALFA survey over 135 deg$^{2}$ ($22{\\rm h}<\\alpha<03{\\rm h}$,  $+26^{\\circ}<\\delta<+28^{\\circ}$). Of these, \\nhvc \\ are identified as HVCs while the remaining \\ngal \\ are extragalactic and all associated with an optical counterpart. These matches are made with great confidence due to the relatively high pointing accuracy and spatial resolution of the survey. Of the extragalactic objects, $62\\%$ are new detections made by ALFALFA and $59\\%$ are new redshift measurements.  Five of the six objects detected with $\\log_{10}(M_{HI})<7.5$ are satellites of either the galaxy pair NGC672/IC1727 or of their neighbor, the dwarf irregular galaxy NGC784. We find no evidence for low-mass optically-dark systems in the vicinity of these large nearby galaxies. We also do not find them in a 460 Mpc$^3$ slice through a nearby local void in the foreground of the Pisces-Perseus Supercluster. The non-detection suggests that a large population of HI-rich yet optically-dark galaxies is unlikely to reconcile the discrepancies between the observed abundance of dwarf galaxies in voids and the predictions made on the basis of numerical simulations.  We have also used the ALFALFA data to study the HVC population. In addition to large bright clouds previously cataloged \\citep{wakker91,deheij} such as the southern extension of Wright's Cloud, we detect a large number of ultra compact HVCs which correspond to the large filaments observed by \\citet{bt04} and associated with a northern extension of the Magellanic Stream. When completed, ALFALFA and its commensal Galactic-ALFA survey will be the most sensitive large scale surveys for High Velocity Clouds, as evidenced by comparing the HVCs found in this catalog to previous detections.  While the data presented here are extracted from only $2\\%$ of the full area ALFALFA will survey, we have been able to show the potential of the survey not only to significantly increase the number of galaxies for which HI redshifts are available, but also to better our understanding of the population of gas-rich galaxies, to put strong constraints on the abundance of HI-rich but optically-dark galaxies, and even to study the population of HVCs."
    },
    "0711/0711.1775_arXiv.txt": {
        "abstract": "We present \\HI\\ observations of the elliptical galaxy NGC 2974, obtained with the Very Large Array. These observations reveal that the previously detected \\HI\\ disc in this galaxy (Kim et al.\\ 1988\\nocite{1988ApJ...330..684K}) is in fact a ring. By studying the harmonic expansion of the velocity field along the ring, we constrain the elongation of the halo and find that the underlying gravitational potential is consistent with an axisymmetric shape.  We construct mass models of NGC 2974 by combining the \\HI\\ rotation curve with the central kinematics of the ionised gas, obtained with the integral-field spectrograph \\sauron. We introduce a new way of correcting the observed velocities of the ionised gas for asymmetric drift, and hereby disentangle the random motions of the gas caused by gravitational interaction from those caused by turbulence. To reproduce the observed flat rotation curve of the \\HI\\ gas, we need to include a dark halo in our mass models. A pseudo-isothermal sphere provides the best model to fit our data, but we also tested an NFW halo and Modified Newtonian Dynamics (MOND), which fit the data marginally worse.  The mass-to-light ratio $M/L_I$ increases in NGC~2974 from 4.3 $M_\\odot/L_{\\odot,I}$ at one effective radius to 8.5 $M_\\odot/L_{\\odot,I}$ at 5 $R_e$. This increase of $M/L$ already suggests the presence of dark matter: we find that within 5 $R_e$ at least 55 per cent of the total mass is dark. \\\\  \\date{} ",
        "introduction": "\\label{sec:introduction}  Although the presence of dark matter dominated haloes around spiral galaxies is well established (e.g. van Albada et al. 1985\\nocite{1985ApJ...295..305V}), there is still some controversy about their presence around early-type galaxies.  Spiral galaxies often contain large regular \\HI\\ discs, which allow us to obtain rotation curves out to large radii, and therefore we can constrain the properties of their dark haloes. But these discs are much rarer in elliptical galaxies (e.g. Bregman, Hogg \\& Roberts 1992\\nocite{1992ApJ...387..484B}), so that for this class of galaxies we are often required to use other tracers to obtain velocity measurements, such as stellar kinematics, planetary nebulae or globular clusters. These tracers however are not available for all early-type galaxies, and give mixed results (e.g.  Rix et al. 1997\\nocite{1997ApJ...488..702R}, Romanowsky et al. 2003\\nocite{2003Sci...301.1696R}, Bridges et al. 2006\\nocite{2006MNRAS.373..157B}).  With the increase in sensitivity of radio telescopes, it has been discovered that many early-type galaxies in the field do contain \\HI\\ gas, though with smaller surface densities than in spiral galaxies (e.g. Morganti et al. 2006\\nocite{2006MNRAS.371..157M}). The average \\HI\\ surface density in the Morganti et al. sample is around $1 M_\\odot$ pc$^{-2}$, which is far below the typical value for spiral galaxies (4 - 8$ M_\\odot$ pc$^{-2}$, e.g. Cayatte et al. 1994\\nocite{1994AJ....107.1003C}). This would explain why previously only the most gas-rich early-type galaxies were detected in \\HI. Morganti et al. find that \\HI\\ can be present in different morphologies: \\HI\\ discs seem to be as common as off-set clouds and tails, though they occur mostly in the relatively gas-rich systems.  Recently rotation curves of \\HI\\ discs in low surface brightness galaxies and dwarf galaxies, complemented with H$\\alpha$ observations, have been used not only to confirm the existence of dark matter haloes, but also to obtain estimates on the inner slope of the density profiles of the haloes (e.g. van den Bosch et al. 2000\\nocite{2000AJ....119.1579V}; Weldrake, de Blok \\& Walter 2003\\nocite{2003MNRAS.340...12W}). Simulations within a cold dark matter (CDM) cosmology yield haloes with cusps in their centres (NFW profiles, see Navarro, Frenk \\& White 1996\\nocite{1996ApJ...462..563N}), but observations suggest core-dominated profiles (e.g. de Blok \\& Bosma 2002\\nocite{2002A&A...385..816D}; de Blok 2005\\nocite{2005ApJ...634..227D}).  Detailed studies of rotation curves of early-type galaxies that contain \\HI\\ discs are sparser, due to lack of spatial resolution: to detect low \\HI\\ surface densities, larger beams are needed.  Also, only few early-type galaxies have \\HI\\ discs that are extended and regular enough to allow for detailed studies.  Comparing $M/L$ values at large radii, derived from \\HI\\ velocities, to $M/L$ at smaller radii measured from ionised gas kinematics, the conclusion is that early-type galaxies also have dark matter dominated haloes (e.g. Bertola et al.  1993\\nocite{1993ApJ...416L..45B}; Morganti et al. 1997\\nocite{ 1997AJ....113..937M}; Sadler et al. 2000\\nocite{2000AJ....119.1180S}).  Franx, van Gorkom \\& de Zeeuw (1994)\\nocite{1994ApJ...436..642F} used the \\HI\\ ring of the elliptical galaxy IC~2006 to determine not only the mass, but also the shape of the dark halo. They concluded that IC~2006 is surrounded by an axisymmetric dark halo, using the geometry of the ring and an harmonic expansion of its velocity map.  In this paper, we present a similar analysis of the regularly rotating \\HI\\ ring around the elliptical (E4) field galaxy NGC~2974. Kim et al. (1988)\\nocite{1988ApJ...330..684K} observed this galaxy before in \\HI\\, but their data had lower spatial resolution than ours, and they found a filled disc instead of a ring. Cinzano \\& van der Marel (1994)\\nocite{{1994MNRAS.270..325C}} found an embedded stellar disc in their dynamical model of this galaxy, based upon long-slit spectroscopic data, but Emsellem, Goudfrooij \\& Ferruit (2003)\\nocite{2003MNRAS.345.1297E} constructed a dynamical model of NGC~2974 based on TIGER integral-field spectrography and long-slit stellar kinematics, that does not require a hidden disc structure. They did report the detection of a two-arm gaseous spiral in the inner 200 pc of NGC~2974 from high resolution WFPC2 imaging.  Krajnovi\\'c et al.  (2005)\\nocite{2005MNRAS.357.1113K} constructed axisymmetric dynamical models of both the stars and ionised gas based upon \\sauron\\ integral-field data. These models require a component with high angular momentum, consisting of a somewhat flattened distribution of stars, though not a thin stellar disc.  Emsellem et al. (2007)\\nocite{2007MNRAS.379..401E} classify NGC~2974 as a fast rotator, which means that it possesses large-scale rotation and that its angular momentum is well defined. Some of the characteristics of NGC~2974 are given in Table~\\ref{tab:ngc2974}.  For our analysis of NGC~2974 we combine kinematics of neutral gas, obtained from our observations with the Very Large Array (VLA), with that of ionised gas, obtained with the integral-field spectrograph \\sauron\\ (Bacon et al.  2001\\nocite{2001MNRAS.326...23B}). This combination of a small scale two-dimensional gas velocity map in the centre of the galaxy, and a \\HI\\ velocity map at the outskirts, allows measurements of a rotation curve ranging from 100 pc within the centre of the galaxy to 10 kpc at the edges of the \\HI\\ ring. We use this rotation curve, together with ground- and space based optical imaging, to determine the dark matter content in NGC~2974, and to constrain the shape of the dark halo.  In section 2, we discuss the two datasets and their reduction, and describe the \\HI\\ ring. We concentrate on the analysis of the velocity maps in section 3. Section 4 is devoted to the rotation curve that we extract from the velocity maps, and in section 5 we show mass models with various halo models, and find the best fit to the rotation curve. Section 6 summarizes our results.   \\begin{table} \\label{tab:ngc2974} \\begin{center} \\begin{tabular}{l|r} \\hline\\hline Parameter & Value \\\\ \\hline Morphological Type & E4 \\\\ $M_B$ (mag) & -20.07 \\\\ Effective $B-V$ (mag) & 0.93 \\\\ PA ($^{\\circ}$) & 41 \\\\ Distance modulus (mag) & 31.60 \\\\ Distance (Mpc) & 20.89 \\\\ Distance scale (pc arcsec$^{-1}$) & 101.3 \\\\ % Effective radius & 24$^{\\prime \\prime}$ \\\\ \\hline \\end{tabular} \\end{center} \\caption{Properties of NGC~2974. The values are taken from the Lyon/Meudon Extragalactic Database (LEDA) and corrected for the distance modulus, which is taken from the surface brightness fluctuation measurements by Tonry et al. (2001)% . Note that 0.06 mag is subtracted to adjust to the Cepheid zeropoint of Freedman et al. (2001)% ; see Mei et al. (2005), % section 3.3, for a discussion. The effective radius is taken from Cappellari et al. (2006).} \\end{table} ",
        "conclusions": "\\label{sec:conclusion}  We obtained \\HI\\ observations of the early-type galaxy NGC~2974 and found that the neutral gas resides in a ring. The ring starts around 50$^{\\prime \\prime}$, and extends to 120$^{\\prime \\prime}$, which corresponds to 12 kpc or 5 effective radii. The total mass of the neutral gas is $5.5 \\times 10^8 M_\\odot$.  We compared the velocity field of the \\HI\\ ring with the kinematics of the ionised gas. We found that both velocity fields are very regular and nicely aligned, indicating that they could form a single disc. A harmonic decomposition of the velocity field showed that at large radii the gravitational potential is consistent with an axisymmetric shape.  We introduced a new way to correct the rotation curve of the ionised gas for asymmetric drift. We found that the correction approaches a constant value and this enabled us to remove the effect of turbulence on the rotation curve, assuming a constant asymmetric drift correction throughout the galaxy. We confirmed that this assumption is valid in NGC~2974, by comparing with the asymmetric drift corrected rotation curve of the stars (which was not affected by turbulence). An interesting question is whether other galaxies show the same behaviour. If this is the case, then with our method we would be able to investigate rotation curves of ionised gas and the effects of turbulence in more detail in other galaxies, and search for connections with e.g. spiral structure and bars. Although in principle we could for NGC~2974 also have used the stellar rotation curve together with the \\HI\\ to constrain the mass models, this will not be the case for all galaxies. For instance, in low surface brightness galaxies stellar kinematics are not easy to obtain, and even in high surface brightness galaxies, the absorption line kinematics need to be binned to higher signal-to-noise than the emission line kinematics, provided that ionised gas is present.  It is clear from the rotation curve of NGC~2974 that dark matter is required to explain the observed velocities.We found that the total mass-to-light ratio increases from 4.3 $M_\\odot/L_{\\odot,I}$ at 1 $R_e$ to 8.5 $M_\\odot/L_{\\odot,I}$ at 5 $R_e$. This last value would correspond to 14 $M_\\odot/L_{\\odot,B}$ in $B$-band. Even in the maximal disc model, 55 per cent of the total mass is dark, and an additional dark halo needs to be included.  We constructed mass models of NGC~2974, where we modeled both the stellar and gaseous contribution to the gravitational potential. The latter is negligible compared to the stars ($M_{\\mathrm{gas}} \\sim 0.001 M_*$), but still included in our models. For the dark halo, we tested two different profiles: the core-dominated pseudo-isothermal sphere and the cuspy NFW profile. We experimented with different values for the stellar mass-to-light ratio, but found that we cannot constrain this value with just the rotation curve: for most $M_*/L$ smaller than the maximal disc value, we could obtain a decent fit for both the pseudo-isothermal sphere and the NFW profile. If we compare models with $M_*/L_I$ from single stellar population models (2.34 $M_\\odot/L_{\\odot,I}$) with maximal disc model ($M_*/L_I = $ 3.8 $M_\\odot/L_{\\odot,I}$), then the first provide better fits to the data. Especially the inner datapoints are better reproduced with $M_*/L_I =$ 2.34 $M_\\odot/L_{\\odot,I}$, but we note that the \\HI\\ data points are better fitted in the models with the larger $M_*/L_I$ value. The pseudo-isothermal sphere fits our data marginally better than the NFW profile, but the difference is not significant. With MOND we can also reproduce the observed rotation curve, but not as well as with models that include a dark matter halo.  The largest uncertainty in our analysis is the stellar mass-to-light ratio. We can only derive an upper limit on this ratio from the maximal disc model or e.g. from Schwarzschild modeling, since $M_*/L$ should always be equal to or smaller than the dynamical $M/L$.  Values for $M_*/L$ from stellar population synthesis models depend significantly on the model assumptions: we mentioned already that going from a Kroupa to a Salpeter IMF can increase $M_*/L$ by as much as 40 per cent. Also, even low-level secondary star formation can affect $M_*/L$ severely.  Furthermore, there is no reason why the $M_*/L$ should remain constant over 5 effective radii, which we assumed when modeling the stellar mass. If the stellar mass-to-light ratio were known, we would be able to determine the dark matter fraction in the galaxy with more accuracy, and either rule out or confirm the maximal disc hypothesis. Also, since $M_*/L$ is the only free parameter when fitting a rotation curve in MOND, knowing this value would provide us with a rotation curve that can be compared to the data directly, providing a clear test for MOND.  We have shown in this paper that it is possible to combine rotation curves of neutral and ionised gas, correcting the latter one for asymmetric drift using the Jeans equations and the higher order velocity moments of the collisionless Boltzmann equations. Our method to correct for the asymmetric drift therefore does not require a cold disc assumption ($\\sigma \\ll V_c$). With more early-type galaxies getting detected in \\HI, and more high quality rotation curves becoming available, we can now study the shape of their dark matter haloes."
    },
    "0711/0711.1633_arXiv.txt": {
        "abstract": "We present results from a study of the orbits of eclipsing binary stars (EBs) in the Magellanic Clouds. The samples comprise $4510$ EBs found in the Large Magellanic Cloud (LMC) by the MACHO project, $2474$ LMC EBs found by the OGLE-II project (of which $1182$ are also in the MACHO sample), $1380$ in the Small Magellanic Cloud (SMC) found by the MACHO project, and $1317$ SMC EBs found by the OGLE-II project (of which $677$ are also in the MACHO sample); we also consider the EROS sample of $79$ EBs in the bar of the LMC. Statistics of the phase differences between primary and secondary minima allow us to infer the statistics of orbital eccentricities within these samples. We confirm the well-known absence of eccentric orbit in close binary stars. We also find evidence for rapid circularization in longer period systems when one member evolves beyond the main sequence, as also found by previous studies. ",
        "introduction": "The two components of a binary star system raise tides in each other. These tides will, in general, cause dissipation in the stars, leading to the phenomenon of \\emph{tidal lag}. These lags exert torques which lead to exchange of angular momentum between the orbit and the spins of the stars. The lowest energy configuration, for a given total angular momentum, is a circular orbit with both spins aligned and synchronous with the orbit. All other configurations evolve towards the synchronous, circular state. Whether or not this is achieved is determined by the rate of dissipation in the tides. \\par The study of tides raised in celestial bodies by mutual attraction dates back at least to \\citet{darwin79}. Darwin considered the tides raised on the Earth, modeled as a homogeneous and deformable \\emph{viscous} body, by a point mass Moon in the weak friction limit; hence his model is also known as \\emph{weak friction model}. The model assumes that viscous dissipation causes a delay in the onset of the tide by a constant amount $\\tau$ so at time $t$ the shape of the Earth is the one that it would have been at time $t-\\tau$ in the absence of dissipation. The axis of the tidal bulge is not therefore aligned with the line of the centers of the two bodies but lags it by a constant amount, resulting in a torque that tends to align the two bodies. Modern discussions of the weak friction model, with emphasis on close binary stars rather than planet-satellite systems, are given by \\citet{alexander73} and \\citet{hut81}. \\par \\citet{zahn66a,zahn66b,zahn66c} proposed that the coupling of the tidal flow with turbulent flows in the envelope of a late type star is chiefly responsible for orbital circularization and synchronization in these stars; evidence of circularization due to tidal interactions in late type giants is found by \\citet{lucy71}. The theory has been more recently revised by \\citet{zahn89a} and compared against observations by \\citet{zahn89b} who found that for late type stars most of the circularization occurs during the Hayashi phase of pre main sequence evolution, as previously suggested by \\citet{mayor84}. \\par The circularization models cited above consider only the \\emph{equilibrium tide} which arises if the star is at all times in hydrostatic equilibrium. If, however, the orbit is not circular or the rotation is not synchronous the star is subjected to a time-varying gravitational potential which excites oscillations, giving rise to a \\emph{dynamical tide} superimposed on the equilibrium tide. Forced oscillations in binary stars were first considered by \\citet{cowling41} and then by \\citet{zahn70}; \\citet{zahn75} considered the damping of the dynamical tide by radiative dissipation in the radiative envelope as a possible circularization mechanism for early type stars; \\citet{giuricin84} studied a sample of $\\sim 200$ early type binary stars and showed the theory to be compatible with the data presented. These mechanisms are all described in \\citet{zahn77}, where time scales of circularization and synchronization are derived. \\par A circularization mechanism in which energy dissipation is due to large scale hydro-dynamical flows resulting from the deformation of the star (meridional circulation) has been proposed by Tassoul \\citep{tassoul87,tassoul95}. This model has been applied by \\citet{claret95} to a homogeneous sample of 45 eclipsing binary stars (EBs) with accurate parameter determination from \\citet{andersen91}: they find a satisfactory agreement with the observations and also find that circularization is still taking place during the main sequence for early-type EBs, a finding consistent with the results of this paper. \\citet{claret97} then apply both the turbulent dissipation mechanism and the radiative damping mechanism to the same \\citep{andersen91} data set. They find that within uncertainties, these formalisms seem able to explain the observed eccentricity distribution, although with a few exceptions. \\par Since tidal forces decrease and the period increases with separation, in a sample of coeval binary stars such as those in star clusters, binaries with a longer period should have orbits with a range of eccentricities since tidal forces have not circularized them yet, whereas binaries with a shorter period should all have circular orbits since tidal forces would have been more effective in circularizing them. Therefore clusters should show a transition between binaries with shorter period and circular orbits and binaries with longer periods and eccentric orbits; the values of this transition period varies with the age of the cluster. Binary stars in cluster are widely studied: for example as a part of the WIYN Open Cluster Study (WOCS) \\footnote{\\url{http://www.astro.ufl.edu/\\~{}ata/wocs/}}, a systematic search of EBs in Milky Way open clusters for testing orbital circularization theories is currently under way \\citep{mathieu04,meibom05}. \\citet{latham02} present spectroscopic orbital solutions for $171$ single-lined stars from their catalog \\citep{carney94} of $1464$ stars selected for high proper motion and find that for the metal poor, high velocity halo binary stars in their sample the transition from circular to eccentric orbits occurs at $\\sim 20~\\mathrm{days}$. \\par Binary systems in which the component stars eclipse each other are a powerful tool for the study of circularization theories since in this case both \\emph{stellar} parameters like radius, mass, and temperature, and \\emph{orbital} parameters like eccentricity, angle of inclination, and longitude of periastron, can be determined to the accuracy necessary to test different theories. To determine these parameters fully one needs to supplement the photometric data with high quality spectra, which can be used to determine orbital velocities. \\par Eclipsing binary stars have been found in large numbers by astronomical surveys studying gravitational microlensing: MACHO\\footnote{\\url{http://www.macho.mcmaster.ca/}}, OGLE\\footnote{\\url{http://sirius.astrouw.edu.pl/\\~{}ogle/}} and EROS\\footnote{\\url{http://eros.in2p3.fr/}}. In particular a sample of $4634$ EBs in the Large Magellanic Cloud (LMC) and of $1509$ EBs in the Small Magellanic Cloud \\citep{faccioli07} have just been published by the MACHO collaboration. Previously a catalogue of $2580$ EBs in the LMC \\citep{wyr03} and of $1351$ EBs in the SMC \\citep{wyr04} were published by the OGLE collaboration, the EBs were selected from their catalogue of variable stars in the Magellanic Clouds \\citep{zebrun01a} compiled from observations taken during the second part of the project \\citep[OGLE-II:][]{udalski97}. A catalogue of 611 EBs in the LMC \\citep{lacy97}, and of $79$ EBs in the bar of the LMC \\citep{gri95} were also published by the MACHO and EROS collaboration respectively. This earlier MACHO sample is contained in the new one. \\par We aim to take advantage of the sheer size of our samples to infer statistically valid conclusions on orbital circularization, even though our data do not have high photometric precision and there are no spectroscopic data. We exploit a simple idea: the phase difference between the primary and secondary minima in an EB light curve, $\\phi_1-\\phi_2$, may be effectively estimated for the light curves in our samples. This object is related to the orbital elements by the following equation, which sets a \\emph{lower bound} to the eccentricity \\citep{milone98}: \\begin{equation}\\label{eq:ecc} |\\phi_1-\\phi_2|-\\frac{1}{2}= \\frac{1}{\\pi}e\\cos\\omega\\left(1+\\frac{1}{\\sin^2i}\\right) \\end{equation} where $\\phi_1-\\phi_2$ is the phase difference between primary and secondary minima, $e$ is the eccentricity, $\\omega$ is the argument of periastron and $i$ is the orbital inclination. For eclipsing systems, especially the wider systems where non zero eccentricities are found, we know that $\\sin i\\sim 1$. The important degeneracy is between $\\omega$ and $e$. It is possible \\emph{in principle} to model high signal-to-noise photometric data to break this degeneracy \\citep{wilson01,wilson02,devor05} but we do not attempt this here. We will show that $\\phi_1-\\phi_2$ alone may be used for our purpose. Although Eq. \\ref{eq:ecc} shows that for an eccentric orbit $\\phi_1-\\phi_2$ can be either $<0.5$ or $>0.5$, depending on $\\cos\\omega$, we will always adopt the convention $\\phi_1-\\phi_2>0.5$ since only the deviation of $\\phi_1-\\phi_2$ from $0.5$ is relevant for the detection of eccentricity. \\par The paper is organized as follows: Section \\ref{sec:obs} describes both the MACHO and the OGLE-II samples for both Clouds; Section \\ref{sec:fit} describes the fits to the MACHO data and its validity (for the OGLE-II data the relevant information is provided by the authors); Section \\ref{sec:results} reports our results for both Clouds both from the MACHO and the OGLE-II samples; Section \\ref{sec:sig} discusses the significance of our results and EROS results are briefly considered in Subsection \\ref{subsec:eros}. Finally Section \\ref{sec:conclusion} states our conclusions. \\par The data presented in this paper can be accessed on line at the Astrophysical Journal website\\footnote{\\url{http://www.journals.uchicago.edu/ApJ/}} and are mirrored at the Harvard University Initiative in Innovative Computing (IIC) /Time Series Center\\footnote{\\url{http://timemachine.iic.harvard.edu/faccioli/CircularizationTables/}}. \\par From now on we will always use the term \\emph{unfolded light curve} to indicate a set of time ordered observations and will reserve the term \\emph{light curve} to indicate a set of time ordered observations \\emph{folded} around a period, omitting for brevity the adjectives ``folded'' or ``phased''. ",
        "conclusions": "\\label{sec:conclusion} We have presented a study of orbital circularization in the Magellanic Clouds using EBs samples compiled both by the MACHO and by the OGLE-II collaboration, as well as a sample for the bar of the LMC compiled by the EROS collaboration. We have shown that in the LMC binary stars with period in the range $1.5\\mathrm{d}-20\\mathrm{d}$ with initially eccentric orbits circularize during post Main Sequence evolution, with a cutoff in color at about $\\vr<0.2~\\mathrm{mag}$: our data are consistent with stars starting their lives with a wide range of eccentricities and circularizing their orbit as they evolve, both on the Main Sequence and in their red giant phase. For the SMC the same conclusion holds, for a period range $1\\mathrm{d}-20\\mathrm{d}$. We have obtained the same results with independently assembled datasets, thus enhancing our confidence in the validity of our conclusions."
    },
    "0711/0711.0686_arXiv.txt": {
        "abstract": "Among various phenomenological $\\Lambda$ models, a time-dependent model $\\dot \\Lambda\\sim H^3$ is selected here to investigate the $\\Lambda$-CDM cosmology. The model can follow from dynamics, underlying the origin of $\\Lambda$. Using this model the expressions for the time-dependent equation of state parameter $\\omega$ and other physical parameters are derived. It is shown that in $H^3$ model accelerated expansion of the Universe takes place at negative energy density, but with a positive pressure. It has also been possible to obtain the change of sign of the deceleration parameter $q$ during cosmic evolution. ",
        "introduction": "Any theoretical model related to cosmology should be supported by observational data. Present cosmological models suggest that the Universe is primarily made of dark matter and dark energy. Various observational evidences, including SN Ia~\\cite{Riess1998,Perlmutter1999,Knop2003,Riess2004,Tegmark2004a,Astier2005,Spergel2007} data, support the idea of accelerating Universe and it is supposed that dark energy is responsible for this effect of speeding up. Now-a-days it is accepted that about two third of the total energy density of the Universe is dark energy and the remaining one third consists of visible matter and dark matter \\cite{Sahni2004}.  Though dark matter had a significant role during structure formation in the early Universe, its composition is still unknown. It is predicted that the dark matter should be non-baryonic and various particle physics candidates and their mixtures are discussed (see e.g. \\cite{khlopov95,khlopovbook1} for review). In particular, time-varying forms of dark matter \\cite{dorr1,Turner:1984nf,Gelmini:1984pe} in Unstable Dark Matter (UDM) scenarios \\cite{dorr1,dorr2,dorr3} are still not fully explored and deserve interest, giving simultaneously clustered and unclustered dark matter components.  The standard cold dark matter (SCDM) model introduced in 1980's which assume $\\Omega_{CDM}=1$ is out of favor today \\cite{Efstathiou1990,Pope2004}. After the emergence of the idea of accelerating Universe, the SCDM model is replaced by $\\Lambda$-CDM (or LCDM) model. This model includes dark energy as a part of the total energy density of the Universe and is in nice agreement with various sets of observations \\cite{Tegmark2004b}. In this connection it is to be noted here that according to $\\Lambda$-CDM model, acceleration of the Universe should be a recent phenomenon. Some recent works \\cite{Padmanabhan2003,Amendola2003} favor the idea that the present accelerating Universe was preceded by a decelerating one and observational evidence \\cite{Riess2001} also support this.  Now, in most of the recent cosmological research, the equation of state parameter $\\omega$ has been taken as a constant. However, its seems that for better result $\\omega$ should be taken as time-dependent \\cite{Chervon2000,Zhuravlev2001,Peebles2003}. Therefore, in the present work $\\Lambda$-CDM Universe has been investigated by selecting a specific time-dependent form of $\\Lambda$, viz., $\\dot \\Lambda\\sim H^3$ along with $\\omega(t)$. A kind of this $\\Lambda$-model was previously studied by Reuter and Wetterich \\cite{Reuter1987} for finding out an explanation of the presently observed small value of $\\Lambda$. This choice found realization in the approach treating time-variation of $\\Lambda$ as Bose condensate evaporation~\\cite{Dymnikova:2001ga} in the framework of model of self consistent inflation~\\cite{Dymnikova:2001ga,Dymnikova:2001jy,Dymnikova:1998ps}. Very recently it is also used by Mukhopadhyay et al. \\cite{Mukhopadhyay2007} for investigating various inherent features of the $\\Lambda$-CDM Universe.  In this context it is important to note that effects of real and virtual particles can play nontrivial role in the origin of $\\Lambda$ and its time-variation. Indeed, $R^2$ term was induced by vacuum polarisation and gave rise to inflation in one of the first inflationary models~\\cite{Starobinsky:1980te}. The back reaction of particle products of Bose-Einstein condensate evaporation slowed down evaporation and damped coherent scalar field oscillations in the approach of self consistent inflation~\\cite{Dymnikova:2001ga,Dymnikova:2001jy,Dymnikova:1998ps}. It resulted in specific time dependence of $\\Lambda$, being determined by the cross section of interaction of evaporated particle. The maximal estimation of this cross section gave the time variation of $\\Lambda$~\\cite{Dymnikova:2001ga,Dymnikova:2001jy}, leading to possible realization of our model in the form $\\dot \\Lambda\\sim H^3$.  Therefore, motivated by the time variation of $\\Lambda$ and using the phenomenological model of the kind $\\dot \\Lambda\\sim H^3$ the expression for the time-dependent equation of state parameter $\\omega$ and various physical parameters are derived in the present investigation. The change in sign of the deceleration parameter q has also been shown in this case which refers to the evolution of the Universe via a phase transition from deceleration to the present acceleration.  Interestingly the expression for energy density obtained in the present model is negative. However, this type of negative density is not at all unavailable in the literature. It can be mentioned here that the negative energy density was first obtained by Casimir \\cite{Casimir1948}. Hawking found the existence of negative energy density at the horizon of a black hole \\cite{Hawking1975}. Davies and Fulling also studied about the negative energy fluxes in the radiation from moving mirrors \\cite{Davies1976,Fulling1977}.  The scheme of the present investigation is as follows: in Section 2 the field equations are provided whereas their solutions as well as the physical features have been sought for in the Sections 3 and 4. In the Section 5 some concluding remarks have been made. ",
        "conclusions": "The objective of this work is to observe the effect of a time-dependent equation of state parameter on a dynamic $\\Lambda$ model selected for dark energy investigation. Assuming $\\dot \\Lambda\\sim H^3$, expressions for time dependent equation of state parameter and matter density have been derived. This assumption can find physical justification through the model of Bose-Einstein condensate evaporation~\\cite{Dymnikova:2001ga,Dymnikova:2001jy}. It is also to be noted that the change of sign of the deceleration parameter has been achieved under this special assumption. Our solutions indicate very nontrivial features of cosmological parameters (see Figs. 1 - 11). They are mostly oscillating and contain singularities. In particular, singular and oscillating nature of $\\omega(t)$ is demonstrated for large $t$ on Fig. 2  Moreover, from the solution set it is clear that the expression for density comes out to be a negative quantity with a positive pressure counter part. Actually, here what is shown is a possible cosmological solution with negative energy density. The mechanism of the accelerated expansion of the present Universe due to this positive pressure and negative density is not yet understood properly. In general, contemporary literature (see e.g. \\cite{Ray2007} and references therein) suggest that via Cosmological parameter $\\Lambda$-dark energy acts as the role for the repulsive pressure which is responsible for the accelerated expansion. Here, in our case, negative energy density seems to possess the same role of repulsive gravitation. Perhaps it's role could be understood in a model with dark matter in the presence of $\\Lambda$ where dark matter will be associated with repulsive nature due to negative density \\cite{Goodman2000}.  The main idea of the result is that there is an acceleration due to negative energy density, but at positive pressure. Though physically it is an awkward situation but not at all unavailable in the literature (as mentioned earlier in the introduction) and also quantum field theory admits it. However in cosmology, we don't see negative energy density phenomena very often. This is because there may be some mechanism restricting negative energies or their interactions with ordinary fields. The idea of anti-gravity (i.e. gravitational repulsion between matter and antimatter) originates from this kind of negative energy \\cite{Hossenfelder2006}. Due to symmetry, each gravitating standard model particle corresponds to an anti-gravitating particle. It is thought that this anti-gravitating particle cancels gravitating particle's contribution to the vacuum energy and hence provides a mechanism for smoothing out of the cosmological constant puzzle \\cite{Quiros2004}.  We would also like to point out further that-  1. The present work has been done by keeping the gravitational constant $G$ as a constant. Therefore, a future work can be carried out along the line of this model with $G$ as a variable.  2. Unless we study {\\it dark matter} + $\\Lambda$, we can not say anything conclusive about $\\Lambda$-CDM cosmology.  3. It is also interesting to consider early Universe for $\\Lambda$ + {\\it radiation}, or more general $\\Lambda$ + {\\it radiation} + UDM, and then to study physics of such Universe and possible observational constraint.  4. It is not very easy to calculate the present age of the Universe from the expressions of various parameters derived in the present work. However, numerical methods may be helpful in calculating the present age from Eq. (15) by using different values of $H_0$ given in  \\cite{Ray2007}. Since the purpose of the present paper is different, so that attempt has not been made here."
    },
    "0711/0711.2156_arXiv.txt": {
        "abstract": "The presence of blue straggler stars (BSs) in star clusters has proven a challenge to conventional simple stellar population (SSP) models. Conventional SSP models are based on the evolution theory of single stars. Meanwhile, the typical locations of BSs in the colour-magnitude diagram of a cluster are brighter and bluer than the main sequence turn-off point. Such loci cannot be predicted by single-star evolution theory. However, stars with such properties contribute significantly to the integrated light of the cluster. In this paper, we reconstruct the integrated properties of the Large Magellanic Cloud cluster ESO~121--SC03, the only cluster populating the well-known age gap in the cluster age distribution, based on a detailed exploration of the individual cluster stars, and with particular emphasis on the cluster's BSs. We find that the integrated light properties of ESO~121--SC03 are dramatically modified by its BS component. The integrated spectral energy distribution (ISED) flux level is significantly enhanced toward shorter wavelengths, and all broad-band colours become bluer. When fitting the fully integrated ISED of this cluster based on conventional SSP models, the best-fitting values of age and metallicity are significantly underestimated compared to the true cluster parameters. The age underestimate is $\\sim40$ per cent if we only include the BSs within the cluster's half-light radius and $\\sim60$ per cent if all BSs are included. The corresponding underestimates of the cluster's metallicity are $\\sim30$ and $\\sim60$ per cent, respectively. The populous star clusters in the Magellanic Clouds are ideal objects to explore the potential importance of BSs for the integrated light properties of more distant unresolved star clusters in a statistically robust manner, since they cover a large range in age and metallicity. ",
        "introduction": "Even using improved observational technologies, the number of galaxies with resolved stellar contents is limited. Comparison of the observed integrated spectral properties with models of ``simple'' stellar populations (SSPs; single-age, single-metallicity stellar populations), i.e. using the so-called population synthesis technique, is therefore a practical approach to studying the formation and evolution processes in unresolved galaxies and their components. As the basic building blocks of the evolutionary population synthesis (EPS) method, conventional SSP models are generally constructed based on the evolution theory of single stars. (We specify the conventional SSP models throughout this paper as the theoretical SSP models based on single-star evolution theory.) Thus far, accuracy assessments of the existing suites of conventional SSP models have not yet adequately considered a potentially significant problem revealed by observations of star clusters.  By assuming that all the original member stars within a cluster are born at the same time from the same progenitor molecular cloud, so that they have the same age and metallicity, star clusters are generally considered as the closest counterparts in the real world to idealised SSPs. The major differences between a star cluster and a conventional model SSP of the same age and metallicity can be illustrated by fitting the observed colour-magnitude diagram (CMD) of the cluster with a theoretical single-star isochrone. Most of the member stars can be well fitted by the isochrone in terms of their positions in the CMD. The number distributions along the isochrone can be retrieved via the adopted stellar initial mass function (IMF), and the number of low-mass member stars can be conserved in the presence of evaporation due to dynamical evolution. Then, the most prominent difference should be attributed to those member stars ``straggling'' away from the isochrone. Yellow and red stragglers (Protegies~Zwart et al. 1997a,b) are not considered in the work, because they are comparatively rare and less luminous than blue stragglers (BSs). The underlopers and faint stars bluer than the main sequence (MS) are not bright enough to be important. Therefore, special attention has only been focussed on the BSs in our previous work (Deng et al. 1999; Xin \\& Deng 2005, hereafter XD05; Xin, Deng \\& Han 2007, hereafter XDH07), in view of their potential non-negligible influence on the integrated properties of star clusters.  BSs have been widely observed in stellar systems of all scales and complexities, as shown by, e.g. Ahumada \\& Lapasset (2007) for Galactic open clusters (OCs), Piotto et al. (2004) for Galactic globular clusters (GCs), Carney et al. (2005) for the Galactic field stellar population, Johnson et al. (1999) for Large Magellanic Cloud (LMC) clusters, Alcaino et al. (2003) for Small Magellanic Cloud (SMC) clusters, and Mapelli et al. (2007) for dwarf spheroidal galaxies. Considering their general locations in the CMD, i.e., brighter and bluer than the MS turn-off point, BSs are remarkably hotter than the most luminous ``normal'' MS stars. Typically for old Population I star clusters, when the red clump giants (RCGs) are populated instead of the blue horizontal branch (HB), BSs can significantly enhance the cluster' integrated spectrum at ultraviolet (UV) and blue wavelengths -- the modifications may lead to the star cluster being predicted as younger or of lower metallicity based on conventional SSP models, and therefore cause uncertainties when simply applying the conventional SSP models in EPS studies (XD05; XDH07). Meanwhile, the most likely formation mechanisms of BSs are all related to stellar interactions, including mass transfer in close binaries (e.g., Tian et al. 2006; and references therein) and stellar collisions in high-density regions (e.g., Sills et al. 2005; and references therein). This provides the theoretical support for this empirical work, since the consequences of the products of stellar interactions are not included in the conventional SSP models.  In order to empirically correct the conventional SSP models for the presence of luminous BSs, we construct the integrated spectrum of an SSP (i.e., a realistic star cluster) by analysing the individual stars, after careful consideration of their membership probability. We assume that the stellar population of a cluster is composed of two components: (i) all member stars that can be well fitted by a single-star isochrone represent a conventional SSP for the cluster's real age and metallicity (e.g., SSP models from Bruzual \\& Charlot 2003; hereafter BC03); and (ii) BSs are responsible for the modifications to the conventional SSPs, and their contributions are considered by including the individual spectra of BSs. By combining these two components we may obtain a better approximation to the true SSP corresponding to the cluster of interest (see for details XD05 and XDH07).  In our previous work (XD05; XDH07), a sample of 100 Galactic OCs spanning an age range $\\in[0.1, 10]$ Gyr and a range in metallicity, $Z\\in[0.005, 0.035]$ (where solar metallicity is Z$_\\odot = 0.020$), was used to detect BS contribution semi-empirically. The general results from our Galactic OC studies show that the integrated spectral properties of the sample clusters are dramatically modified by their BS components. A preliminary assessment of the uncertainties inherent to the conventional SSP models was made. Using either spectra or broad-band colours, the resulting ages and/or metallicities will be underestimated significantly: conservatively, the age underestimates are $\\sim50$ per cent, as are the likely underestimates of the metallicity (although they are less accurately constrained because of our lower resolution in metallicity).  The results from the Galactic OCs clearly showed the non-negligible contributions of BSs to the integrated properties of SSPs. However, the numbers and distributions of BSs in OCs are affected by significant stochastic uncertainties (Ahumada \\& Lapasset 1995, 2007). To make the results and conclusions more reliable and convincing, we expand our working sample by including star clusters in the Magellanic Clouds (MCs). MC clusters hold significant advantages over Galactic OCs: (i) they span a wider range in both age and metallicity, and (ii) they are significantly more massive than their counterparts in the Galaxy, and therefore provide better statistics on an individual cluster basis (e.g., Mackey \\& Gilmore 2003a,b; de~Grijs \\& Anders 2006).  In this paper, the LMC cluster ESO~121--SC03 is analyzed using its integrated spectral properties. This serves as the first example to detect BS contributions to the conventional SSP models using a massive intermediate-age cluster in a low-metallicity environment, based on a novel approach and better statistics than we were afforded by our use of Galactic OCs. ESO~121--SC03 ($\\alpha_{2000}=06^{\\rm h}02^{\\rm m}01^{\\rm s}.36$, $\\delta_{2000}=-60^\\circ31'22''.6$) is a distant northern LMC cluster, lying at a projected angular separation of $\\sim10^\\circ$ from the LMC centre. It is described as a ``unique'' LMC cluster by Mackey, Payne \\& Gilmore (2006; hereafter MPG06), because it is the only known cluster to lie in the LMC age gap. A significant number of previous studies (e.g., Mateo et al. 1986; Bica et al. 1998; MPG06) claim an absolute age of ESO~121--SC03 in the range of $8-10$ Gyr. Mateo et al. (1986) obtained [Fe/H]~$=-0.9\\pm0.2$ combined with a reddening of $E(B-V)=0.03$ mag. MPG06 derive [Fe/H]~$=-0.97\\pm0.01$ and $E(V-I)=0.04\\pm0.02$ for the cluster. MPG06 also mark a region in the CMD used to define BS candidates in the cluster, but they do not study the BS population in detail.  We present our photometric data reduction steps in Section 2. In the subsequent sections we present the details of the model construction and our main results. The CMD and the best-fitting isochrone of the cluster are discussed in Section 3. The identification of BSs in the CMD is described in Section 4. A discussion of the contributions of BSs to the integrated spectral properties of the cluster is presented in Section 5. In Section 6, the synthetic integrated spectral energy distribution (ISED) of the cluster is fitted with conventional SSP models in order to assess the uncertainties. Finally, a summary and the conclusions of this study are presented in Section 7. ",
        "conclusions": "In this paper, the LMC star cluster ESO~121--SC03 is taken as a specific cluster to show in detail the technique and procedure we have developed to assess the potential importance of BSs to the integrated properties of unresolved star clusters: data reduction, isochrone fitting, identification of the BS population, synthetic ISED construction, and the concept of detecting BS contributions to conventional SSP models using integrated light properties.  We construct the synthetic ISED of the cluster by decomposing the stellar population of the cluster into two components: the SSP component composed of all the member stars that can be well fitted by an isochrone of single stars; and the BS component defined by the BSs in the cluster. The synthetic ISED of the sum of these two components is regarded as the real ISED of the cluster. The contributions to the ISED of the cluster owing to the BS component are considered for two cases. For case (i), only BSs inside $R_{\\rm HL}$ are used, which provides a conservative estimate of the modification caused by the presence of BSs; while for case (ii), all BSs in the cluster are included, which renders the maximum likely alteration to the conventional results. These two cases bracket a range of changes to the conventional SSP models caused by the BS population in a real cluster.  Integrated spectral properties are used to measure the modifications to the conventional SSP due to BSs. By the inclusion of BSs, the synthetic ISED of ESO~121--SC03 is greatly enhanced toward shorter wavelengths, and all UV and optical broad-band colours become bluer. Both of these observations can make the population appear a lot younger or more metal poor than would be derived from fitting conventional SSP models. The modifications are quantified by fitting the synthetic ISED with conventional SSP models. For the synthetic cluster ISED including only BSs inside $R_{\\rm HL}$, the age underestimate is at the level of $\\sim40$ per cent, and the metallicity underestimate is $\\sim30$ per cent. For the synthetic ISED that includes all BSs identified in the cluster, the underestimates of both age and metallicity increase to $\\sim60$ per cent.  In view of the common occurrence of BSs in MC star clusters (e.g., Johnson et al. 1999; Alcaino et al. 2003), and given that our results for the comparatively rich MC clusters (at least, with respect to the Galactic OCs) are statistically robust, we have reason to believe that the BS effect discussed in this paper for ESO~121--SC03 may not be an exception. BS contributions (to the conventional SSPs) measured based on other MC star clusters are expected to be at a similar level as for ESO~121--SC03, and likely follow the same trend as for our previous results based on Galactic OCs covering a wide age range (Deng et al. 1999; XD05; XDH07): the clusters' age and/or metallicity estimates will be underestimated significantly, by a factor of 2--4, based on fits of SSPs to either spectra or broad-band colours. In a subsequent paper, we will discuss the equivalent results for the entire MC star cluster sample, i.e., based on the application of the method discussed here, which will significantly extend the parameter coverage and the robustness of this conclusion.  In order to constrain the ISEDs of real stellar populations in a reliable way and to provide the community with a set of modified SSP models, extending our work on cluster samples to the MC rich star clusters is now in progress. In the context of our ultimate goal, we are currently still at the stage of information collection. To correct the conventional SSP models empirically, a sufficient working sample of clusters covering large ranges of age and metallicity, and of different dynamical environments, is definitely needed. Such a sample will provide statistical results of the BS contributions, and reliable information of BS populations in SSPs, such as the specific frequency of BSs (BS numbers in an SSP), and luminosity and colour functions of BSs (BS positions in a CMD). These will, we expect, eventually lead to the final solution to this problem, when we can construct the BS-corrected SSP models which can then be used in EPS studies in unresolved conditions."
    },
    "0711/0711.0962_arXiv.txt": {
        "abstract": "Photometric redshift (photo-z) estimates are playing an increasingly important role in extragalactic astronomy and cosmology. Crucial to many photo-z applications is the accurate quantification of photometric redshift errors and their distributions, including identification of likely catastrophic failures in photo-z estimates. We consider several methods of estimating photo-z errors and propose new training-set based error estimators based on spectroscopic training set data. Using data from the Sloan Digital Sky Survey and simulations of the Dark Energy Survey as examples, we show that this method provides a robust, relatively unbiased estimate of photo-z errors. We show that culling objects with large, accurately estimated photo-z errors from a sample can reduce the incidence of catastrophic photo-z failures. ",
        "introduction": "While spectroscopic redshifts have now been measured for over one million galaxies, in recent years digital sky surveys have obtained multi-band imaging for over a hundred million galaxies. Deep, wide-area surveys planned for the next decade will increase the number of galaxies with multi-band photometry to a few billion. Over the last decade, substantial effort has gone into developing photometric redshift (photo-z) techniques, which use multi-band photometry to estimate approximate galaxy redshifts \\citep{con95a,bol00,ben00,col04,wad04}. For many applications in extragalactic astronomy and cosmology, the precision achieved by photometric redshifts is sufficient, provided one can accurately characterize the uncertainties in the photo-z estimates, i.e., the photo-z errors. A number of recent papers have considered the effects of photo-z errors on cosmological probes including baryon acoustic oscillation \\citep{zha06}, weak lensing tomography \\citep{hut06,ma06}, supernovae \\citep{hut04} and galaxy clusters \\citep{hut04,lim07}.   A number of methods have been proposed to characterize photometric redshift errors to date. They can be roughly divided into two categories: methods based on estimating statistical errors in template fitting, e.g., the $\\chi^{2}$ method and its Bayesian counterparts \\citep{bol00,ben00}; and methods that explicitly propagate errors in the input parameters, typically magnitudes or colors, through the photo-z estimator \\citep[e.g.,][]{bru99b,hsi05,col04}.  The error in a photometric redshift estimate \\zphot~ is simply the difference between the photo-z estimate and the true (hereafter, spectroscopic) redshift, $\\Delta z =$\\zphot $-$\\zspec. In practice, the errors for the vast majority of objects in a deep photometric sample are unknown, since the spectroscopic redshifts are not measured. Our goal is to devise an estimator of $\\Delta z$ that has desirable statistical properties, e.g., minimum bias and variance, based on whatever information is at hand. Given a photo-z estimate, an error estimator should give the range of redshifts over which the true redshift will be found at some confidence level.  In most cases, spectroscopic redshifts are available for a small subset of the photometric sample. Such spectroscopic samples are often used as training sets for empirical or machine-based learning photo-z estimators. In this paper, we develop methods of photo-z {\\it error} estimation that are based on the use of spectroscopic training sets to accurately characterize the error distribution. We show that training-set based error estimators outperform other commonly used methods when a representative training set is available and that they are competitive even when the training set is not fully representative of the photometric sample. In cases where the magnitude errors are not well determined, we show that the relative advantages of the new training-set based methods are further increased.  This paper is organized as follows. In \\S \\ref{section:test}, we describe the data sets that we use in this work. In \\S \\ref{section:train}, we introduce the training-set based error estimators and their implementations, as well as their advantages and disadvantages. For comparison, we review the traditional error estimators in \\S \\ref{section:comparison} and highlight the key differences between them and the training-set based error estimators. We show in \\S \\ref{section:culling} that the over-all photo-z scatter and outlier fraction can be significantly improved by culling objects with high estimated photo-z errors, possibly leading to improved results in analyses that rely on photo-z's. Finally, we offer concluding remarks in \\S \\ref{section:conclusion}. ",
        "conclusions": "\\label{section:conclusion}  In this paper, we have introduced a new approach to estimating photometric redshift errors using a spectroscopic training set. We presented two implementations of the training-set approach, Kd-tree and Nearest Neighbor Error (NNE), and found that NNE is the best error estimator when a representative training set is available. Compared to the $\\chi^{2}$ error and the MDE estimators, training-set based error estimators are less sensitive to systematic errors in magnitude error estimates. They incorporate both the bias and scatter of the photo-z's, important features given the often substantial biases in photo-z estimates. Comparison of NNE and Kd-tree errors with error estimators from the literature shows that these training-set error estimators are in general more accurate and better behaved (in the sense that the error residual distribution is closer to a Gaussian).  Since a fully representative spectroscopic training set is not always available, we explored the impact on these error estimates of non-representative training sets. We found that this does not substantially degrade the accuracy of the training-set error estimates. In fact, we showed that even for training sets with very different magnitude and redshift distributions from the photometric sample, the training-set error estimates remain accurate at the 10\\% level.  Finally, we demonstrated that one can cull galaxies with large estimated errors from a sample and thereby significantly improve the overall scatter and bias of the photo-z estimates. Because the training-set error estimators are more accurate than other error estimators, and because the photo-z error residuals are nearly Gaussian distributed for these methods, culling objects using NNE or Kd-tree results in greater performance improvement than culling with other error estimators."
    },
    "0711/0711.1013_arXiv.txt": {
        "abstract": "We present intrinsic extinction curves for 14 AGNs.  The AGNs have reddenings, E(B-V), of up to 0.36 mag.  The majority (13 out of 14) of the extinction curves are not steep in the UV. Of the seven best determined extinction curves, five have extinction curves that are as flat as the standard Galactic curve in the optical and near UV, but flatter in the far UV, and without the $\\lambda$2175 feature. One AGN, B3 0754+394, has a steep SMC-like extinction curve, and another, Mrk 304, has an LMC-like extinction curve, including a probable $\\lambda$2175 bump.  The remaining seven, lower-quality, extinction curves have overall shapes that are consistent with an LMC-like shape or a flatter shape.  Two have possible $\\lambda$2175 features, and one might be identical to the Galactic curve. The flatter curves that predominate in our best determined extinction curves are not as flat as the Gaskell et al.\\@ (2004) extinction curve for radio-loud AGNs. This suggests that the previous radio-loud extinction curve might be slightly too flat in the range $4 < 1/\\lambda < 6.5$ $\\mu$m$^{-1}$ because of luminosity-dependent reddening biases in the composite spectra, but further investigation is needed. We present a parameterized average AGN extinction curve. Observed variations in the continuum properties of the AGNs are inconsistent with intrinsic object-to-object variations because observed differences are least in the far UV where changes in the accretion disk spectrum should be greatest. We suggest that the steepening of AGN spectra around Lyman $\\alpha$ is the result of a small amount of SMC-like dust (E(B-V) $\\sim 0.03$). We find the largest object-to-object differences in spectral shape to be in the \\ion{Fe}{2} emission of the ``small blue bump''. ",
        "introduction": "Knowing the reddening of the continua and emission lines of AGNs is vitally important for understanding the energy-generation mechanism of AGNs, the physical conditions of the line-emitting gas, the differences between AGNs, and AGN demographics.  Gaskell et al.\\@ (2004) showed that reddening is the major cause of differences in optical to UV spectral energy distributions (SEDs), and that the mean extinction in AGNs increases with decreasing luminosity. Gaskell et al.\\@ estimated mean AGN extinction curves from composite spectra based on samples of AGNs with different mean extinctions. These extinctions curves differ from the well known extinction curve for dust in the solar neighborhood by lacking the $\\lambda$2175 carbon feature. The mean extinction curve for the Small Magellanic Cloud (SMC) also lacks the $\\lambda$2175 feature, but the mean AGN extinction curves found by Gaskell et al.\\@ (2004) are considerably flatter than the SMC curve in the ultra-violet. Czerny et al.\\@ (2004) also derived a flat extinction curve with no $\\lambda$2175 bump from SDSS composite AGN colors.  The extinction curves derived by Gaskell et al.\\@ (2004) and Czerny et al.\\@ (2004) refer to samples of AGNs. Using composite spectra has the advantage of minimizing the effects of possible real differences in the spectral-energy distributions, but the composite spectra necessarily combined quasars of different redshifts.  For an apparent-magnitude-limited sample the objects at higher redshift have higher mean luminosities.  The luminosity-dependence Gaskell et al.\\@ find for the mean extinction therefore introduces the possibility of bias in deriving a extinction curve for AGNs from composite spectra, since the UV part of the extinction curve will tend to come from AGNs with lower mean extinction (Willott 2005). In this paper we avoid this bias by deriving extinction curves for individual AGNs.  The derivation of the extinction curves is discussed in sections 2 and 3, and the curves are presented in section 4.  We consider luminosity dependence in section 5 and in section 6 we show that the extinction curves we derive cannot be the result of intrinsic variation of the underlying quasar spectral energy distributions or of varying amounts of host galaxy light. We discuss our results in section 7. ",
        "conclusions": "In agreement with Gaskell et al.\\@ (2004) we argue that the UV-optical continuum shape is profoundly affected by reddening for all but the bluest AGNs.  We believe that all apparent differences in the UV-optical continua of AGNs are primarily due to differences in extinction and in the forms of the extinction curves.  Until the effects of reddening are carefully allowed for it is not possible to investigate possible remaining differences in the spectral energy distributions of AGNs as a function of such things as luminosity, black-hole mass, accretion rate, or radio power. At present all AGNs have optical-UV continuous spectra that appear to be consistent with a single shape after allowance for the effects of reddening and the small blue bump.  Even in a sample, such at the {\\it HST/FUSE} sample, biased towards low reddening, most of the objects show detectable reddening.  Our results support the interpretation of Shang et al.\\@ (2003) that the second eigenvector in their spectral principal component analysis is due to dust reddening, and disagree with the interpretation of Yip et al.\\@ (2004) that the main continuum eigenvectors are due to changing host galaxy contributions and variations in the intrinsic spectral index.  We find rising SMC-like reddening curves to be rare in the relatively low reddening {\\it HST/FUSE} sample we consider and believe that they are not the most common reddening curve for AGNs. Flatter reddening curves imply that the visual extinction of AGNs is greater than if the reddening curves were SMC-like.  We confirm that the $\\lambda$2175 extinction feature is mostly absent in AGN extinction curves. The rise into the far UV seen in the Galactic, LMC, and SMC reddening curves seems also to be mostly absent.  These differences imply that the dust has been modified by the AGN environment.  Our new mean AGN extinction curve, while flat in the far UV (1/$\\lambda > 6.5$ $\\mu$m$^{-1}$) is steeper in the range $3.5 < 1/\\lambda < 6.5$ $\\mu$m$^{-1}$ than the Gaskell et al.\\@ (2004) extinction curve from composite spectra of radio-loud AGNs. Although the objects are not the same, the difference suggests that luminosity-dependent extinction (Willott 2005) might have caused the Gaskell et al.\\@ curve to be too flat in the UV.  In disagreement with Willott (2005), however, we do not find the majority of AGN extinction curves to be SMC-like."
    },
    "0711/0711.3234_arXiv.txt": {
        "abstract": "The modified Gauss-Bonnet gravity can be motivated by a number of physical reasons, including: the uniqueness of a gravitational Lagrangian in four and higher dimensions and the leading order $\\alpha^\\prime$ corrections in superstring theory. Such an effective theory of scalar-tensor gravity has been modeled in the recent past to explain both the initial cosmological singularity problem and the observationally supported cosmological perturbations. Here I present an overview of the recent developments in the use of modified Gauss-Bonnet gravity to explain current observations, touching on key cosmological and astrophysical constraints applicable to theories of scalar-tensor gravity. The Gauss-Bonnet type modifications of Einstein's theory admits nonsingular solutions for a wide range of scalar-curvature couplings. It also provides plausible explanation to some outstanding cosmological conundrums, including: the transition from matter dominance to dark energy and the late time cosmic acceleration. The focus is placed here to constrain such an effective theory of gravity against the recent cosmological and astrophysical observations. ",
        "introduction": "Einstein's general relativity has been very successful as a classical theory of gravitational interactions, especially, in a non-accelerating (or non-expanding) spacetime. In a cosmological background, the theory predicts spacetime singularities, so its modification is inevitable at high energy scales. Further the recently observed accelerating expansion of the universe~\\cite{Supernova} provides some insight to the possibility that general relativity together with ordinary matter and radiation, described by the standard model of particle physics, cannot fully explain the current observations. The question arises because the current observations~\\cite{WMAP03} require in the fabric of the cosmos the existence of a dark energy component of magnitude about $73\\%$, which does not `clump' gravitationally. Another $23\\%$ of the mass-energy is in the form of mysterious non-baryonic dark matter, which `clumps' gravitationally.  The bulk of the universe appears to be dark energy and dark matter. So far there is no fully consistent explanation of these energy components supported by a fundamental theory. The main focus of this meeting is obviously to update our knowledge on DARK matter and DARK energy searches and the physics behind these. The focus of my presentation will be on a possible resolution of dark energy problem within some string-inspired theories of scalar-tensor gravity. ",
        "conclusions": "The important ingredient of the present approach to dark energy cosmology is the treatment of gravitational coupling between the dynamical scalar field $\\varphi$ and the quadratic Riemann invariant of the Gauss-Bonnet form, which gives rise to nonsingular cosmologies for a wide range of scalar-curvature couplings. The model also provides plausible explanation to some outstanding cosmological conundrums, including: the transition from matter dominance to dark energy and the late time cosmic acceleration. Furthermore, the scalar-curvature coupling can easily trigger onset of a late dark energy domination. Despite these promising signs, it remains to be checked whether the Gauss-Bonnet modification of Einstein's theory will lead to genuine contact between observations and string theory.  String theory is known to be free from ghosts and superluminal modes, at least, in a flat ten-dimensional Minkowski background. This is perhaps not essentially the case in a four-dimensional FRW background. The effective string actions in four dimensions may well exhibit some unwarranted features, such as, short scale instabilities and superluminal propagation of tensor or scalar modes, under inhomogeneous (cosmological) perturbations. The model discussed here is perhaps not an exception; at least, for the potential and the Gauss-Bonnet coupling in simplest exponential forms, one could see that the tensor or scalar modes propagate at a speed faster than light at some stage, especially, for large slope parameters. The appearance of a superluminal mode is generic, and perhaps also acceptable, if such an effect is only transient.  In the present proposal for explaining a crossing of cosmological (dark energy) equation of state, $w\\Z{DE}=-1$, and a superluminal propagation of scalar or tensor modes, a number of important, physically falsifiable predictions can be made. These include a transient violation of Lorentz symmetry and the weak equivalence principle, associated with the microscopic effects of the coupling between $\\varphi$ and background matter in high density regions.  Ninety years after Einstein's proposition of general relativity with a cosmological constant, a modified cosmological scenario with its natural generalization is close to experimental test and possibly an outlet compatible with present experimental data. The coming generation of cosmological experiments, including Dark Energy Survey, will probably rule out the great majority of string-derived models, as well as exclude those class of scalar-tensor theories which give rise to unphysical states. {\\sl Time will tell}.   \\medskip  {\\it Acknowledgements}: This research is supported by the Foundation for Research, Science and Technology (New Zealand) under Research Grant No. E5229 and also by Elizabeth Ellen Dalton Research Award (E5393)."
    },
    "0711/0711.4622_arXiv.txt": {
        "abstract": "The character of the first galaxies at redshifts $z\\ga 10$ strongly depends on their level of pre-enrichment, which is in turn determined by the rate of primordial star formation prior to their assembly. In order for the first galaxies to remain metal-free, star formation in minihaloes must be highly suppressed, most likely by H$_2$-dissociating Lyman-Werner (LW) radiation. We show that the build-up of such a strong LW background is hindered by two effects. Firstly, the level of the LW background is self-regulated, being produced by the Population~III (Pop~III) star formation which it, in turn, suppresses. Secondly, the high opacity to LW photons which is built up in the relic H~II regions left by the first stars acts to diminish the global LW background. Accounting for a self-regulated LW background,  we estimate a lower limit for the rate of Pop III star formation in minihaloes at $z$ $\\ga$ 15. Further, we simulate the formation of a 'first galaxy' with virial temperature $T_{\\rm{vir}}\\ga 10^{4}~\\rm{K}$ and total mass $\\ga$ 10$^8$ M$_{\\odot}$ at $z\\ga 10$, and find that complete suppression of previous Pop~III star formation is unlikely, with stars of $\\ga 100~\\rm{M_{\\odot}}$ (Pop III.1) and $\\ga 10~\\rm{M_{\\odot}}$ (Pop~III.2) likely forming. Finally, we discuss the implications of these results for the nature of the first galaxies, which may be observed by future missions such as the {\\it James Webb Space Telescope}. ",
        "introduction": "With the goal of estimating the degree to which Pop III star formation in minihaloes affects the properties of the gas from which the first galaxies form at $z$ $\\ga$ 10, we have carried out two cosmological hydrodynamics simulations which track the detailed evolution of the primordial gas under the influence of a LW background radiation field. The first simulation is designed to test the validity of our self-regulated model for the build-up of the LW background, and is described further in Section 4.  The second simulation follows the assembly of a galaxy virializing at $z$ $\\ga$ 10, assuming a critical LW background, $J_{\\rm LW, crit}$. This simulation is described in Section 5.  As with previous work, for our three-dimensional numerical simulations we employ the parallel version of GADGET (version 1), which includes a tree (hierarchical) gravity solver combined with the smoothed particle hydrodynamics (SPH) method for tracking the evolution of gas (Springel et al. 2001; Springel \\& Hernquist 2002).  Along with H$_2$, H$_2^{+}$, H, H$^-$, H$^+$, e$^-$, He, He$^{+}$, and He$^{++}$, we have included the five deuterium species D, D$^+$, D$^-$, HD and HD$^+$, using the same chemical network as in Johnson \\& Bromm (2006, 2007).  The specific initial conditions for our simulations are described separately in Sections 4 and 5.  \\subsection{Sink Particles} The primary result obtained from the present simulations is the number of Pop~III stars which are able to form in minihaloes at $z$ $\\ga$ 15 during the assembly of a galaxy virializing at $z$ $\\ga$ 10. To track the number and location of the formation sites for Pop III stars, we allow sink particles to form when the gas has collapsed to high densities. Specifically, sinks are created when the gas has collapsed to a threshold density $n_{\\rm res}$, above which the mass resolution of our simulation is insufficient to reliably follow the continued collapse of the gas, as explained further in Johnson et al. (2007).  We note that the time at which a sink particle forms will precede the time at which a star is able to form by roughly a free-fall time in gas with a density $n_{\\rm res}$; while this can introduce an error of $\\la$ 10 Myr in the precise timing of star formation, the conclusions that we derive from the results presented here are not significantly impacted. Once having formed, the sink particles act as tracers of star formation sites, each sink representing the formation site of a single Pop~III star. We focus our attention on when and where these sinks emerge in order to estimate the number of stars formed, and to gauge the feedback effects from these stars.  \\subsection{Implementation of Background Radiation} Motivated by the considerations presented in Section 2 and further justified in Section 4, we have included in our simulations the self-regulated model of the LW background radiation field, in the form of a global H$_2$ dissociation rate.  The LW background is taken to be a function of redshift, its evolution given by $J_{\\rm LW, crit}(z)$, as shown in Fig.~2. The dissociation rate is  \\begin{equation} k_{\\rm diss}(z) = 10^{-12} J_{\\rm LW, crit}(z) {\\rm s}^{-1}\\mbox{\\ .} \\end{equation}  In addition to the photodissociation of H$_2$, the photodestruction of other species can, in principle, affect the cooling properties of the primordial gas.  The main reaction producing H$_2$ being H$^-$ + H $\\to$ H$_2$ + e$^-$ (e.g. Peebles \\& Dicke 1968; Lepp \\& Shull 1984), the photodestruction of H$^-$ is a process that might lead to suppression of the H$_2$ fraction of the gas (Chuzhoy et al. 2007; see also Glover et al. 2006).  However, we find the photodissociation of H$^-$ to be unimportant for the suppression of H$_2$ in collapsing minihaloes.  As discussed in Section 2.1, the gas in such minihaloes will adiabatically collapse to densities of $\\ga$ 1 cm$^{-3}$ and temperatures of $\\ga$ 10$^3$ K, even if the fractions of H$_2$ or H$^-$ are low.  Following this adiabatic collapse, the fate of the gas does depend sensitively on the abundances of these species, and it is only at this stage at which the photodissociation of H$^-$ might begin to have an effect.  We find that the chemical species at this stage, absent strong shocks or photoionization, are roughly in equilibrium when we include the effect of the LW background, and therefore suppression of the H$_2$ abundance can be directly tied to the abundance of H$^-$ at this stage.  At this temperature and density, we find that the H$^{-}$ formation time is of the order of years.  We find that the minimal photodissociation timescale of H$^-$, produced for the case of the maximal LW background $J_{\\rm LW, max}$ (see Section 2.2), is of the order of $\\ga$ 10$^2$ yr, roughly two orders of magnitude higher than the formation timescale.  Thus, the H$^-$ fraction will not be significantly affected by a background radiation field, and in turn the abundance of H$_2$ will not be significantly affected, allowing us to neglect the photodissociation of H$^-$ as a factor in delaying star formation. Furthermore, we neglect the photodissociation of H$^-$ due to radiation from local star formation for similar reasons, as explained in Johnson et al. (2007).  We have also neglected Ly$\\alpha$ heating as a suppressant of minihalo star formation (see Chuzhoy \\& Shapiro 2007; Ciardi \\& Salvaterra 2007). For the Pop III star formation rates, and corresponding production rates of UV photons, found in our simulations, Ly$\\alpha$ heating will not raise the temperature of the IGM appreciably, even for the maximal SFR considered here, ${\\rm SFR}_{\\rm III, max}$ (see Section 2.2). Finally, we do not include a photoionizing background in our calculations. At these high redshifts, the IGM is still predominantly neutral. Photoionization will thus generally occur only locally, when isolated Pop~III stars form, an intermittent process which is not likely to dramatically affect the overall SFR (see e.g. Ahn \\& Shapiro 2007; Johnson et al. 2007; Whalen et al. 2007; Wise \\& Abel 2007a; but see Ricotti et al. 2002a,b). ",
        "conclusions": ""
    },
    "0711/0711.3528_arXiv.txt": {
        "abstract": "We present a possible explanation of the recently observed 511 keV $\\gamma$-ray anomaly with a new ``millicharged'' fermion. The new fermion is light (${\\cal O}({\\rm MeV})$) but has never been observed by any collider experiments mainly because of its tiny electromagnetic charge $\\varepsilon e$. We show that constraints from its relic density in the Universe and collider experiments allow a parameter range such that the $511$ keV cosmic $\\gamma$-ray emission from the galactic bulge may be due to positron production from this millicharged fermion. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.1377_arXiv.txt": {
        "abstract": "We consider the radiative feedback processes that operate during the formation of the first stars, including the photodissociation of H$_2$, \\lal\\ radiation pressure, formation and expansion of an \\ion{H}{2} region, and disk photoevaporation. These processes may inhibit continued accretion once the stellar mass has reached a critical value, and we evaluate this mass separately for each process. Photodissociation of H$_2$ in the local dark matter minihalo occurs relatively early in the growth of the protostar, but we argue this does not affect subsequent accretion since by this time the depth of the potential is large enough for accretion to be mediated by atomic cooling. However, neighboring starless minihalos can be affected. Ionization creates an \\ion{H}{2} region in the infalling envelope above and below the accretion disk. \\lal\\ radiation pressure acting at the boundary of the \\ion{H}{2} region is effective at reversing infall from narrow polar directions when the star reaches $\\sim 20-30 M_\\odot$, but cannot prevent infall from other directions. Expansion of the \\ion{H}{2} region beyond the gravitational escape radius for ionized gas occurs at masses $\\sim 50-100\\sm$, depending on the accretion rate and angular momentum of the inflow. However, again, accretion from the equatorial regions can continue since the neutral accretion disk has a finite thickness and shields a substantial fraction of the accretion envelope from direct ionizing flux. At higher stellar masses, $\\sim 140\\sm$ in the fiducial case, the combination of declining accretion rates and increasing photoevaporation-driven mass loss from the disk act to effectively halt the increase in the protostellar mass. We identify this process as the mechanism that terminates the growth of Population III stars (i.e., stars with primordial composition) that have not been affected by prior star formation (Population III.1 stars).  We discuss the implications of our results for the initial mass function of these stars. In the Appendix we develop approximate solutions to a number of problems relevant to the formation of the first stars: the effect of Rayleigh scattering on line profiles in media of very large optical depth; the intensity of Lyman-$\\alpha$ radiation in very opaque media; an approximate determination of the radiative acceleration in terms of the gradient of a modified radiation pressure; the determination of the flux of radiation in a shell with an arbitrary distribution of opacity; and the vertical structure of an accretion disk supported by gas pressure with constant opacity. ",
        "introduction": "}  There has been substantial recent progress in our theoretical understanding of how the first stars formed (Bromm \\& Larson 2004). In marked contrast to the case for contemporary star formation, the initial conditions for the formation of the first stars are believed to be relatively well understood: they are determined by the growth of small-scale gravitational instabilities from cosmological fluctuations in a cold dark matter universe. The first stars are expected to form at redshifts $z\\sim 10-50$ in dark matter halos of mass $\\sim 10^6 M_\\odot$ (Tegmark et al. 1997).  In the absence of any elements heavier than helium (other than trace amounts of lithium) the chemistry and thermodynamics of the gas are very simple (Abel, Bryan, \\& Norman 2002, hereafter ABN; Bromm, Coppi, \\& Larson 2002, hereafter BCL02). There are no dust grains to couple the gas to radiation emitted by the protostar.  There are no previous generations of stars to roil the gas out of which the stars form, nor is there any radiation other than the cosmic background radiation.  Existing calculations have assumed that there were no significant primordial magnetic fields, thereby eliminating a major complication that occurs in contemporary star formation.  However, even in the absence of significant primordial fields, it is possible that magnetic fields could have been generated in the accretion disk surrounding a primordial protostar (Tan \\& Blackman 2004), although even in this case the magnetic fields become dynamically important only after the star formation process is well underway, and they do not affect the initial conditions.  Given this relative simplicity, there is some confidence in the results of numerical simulations that have followed the collapse of cosmological scale perturbations down to almost stellar dimensions (ABN; Bromm, Coppi, \\& Larson 1999; Yoshida et al. 2006; O'Shea \\& Norman 2007). This confidence is strengthened by the fact that it appears to be the microphysics of $\\rm H_2$ cooling that determines the types of baryonic structures that are formed, and not, for example, the details of the initial power spectrum of fluctuations in dark matter density. The results of these simulations suggest that the initial gas cores out of which stars form are quite massive, $\\mcore \\sim 100 - 1000\\sm$.  The observational imprint of the first stars depends on their mass: These stars were likely to be of critical importance in reionizing the universe, in producing the first metals, and in creating the first stellar-mass black holes.  The number of ionizing photons emitted per baryon depends on the stellar mass for $m_*\\la 300 M_\\odot$ (Bromm, Kudritzki, \\& Loeb 2001). The hardness of the radiation field is also sensitive to the mass (e.g. Tumlinson \\& Shull 2000; Schaerer 2002), so that He reionization can be affected. The effectiveness of the first stars in enriching the intergalactic medium (IGM) with metals and in producing the first stellar-mass black holes also depends sensitively on the mass of the star. A potential simplification in assessing these effects is that massive primordial stars are thought to have much smaller mass-loss rates than contemporary massive stars (Kudritzki 2002), so that the mass at core collapse should be quite similar to the initial mass.  However, Meynet, Ekstr\\\"om, \\& Maeder (2006) have argued that if rotation is allowed for, then mass loss can be significant.  Heger \\& Woosley (2002) showed that stars exploding as supernovae above about $260 M_\\odot$ and between 40 and $140 M_\\odot$ should collapse directly to black holes, and they argued that such stars would providing relatively little metal enrichment.  However, Ohkubo et al. (2007) followed the collapse and explosion of $500 M_\\odot$ and $1000 M_\\odot$ stars in two dimensions and concluded that substantial amount of metals could be ejected; they proposed that such supernovae could produce intermediate-mass black holes.  Heger \\& Woosley (2002) also showed that for $140\\la m_*\\la 260 M_\\odot$, the pair instability leads to explosive O and Si burning that completely disrupts the star, leaving no remnant and ejecting large quantities of heavy elements. Such supernovae produce a dramatic odd-even effect in the nuclei produced.  Stars below $\\sim 40 \\sm$ are expected to form neutron stars, with more normal enrichment rates. In principle, metallicity determinations from high-redshift absorption line systems (Schaye et al. 2003; Norman, O'Shea, \\& Paschos 2004) or from very metal poor local stars (Beers \\& Christlieb 2005) can constrain the initial mass function (IMF) of the early generations of stars. Indeed, based on observations such as these, Daigne et al. (2004) argue that the stars responsible for reionizing the universe mostly likely had masses $\\la 100\\sm$, and Tumlinson (2006) concludes that stars above $140 M_\\odot$ could have contributed at most 10\\% of the iron observed in extremely metal poor stars (those with [Fe/H]~$<10^{-3}$---Beers \\& Christlieb 2005).  In discussing the formation of the first stars, the terms ``first stars'' and ``Population III stars'' are often used interchangeably, but this can lead to confusion. To be precise, we shall follow the conventions suggested by one of us at the First Stars III conference (O'Shea et al. 2008) and define Population III stars as those stars with a metallicity sufficiently low that it has no effect on either the formation or the evolution of the stars. The value of the critical metallicity for star formation---i.e., the value below which the metals do not influence star formation---is uncertain, with estimates ranging from $\\sim 10^{-6}Z_\\odot$ if the cooling is dominated by small dust grains that contain a significant fraction of the metals (Omukai et al. 2005) to $\\sim 10^{-3.5}Z_\\odot$ if the cooling is dominated by the fine structure lines of C and O and there is negligible \\htwo\\ (Bromm \\& Loeb 2003); if \\htwo\\ cooling is included, Jappsen et al. (2007) argue that there is no critical metallicity for gas-phase metals.  It is possible that values of the metallicity even less than $10^{-6}Z_\\odot$ could significantly affect the evolution of primordial stars (Meynet, private communication). Among Population III stars, we distinguish between the first and second generations, termed Population III.1 and III.2, respectively: The initial conditions for the formation of Population III.1 stars are determined solely by cosmological fluctuations whereas those for Population III.2 stars are significantly affected by other stars.  It is likely that Population III.1 stars have a primordial composition, since it is hard to see how the gas out of which they form could be contaminated by even small trace amounts of metals without having been affected by radiation from the star that produced the metals.  Stars affect the initial conditions for the formation of Population III.2 stars primarily by radiation, both ionizing radiation and Lyman-Werner band radiation that destroys \\htwo\\ molecules. The latter reduces the cooling efficiency of the gas, allowing compression to heat up to the point that it begins to become collisionally ionized; shocks associated with \\ion{H}{2} regions and supernova remnants can also ionize the gas. Once the gas has been partially ionized, HD cooling can become important, reducing the characteristic star-formation mass (Uehara \\& Inutsuka 2000). Greif \\& Bromm (2006) use the term ``Population II.5\" to describe stars that form from gas in which HD cooling is important, whereas in our terminology these would be Population III.2 stars, but we believe that it is better to describe all essentially metal-free stars as ``Population III.\" It should be noted that our definition of Population III.2 stars includes all Population III stars that were significantly affected by previous generations of star formation, even if that did not result in significant HD production.  Those Population III.2 stars that form out of gas that is enriched in HD will typically be less massive than Population III.1 stars---by about a factor 10 according to Greif \\& Bromm (2006). They infer that Population III.1 stars are relatively rare, constituting about 10\\% by mass of all Population III stars.   In this paper we wish to estimate the characteristic mass of the first generation of stars (Population III.1). Even if they are relatively rare, they are critical in setting the initial conditions for the star formation that followed, and therefore in determining the reionization of the universe and the production of the first metals and the first stellar mass black holes. For contemporary star formation, it is believed that the IMF is set by a combination of gravitational fragmentation in a turbulent medium (Elmegreen 1997; Padoan \\& Nordlund 2002) and feedback effects.  The characteristic stellar mass is of order the Bonnor-Ebert mass, $m_{\\rm BE}\\propto T^{3/2}/\\rho^{1/2}$.  However, not all of the initial core mass is incorporated into the final star, since contemporary protostars have powerful outflows that eject some of the core mass (Nakano, Hasegawa, \\& Norman 1995; Matzner \\& McKee 2000). There are a number of feedback effects that occur in contemporary massive star formation (Larson \\& Starrfield 1971), particularly radiation pressure on dust and photoionization associated with the growth of an \\ion{H}{2} region. It remains unclear whether the upper limit on the contemporary IMF is set by feedback or by instabilities that afflict very massive stars.  The clumps out of which the first stars form have total masses, including dark matter, of order $10^6 M_\\odot$; these objects are typically referred to as minihalos.  Cooling by trace amounts of \\htwo\\ generally leads to the formation of a gravitationally unstable core of baryonic mass $\\sim 10^{2-3} \\sm$ (ABN). In contrast to contemporary star-forming regions, the turbulence in this gas is subsonic, due to the higher temperatures and the lack of internal and external sources of turbulence.  As a result, gravitational fragmentation is much less effective---indeed, numerical simulations show no evidence for it (e.g., ABN, Yoshida et al. 2006, but see Clark, Glover, \\& Klessen 2008), and analytic calculations (Ripamonti \\& Abel 2004), including those that consider disk fragmentation (Tan \\& Blackman 2004), show no evidence for fragmentation either. It therefore appears that the mass of the first stars is likely to be set by feedback effects.  Feedback effects can be classified as either radiative or kinetic. Kinetic feedback includes protostellar outflows and main-sequence winds. In contemporary star formation, outflows are believed to be hydromagnetically driven; in this paper, we assume that the magnetic fields associated with the protostar are too weak or too tangled to drive a strong outflow (see Tan \\& Blackman 2004 for a more extensive discussion of the effect of these outflows). Due to the absence of metals, main-sequence winds are very weak in the absence of rotation (Kudritzki 2002), although they could become important in the later stages of evolution if the star is rapidly rotating. Since we are primarily interested in the early stages of evolution of the star, we neglect kinetic feedback (i.e., outflows and winds).   Radiative feedback includes radiation pressure, photoionization heating, and photodissociation. Radiation pressure can be due to continuum radiation or to resonance line scattering; the continuum radiation pressure can be due to electron scattering or to photoionization. Photoionization also leads to heating, which unbinds gas beyond the gravitational radius $r_g\\equiv Gm_*/c_s^2$, where $c_s$ is the isothermal sound speed of the gas. If the gas is initially in a disk, this process is termed photoevaporation. Finally, photodissociation destroys \\htwo\\, the dominant coolant in neutral primordial gas.  Most previous work has focused on the effects of continuum radiation pressure and photoionization heating in limiting the mass of primordial stars.  Omukai \\& Palla (2001; 2003) focused on electron scattering, which leads to the Eddington limit on an accreting mass. For the case of spherical accretion at a rate of $4.4\\times 10^{-3}\\smyr$, radiation pressure first becomes important at around $80\\sm$, leading to a dramatic swelling of the stellar surface. This, however, is a transient effect because an important part of the luminosity is due to accretion, which is reduced by the increase in the stellar radius.  Only at masses around $300\\sm$ does the internal luminosity become very close to the Eddington value, leading to runaway expansion of the star and, presumably, the end of accretion. Omukai \\& Palla (2003) considered a range of accretion rates. They found that if the accretion rate is smaller than $4.4\\times10^{-3}\\smyr$, then the total luminosity remains sub-Eddington and the star continues to grow along the main sequence to arbitrarily large masses. On the other hand, if the accretion rate is somewhat larger than this critical rate, the Eddington limit becomes important at around $100\\sm$. Accretion at a rate based on the settling motions in the core of ABN is slow enough that the Eddington limit does not affect the final mass. However, these models ignored the influence of other protostellar feedback mechanisms on the infalling envelope. These models also assumed spherical symmetry, which leads to much larger photospheric radii and thus a softer radiation field than in the more realistic case of disk accretion (see \\S 7 in Paper I).   Omukai \\& Inutsuka (2002) considered the the combined effects of photoionization heating and continuum radiation pressure due to photoionization in the infalling envelope.  They show that there is a critical stellar mass at which the hydrogen ionizing luminosity is sufficient to create an \\ion{H}{2} region, which rapidly spreads to large distances where its thermal pressure becomes dynamically important in slowing infall. However, the ionizing radiation force decelerates the inflowing gas, raising the gas density and therefore reducing the radius of the \\ion{H}{2} region. For spherical inflow, this mechanism is so effective that the radius of the \\ion{H}{2} region remains well below the gravitational radius $r_g$, stopping any mass loss.  They concluded that with this effect, there was no limit on protostellar masses below $1000\\sm$. Without this radiation force, they predicted a mass limit of order $300\\sm$. As we shall see below, these conclusions are sensitive to the assumption of spherical accretion.   In Paper I (Tan \\& McKee 2004) we modeled the growth of a primordial protostar from very small masses to large. We included the effects of rotation of the infalling gas, which led to the formation of an accretion disk around the protostar.  The goal of this paper is to determine when the energy output from the protostar is sufficient to halt accretion and set the final stellar mass.  This is an extremely complicated problem, the full solution of which requires three dimensional hydrodynamical simulations that include the generation, propagation, and dynamical influence of radiation. Furthermore these simulations must resolve a large range of scales: the protostar is of order $10\\:R_\\odot$, while the size of the quasi-hydrostatic core that encloses $\\sim 1000\\sm$ is of order 1~pc, several million times greater. The demands on the time scale are even greater: the simulation may have to follow the evolution of the star over its lifetime $\\sim 2 - 3$ Myr (Schaerer 2002) while at the same time following the dynamics of an accretion disk with a characteristic time scale as short as $10^4$ s. The numerical simulations of ABN were able to resolve an even greater range of radii, but it will be some years before it is possible to meet the required dynamical range in time scales. As a result, we shall develop simple analytic models for the feedback interaction that we hope will provide a useful first step.  We begin our discussion with a review the results of Paper I in \\S \\ref{S:review}.  Feedback effects are then considered in the approximate order in which they become manifest. In \\S\\ref{S:dissoc}, we briefly discuss the effects of photodissociation.  \\lal\\ radiation pressure feedback, is considered in \\S\\ref{S:lya} and in several Appendixes.  The feedback from ionizing photons that can create an \\ion{H}{2} region is considered in \\S\\ref{S:ion}.  After discussing shadowing by accretion disks in \\S\\ref{S:shadow} and an Appendix, the feedback due to disk photoevaporation is considered in \\S\\ref{S:diskevap}.  Figure \\ref{fig:overview} gives an overview of these feedback processes occurring near the protostar.  Finally, our conclusions are summarized in \\S\\ref{S:conclusions}.   \\begin{figure}[h] \\begin{center} \\epsfig{ file=f1.eps, angle=0, width=6in} \\end{center} \\caption{ \\label{fig:overview} Overview of the accretion geometry and feedback processes involved in primordial star formation. (a) Top left: Cross section of the accretion geometry: the dashed lines show streamlines of the rotating, infalling gas, with figure of revolution from each streamline separating 10\\% of the total infall from this hemisphere. The aspect ratio of the accretion disk is realistic, while the size of the star has been exaggerated for clarity. (b) Top right: The shaded region around the star shows the extent of the \\ion{H}{2} region, which at this relatively early stage is still confined inside the gravitational radius for the escape of ionized gas, $r_g$. \\lal\\ radiation pressure feedback should be strong enough to prevent accretion in the polar directions. (c) Bottom left: The stellar mass and ionizing luminosity have grown, and the \\ion{H}{2} region is just in the process of breaking out of the accretion flow. Once a significant volume beyond $r_g$ is ionized, accretion from these directions is expected to be shut off. (d) Bottom right: Final stage of accretion involves shadowing of the equatorial region from stellar ionizing flux by the disk, which at the same time is photoevaporated. The competition between this photoevaporative outflow and the residual accretion rate sets the final mass of the star.} \\end{figure} ",
        "conclusions": "\\label{S:conclusions}  Recent numerical studies have indicated that the initial conditions for primordial star formation are dense, massive gas cores in approximate hydrostatic and virial equilibrium. These physical properties are set mostly by the microphysics of $\\rm H_2$ cooling and not by the initial cosmological density perturbations.  We have described the rate of collapse of these gas cores as a function of the entropy parameter of the gas, $K$, and the amount of mass that already collapsed. This accretion rate is very large, so that once an optically thick protostellar core forms the star grows very quickly.  We have developed a simplified method for modeling protostellar evolution and applied the appropriate accretion rate for primordial protostars. The method allows for the treatment of accretion of gas with angular momentum, so that part of the accretion occurs via a disk. Using a realistic degree of rotation for the initial gas core we find that, after the protostar has grown to about a solar mass, essentially all of the accretion flow is via the disk and conditions at the protostar are optically thin, in contrast to the spherical case. This means that the radiation field that the accretion envelope is exposed to is significantly hotter so that ionization and FUV radiation feedback can become important.  We considered the impact of the protostellar feedback on the infalling envelope. Again rotation is important because it modifies the density distribution in the vicinity of the star. First we discussed the effects of photodissociation of \\htwo, the primary coolant. We showed that this does not stop accretion if the protostar has already begun to form, but it can suppress star formation in the vicinity (c.f. Glover \\& Brand 2001).  Next we considered radiation pressure feedback due to resonant scattering of FUV radiation in the \\lal\\ damping wings. As a result of the high column densities of neutral gas around the \\ion{H}{2} region, this radiation is trapped and the pressure amplified by large factors. This radiation pressure becomes larger than the ram pressure of the infalling gas in the polar directions for stellar masses of order 20$\\sm$. However, once the infall is reversed at the poles, the \\lal\\ photons can escape and the accretion in other directions proceeds unimpeded.  We then considered the growth of the ionized region.  Once the expansion velocity of this region exceeds the free-fall velocity, the accretion is halted.  This typically occurred at about 50 to 100$\\sm$, although it took much larger masses for cases with little angular momentum.  The ionized gas is confined to the region above the disk, however, so accretion can continue in the shadow of the disk. Evaluating this, we found that shadowing permitted accretion to continue at a rate of about 20-30\\% of that in the absence of the \\ion{H}{2} region.  Allowing for photoevaporation of the disk, we found that the final stellar mass is about $140 M_\\odot$ in the fiducial case.   Table \\ref{tab:m} summarizes how the mass scales set by protostellar feedback depend on model parameters.  The final mass of a Pop III.1 star depends fairly sensitively on the entropy parameter of the accreting gas (i.e., approximately as $(K^\\prime)^{1.3}$), which in turn determines the overall accretion rate to the star$+$disk, but not very much on core rotation (for $f_{\\rm Kep}\\gtrsim 0.25$) or ionized gas temperature ($T_i$). At very low values of core rotation, \\ion{H}{2} region breakout is delayed until high protostellar masses, at which point the disk photoevaporation rate soon exceeds the residual disk-shadowed accretion rate. However, these small values of $f_{\\rm Kep}$ are not very likely to occur in nature.  The final masses predicted by our model overlap the range of masses necessary to produce pair instability supernovae, $140-260\\;M_\\odot$ (Heger \\& Woosley 2002).  Rotation may lower these limits (S. Woosley, private comm.). The lack of the expected nucleosynthetic signature of such supernovae in the abundance patterns of very metal poor halo stars (Tumlinson et al. 2004), could indicate that such massive Pop III.1 stars were relatively rare or that they tended to enrich regions not probed by typical halo stars, perhaps the centers of larger galactic halos.  The conclusion by Scannapieco et al. (2006) that Pop III star formation should be fairly widespread in regions now probed by Galactic halo stars, can be reconciled with the abundance pattern observations if most of this star formation leads to either Pop III.1 stars from relatively low entropy ($K^\\prime \\lesssim 1$) gas cores or Pop III.2 stars that also have a mass scale below the pair instability threshold (see also the study by Greif \\& Bromm 2006). Further work is required to determine the range of pre-stellar core parameters, primarily $K^\\prime$ and $f_{\\rm Kep}$, exhibited in cosmological simulations, in order to predict the frequency of pair instability supernovae.  \\begin{deluxetable}{cccccc} \\tablecaption{Mass Scales of Population III.1 Protostellar Feedback\\label{tab:m}} \\tablewidth{0pt} \\tablehead{ \\colhead{$K^\\prime$} & \\colhead{$f_{\\rm Kep}$} & \\colhead{$T_i/(10^4\\:{\\rm K})$} & \\colhead{$m_{\\rm *,pb}$ ($\\sm$)\\tablenotemark{a}} & \\colhead{$m_{\\rm *,eb}$ ($\\sm$)\\tablenotemark{b}} & \\colhead{$m_{\\rm *,evap}$ ($\\sm$)\\tablenotemark{c}} } \\startdata 1 & 0.5 & 2.5 & 45.3 & 50.4 & 137\\tablenotemark{d}\\\\ \\hline 1 & 0.75 & 2.5 & 37 & 41 & 137\\\\ 1 & 0.25 & 2.5 & 68 & 81 & 143\\\\ 1 & 0.125 & 2.5 & 106 & 170 & 173\\\\ 1 & 0.0626 & 2.5 & 182 & 330\\tablenotemark{e} & 256\\\\ \\hline 1 & 0.5 & 5.0 & 35 & 38 & 120\\\\ 1 & 0.25 & 5.0 & 53.0 & 61 & 125\\\\ \\hline 0.5 & 0.5 & 2.5 & 23.0 & 24.5 & 57\\\\ \\hline 2.0 & 0.5 & 2.5 & 85 & 87 & 321\\\\ \\enddata \\tablenotetext{a}{\\footnotesize Mass scale of \\ion{H}{2} region polar breakout.} \\tablenotetext{b}{\\footnotesize Mass scale of \\ion{H}{2} region near-equatorial breakout.} \\tablenotetext{c}{\\footnotesize Mass scale of disk photoevaporation limited accretion.} \\tablenotetext{d}{\\footnotesize Fiducial model.} \\tablenotetext{e}{\\footnotesize This mass is greater than $m_{\\rm *,evap}$ in this case because it is calculated without allowing for a reduction in $\\dot{m}_*$ during the evolution due to polar \\ion{H}{2} region breakout.} \\end{deluxetable}     One may ask how the feedback mechanisms we have considered relate to those that operate in contemporary massive star formation. We note that the maximum mass attained in our fiducial model of PopIII.1 star formation is very similar to that inferred observationally in local massive star clusters (e.g. Figer 2005). However, after decades of study, it remains unclear whether the maximum mass of stars forming today is set by feedback or instabilities in very massive stars (Larson \\& Starrfield 1971).  Here, we have argued that the maximum mass of primordial stars is set by feedback. The primary differences in the feedback processes then and now are:   (1) Dust.  In contemporary star-forming regions, dust destroys \\lal\\ photons, eliminating them as a significant pressure. On the other hand, the dust couples the pressure of the UV continuum radiation to the gas very effectively, and it remains to be determined whether this limits the final mass of the star; e.g., Yorke \\& Sonnhalter (2002) find that it does, whereas Krumholz et al. (2005a) have not found evidence that it does.  Dust also affects the evolution of \\ion{H}{2} regions, absorbing a significant fraction of the ionizing photons in dense \\ion{H}{2} regions (Spitzer 1978), thereby reducing the \\ion{H}{2} feedback discussed in \\S 5 and the photoevaporation in \\S 7.  (2) Magnetic fields. In contemporary protostars, magnetic fields drive powerful protostellar winds that drive away a significant fraction of the core out of which the star is forming (Matzner \\& McKee 2000). The cavities created by these winds allow radiation to escape from the vicinity of the protostar, significantly reducing the radiation pressure (Krumholz et al. 2005b).  Tan \\& Blackman (2004) considered the influence of such outflows on Pop III.1 cores, concluding the instantaneous efficiency of accretion could be reduced by a factor of about 2 from the no-feedback case in an isotropic core by the time the star reached 100$\\sm$. However, these outflows would not confine ionizing feedback from the star at these masses, so much of the gas that could be expelled by outflows would have already been disrupted by \\ion{H}{2} region breakout. We conclude that outflows would have a relatively minor effect on the results presented here, and that the masses of the first stars are mostly influenced by radiative feedback. See Tan \\& McKee (2008) for further discussion.  (3) Stellar temperatures and luminosities. Primordial stars were significantly hotter than contemporary stars, resulting in significantly greater ionizing luminosities. In addition, the accretion rates of primordial massive stars are much greater, at least initially, than those of contemporary massive stars (McKee \\& Tan 2002; 2003). Future calculations will show whether feedback can be as effective in setting the maximum mass of contemporary stars as we have argued that it is for primordial stars."
    },
    "0711/0711.3144_arXiv.txt": {
        "abstract": "{} { We investigate potential models that could explain why multiple proto-stellar systems predominantly show single jets. During their formation, stars most frequently produce energetic outflows and jets. However, binary jets have only been observed in a very small number of systems.} { We model numerically 3D binary jets for various outflow parameters. We also model the propagation of jets from a specific source, namely L1551 IRS 5, known to have two jets, using recent observations as constraints for simulations with a new MHD code. We examine their morphology and dynamics, and produce synthetic emission maps.  } { We find that the two jets interfere up to the stage where one of them is almost destroyed or engulfed into the second one. We are able to reproduce some of the observational features of L1551 such as the bending of the secondary jet.  } { % While the effects of orbital motion are negligible over the jets dynamical timeline, their interaction has significant impact on their morphology. If the jets are not strictly parallel, as in most observed cases, we show that the magnetic field can help the collimation and refocusing of both of the two jets.  } ",
        "introduction": "Well-collimated jets occur across a broad range of mass scales from young stellar objects to active galactic nuclei. In young stellar objects (YSOs), jets and outflows are more easily observed at optical, infrared and UV wavelengths and the properties are relatively well known \\citep{2001ARA&A..39..403R,2007prpl.conf..231R}. YSO jets are believed to be launched by means of a magnetic field which is anchored or frozen into a circumstellar disk, and pinched and twisted by the disks rotation. Once the pinched-in field reaches a critical angle, fast-rotating gas from the disk may be loaded onto the field lines and thus launched into the parent cloud in an outflow \\citep{1982MNRAS.199..883B}. The most accepted schools of thought diverge on whether the wind is launched primarily from the inner disk region \\citep[disk wind, ][]{1997A&A...319..340F,2000prpl.conf..759K} or the X-annulus where the young star's magnetosphere interacts with the disk (X-wind, \\citep{2000prpl.conf..789S}). However in either case the jet is believed to brake the rotation of the disk and is thus essential to the star formation process, as it allows the star to accrete up to its final mass.  In all such models a single star with its associated circumstellar disk produces a bipolar outflow or jet. Estimates derived from studies of stellar populations, however, show that large numbers of binary and multiple star systems are expected \\citep{2002ApJ...581..654P,1995ApJ...443..625S,1993AJ....106.2005G,1991A&A...248..485D}. Do such systems present associated disks and jets ? It has been argued that the disk configuration in a binary system may either be in circumstellar disks or circumbinary disks depending on the star separation distance \\citep{2000prpl.conf..841H}. On the other hand, it is well established that multiple sources can be the source of HH objects, e.g. T Tau, IRAS 04325-1419, Z Cma, Sz 68, SR 24, S Cr A, AS 353 \\citep{1993ApJ...408L..49R}. Now, a multiple system may produce a single jet or outflow or a set of multiple outflows. Although there should be ``many'' visible binary jets, only a few binary jets from binary protostars have been observed. The examples known to the authors are compiled in Table \\ref{BinaryOutflowTable}. The frequency of existing binary jets is low compared to the large number of protostellar binary sources.    \\begin{table*} \\begin{center} \\caption{Evidence of Binary Jets and Outflows?} \\label{BinaryOutflowTable} \\begin{tabular}{l l l l l} \\hline \\hline Outflow (Source) & Natal Cloud & RA (J2000) & Dec (J2000) & Reference \\\\ \\hline HH154 (L1551 IRS 5) & Tau-Aur & 04 31 34.20 & +18 08 04.8 & \\citet{2005ApJ...L}\\\\ HH1-2 - HH144 & Orion & 05 36 22.85 & -06 46 06.6 & \\citet{1993ApJ...408L..49R} \\\\ HH111 - HH121 & Orion & 05 51 46.07 & +02 48 30.6 & \\citet{1994AA...289L..19G} \\\\ L723 (IRAS 19156+1906) & Cep & 19 17 53.16 & +19 12 16.6 & \\citet{2004RMxAC..21..100A} \\\\ HH288 (IRAS 00342+6347) & Cep& 00 37 11.07 & +64 03 59.8 & \\citet{2001AA...375.1018G} \\\\ HH377 (IRAS 23011+6126) & Cep E & 23 03 13.9 & +61 42 21 & \\citet{1997ApJ...474..749L} \\\\ IRAS 16293-2422 & $\\rho$ Oph E & 16 32 22.8 & -24 28 33 D & \\citet{2001ApJ...547..899H} \\\\ IRAS 20050+2720 & Cyg Rift & 20 07 06.7 & +27 28 53 & \\citet{1995ApJ...445L..51B} \\\\ \\hline \\end{tabular} \\end{center} \\end{table*}      Observations have also shown that multiple or quadrupolar jets may occur. \\citet{1993ApJ...408L..49R} discovered a second flow HH144 from the same source as HH1-2. \\citet{1994A&A...289L..19G} observed a second flow (HH121) from the same source as the well-known HH111 outflow. In both cases there are large angles of separation between the two jets. This can be explained if the disks are not coplanar \\citep{1994ARA&A..32..465M}. However, observations of apparent double jets can also be explained by other means: \\citet{1990ApJ...357..524A} imaged the outflow in L723 and found a distinct multi-polar morphology. They concluded that the lobes of the jets were in fact shell walls of two cavities swept clear by a single bipolar outflow. \\citet{1991ApJ...376..615A} discovered a double radio source at the centre of the outflow structure and this led to a reappraisal of the of the double outflow theory \\citep{1996ApJ...473L.123A,1997ApJ...489..734G,1998ASPC..132..303A,1998ApJ...504..334H,1999A&A...346..233P,2002ApJ...575..337S,2003RMxAC..15..135E,2004RMxAC..21..100A}. Other example of multiple molecular outflows include IRAS 16293-2422, IRAS 20050+2720.   In order to explore these possibilities, and to investigate about the nature of jets from multiple systems, we perform numerical simulations of binary jets with typical physical quantities. Our aim is to demonstrate the effects of their relative sizes and speeds on their interaction, in order to test if the jet evolution can lead to situations compatible with the few binary jets observed, and with the corresponding observed properties.  As a starting point, we focus on jets emanated independently by two sources that launch the two outflows with a small angle of separation. A suitable test case for this scenario is the the binary jet HH 154, emanating from the source L1551 IRS 5.      Numerical models of optical jets have evolved over the past twenty years, from the original models by \\citet{1981ApJ...247...52N} of the \\citet{1974MNRAS.169..395B} twin-exhaust model, to later models by \\citet{1984ApJ...276..560H,1987ApJ...316..323H,1988ApJ...326..323R} which were 1.5D bow shock models which explicitly tried to model observed features, to axisymmetric and 3D geometries and including such physics as atomic radiative cooling \\citep{1990ApJ...360..370B,1998AJ....116.2943R,2000ApJ...540..192S,2000A&A...364..763R,2001A&A...367..959R}.  In our case we model a fully three-dimensional binary jet system using ATLAS, a new shock-capturing, multi-dimensional, constrained transport, adaptive-grid, directionally unsplit, higher-order Godunov astrophysical MHD code, ATLAS and examine its morphology and propagation dynamics.  The paper is structured as follows: in Section \\ref{Observation_Section} we describe the observed properties of the binary jet L1551 IRS 5, that constitutes our test case. In Section \\ref{Simulation_And_Code} we describe the methodology, the set of equations used, the microphysics included, the numerical code and the initial and boundary conditions on the computational grid. In Sect. \\ref{Results} we present the results of the simulation, and finally in Sect. 5 we discuss our results and the insights brought forward by our model. ",
        "conclusions": "\\label{Discussion and Summary}   The L1551 IRS 5 jet pair can be seen as a pathological object unlike any other observed at such close range. In fact many effects, like binary source orbiting, peculiar environment, density parameters, magnetic fields play a role in sculpting this most unusual object. To attempt an explanation of the observed phenomena that can take into account all the above mentioned elements, a full 3D numerical treatment is necessary. In this paper we have presented 3D HD and MHD simulations of interacting binary jets that may help explain the L 1551-IRS5 case. We have used the code ATLAS including radiative cooling and non-equilibrium ionisation treatment. The main results can be summarised as follows:  \\begin{itemize} \\item We can reproduce the kink and bending of the secondary jet and the knotty structure of the two components as observed for L1551 IRS 5. \\item Over a long time-scale, we see precession induced by the orbital motion of the source. On the short lifetime of the jets from L1551, this effect is negligible. \\item If the jets are not strictly parallel, as in most observed cases, we show that the magnetic field can help the collimation and refocusing of both the jets. In fact the toroidal field affects the motion of the southern jet; and this magnetically driven change in direction together with the interaction of the bow shock of the fast jet with the beam of the slow jet contribute to the distinctive morphology. \\item We have produced emission maps in the [OI] lines which can be compared with observations. We show that the kink structure of interacting jets is still apparent in the synthetic observations. \\end{itemize}  Some aspects however remain to be clarified. For example, one of them is the nature of the the source of the X-ray emission. \\citet{2003ApJ...584..843B} find a source of X-rays in L1551 IRS 5 which they attribute to either fast shocks or reflected x-rays scattered out through the outflow cavity. The fast shocks may be caused by jet collimation which could be magnetic in nature. In our model there is a peak in emission (Knot Z in Figure \\ref{fig:4-1}) which appears to come from the colliding winds. The Bally model suggests that the X-rays come from a moving source at the base of the jets. If the density ratio is $n_{ambient}/n_{jet}=10$ then the shock speed will be too slow to achieve X-ray emission temperatures. If on the other hand the density ratio is lower i.e. 0.1 the jet is moving into a less dense region swept out by the slower jet the shock velocities will be higher. To get up to the observed shock velocity the density contrast would need to be much lower than observed e.g. 0.01 would give a velocity in the range required. We also could assume that it is caused by magnetic reconnection in the interval between the two jets where the field is compressed by the pair of bow shocks.  One other interesting aspect to be discussed is the implications for jet launching of the binarity of the source. The binary provides a vastly different natal environment for a protostar and its attendant jets and outflows than that envisaged in a single star formation. We identify three main differences for jet launching for binaries, orbital motion, circumbinary material and interaction with a jet. In this work we showed that the impact upon the jet beam of the bow shock of its neighbour is large enough to be observable as a knot of emission moving along along the jet. The orbital motion is of too long a period to affect the jet propagation. The role of the circumbinary material may be to twist up the magnetic field into the toroidal shape which appears in the polarized light. In conclusion, our simulations have demonstrated to be able to reproduce the observed signatures in this source and to give powerful constraints on the physical effects at play. A larger number of observed cases has however to be modeled with our code in order to give a complete view of binary jet production."
    },
    "0711/0711.4078_arXiv.txt": {
        "abstract": "We study the structure and evolution of ``quasistars,\" accreting black holes embedded within massive hydrostatic gaseous envelopes. These configurations may model the early growth of supermassive black hole seeds. The accretion rate onto the black hole adjusts so that the luminosity carried by the convective envelope equals the Eddington limit for the total mass, $M_* + M_{\\rm BH} \\approx M_*$. This greatly exceeds the Eddington limit for the black hole mass alone, leading to rapid growth of the black hole. We use analytic models and numerical stellar structure calculations to study the structure and evolution of quasistars. We show that the photospheric temperature of the envelope scales as $T_{\\rm ph} \\propto M_{\\rm BH}^{-2/5}M_*^{7/20} $, and decreases with time while the black hole mass increases. Once $T_{\\rm ph} < 10^{4}$~K, the photospheric opacity drops precipitously and $T_{\\rm ph}$ hits a limiting value, analogous to the Hayashi track for red giants and protostars, below which no hydrostatic solution for the convective envelope exists. For metal-free (Population~III) opacities this limiting temperature is approximately 4000 K. After a quasistar reaches this limiting temperature, it is rapidly dispersed by radiation pressure. We find that black hole seeds with masses between $10^3 \\ M_\\odot$ and $10^4 \\ M_\\odot$ could form via this mechanism in less than a few Myr. ",
        "introduction": "The formation mechanism for supermassive black holes in galactic nuclei remains unknown. Variations on most of the formation channels identified by \\cite{begelman78} --- which include instabilities in clusters of stars or stellar remnants, and the collapse of supermassive stars or massive discs --- are still under consideration today \\citep{umemura93,freitag06,shibata02,lodato06}. A more recent idea holds that supermassive black holes result from sustained accretion onto, or mergers of, the remnants of Population~III (i.e. metal-free) stars \\citep{volonteri03}, at least some of which seem likely to be massive, short-lived progenitors of stellar mass black holes \\citep{carr84,bromm99,abel02,heger03,tumlinson04}.  No current observation directly constrains the different classes of model for black hole formation. However, there are clues. The existence of massive black holes in quasars at redshifts $z > 6$ strongly suggests that these black holes, at least, started forming at redshifts high enough ($z \\sim 20$) to predate extensive metal enrichment by the first stars. Motivated by this, we develop in this paper a model proposed by \\cite{Be06} (BVR), in which supermassive black holes form, not from Pop~III stars themselves, but rather from the evolution of a new class of Pop~III objects that might form in metal-free haloes too massive to yield individual stars. In outline, the BVR model envisages a three-stage process for black hole formation. First, gas in metal-free haloes with a virial temperature above $T \\simeq 10^4 \\ {\\rm K}$ flows toward the center of the potential as a result of gravitational instabilities, forming a massive, pressure-supported central object. Nuclear reactions may start, but the very high infall rate continues to compress and heat the core, precluding formation of an ordinary star. Eventually, when the core temperature attains $T \\sim 5 \\times 10^8 \\ {\\rm K}$, neutrino losses result in a catastrophic collapse of the core to a black hole. We dub the resulting structure --- comprising an initially low mass black hole embedded within a massive, radiation-pressure supported envelope --- a {\\em quasistar}. Initially, the black hole is much less massive than the envelope. Over time, the black hole grows at the expense of the envelope, until finally the growing luminosity succeeds in unbinding the envelope and the seed black hole is unveiled. The key feature of this scenario is that while the black hole is embedded within the envelope, its growth is limited by the Eddington limit for the {\\em whole quasistar}, rather than that appropriate for the black hole mass itself. Very rapid growth can then occur at early times, when the envelope mass greatly exceeds the black hole mass.  The structure of quasistars, illustrated in Fig.~\\ref{fig_cartoon}, resembles that of more familiar objects. The outer regions have qualitative similarities to Thorne-\\.Zytkow (1977) objects, although in our case the luminosity derives entirely from black hole accretion. The accretion physics is related to the ``hypercritical\" regime of accretion \\citep{begelman79,blondin86}, which may be observed in X-ray binaries such as SS433 \\citep{begelman06}. However, these are very rough analogues. The large difference between Pop~III opacity and that of present-day, metal-enriched gas, together with the dominant role of radiation pressure and the presence of ongoing accretion in quasistars, result in a unique quasistar structure.  \\begin{figure} \\psfig{figure=fig1.eps,width=0.48\\textwidth} \\caption{Schematic illustration of the quasistar structure that we consider in this paper. A seed black hole of mass $M_{BH}$ accretes gas from a massive, radiation pressure-supported envelope at a rate set by the conditions outside the Bondi radius. The luminosity liberated by the accretion process is transported convectively in the inner regions of the envelope, with a transition to a radiative zone once convection becomes inefficient. In this paper we consider isolated, spherically symmetric models of quasistars, but the more physical situation would also include partial rotational support leading to flattening of the quasistar, and ongoing disc accretion at a fraction of a Solar mass per year.} \\label{fig_cartoon} \\end{figure}  In this paper, we study the structure and evolution of quasistars. Using Pop~III opacities, we construct simplified analytic models which we compare with numerical integrations of the stellar structure equations. Our model quasistars are assumed to be spherically symmetric and in hydrostatic equilibrium. Real quasistars, if they exist, may not obey these strictures, and thus our results cannot be definitive. Rather, our calculations are intended to provide a first estimate of the maximum black hole mass that this mechanism can yield, together with a determination of the photospheric temperature and luminosity of quasistars that can be used to assess their detectability with next-generation observatories such as the {\\em James Webb Space Telescope}. ",
        "conclusions": "In this paper, we have studied the structure and evolution of quasistars, rapidly accreting black holes embedded within massive gas envelopes. We find that, for any black hole mass, there is a minimum photospheric temperature below which rapid dissolution of the envelope is inevitable. Both analytic and numerical models, computed using Pop~III opacities, suggest that this minimum temperature is around 4000-5000~K. If quasistars are implicated in the formation of seeds for supermassive black holes in pre-galactic haloes, as suggested by \\cite{Be06}, this floor temperature implies that the most luminous quasistars would emit most strongly in the rest-frame near-IR. At typical redshifts of $z \\simeq 10$, the observed spectrum would peak at $\\lambda \\sim 10 \\mu {\\rm m}$. We defer discussion of the possible cosmological density of such sources, and hence their observability with future facilities such as the {\\em James Webb Space Telescope}, to a subsequent paper.  We have also studied the evolution of quasistars in simple scenarios for the mass growth of their envelopes. Generically, as the black hole mass grows the photospheric temperature falls, until eventually the limit imposed by the behavior of the Pop~III opacity is reached.  Unveiling of the black hole appears inevitable long before it succeeds in accreting much of the envelope. Both analytic and numerical estimates suggest that seed black holes with masses between a few $10^3 \\ M_\\odot$ and $10^4 \\ M_\\odot$ are plausible outcomes of this scenario. The efficiency of the black hole accretion process that powers quasistars (parameterized here by $\\alpha$) is unknown --- if the efficiency is low (due, for example, to polar outflows that couple poorly to the envelope) then larger seed black holes are possible.  Although our numerical integrations of quasistar structure confirm many aspects of the analytic model, they also reveal complex behavior in the radiative zone immediately beneath the photosphere. Narrow regions in which the flux is locally super-Eddington develop for temperatures significantly in excess of the minimum temperature at which the entire radiative zone becomes super-Eddington. In our hydrostatic models, these zones are characterized by density inversions. Interpreting this structure is tricky, since such density inversions may well be subject to one- or two-dimensional instabilities whose ultimate resolution is unknown. Hydrodynamic studies will be needed to determine the extent of mass loss that may occur at temperatures above the theoretical floor value.  Finally, we note that here we have focused on black holes embedded within truly primordial gas at high redshift. {\\em Small} amounts of metal pollution would act to increase the opacity in the radiative zone, altering the minimum temperature and possibly increasing the likelihood of mass loss due to the formation of dust. Changes to the interior structure would be smaller. In particular, the density at the base of the radiative zone is so low ($\\rho \\sim 10^{-9} \\ {\\rm g cm}^{-3}$) that electron scattering will continue to dominate the opacity for $T > 10^5 \\ {\\rm K}$, even in the presence of pollution. Of greater concern is the fact that metal enriched gas in the halo {\\em outside} the quasistar would be more susceptible to fragmentation and star formation. Too much star formation would reduce the rate of mass accretion onto the quasistar below the values ($\\sim 0.1 \\ M_\\odot {\\rm yr}^{-1}$ or higher) that we have assumed, resulting in smaller final black hole masses. It is also possible that generically similar structures could form at lower redshift whenever small black holes encounter very high rates of gas inflow. Structures similar to those we have described could allow stellar remnants at the center of merging galaxies to grow significantly via a transient quasistar stage."
    },
    "0711/0711.1970_arXiv.txt": {
        "abstract": "We report on accretion- and outflow-related X-rays from T Tauri stars, based on results from the \"XMM-Newton Extended Survey of the Taurus Molecular Cloud.\" X-rays potentially form in shocks of accretion streams near the stellar surface, although we hypothesize that direct interactions between the streams and magnetic coronae may occur as well. We report on the discovery of a \"soft excess\" in accreting  T Tauri stars supporting these scenarios. We further discuss a new type of X-ray source in jet-driving T Tauri stars. It shows a strongly absorbed coronal component and a very soft, weakly absorbed component probably related to shocks in microjets. The excessive coronal absorption points to dust-depletion in the accretion streams. ",
        "introduction": "Classical and weak-lined T Tauri stars (CTTS and WTTS) are pre-main sequence stars showing vigorous X-ray emission with X-ray luminosities ($L_{\\rm X}$) near the empirical saturation limit found for main-sequence (MS) stars, $L_{\\rm X}/L_{\\rm bol} \\approx 10^{-3.5}$. Consequently, X-ray emission from T Tauri stars has been attributed to solar-like coronal activity. However, both accretion and outflows/jets may  contribute to or alter  X-ray production in the magnetic environment of strongly accreting young stars. X-rays thus provide an important diagnostic for the detection and study of accretion and outflow activity.   Support for X-ray {\\it suppression} in CTTS has been reported from X-ray photometry. CTTS are, on average, less X-ray luminous than WTTS (e.g., \\cite[Strom \\& Strom 1994]{strom94}, \\cite[Neuh\\\"auser \\etal\\ 1995]{neuhauser95}). This finding has  been partly supported by recent, deep surveys of the Orion Nebula Cluster (\\cite[Getman \\etal\\ 2005]{getman05}, \\cite[Preibisch \\etal\\ 2005]{preibisch05}).  On the other hand, X-rays may be {\\it generated} by surface accretion shocks (\\cite[Lamzin 1999]{lamzin99}). If gas collides with the stellar photosphere in free fall, shocks heat it to a few million K. Given the appreciable accretion rates, high shock densities of order $10^{12}- 10^{14}$~cm$^{-3}$ are expected, as first reported for the CTTS TW Hya (\\cite[Kastner \\etal\\ 2002]{kastner02}). X-rays may similarly be produced in shocks forming in jets (\\cite[Raga \\etal\\ 2002]{raga02}).  The {\\it XMM-Newton Extended Survey of the Taurus Molecular Cloud (XEST)} (\\cite[G\\\"udel \\etal\\ 2007a]{guedel07a}) has provided new insights into these issues. XEST covers the most populated  $\\approx$5~sq. deg of the Taurus star-forming region. The average on-axis detection limit is $\\approx 10^{28}$~erg~s$^{-1}$ for lightly absorbed objects, sufficient to detect about half of the observed brown dwarfs (\\cite[Grosso \\etal\\ 2007a]{grosso07a}). We discuss X-ray results relevant for  the star-disk interface in which accretion occurs and where (parts of) the jets may be accelerated. ",
        "conclusions": ""
    },
    "0711/0711.2142_arXiv.txt": {
        "abstract": "The hot Jupiter HD 189733b was observed during its primary transit using the Infrared Array Camera on the \\spitzer\\ \\emph{Space Telescope}. The transit depths were measured simultaneously at 3.6 and 5.8 $\\mu$m. Our analysis yields values of $2.356 \\pm 0.019$\\% and $2.436 \\pm 0.020$\\% at 3.6 and 5.8 $\\mu$m respectively, for a uniform source. We estimated the contribution of the limb-darkening and star-spot effects on the final results. We concluded that although the limb darkening increases by $\\sim $ 0.02-0.03 \\% the transit depths, and the differential effects between the two IRAC bands is even smaller, 0.01 \\%. Furthermore, the host star is known to be an active spotted K star with observed photometric modulation. If we adopt an extreme model of 20 \\% coverage with spots 1000K cooler of the star surface, it will make the observed transits shallower by 0.19 and 0.18 \\%. The difference between the two bands will be only of 0.01 \\%, in the opposite direction to the limb darkening correction. If the transit depth is affected by limb darkening and spots, the differential effects between the  3.6 and 5.8 $\\mu$m  bands are very small. The differential transit depths at 3.6 and 5.8 $\\mu$m and the recent one published by Knutson et al.( 2007) at 8 \\microns\\ are in agreement with the presence of water vapour in the upper atmosphere  of the planet. This is the companion paper to Tinetti et al. (2007b), where the detailed atmosphere models are presented. ",
        "introduction": "\\label{sec:intro}  Over 240 planets are now known to orbit stars different from our Sun (Schneider 2007; Butler et al. 2007). This number is due to increase exponentially in the near future thanks to space-missions devoted to the detection of exoplanets and the improved capabilities of the ground-based telescopes. Among the exoplanets discovered so far, the best known class of planetary bodies are giant planets (EGPs) orbiting very close-in (hot-Jupiters). In particular, hot-Jupiters that transit their parent stars offer a unique opportunity to estimate directly key physical properties of their atmospheres (Brown, 2001). In particular, the use of the primary transit (when the planet passes in front of its parent star) and transmission spectroscopy to probe the upper layers of the transiting EGPs, has been particularly successful in the UV and visible ranges (Charbonneau et al. 2002; Richardson et al. 2003, 2003a; Deming et al. 2005a; Vidal-Madjar et al. 2003, 2004; Knutson et al. 2007; Ballester, Sing, \\& Herbert 2007; Ben-Jaffel, 2007) and in the thermal IR (Richardson et al. 2006; Knutson et al. 2007).  The hot-Jupiter HD 189733b (Bouchy et al. 2005) has a mass of $M_p=1.15 \\pm 0.04$~$M_{\\rm Jup}$, a radius of $R_p=1.26 \\pm 0.03$~$R_{\\rm Jup}$, and orbits a main sequence K-type star at a distance of 0.0312 AU. This exoplanet is orbiting the brightest and closest star discovered so far, making it one of the prime targets for observations (Bakos et al. 2006; Winn et al. 2007;  Deming et al. 2006; Grillmair et al. 2007; Knutson et al. 2007).  Tinetti et al. (2007a) have modeled the infrared transmission spectrum of the planet HD 189733b during the primary transit and have shown that \\spitzer observations are well suited to probe the atmospheric composition and, in particular, constrain the abundances of water vapor and CO. Here we analyze the observations of HD 189733b with the Infrared Array Camera (IRAC Fazio et al., 2004) on board the \\spitzer Space Telescope in two bands centered at 3.6 and 5.8 $\\mu$m. We report the data reduction, and discuss our results in light of the theoretical predictions and the recent observation at 8 $\\mu$m (Knutson et al. 2007). ",
        "conclusions": "\\label{sec:conclusions}  We estimated accurately the radius of the extrasolar planet HD\\,189733b using its primary transit, at 3.6 and 5.8~\\microns. The small error bars are the result of a high signal-to-noise ratio and weak influence of the limb-darkening effect in the IR. The planetary radius appears $(1.6 \\pm 0.5)$\\% larger at 5.8~\\microns\\ than at 3.6~\\microns. The observations match the predictions by Tinetti et al. (2007a).  Detailed interpretation of these results (Tinetti et al. 2007b) combined with the 8 ~\\microns\\ observations (Knutson et al. 2007) confirm that water vapor is the most likely explanation for the observed photometric signature in the IR. The comparison with the visible is more complex because of the possibly important role of star spots.  Our observations show that the combination of the primary transit technique and comparative band photometry at multiple wavelengths is an excellent tool to probe the atmospheric constituents of transiting extrasolar planets. Similar studies and observations should be considered for other targets, especially with the foreseen {\\it James Webb Space Telescope}, which could observe more distant and smaller transiting planets."
    },
    "0711/0711.0001_arXiv.txt": {
        "abstract": "Star-forming galaxies constitute the majority of galaxies with stellar masses $\\gtrsim10^{10}\\,h^{-2} M_\\Sun$ at $z\\sim2$.  It is thus critical to understand their origins, evolution, and connection to the underlying dark matter distribution.  To this end, we identify the dark matter halos (including subhalos) that are likely to contain star-forming galaxies at $z\\sim2$ (z2SFGs) within a large dissipationless cosmological simulation and then use halo merger histories to follow the evolution of z2SFG descendants to $z\\sim1$ and $z\\sim0$.  The evolved halos at these epochs are then confronted with an array of observational data in order to uncover the likely descendants of z2SFGs.  Though the evolved halos have clustering strengths comparable to red galaxies at $z\\sim1$ and $z\\sim0$, we find that the bulk of z2SFGs {\\it do not} evolve into red galaxies, at either epoch. This conclusion is based primarily on the fact that the space density of z2SFGs is much higher than that of lower redshift red galaxies, even when accounting for the merging of z2SFG descendants, which decreases the number density of z2SFG descendants by at most a factor of two by $z\\sim0$.  Of the $\\sim50$\\% of z2SFGs that survive to $z\\sim0$, $\\sim70$\\% reside at the center of $z\\sim0$ dark matter halos with $M>10^{12} \\,h^{-1} M_\\Sun$.  Halo occupation modeling of $z\\sim0$ galaxies suggests that such halos are occupied by galaxies with $M_r\\lesssim-20.5$, implying that these z2SFGs evolve into ``typical'' $\\sim L^\\ast$ galaxies today, including our own Galaxy.  The remaining $\\sim30$\\% become satellite galaxies by $z\\sim0$, and comparison to halo occupation modeling suggests that they are rather faint, with $M_r\\lesssim-19.5$.  These conclusions are at least a partial departure from previous work due primarily to the increased accuracy of observational data at $z\\lesssim 1$, and to higher resolution $N$-body simulations that explicitly follow the evolution of dark matter subhalos, whose observational counterparts are likely satellite galaxies.  These conclusions are qualitatively generic in the sense that any halo mass-selected sample of galaxies at one epoch will evolve into a more complex and heterogeneous sample of galaxies at a later epoch. This heterogeneity is driven largely by the fact that some galaxies will continue to accrete matter and form stars throughout their evolution, while others will become satellites and thus have their growth suppressed relative to galaxies in the field. ",
        "introduction": "\\label{s:intro}  The past decade has witnessed an enormous increase in our knowledge of the high-redshift ($z>2$) Universe.  The unique rest-frame UV spectral signatures of star-forming galaxies have allowed observers to select galaxies at $z\\gtrsim2$ based solely on optical photometry with very high efficiency \\citep[e.g.][]{Steidel03, Steidel04}.  Extensive multi-wavelength follow-up of these objects, including optical and near-IR spectroscopy, and imaging at wavelengths from X-ray through radio, has yielded estimates of rest-frame UV luminosities and colors \\citep{Steidel99, Adelberger00, Reddy07}, stellar masses \\citep{Shapley05, Erb06b}, star-formation rates \\citep{Erb06a, Reddy06a}, chemical abundances \\citep{Pettini01, Erb06c}, and clustering strengths \\citep{Adelberger03, Adelberger05, Giavalisco01, Ouchi05, Lee06, Hamana06, Quadri07} of a large, well-defined population of galaxies that dominate the star-formation rate density \\citep{Reddy07} at an epoch when the Universe was only a few billion years old. Complementary techniques tuned to the rest-frame optical properties of galaxies at similar epochs have identified objects with typically redder colors and higher mass-to-light (M/L) ratios, but the corresponding spectroscopic confirmation has only been obtained for the brightest, rare objects \\citep{Franx03, Daddi04, vanDokkum06}.  At the same time, models connecting galaxies to the underlying dark matter distribution have been put forth which either focus on quantifying the statistical relation between galaxies and dark matter halos \\citep[e.g.][]{Berlind02, Bullock02, Yan03, Seljak00, Scoccimarro01, Zheng04, Kravtsov04, Conroy06a}, or predicting the connection based on physical principles \\citep[e.g.][]{White91, Somerville99, Cole00, Hatton03, Springel01, Croton06, Bower06}. Essential to these latter efforts has been the vast increase in size and resolution of $N$-body simulations, which has only recently allowed for the construction of detailed merger trees that follow not only the evolution of distinct halos, i.e. those halos not contained within any larger halo, but also subhalos --- the halos contained within the virial radii of distinct halos --- whose observational counterparts are likely satellite galaxies \\citep[e.g.][]{Springel01}.  Throughout we refer to both distinct halos and subhalos generically as halos.  Despite impressive advances in our understanding both of the properties of halos themselves and of the relation between galaxies and halos, fundamental questions remain.  The origin and fate of z2SFGs is one such question.  Herein we define z2SFGs as those galaxies selected to lie at $z\\sim2$ based on optical $U_nG\\mathcal{R}$ photometry that are brighter than $\\mathcal{R}=25.5$, which corresponds to a rest-frame UV magnitude limit at this epoch.  Historically, UV-selected high-redshift galaxies have been described as either low-mass, merger-induced starbursts or more massive galaxies ``quiescently'' forming stars.  With early data, it was not possible to distinguish between these opposing ideas \\citep{Lowenthal97, Coles98, Mo99, Giavalisco01, Kolatt99, Somerville01, Wechsler01}.  However, recent modeling in conjunction with more recent data on the abundance and spatial clustering seems to favor the latter picture \\citep{Conroy06a, Adelberger05}. The distribution of stellar populations of these objects lends support to the massive, quiescent scenario as well \\citep{Shapley01, Erb06b}.  Once identified at high redshift, it is also important to determine what type(s) of galaxies z2SFGs evolve into at later epochs.  Most previous analyses have focused on the so-called Lyman-Break Galaxies, which are UV-selected star-forming galaxies at $z\\sim3$.  We argue in later sections that these galaxies are qualitatively similar to z2SFGs and thus here we make no distinction between the two. Some analyses focusing on the strong observed clustering of z2SFGs favored the scenario where z2SFGs evolved into the observed massive red galaxies at $z\\sim1$, and by $z\\sim0$ were at the centers of rich groups and clusters \\citep{Mo96b, Baugh98, Governato98, Governato01, Blaizot04, Adelberger05}, though other models incorporating the high observed comoving number density of z2SFGs suggested a more nuanced history \\citep{Moustakas02}.  Subsequent to these efforts, there has been a vast increase in our understanding of the number densities and clustering of galaxies as a function of luminosity and color both at $z\\sim1$ \\citep[e.g.][]{Coil06b, Pollo06, Meneux06, Coil07} and $z\\sim0$ \\citep[e.g.][]{Zehavi05}. Better constraints on the statistics of galaxy populations at lower redshift are crucial for a robust identification of the descendants of high-redshift galaxies. These recent data at $z\\leq 1$, in conjunction with new high-resolution simulations and dark-matter-halo merger trees, motivate a re-evaluation of the nature and subsequent evolution of z2SFGs.  In the present work we focus on the fate(s) of observed z2SFGs.  This redshift range is our focus here rather than $z\\geq 3$, as there has been much recent attention to this lower redshift window with several complementary galaxy selection techniques. Plus, a more extensive set of multi-wavelength imaging and spectroscopy exists for $z\\sim 2$ galaxies, allowing the more robust estimate of a number of physical properties such as stellar, gas, and dynamical masses, and extinction-corrected star-formation rates \\citep{Erb06b,Shapley05,ForsterSchreiber06, Reddy06a}, which are important to consider along with the inferred host dark-matter halo properties. Such a detailed comparison of halo and galaxy masses and formation histories is currently not possible at higher redshifts.  Under the assumption of a tight correlation between galaxy light and halo mass, the observed clustering of z2SFGs puts a constraint on the minimum halo mass hosting such galaxies ($\\S$\\ref{s:p}).  Halo merger trees extracted from a cosmological $N$-body simulation are then used to follow the halos hosting z2SFGs to later epochs.  Comparing the clustering, satellite fraction, and number density of these evolved halos to observed galaxy populations at both $z\\sim1$ and $z\\sim0$ then provides constraints on the fates of z2SFGs ($\\S$\\ref{s:r}).  Throughout we assume a $\\Lambda$CDM cosmology with $(\\Omega_m,\\Omega_{\\Lambda},\\sigma_8) = (0.25,0.75,0.9)$, consistent with the {\\it WMAP1} data \\citep{Spergel03}.  Where applicable, we leave results in terms of $h$, the Hubble parameter in units of $100$ km s$^{-1}$ Mpc$^{-1}$.  In several sections we discuss how our results are affected by adopting the cosmological parameters favored by the {\\it WMAP3} data \\citep{Spergel07}, where the primary difference is a lower $\\sigma_8$. Halo masses are measured as the mass interior to a region with mean enclosed density equal to $200$ times the critical density.  A \\citet{Chabrier03} IMF is assumed when quoting stellar masses.  All magnitudes are quoted in the AB system; we omit the factor of $5{\\rm log} (h)$ when quoting absolute magnitudes for brevity. ",
        "conclusions": "\\label{s:disc}  In the preceding section we found that the observed abundance and clustering of z2SFGs implied that every halo at $z\\sim2$ above a mass of $10^{11.4}\\, h^{-1} \\, M_\\Sun$ contains one z2SFG.  By $z\\sim0$, the overall population of halos that once hosted z2SFGs, and survived to $z\\sim0$, has a comoving number density comparable to that of $\\sim L^\\ast$ galaxies ($M_r\\lesssim-20$; while $M_r^\\ast=-20.4$), a clustering strength comparable to luminous ($M_r\\lesssim-21$) red galaxies, and a satellite fraction comparable to $\\sim L^\\ast$ or fainter red galaxies.  We found that a subsample of z2SFG descendant halos where 60\\% of the satellites were removed resulted in much better agreement with the clustering strength, satellite fraction, and number density of $\\sim L^\\ast$ galaxies at $z\\sim0$.  Comparison to halo occupation modeling of observed galaxy samples revealed that the z2SFG descendants that are central galaxies most closely resemble observed galaxies brighter than $M_r\\sim-20.5$ while the satellite properties compare more favorably to galaxies brighter than $M_r\\sim-19.5$.  The average redshift at which z2SFG descendant satellites at $z\\sim 0$ were accreted onto their parent halo was $z\\approx0.5$, and the mass of the halo associated with these satellites when they were accreted was only 55\\% of the mass of distinct halos at $z\\sim0$, on average.  This indicates that the growth of subhalos (and the galaxies within them) was significantly less than that of distinct halos.  These results have straightforward implications regarding the connection between z2SFGs and their descendants at $z\\sim0$.  First, it is abundantly clear that z2SFGs do not evolve into any single sample of galaxies defined either according to a luminosity cut or a luminosity and color cut.  In particular, they do not evolve exclusively into massive red galaxies, either by $z\\sim1$ or $z\\sim0$. Rather, the subsequent evolution of z2SFGs is more nuanced.  Merger trees extracted from simulations indicate that the number density of z2SFG descendants at $z\\sim0$ is $\\sim50$\\% of the number density of z2SFGs because some descendants merge together between $z\\sim2$ and $z\\sim0$.  This is an upper bound on the evolution of the number density because halos in our simulations can merge or disrupt without the subsequent merger or disruption of the galaxy within it due to simulation resolution effects or possible baryonic effects.  Of the remaining descendants, $\\sim70$\\% are centrals and $\\sim30$\\% are satellites.  The central galaxies have properties similar to observed galaxies with $M_r\\lesssim-20.5$.  In particular, they reside predominantly in halos with mass $M\\gtrsim10^{12}\\, h^{-1} M_\\Sun$.  An anecdotal consequence is that our own Galaxy, with halo mass $\\sim 10^{12} \\, h^{-1} \\,M_\\Sun$ \\citep{Klypin02}, was once likely a z2SFG. This conclusion is in accord with various properties of the bulge of our Galaxy.  For example, the bulge has a stellar mass of $\\sim 10^{10}\\,M_\\Sun$ that is thought to have formed over a short period of time ($<1$ Gyr) roughly 10 Gyr ago \\citep{Zoccali03, Ferreras03}. These facts imply that the star-formation rate during the formation of the bulge was $>10\\, M_\\Sun$ yr$^{-1}$.  The epoch of formation of the bulge and its high star-formation rate imply that it would have likely been detected as a z2SFG \\citep[see also][who reached similar conclusions]{Pettini06}.  In contrast, the z2SFG descendants that become satellites by $z\\sim0$ are much more similar to a fainter sample of observed galaxies ($M_r\\lesssim-19.5$).  This is not surprising in light of the different histories of satellite compared to central galaxies as inferred from the histories of their dark matter halos.  Since satellites fell into their parent halo on average at $z\\approx0.5$ ($3.5\\,h^{-1}$ Gyr ago), their growth was significantly retarded relative to central galaxies.  Satellites grew less not only because they could not accrete new gas but also because the gas they possessed at infall was likely prohibited from cooling to form stars (or removed altogether) due to one or more quenching processes, such as gas strangulation, harassment, or ram-pressure stripping.  Not all satellites are required to be so faint; our results indicate that $\\sim40$\\% are brighter than $M_r\\sim-20.5$.  It is plausible that the $\\sim60$\\% required to have $-19.5\\lesssim M_r\\lesssim-20.5$ were the satellites accreted at the earliest epochs, because such satellites would have the least amount of growth between $z\\sim2$ and $z\\sim0$.  Simulated star-formation histories coupled with a stellar population synthesis code \\citep{Bruzual03} confirm that galaxies with properties similar to z2SFGs can easily evolve into galaxies with $M_r\\sim-19.5$ if their star-formation is truncated at $z\\sim0.5-1$. For example, a galaxy at $z\\sim2$ with stellar mass $10^{10}\\,h^{-2}\\, M_\\Sun$ and a star-formation rate of 10 $M_\\Sun$ yr$^{-1}$ that steadily declines to zero at $z\\sim1$ fades from $M_r=-21.5$ to $M_r=-19.2$ by $z\\sim0$.  Similar conclusions hold for the descendant halos at $z\\sim1$.  In particular, from analysis of the Millennium simulation merger trees we find that $\\sim25$\\% of z2SFG descendants have merged away by this time.  Of the remaining halos, $\\sim25$\\% are satellites and $\\sim75$\\% are central galaxies.  If half of the satellites are fainter than $\\sim L^\\ast$, then the remaining descendant halos that are satellites and centrals have a number density, clustering strength, and satellite fraction comparable to observed $\\sim L^\\ast$ galaxies at $z\\sim1$.  There has been some debate in the literature over whether or not the optical color selection technique used to select z2SFGs misses a significant population of redder (i.e. older and/or dustier) galaxies \\citep[e.g.][]{Franx03, Daddi04, vanDokkum06, Smail02}.  Under our assumption of a tight correlation between rest-frame UV luminosity and halo mass, we have demonstrated that at $z\\sim2$ roughly every halo with mass above $10^{11.4} \\,h^{-1}\\, M_\\Sun$ contains one z2SFG.  It would be rather surprising if each halo also contained another comparably massive galaxy not detectable with the optical color selection technique, if only by analogy with lower redshift, where the vast majority of $\\sim L^\\ast$ galaxies live alone in their dark matter halos \\citep[e.g.][]{Zheng07}.  Thus, while there may well be a population of massive ($M_\\ast\\gtrsim 10^{11} \\,h^{-2}\\,M_\\Sun$), red galaxies with a correspondingly low space density ($\\sim10^{-4}\\, h^3$ Mpc$^{-3}$), that is missed with the optical color selection technique \\citep{Franx03, Daddi04, vanDokkum06}, it is hard to imagine that this could be the case at lower galaxy masses.  Since our conclusions are most sensitive to lower mass galaxies, they are not influenced by this possible distinct population of more massive, less numerous, red galaxies.  There are two basic facts about dark matter halos and their evolution that, when combined, provide both a qualitative understanding of our results and more general insights into the evolution of populations of galaxies across time.  The first fact is that, to with a factor of $\\sim2$, the space density of a set of halos above a mass threshold does not change from high redshift to the present.  For example, of the halos with $M\\geq10^{11.4}\\, h^{-1}\\, M_\\Sun$ at $z\\sim2$, 50\\% survive to $z\\sim0$.  For halos more massive than $10^{11.9}\\, h^{-1}\\, M_\\Sun$ and $10^{12.4}\\, h^{-1}\\, M_\\Sun$ at $z\\sim2$, the fraction that survive to $z\\sim0$ increases to 60\\% and 70\\%, respectively.  The second fact is that, by and large, the ranking of halos by mass is preserved as the Universe evolves.  For example, the most massive halos at $z\\sim2$ evolve to the most massive halos today. If, as we have argued, the z2SFGs are an at least approximately halo mass selected sample, then these two basic facts make less surprising the conclusion that z2SFGs evolve largely into a roughly halo mass selected sample of abundant galaxies at $z\\sim0$.  As described above, it is primarily the uncertain evolution of satellite galaxies that complicates this simple description.  These facts can also shed light on the likely descendants of the massive red galaxies discussed above.  If such galaxies are in the most massive halos \\citep[as suggested by their high stellar mass and clustering strength;][]{Quadri07}, then it is these galaxies that will evolve, largely intact, to the centers of rich groups and clusters by $z\\sim0$.  Better constraints on the clustering and abundances of these observed massive red galaxies are required to make more detailed statements.  Finally, these basic facts can connect our discussion herein to the fates of star-forming galaxies at $z>2$.  Indeed, much of the previous high-redshift modeling efforts discussed in the Introduction focused on star-forming galaxies at $z\\sim3$, the so-called Lyman-Break Galaxies (LBGs).  If we apply our methods to match the observed clustering of LBGs \\citep[$r_0=4.0\\pm0.6\\,h^{-1}$ Mpc;][]{Adelberger05} with a set of halos in the Millennium simulation above some minimum mass, we find a best-fit $\\Mmin = 10^{11.1\\pm0.2}\\,h^{-1}\\,M_\\Sun$.  Analysis of the z2SFG halo merger trees indicates that every z2SFG halo contains one or more of these LBG halos in their history and approximately 70\\% of LBGs evolve into z2SFGs \\citep[in agreement with earlier work; e.g.][]{Adelberger05}. While the link between LBGs and z2SFGs appears strong, there are important subtleties one must keep in mind when comparing LBGs to z2SFGs.  In particular, the clustering of these two populations has been measured for all objects above a common apparent magnitude limit ($\\mathcal{R}=25.5$), and thus the LBGs, being more distant, are an intrinsically more luminous sample compared to the z2SFGs.  However, the luminosity functions of these two classes of galaxies are nearly the same \\citep{Reddy07}, suggesting that if the two samples were defined with respect to the same absolute luminosity limit, their evolutionary bond would be even stronger.  Nonetheless, the clear link between LBGs and z2SFGs implies that LBGs share the same varied fates as those we have outlined for z2SFGs.  The most significant uncertainty in our analysis is the clustering strength of z2SFGs and the subsequent merging of their descendants.  A more quantitative understanding of the fate of these z2SFGs will thus require a more accurate accounting of their clustering strengths, on the observational side, and a more sophisticated treatment of z2SFG mergers on the theoretical side.  In addition, clustering measurements on smaller scales ($\\lesssim1$Mpc) and over a range of rest-frame luminosities will allow a more detailed modeling of the connection between z2SFGs and halos at $z\\sim2$ than has been presented here and will provide a critical test of our assumption of a tight correlation between rest-frame UV luminosity and halo mass at $z\\sim2$.  With more accurate data over the interval $0\\lesssim z\\lesssim 2$, the goal will ultimately be to model the galaxy-halo connection as a function of galaxy properties such as color and luminosity, or star-formation rate and stellar mass, in narrow redshift intervals. The galaxy-halo connections at higher redshift can then be evolved into the connections at lower redshift with the aid of halo merger trees.  The merger trees are the key to this approach, as they put the static (fixed redshift) halo occupation models into motion.  Such an approach has only recently become feasible, due both to the vastly increased data at these epochs and to a converging set of cosmological parameters that removes significant uncertainties in the properties and histories of dark matter halos \\citep[see e.g.][]{Zheng07, MWhite07, Conroy07b}.  A simplified version of this approach has been explored in the present work.  A more detailed version will be pursued elsewhere."
    },
    "0711/0711.1627_arXiv.txt": {
        "abstract": "We announce the discovery of six $\\beta$~Cephei stars and many other variable stars in the young open cluster h Persei (NGC 869). The cluster seems to be very rich in variable B-type stars, similarly to its twin, $\\chi$~Persei (NGC\\,884). ",
        "introduction": "The young open clusters h and $\\chi$ Persei (NGC 869 and 884) form a well-known bright double cluster in Perseus. The clusters have similar age, distance and contain comparable number of stars. Recently, $\\chi$~Persei was the target of a large international campaign aimed at the discovery and subsequent asteroseismology of pulsating stars, in particular $\\beta$ Cephei stars \\cite{vienna,sophie}. Preliminary reduction led to the discovery of seven $\\beta$ Cephei stars in this cluster which makes it one of the richest in this kind of variable. It can be suspected that h Persei, the twin of $\\chi$~Persei, also contains a large number of early-type pulsating stars. For this reason, we reduced and analyzed CCD observations of h~Persei obtained during 8 observing nights in 2001 at Piszk\\'estet\\H{o} station of the Konkoly Observatory (Hungary). We present here the preliminary results of this study. ",
        "conclusions": ""
    },
    "0711/0711.1545.txt": {
        "abstract": " ",
        "introduction": "}  Binaries are a well-established feature of the transneptunian population; more than 50 systems are known (Noll et al.~2008).  The number of transneptunian binaries (TNBs) is most usefully expressed as a fraction of the population; this fundamental quantity is among the first to be determined by large surveys.  Any such fraction must be qualified by the detection limits of the instrument configuration used to carry out the observations.  As the number of detected binaries increases, it is becoming clear that the fraction of detectable TNBs is a function of factors beyond instrument sensitivity (Stephens and Noll 2006).  One of the surprising facets of the transneptunian population is its dynamical complexity, a complexity that has fueled a proliferation of classification schemes.  While there is no generally agreed-upon system of classification ({\\it e.g.} Elliot et al.~2005; Gladman et al.~2008; Lykawka and Mukai 2007), all of the proposed schemes recognize objects in stable, non-resonant, non-scattering, low-eccentricity ($e_\\odot$ ${<\\atop\\sim}$ $0.24$) orbits as a distinct class.  These objects are generally referred to as ``Classical\" transneptunian objects (TNO) and are mostly located between 42-48 AU.  The Classicals are sometimes further subdivided into two subgroups, Hot and Cold, based on their inclination.   The high-inclination Hot Classicals differ in measurable ways (color, absolute magnitude) from the low-inclination Cold Classicals and may have a distinct evolutionary history  (Tegler and Romanishin 2000; Levison and Stern 2001; Brown 2001).   As we show in this work, an additional distinction between these two subclasses is the fraction of objects that are binary.  We have been conducting a series of surveys of transneptunian objects using the Hubble Space Telescope since 2001.  In this paper we report on the result of searches for binary companions of 101 Classicals observed with the Near Infrared Camera Multi-Object Spectrometer (NICMOS) or the Advanced Camera for Surveys (ACS).  We find a remarkable concentration of binaries among the low inclination Classicals and, conversely, a marked absence of these systems among higher inclination Classicals.  The variation in the fraction of binaries with inclination appears to be another indicator that the Classicals are composed of two partially overlapping populations that may have originated in distinctly different parts of the early Solar System. ",
        "conclusions": "Binaries appear to be nearly absent among 100-km-sized Hot Classicals,  {\\it i.e.}~low-eccentricity, non-resonant Kuiper Belt objects with absolute magnitudes of $H_{\\rm V} \\geq 5$ and inclinations of $i \\geq 5.5^{\\circ}$.  Only one such pair has been found, the binary 2001 QC$_{298}$ (which is classified as a Classical by only one of the two systems we considered).  Conversely, at inclinations of $i < 5.5^{\\circ}$ binaries are common.  A large percentage, 29$\\pm {7\\atop 6}$ \\%, of low-inclination Classicals have detectable binaries that are separated by $s > 0.06$ arcsec and with a magnitude difference of $\\Delta_{\\rm mag} < 2$.  The largest Classicals, those with absolute magnitudes of $H_{\\rm V} < 5$, are different.  Of the 11 objects that fit this description, only 1 has an inclination less than 5.5$^{\\circ}$.  Satellites have been found around 3 of the 11.  However, all 3 systems have small secondaries compared to the primaries, $\\Delta_{\\rm mag} > 2$, unlike any of the more numerous, smaller Classical binaries.  One of these, (120347) 2004 SB$_{60}$ has heliocentric orbital elements that are similar to 2003 EL$_{61}$ and its possible collision family (Brown et al.~2007), although its neutral infrared colors (Kern, private communication 2007) make an association with this group unproven.  All of these factors suggest that the satellites of the largest Classicals may be collisional in origin as has been hypothesized for the largest TNOs generally (Brown et al.~2006).  While this leads to an unsatisfying complexity in origin of binaries among the Classicals, it may be an unavoidable byproduct of the fact that the Classicals are a composite population.  Of the two logical possibilities for the origin of an inclination dependence for binaries, only the two-populations hypothesis appears to be consistent with available evidence.  Several of the scenarios for Kuiper Belt formation have natural explanations for the existence of these two populations.  In all these models, these two parent populations are formed at different locations in the protoplanetary disk.  We are then left with the unsolved question of how, exactly, binary fraction, color, size, and possibly other physical properties varied as a function of heliocentric distance in the early Solar System.  Because the same processes could apply in other circumstellar disks, further investigations into these questions have an intrinsic value that goes beyond our own local example of a planetary system.  \\bigskip \\bigskip"
    },
    "0711/0711.4833_arXiv.txt": {
        "abstract": "We use a sample of 43,690 galaxies selected from the Sloan Digital Sky Survey Data Release 4 to study the systematic effects of specific star formation rate (SSFR) and galaxy size (as measured by the half light radius, $r_h$) on the mass-metallicity relation.  We find that galaxies with high SSFR or large $r_h$ for their stellar mass have systematically lower gas phase-metallicities (by up to 0.2 dex) than galaxies with low SSFR or small $r_h$.  We discuss possible origins for these dependencies, including galactic winds/outflows, abundance gradients, environment and star formation rate efficiencies. ",
        "introduction": "The well-established luminosity-metallicity relation (LZR) that exists over at least 10 magnitudes in M$_B$ (e.g., Skillman, Kennicutt \\& Hodge 1989; Zaritsky, Kennicutt \\& Huchra 1994; Salzer et al. 2005; Lee et al. 2006) has been recently confirmed to be a manifestation of a more fundamental stellar mass-metallicity relation (MZR; Lequeux et al. 1979; Tremonti et al. 2004).    Although there is evidence that galaxy interactions may affect the normalization of the LZR and MZR (Kewley et al. 2006; Ellison et al. 2007; Rupke et al. 2007), the MZR appears to be independent of large-scale environment (Mouhcine, Baldry \\& Bamford 2007) and remains intact out to $z \\sim 2.5$ (Savaglio et al. 2005; Erb et al. 2006).  Several origins of this apparently fundamental relation have been proposed.  Many of these models invoke winds as the fundamental driver of the MZR (e.g., Kobayshi, Springel \\& White 2007) and observational evidence indicates that gas ejection and/or accretion could be important (Tremonti et al. 2004; Gallazzi et al. 2005). However, the effects of supernova feedback on star formation efficiency and metal-poor gas infall have also been cited as important factors (Brooks et al. 2007; Finalator \\& Dav\\'e 2007). In this paper, we investigate the dependence of the MZR on various physical parameters that may shed light on the underlying mechanism that shapes it.  We adopt a concordance cosmology of $\\Omega_{\\Lambda} = 0.7$, $\\Omega_M = 0.3$, $H_0 = 70$ km/s/Mpc where applicable. ",
        "conclusions": "We have investigated how the galactic MZR depends on observable parameters such as the half light radius, SSFR and morphology and discuss various scenarios to explain the dependencies. We find that, at a given stellar mass, the MZR is offset towards lower metallicities (by up to 0.2 dex) for galaxies with larger half light radii and higher specific star formation rates, but that there is no significant dependence on bulge fraction.  These dependencies have little impact on the overall scatter in the MZR and the basic shape of the relation exists for all subsets of $r_h$ and SSFR.  We conclude that environment and abundance gradients are unlikely to account for (all) of the $r_h$ and SSFR dependence of the MZR.  Infall of metal-poor gas or metal-enriched outflows also seem unlikely explanations based on timescale arguments.  Of the possibilities considered here, a sensitivity to star formation efficiency is the most plausible reason for the dependence of the MZR on $r_h$ and SSFR."
    },
    "0711/0711.4372_arXiv.txt": {
        "abstract": "{In this work we investigate the stellar population, metallicity distribution and ionized gas in the elliptical galaxy NGC\\,5044, using long-slit spectroscopy and a stellar population synthesis method. We found differences in the slope of metal-line profiles along the galaxy which suggests an enhancement of $\\alpha$ elements, particularly towards the central region.   The stellar population synthesis showed that the component with $[Z/Z_\\odot]\\sim0.0$ dominates the $\\lambda5870$\\,\\AA\\ flux in the central region of NGC\\,5044, contributing with $\\sim42\\%$ of the total flux, while in the external regions the contribution decreases to $\\sim8.0\\%$. The component with $[Z/Z_\\odot]\\sim-0.4$ contributes with $\\sim32\\%$ in the central region, and $\\sim55\\,\\%$ in the external regions. The component with $[Z/Z_\\odot]\\sim-1.1$ contributes with $\\sim26\\,\\%$ in the central region, and $\\sim 37\\,\\%$ in the external regions.  The three components have $\\sim\\,10^{10}$ years. The presence of a non-thermal ionization source, such as a low-luminosity AGN and/or shock ionization, is implied by the large values of the ratio $\\frac{\\ion{[N}{ii]}}{H\\alpha}$ observed in all sampled regions.  However, the emission lines observed in the external regions indicate the presence of an additional ionization source, probably hot, post-AGB stars. ",
        "introduction": "A great effort has been made in the last few years to understand how early-type galaxies form and evolve. The relatively recent discoveries that ellipticals are not merely a one-parameter family and that many ellipticals show signatures of interaction with the environment suggest that these galaxies may have different star formation histories, with stellar populations differing in metallicity and/or age (Worthey, Faber \\& Gonzales 1992).  Crucial information on the above issues has been derived from metal line-strength indices (e.g. Davies et al. 1987) and their radial variation inside the galaxies. The $Mg2$ line-strength distribution in early-type galaxies, for example, can vary considerably, ranging from essentially featureless to structured profiles showing, e.g., changes of slope possibly associated with kinematically decoupled cores, or anomalies in the stellar population of some ellipticals  (Carollo et al. 1993).  NGC\\,5044 is the central and brightest member of a rich group that contains many dE dwarf members and a few Im and Sm dwarf candidates. The mean radial velocity of the group is $\\sim2048$ km\\,s$^{-1}$ (Ferguson \\& Sandage 1990). This galaxy presents a very bright ionized gas emission in the form of extended filaments up to $40\\arcsec$ from the center, being larger in the southern part of the galaxy (Macchetto et al. 1996; Goudfrooij 1991). The dust distribution has an irregular morphology, concentrated in the inner $10\\arcsec$, where two central dark clouds can be seen (Ferrari et al. 1999). NGC\\,5044 has also been  detected by IRAS at 60 and $100\\mu\\,m$ and also at 6, 12, and $18\\mu\\,m$ by ISO (Ferrari et al. 2002). Its IRAS luminosity is $14\\times 10^{8}\\,L_\\odot$. The galaxy has a nuclear radio source and extended X-ray emission (Fabbiano et al. 1992) and shows a very complex gas and stellar kinematics. The gas velocity profile is very irregular, with many  humps and dips; its radial velocity is systematically blueshifted with respect to the stellar systemic velocity by $\\sim60 - 100$ km\\,$\\mathrm{s^{-1}}$.  Furthermore, the stellar velocity curves observed on a fairly large area near the galaxy center counter-rotate with respect to the outer regions. Very slow or no stellar rotation is detected along the galaxy's minor axis. The velocity dispersion of the gas peaks at $\\sim200 - 230$ km\\,s$^{-1}$ in the center, decreasing rapidly outwards (Caon et al. 2000). The peculiar kinematics and gas morphology described above may be signatures either of interactions of NGC\\,5044 with the environment or a post-merger event.  The goal of this paper is to investigate the stellar population, metallicity distribution and the presence of ionized gas in NGC\\,5044, which are fundamental parameters to understand the formation and evolution of this galaxy.  This paper is organized as follows: Sect.~2 presents the observations and data reduction; Sect.~3 deals with the equivalent width analysis of the absorption lines and their radial dependence; Sect.~4 describes the stellar population synthesis, whose results are discussed in Sect.~5; Sect.~6  presents an analysis of the emission lines and the nature of the ionization source. Finally, in Sect.~7 a general discussion of the results and conclusions will be drawn. ",
        "conclusions": "The stellar population, metallicity distribution and ionized gas in NGC\\,5044 have been investigated in this paper by means of long-slit spectroscopy and stellar population synthesis. With respect to the spatial distribution of metal-strength indices (corrected for velocity dispersion), the \\ion{Mg}{2} profile increases towards the central region, contrary to what is observed for the profile of \\ion{Fe}{i}$_{5270}$. This difference in slopes may be accounted for by an enhancement of $\\alpha$-elements in general, Mg in particular. The \\ion{Mg}{2} gradient slope compared with the mass of NGC\\,5044 indicates a monolitic collapse.  However, the measure [Fe/H] gradient slope is not characteristic of one monolitic collapse.  The stellar population synthesis gives that most metallic component ($[Z/Z_\\odot]\\sim0.0$) dominates the $\\lambda5870$\\,\\AA\\ flux in the central region of NGC\\,5044. In the nucleus, this component contributes with $\\sim42\\%$ of the total flux, while in the external regions the contribution decreases to $\\sim8.0\\%$. The component with $[Z/Z_\\odot]\\sim-0.4$ contributes with $\\sim32\\%$ in the central region, and $\\sim55\\,\\%$ in the external regions. The metal-poor component ($[Z/Z_\\odot]\\sim-1.1$) contributes with $\\sim26\\,\\%$ in the central region, and $\\sim 37\\,\\%$ in the external regions. The spatial distribution of the internal reddening E(B$-$V)$_{\\rm{i}}$ along the north and south directions does not characterize a gradient.  The large values of the ratio $\\frac{\\ion{[N}{ii]}}{H\\alpha}$ observed in all sampled regions characterizes the presence of a non-thermal ionization source, such as a low-luminosity AGN and/or shock ionization.  However, in the external regions an additional ionization source is necessary to explain the emission lines, which might be hot, post-AGB stars."
    },
    "0711/0711.4975_arXiv.txt": {
        "abstract": "{} {The paper investigates the milli-arcsecond scale structure of the present-day mass-loss of the post-Red Supergiant \\IR.} {We use three telescopes of the VLT Interferometer in combination with the AMBER near-infrared beam combiner to measure spectrally dispersed correlated fluxes in the $K$-band around the Br$\\gamma$ transition. The resulting visibilities are compared to the predicted visibilities of emission structures with various simple models in order to infer the size  of the observed emission region.} {The Br$\\gamma$ line is resolved by VLTI+AMBER on all three baselines, with the maximum projected baseline extending 69 meter and a P.A. ranging between 10\\degr~and 30\\degr. A differential phase between line and continuum is detected on the longest baseline. The Br$\\gamma$ emission region is found to have a diameter of $3.3$ milli-arcseconds (FWHM), when compared to a Gaussian intensity distribution. A uniform disk and a ring-like intensity distribution do not fit the line visibilities. The size of the continuum emission is not constrained by the observations. Comparing the AMBER equivalent width of \\Brg with measurements from various epochs, we find that the stellar photosphere contributes about 60\\% of the total continuum light at 2.2$\\mu$m. The remaining 40\\% continuum emission is found on scales larger than the 66\\,mas AMBER field of view. This independently confirms  similar results made by previous studies. If the \\Brg emission is optically thin, then the observations do not allow to make any inferences about the shape of the line forming region. However, there is indirect evidence that the hydrogen recombination line emission is optically thick. In that case, using simple arguments, we find that the line emitting region is elongated. This is because the spectrum indicates that the projected line emitting area is about twice that of the stellar surface. This value is an order of magnitude less than a circular \\Brg line emitting area would have using the measured size of the emission region along our baseline. We briefly discuss the possibilities whether such a structure is due to a bi-polar flow or a circumstellar disk.} {} ",
        "introduction": "\\label{intro} Throughout the evolution of a high-mass star, stellar winds carry off a significant fraction of the initial star mass. This process dominates their evolution and determines their final fate. In the late stages of massive star evolution, important information of the mass loss process due to stellar winds is contained in the complicated circumstellar nebulae, like the ones found near e.g. LBVs and Wolf-Rayet stars (Nota et al. 1995\\nocite{1995ApJ...448..788N}; Nota \\& Clampin 1997\\nocite{1997ASPC..120..303N}; Van der Sluys \\& Lamers 2003\\nocite{2003A&A...398..181V}). The nebulae may disclose the mass-loss history, identifying phases with varying stellar wind properties like the wind momentum, geometry and total mass lost. One explanation for the observed nebula geometry is that the different mass loss episodes lead to wind-wind interactions, in the sense that a more recent fast wind collides with a pre-existing slower expanding wind. Crucially, at least one of these winds should be asymmetric, possibly induced by mass loss anisotropies at the stellar surface (Maeder 2002\\nocite{2002A&A...392..575M}). During which evolutionary phase this would actually occur is not known however, and various ideas have been put forward (see e.g. Mellema 1997\\nocite{1997A&A...321L..29M}; Heger \\& Langer 1998\\nocite{1998A&A...334..210H}; Dwarkadas \\& Owocki 2002\\nocite{2002ApJ...581.1337D}).  \\IR is  one of the few stars that evolve from the Red Supergiant (RSG) phase back towards the blue, witnessed by an extremely rapid increase in its surface temperature of $\\sim$2200\\,K over an interval of 30 years currently having a mid-A spectral type (see Oudmaijer et al. 1996\\nocite{1996MNRAS.280.1062O}; Klochkova et al. 2002\\nocite{2002ARep...46..139K}). Its estimated distance of 3.5 to 5\\,kpc implies a luminosity typical for a star with an initial mass of around $\\rm 40\\,M_{\\odot}$, putting it close to the Humphreys-Davidson limit in the Hertzsprung-Russell diagram (Jones et al. 1993\\nocite{1993ApJ...411..323J}). The star is bright at 2.2$\\mu$m $(K_s=3.6)$ and exhibits a huge infrared excess due to warm dust indicating that it was only recently in an extreme mass-losing Red Supergiant phase (Oudmaijer et al. 1996\\nocite{1996MNRAS.280.1062O}). These properties make the star a prime candidate to eventually become a Wolf-Rayet star, surrounded by an intricate nebula (Jones et al. 1993\\nocite{1993ApJ...411..323J})  The circumstellar environment of \\IR has various distinctive components. We list the basic features on the various size scales. Based on HST imaging, \\IR's optical reflection nebula portrays a spherical distribution on scales larger than 2\\arcsec~ out to 6 to 8\\arcsec~ (Humphreys et al. 1997\\nocite{1997AJ....114.2778H}). Sphericity of the circumstellar material on this scale is seen at other wavelengths too, e.g. in the $J$-band by Kastner \\& Weintraub (1995)\\nocite{1995ApJ...452..833K}, and in CO transitions by Castro-Carrizo et al. (2007\\nocite{2007A&A...465..457C}). On a scale of about 1\\arcsec, one finds a bipolar geometry, with two patches of $K$-band and mid-IR emission to the NE and SW direction of the star (Humphreys et al. 1997). Given the relatively large distance from the central star, this emission should be scattered radiation. The location of these patches corresponds to an approximate position angle of 40\\degr. Closer in and down to 0.3\\arcsec, HST reveals a complicated and non-uniform distribution of various jet-like, ray and arc features (Humphreys et al. 1997).  \\begin{table*} { \\begin{center} \\caption[]{Predicted and measured continuum visibilities for \\IR and the calibrator star HR\\,7648. \\IR has an observed upper limit to its diameter of 3\\,mas (Monnier et al. 2004), whereas the calibrator has a uniform disk diameter of 1.95\\,mas (Richichi et al. 2005). Corresponding theoretical visibilities are given in cols.(2) and (6). The measured raw values of $\\rm V_{cont}^{2}$ in cols. (3) and (7) correspond to the average continuum value for a SNR frame selection of 20\\% (see text), these values are subject to systematic uncertainties. The uncertainties on the raw $\\rm V_{cont}^{2}$ given reflect the internal variation. Col.(8) is the calibrated $\\rm V_{cont}^{2}$. Col.\\,(9) gives the ratio between the line and continuum (see Fig.\\,\\ref{select}). Col.\\,(10) is the range for $\\rm V_{Br\\gamma}$ as derived in Sect.\\,\\ref{brgsize}, taking into account the various uncertainty contributions.} \\begin{tabular}{cccc|ccccccc} \\hline \\hline &  HR7648     &              &                    & \\IR      &       &         &                     &                      &                  & \\\\ &  Baseline   & $\\rm V_{UD}$ & $\\rm <V_{cont}^2>$ & Baseline & P.A.  &  V      & $\\rm <V_{cont}^2>$   &   $\\rm <V_{cont}^2>$  & V$^2_{\\rm cont.}$/V$^2_{Br\\gamma}$  &   $V_{Br\\gamma}$ \\\\ &  (m)        & theo.         & raw                & (m)      & \\degr &  theo.  & raw  & calib.   &&\\\\ &  (1)  & (2) & (3) & (4) & (5) & (6) & (7) & (8) & (9)  & (10) \\\\ \\hline U1-U2 &  43.9 &  0.96  & $0.111\\pm0.002$          & 39.8     & 12.8  & $>0.92$ & $0.169\\pm0.004$ & $1.46\\pm0.03$& $1.36\\pm0.07$&0.60--0.74\\\\ U2-U3 &  35.8 &  0.97  & $0.149\\pm0.003$          & 30.3     & 32.7  & $>0.95$ & $0.210\\pm0.005$ & $1.37\\pm0.03$& $1.20\\pm0.07$&0.74--0.89 \\\\ U3-U1 &  78.7 &  0.87  & $0.201\\pm0.004$          & 69.1     & 21.5  & $>0.76$ & $0.265\\pm0.005$ & $1.15\\pm0.03$& $1.68\\pm0.06$&0.39--0.57\\\\ \\hline \\end{tabular} \\label{tabvs} \\end{center} } \\label{obs} \\end{table*}  At even smaller angular scales, the dusty environment giving rise to the near-IR excess has been resolved with speckle interferometry at 2.11$\\mu$m (Bl\\\"{o}cker et al. 1999\\nocite{1999A&A...348..805B}). Modelling the spectral energy distribution (SED) and visibility curves with DUSTY leads these authors to conclude that \\IR is surrounded by two separate spherical shells with diameters of $0.070\\arcsec$ and $0.310\\arcsec$ centred on the star. Finally, the geometry on a size scale comparable to the stellar radius ($\\sim$ milli-arcseconds) can be probed with the various emission lines, like H$\\alpha$ and \\ion{Fe}{ii}, that are reflected off the large scale blue nebula (Humphreys et al. 2002\\nocite{2002AJ....124.1026H}; Davies et al. 2007\\nocite{2007arXiv0708.2204D}). With IFU observations of the nebula, Davies et al. find evidence for a present-day wind which has an axi-symmetric, rather than spherically symmetric geometry. The H$\\alpha$ emission symmetry axis is found to have a position angle of 33\\degr.  In this paper on \\IR we present complementary observations of the subarcsecond ionised and continuum emitting material inside the 70\\,mas dust shell found by Bl\\\"{o}cker et al. (1999\\nocite{1999A&A...348..805B}). These authors found that about 60\\% of the $K$-band continuum emission remains unresolved (see also Monnier et al. 2004), presumably corresponding to the star itself. We use the VLT interferometer in conjunction with the near-IR three telescope beam combiner AMBER, attaining a maximum spatial resolution of 6.5\\,mas. The ultimate objective is to measure the size and shape of the Br$\\gamma$~emission line region, giving a direct measurement of the present-day mass-loss geometry, an important ingredient to the mass-loss puzzle. In Sect.\\,\\ref{observ} we present the AMBER data, and provide a list of steps and considerations taken in the reduction process. Results regarding the final visibility, flux and differential phase spectra of \\IR are presented in Sect.\\,\\ref{results}. In this section we derive the size for the continuum and Br$\\gamma$~emission. We discuss the shape of the Br$\\gamma$~emission region in Sect.\\,\\ref{discu}. The latter is found to be elongated along the VLTI baseline of the observation with a P.A. of $\\sim 20\\degr$. We conclude our findings in Sect.\\,\\ref{concl}. ",
        "conclusions": "\\label{concl} We have presented VLTI/AMBER observations in the $K$-band of the post-Red Supergiant transition object \\IR, with the aim to spatially resolve the Br$\\gamma$ emission region. The observations were done with three UT telescopes and consist of intermediate spectral resolution spectra of the coherent and incoherent flux of \\IR within the AMBER field of view of 66\\,mas and with an angular resolution of 6.5\\,mas. The telescope configuration probed a range of 20\\degr~in position angle. The conclusions of this study are as follows.   \\begin{itemize}  \\item The AMBER visibility spectra show a decreased \\Brg line visibility for all three baselines.  It indicates that the line emitting region has been resolved. The continuum visibilities on the other hand are not constrained by this dataset. We fit the \\Brg line visibilities with simple intensity distributions, and find that a Gaussian distribution with a FWHM of 3.3\\,mas fits the data best. The derived diameter corresponds to a distance of approximately $\\rm 5\\,R_{*}$ from the central star.  \\item The AMBER differential phase spectra show a change in phase for the U3-U1 baseline at the \\Brg transition.  This could indicate an offset of the photocentre projected on this baseline. New measurements of better quality would help to characterize the nature of this signal.   \\item The AMBER flux spectra show the \\Brg line of \\IR in emission.  It is stronger than reported by two previous studies. We have argued that this is due to the fact that most of the $K$-band excess emission is located outside the AMBER field of view. This result is consistent with the speckle-interferometric results presented previously by Bl\\\"ocker et al. (1999), where they find the inner most dust sphere to have a diameter of 70\\,mas.  \\end{itemize} We have speculated on the geometry of the \\Brg emitting region, concluding that it would be elongated along the interferometer baseline if the emission is optically thick. Optically thick emission seems however to be consistent with observed line ratios. Additional measurements perpendicular to the baselines presented here will help in further constraining the processes that shape the environment of this massive star in transition to its final fate."
    },
    "0711/0711.0551_arXiv.txt": {
        "abstract": "In the standard supernova picture, type Ib/c and type II supernovae are powered by the potential energy released in the collapse of the core of a massive star.  In studying supernovae, we primarily focus on the ejecta that makes it beyond the potential well of the collapsed core.  But, as we shall show in this paper, in most supernova explosions, a tenth of a solar mass or more of the ejecta is decelerated enough that it does not escape the potential well of that compact object.  This material falls back onto the proto-neutron star within the first 10-15 seconds after the launch of the explosion, releasing more than $10^{52}$\\,erg of additional potential energy. Most of this energy is emitted in the form of neutrinos and we must understand this fallback neutrino emission if we are to use neutrino observations to study the behavior of matter at high densities.  Here we present both a 1-dimensional study of fallback using energy-injected, supernova explosions and a first study of neutrino emission from fallback using a suite of 2-dimensional simulations. ",
        "introduction": "The collapse of a massive star down to a neutron star releases over $10^{53}$\\,erg of potential energy.  Just one percent of this energy is required to power the observed type Ib/c and type II supernovae. Most of the energy is emitted in the form of neutrinos.  How this small fraction of the energy is converted into explosion energy is still a matter of debate, but it is believed by most that neutrinos play a role in depositing energy above the collapsed core to drive this explosion (see Fryer 2003 for a review).  Whether or not neutrinos are important for the explosion, they do provide astronomers a window into the explosion mechanism behind supernovae.  Neutrinos also provide a window into the material conditions at the core of a supernova explosion.  With temperatures above 10\\,MeV and densities above nuclear densities, core-collapse supernovae make ideal laboratories for nuclear physics.  To extract information about nuclear physics from such laboratories, we must understand the systematics in our supernova experiment.  When the explosion engine is active, convective motions in the engine (Herant et al. 1994) make it very difficult to interpret the neutrino signal.  An observed elevated neutrino luminosity could be a change in the neutrino opacity or it could be a difference in the convective engine.  One approach to eliminate issues with convection would be to wait until all convective activity ceases.  A number of studies have now been presented following the evolution of a collapsing star over 0.5-1.5\\,s after bounce (Fryer \\& Heger 2000; Burrows et al. 2006; Scheck et al. 2006; and, in 3-dimensions, Fryer \\& Young 2007).  In all these calculations, hydrodynamic motions near the proto-neutron star continue to dramatically affect the neutrino emission through the end of the simulations.  To achieve a clean neutrino signal, we must wait until after the launch of the explosion.  We must also wait a few seconds after the launch of the explosion because the proto-neutron star may experience deep convection (Keil et al. 1996).  After this time (roughly a few seconds to 10-20\\,s), the supernova finally seems to have achieved an ideal condition as a physics laboratory (Reddy et al. 1999).  After 10-20\\,s, the neutrino signal will be too weak to do much experimental science, so we truly are limited to this narrow time window.  Unfortunately, even at these late times, Nature does not allow completely pristine conditions.  Material falling back from the supernova explosion (``fallback'') may well produce a new round of convection and confusion to our neutrino signal.  This fallback has been studied both in its important role in calculating the initial mass of the neutron star formed in a supernova explosion (e.g. Fryer \\& Kalogera 2001) and in estimating the r-rpocess yields in supernovae Fryer et al. (2006). In this paper, we study its role in determining the neutrino luminosity arising after the first few seconds of a supernova.  In \\S 2, we review the history of supernova fallback and present new calculations of supernova explosions estimating the fallback for a range of explosion energies and stellar masses.  \\S 3 describes the 2-dimensional code used to model the neutrino emission from the this fallback and shows results for the suite of simulations run for this paper.  Different than the multi-dimensional simulations modeling stellar collapse, these simulations do not start at the onset of collapse (or bounce) and end at the launch of the explosion. Instead, they start a 2-10\\,s (depending on the explosion energy, etc. from \\S 2) after the launch of the explosion when material begins to fall back onto the neutron star.  We conclude with a brief discussion on how neutrino signals can be used to help better understand the supernova explosion, neutron star birth masses and neutrino cross-sections. ",
        "conclusions": "Fallback (of at least 0.1\\,M$_\\odot$) is likely to occur in normal ($10^{51}$\\,erg) supernova explosions from stars with initial masses of 12\\,M$_\\odot$ or greater.  With energy-injected (more realistic than piston-driven) explosion calculations, this fallback occurs quickly, generally in the first 15\\,s and peak accretion rates in the 0.01-0.1\\,M$_\\odot$\\,s$^{-1}$ range should be expected.  If the engine stays active long after the launch of the supernova shock, the total fallback will be lower.  But neutrino engines weaken significantly quickly after the shock is launched and the material above the neutrinosphere (absorbing the energy) is ejected. Magnetic-driven explosions could be very different.  The mechanism for this fallback is essentially that proposed by Colgate (1971) whereby the outflowing material decelerates as it pushes against the material above it, ultimately decreasing its velocity below the escape velocity and causing it to fall back onto the proto-neutron star.  This means that the bulk of the fallback is fairly insensitive to the structure on the outer envelope of the star (so fallback is roughly the same whether or not the star is in a binary).  Our multi-dimensional simulations of this fallback suggest that not all of this fallback material is actually incorporated into the proto-neutron star.  Some of it flows out of the system.  The corresponding neutrino luminosity is also lower, both because less material is accreted and some of the energy released goes toward driving outflows.  But the neutrinos from fallback could still dominate the neutrino emission from a proto-neutron star 10\\,s into the explosion.  At this time, observations of neutrinos from supernovae are limited to those of supernova 1987A (Hirata et al. 1987, Bionta et al. 1987). With only 20 neutrinos over $\\sim$15\\,s, it is difficult to place many constraints on the fallback.  Some authors have claimed that claimed that the late-time neutrinos could not be easily explained by a cooling neutron star (e.g. Suzuki \\& Sato 1987), but others have argued that the observed neutrino signal is consistent with neutrinos diffusing out of a hot proto-neutron star (Burrows \\& Lattimer 1987; Bruenn 1987) .  In our neutron-star forming models, fallback ends within the first 10-15\\,s, comparable to the duration of the observed neutrino burst and the flux is consistent with our more standard accretion rates.  The observations are consistent with fallback, but could easily be explained by a neutron star without any fallback. Given that the progenitor star is believed to be greater than 15\\,M$_\\odot$, fallback is likely to have occured.  But with the current errors in the time-dependent luminosity, it is difficult to determine whether or not accretion is taking place.  If we assume fallback accretion did occur, the fact that the mean neutrino energy appears to be dropping with time in the observations suggests that the accretion rate is dropping at 10\\,s.  But neutrinos are still observed at 10\\,s.  Unless the fallback material has considerable angular momentum, the compact remnant is not a black hole at 10\\,s.  In addition, because the accretion rate is dropping dramatically at this time, it is unlikely that the remnant will accrete much additional mass.  It will remain a neutron star.  In our fallback calculations, systems that accrete enough material to form a black hole are either already a black hole at 10\\,s or still accreting rapidly at 10\\,s, neither of which is supported by the observations.  But we really need a better neutrino signal to say anything definitive.  With such high accretion rates, fallback can easily dominate the neutrino luminosity after a few seconds, more than doubling the emission from the neutron star at thearly times.  If we wish to study neutrino opacities in the supernova explosion, we will have to be able to calculate this fallback.  It may be possible to estimate the fallback rate from detailed study of supernova light-curves. Fortunately, any system that is detected in neutrinos will have a wealth of data in photons of all wavelengths.  If the fallback rate is high, the infalling material will alter the position of the neutrinosphere, and the problem becomes nearly as complex as the supernova explosion modeling itself.  We have only touched the surface on the difficulties in modeling such systems.  In this paper, we have not addressed the production or feedback of magnetic fields, which may definitely alter the fate of the fallback. We have also not discussed the nuclear yields from the ejecta (a first paper on this is by Fryer et al.  2007 and we plan future projects studying this nucleosynthesis).  However, note that the old ``wind'' r-process picture will not work in any system with fallback (stars more massive than $\\sim 12$M$_\\odot$).  The matter trajectories just can't be explained by the wind solution in the cases where fallback dominates the motion near the neutron star.  We also have not studied the actual accretion onto the neutron star in detail.  Note that the final neutron star masses in Table~\\ref{tab:fallback} assumed no mass ejecta.  But as we have found in this study, over half of the fallback material may be ejected, leading to smaller neutron star masses and a narrower mass range for these neutron stars."
    },
    "0711/0711.2886_arXiv.txt": {
        "abstract": "We present a complete spectroscopic survey of 2414 2MASS-selected blue horizontal branch (BHB) candidates selected over 4300 deg$^2$ of the sky. We identify 655 BHB stars in this non-kinematically selected sample. We calculate the luminosity function of field BHB stars and find evidence for very few hot BHB stars in the field. The BHB stars located at a distance from the Galactic plane $|Z|<4$ kpc trace what is clearly a metal-weak thick disk population, with a mean metallicity of [Fe/H] $=-1.7$, a rotation velocity gradient of $d v_{rot}/d|Z|=-28\\pm3.4$ km s$^{-1}$ in the region $|Z|<6$ kpc, and a density scale height of $h_Z=1.26\\pm0.1$ kpc. The BHB stars located at $5<|Z|<9$ kpc are a predominantly inner-halo population, with a mean metallicity of [Fe/H] $=-2.0$ and a mean Galactic rotation of $-4\\pm31$ km s$^{-1}$.  We infer the density of halo and thick disk BHB stars is $104\\pm37$ kpc$^{-3}$ near the Sun, and the relative normalization of halo to thick-disk BHB stars is $4\\pm1$\\% near the Sun. ",
        "introduction": "Theoretical simulations show that the remnants of hierarchical galaxy formation in the Milky Way should still be visible as star streams in the stellar halo \\citep{johnston96, harding01, abadi03b, bullock05, font06}. Star counts and color-magnitude diagrams have proven very effective in identifying structures in the halo, including the Sagittarius stream wrapping around the sky \\citep{majewski03} and overdensities in Monoceros \\citep{newberg02,ibata03,yanny03}, Canis Major \\citep{martin04}, Triangulum-Andromeda \\citep{rochapinto04}, Virgo \\citep{duffau06,vivas06,newberg07}, and elsewhere \\citep{belokurov06, grillmair06a, grillmair06b}.  Stellar spectroscopy opens up the possibility of finding structures in velocity \\citep[such as the Sagittarius dwarf galaxy, e.g.][]{ibata94}, in metallicity, and in distance.  The major difficulty in mapping the stellar halo is finding tracer stars that are luminous enough to observe at great depths yet common enough to densely sample the halo.  In \\citet[hereafter Paper I]{brown03}, we introduced the Century Survey Galactic Halo Project, a photometric and spectroscopic survey from which we selected Blue Horizontal-Branch (BHB) stars as probes of the Milky Way halo. BHB stars are evolved, helium core-burning stars that serve as excellent ``standard candles.'' In \\citet[hereafter Paper II]{brown05b}, we explored the Two Micron All Sky Survey \\citep[2MASS]{skrutskie06} and the Sloan Digital Sky Survey \\citep[SDSS]{adelman06} as the basis for a large spectroscopic survey of BHB stars.  In Paper II we calculated the first field BHB luminosity function, and concluded that field BHB stars are consistent with populations seen in most globular clusters, but inconsistent with globular clusters that have substantial extended BHBs.  Here we describe a complete, non-kinematically selected sample of BHB stars covering 10\\% of the entire sky.  Our survey is inspired by \\citet{brown04}, in which we photometrically selected BHB candidates from the completed 2MASS catalog. We have now obtained spectroscopy for 2414 2MASS-selected BHB candidates, allowing us to measure velocities and metallicities for stars to a depth of 8 kpc over a 4300 deg$^2$ region. Such a large-area survey is necessary to unambiguously identify halo structure: theoretical simulations predict that star streams cover hundreds of square degrees on the sky \\citep{bullock05, font06}.  Previous spectroscopic surveys of field BHB stars \\citep{pier83, sommer89, arnold92, kinman94, wilhelm99b, kinman04, clewley04, clewley05, kinman07} have identified BHB stars over relatively small fractions of the sky (10$^2$ - 10$^3$ deg$^2$) compared to the Century Survey Galactic Halo Project.  The exception is the sample of 1170 BHB stars observed by the SDSS as mis-identified quasars or as filler objects in low density regions \\citep{sirko04a, sirko04b, clewley06}.  In comparison, our spectroscopic survey of BHB stars is cleanly selected and is 100\\% complete within the selection limits.  Although our large-area spectroscopic survey is a rich source for general studies of the thick-disk and inner-halo populations, here we focus our attention on the properties of the BHB stars.  Our goal is to lay the groundwork for a structure analysis to be presented in later paper (in preparation).  In \\S 2 we describe spectroscopic observations of stars in the new 4300 deg$^2$ region, and discuss our selection efficiency for BHB stars.  In \\S 3 we discuss the global properties of the sample, including the mean galactic rotation and metallicity distribution of the stars.  In \\S 4 we calculate the luminosity function of our clean sample of field BHB stars.  In \\S 5 we fit for the density distribution of the BHB stars and determine the relative normalization of thick disk to inner halo stars.  We conclude in \\S 6. ",
        "conclusions": "We discuss a complete spectroscopic survey of 2414 2MASS-selected BHB candidates over 4300 deg$^2$ of sky.  We identify 655 BHB stars in this non-kinematically selected sample.  The luminosity function of the field BHB stars has a median value of $M_V=0.65$ and a small tail extending to $M_V>1$, but shows very few extended BHB stars.  The BHB stars located at $|Z|<4$ kpc have a mean Galactic rotation and density distribution remarkably consistent with a metal-weak thick-disk population. The $|Z|<4$ kpc BHB stars have a mean metallicity of [Fe/H] $=-1.7$, a velocity gradient of $d v_{rot}/d|Z|=-28\\pm3.4$ km s$^{-1}$, and a vertical scale height of $h_Z=1.26\\pm0.1$ kpc.  We infer a space density of $100\\pm36$ kpc$^{-3}$ thick disk BHB stars near the Sun.  RR Lyrae stars, by comparison, have a much less prominent disk component near the Sun \\citep{martin98}.  The existence of metal-poor BHB stars with thick-disk kinematics and scale heights present another clue for formation scenarios of the Milky Way.  The BHB stars located at $5<|Z|<9$ kpc have a mean Galactic rotation and density distribution consistent with a predominantly halo population. The halo BHB stars have a mean metallicity of [Fe/H] $=-2.0$, a mean Galactic rotation of $-4\\pm31$ km s$^{-1}$, and relative normalization of $4\\pm1$\\% with respect to the thick disk near the Sun $(R, Z)=(8, 0)$ kpc.  This is the best currently available measurement of the relative normalization of the halo and thick disk, an important quantity that enters into all models of the local structure of the Galaxy.  In the future, having established the global properties of our survey, we hope to analyze the BHB sample for structure in space, velocity, and metallicity."
    },
    "0711/0711.2281_arXiv.txt": {
        "abstract": "The ISO/LWS far-infrared spectrum of the ultraluminous galaxy Mkn 231 shows OH and H$_2$O lines in absorption from energy levels up to 300 K above the ground state, and emission in the [O I] 63 $\\mu$m and [C II] 158 $\\mu$m lines. Our analysis shows that OH and H$_2$O are radiatively pumped by the far-infrared continuum emission of the galaxy. The absorptions in the high-excitation lines require high far-infrared radiation densities, allowing us to constrain the properties of the underlying continuum source. The bulk of the far-infrared continuum arises from a warm ($T_{{\\rm dust}}=70-100$ K), optically thick ($\\tau_{100\\mu{\\rm m}}=1-2$) medium of effective diameter 200-400 pc. In our best-fit model of total luminosity $L_{{\\rm IR}}$, the observed OH and H$_2$O high-lying lines arise from a  luminous ($L/L_{{\\rm IR}}\\sim0.56$) region with radius $\\sim100$ pc. The high surface brightness of this component suggests that its infrared emission is dominated by the AGN. The derived column densities $N({\\rm OH})\\gtrsim10^{17}$ cm$^{-2}$ and $N({\\rm H_2O})\\gtrsim6\\times10^{16}$ cm$^{-2}$ may indicate XDR chemistry, although significant starburst chemistry cannot be ruled out. The lower-lying OH, [C II] 158 $\\mu$m, and [O I] 63 $\\mu$m lines arise from a more extended ($\\sim350$ pc) starburst region. We show that the [C II] deficit in Mkn 231 is compatible with a high average abundance of C$^+$ because of an extreme overall luminosity to gas mass ratio. Therefore, a [C II] deficit may indicate a significant contribution to the luminosity by an AGN, and/or by extremely efficient star formation. ",
        "introduction": "\\label{sec:intro}  The peculiar ultraluminous infrared galaxy (ULIRG, $L_{{\\rm IR}}\\ge10^{12}$ \\Lsun) Markarian 231 (Mkn 231, 12540+5708) is the most luminous infrared galaxy in the local universe, with a 8-1000 $\\mu$m luminosity of $3.2\\times10^{12}$ \\Lsun\\ \\citep{san03}, and may be a representative example of the link between AGNs and nuclear starbursts \\citep{sco04}. A QSO-like nucleus is evident from many observations:  optically it is classified as a Type 1 Seyfert \\citep*{bok77,cut84,baa98},  it exhibits UV through IR polarization and broad absorption lines \\citep{smi95}, it has compact X-ray emission \\citep[e.g.,][]{gal02} and extremely compact mid-infrared emission \\citep{soi00}, and in the radio it is variable and possesses a parsec scale jet \\citep*{ulv99,tay99}. Nevertheless, there is also evidence of a compact starburst in these results as well as in VLA observations of H I 21 cm absorption \\citep*{car98}, near-infrared observations \\citep{tac02}, and millimeter CO interferometry \\citep[][hereafter DS98]{bry96,dow98}. Estimates for the starburst luminosity range from 1/3 to 2/3 of the bolometric luminosity \\citep[][DS98]{dav04}.  Molecular observations have provided important clues about the concentration and kinematics of the gas in Mkn 231. DS98 showed the presence of an inner nuclear disk of radius $\\sim460$ pc in CO (2-1), and a more extended disk with lower brightness. Most of the molecular gas has been found to be dense ($\\sim10^4$ \\cmt) and warm ($\\sim70$ K) from recent observations of CO and HCN submillimeter lines \\citep*[][hereafter PIW07]{pap07}. \\cite{lah07} have inferred embedded starburst chemistry in Mkn 231 and other ULIRGs based on mid-IR Spitzer observations of ro-vibrational bands of warm/hot HCN and C$_2$H$_2$, while \\cite{gra06} and \\cite{aal07} have inferred XDR chemistry and/or radiative pumping based on anomalous intensity ratios of millimeter lines of HCN, HNC, and HCO$^+$.   The bulk of the luminosity in ULIRGs is emitted at far-infrared (FIR) wavelengths, where a number of molecular tracers are detected, mostly in absorption. Prominent lines of OH and \\hdo\\ were detected using ISO/LWS in the FIR spectrum of Arp 220, along with absorption features by radicals such as NH and CH, revealing a chemistry that may be indicative of PDRs with plausible contribution by shocks and hot cores \\citep[][hereafter Paper I]{gon04}. However, those species are also expected to be enhanced in XDRs \\citep{mei05}, so that the dominant chemistry in the nuclear regions of ULIRGs remains uncertain. In Paper I, the ISO/LWS FIR spectrum of Arp 220 was analyzed by means of radiative transfer calculations, which included a non-local treatment of the molecular excitation by absorption of FIR photons. Paper I showed that the population of high-excitation OH and \\hdo\\ rotational levels, in evidence from absorption in high-lying lines, is pumped through absorption of FIR continuum photons, a process that requires high FIR radiation densities. The detection of these lines thus not only reveals the chemical and excitation conditions in the absorbing regions, it also sheds light on the size and characteristics of the underlying continuum FIR source in spite of the low angular resolution currently available at these wavelengths.  In this paper we extend our approach of Paper I to the ISO/LWS FIR spectrum of Mkn 231, and show that this galaxy spectrum presents striking similarities to that of Arp 220.  Specifically, strong absorption in the high-excitation OH and \\hdo\\ lines is also seen in Mkn 231. Rotationally excited OH in Mkn 231 has been previously detected via the $^2\\Pi_{1/2}\\,\\,\\Lambda$-doublet transitions \\citep*{hen87}. VLBI observations of the mega-maser OH emission at 18 cm wavelength trace an inner torus or disk of size $\\sim100$ pc around the AGN \\citep*{klo03}, and MERLIN observations were able to map essentially the whole single-dish mega-maser OH emission with angular resolution of $\\approx0.3''$ \\citep{ric05}. We analyze here both the FIR continuum emission and the high-excitation OH and \\hdo\\ lines, as well as the [C II] 158 $\\mu$m and [O I] 63 $\\mu$m emission lines. In \\S\\ref{sec:obser} we present the ISO spectroscopic observations of Mkn 231.  In \\S\\ref{sec:models} we first analyze simple models for the FIR continuum emission from Mkn 231, and then examine how well those models reproduce the observed FIR emission and absorption lines.  \\S\\ref{sec:discussion} summarizes our results. We adopt a distance to Mkn 231 of 170 Mpc ($H_0=75$ km s$^{-1}$ Mpc$^{-1}$ and $z\\approx0.042$). ",
        "conclusions": "\\label{sec:discussion}  Simultaneous modeling of the FIR continuum and line absorptions is the only consistent way to derive from FIR observations the physical and chemical properties of bright-infrared galaxies, where due to high radiation densities, the OH and \\hdo\\ lines are pumped through absorption of continuum photons. In Mkn 231, the high-lying 65-68 $\\mu$m OH and \\hdo\\ lines indicate the presence of a compact ($\\sim100$ pc) and warm ($\\sim100$ K) region, presumably located around the AGN, with high continuum opacity ($\\tau_{{\\rm 100\\mu m}}\\sim1$). The size of this region is comparable to (though somewhat larger than) the inner nuclear disk where the OH mega-masers have been observed from VLBI observations \\citep{klo03}, and is similar to the spatial extent of the regions where the bulk of the OH mega-maser emission is generated \\citep[regions C, NE, and SW in Fig. 5 of][]{ric05}. This strongly suggests that we are observing the FIR counterpart of the 18 cm OH emission \\citep[e.g.,][]{ski97}. \\cite{ric05} estimate $N({\\rm OH})\\sim1.5\\times10^{15}$ \\cmd\\ in 1 pc clouds where $\\Delta V\\approx1.7$ \\kms; therefore, $N({\\rm OH})/\\Delta V \\sim 10^{15}$ \\cmd/(\\kms). In both the warm and cold components of model $D$, we have obtained $N^{{\\rm scr}}({\\rm OH})/\\Delta V \\sim 1.5\\times 10^{15}$ \\cmd/(\\kms), in rough agreement with the estimates by \\cite{ric05}. The ``screen'' case is most appropriate for comparison here because the OH megamasers amplify the background radio continuum emission.   The nuclear region is surrounded by a more extended ($R_{eq}\\sim350$ pc) and colder ($\\sim45$ K) region, which produces the bulk of the [C II] and [O I] line emissions and contributes to the low-lying OH lines. Allowing for an area filling factor of 0.6-0.7, the total extent of the FIR emission is then comparable to the size of the CO (inner) disk \\citep[][DS98]{bry96}. The inferred densities ($n({\\rm H_2})\\sim10^4$ \\cmt) may explain the submillimeter CO emission detected by \\cite{pap07}, and somewhat lower densities (${\\rm a \\,\\, few}\\times10^3$ \\cmt) can account for the [O I] line emission. We do not however rule out the possibility that some of the FIR emission arises from a more extended region, as the outermost disk observed by DS98 in the millimeter CO lines or the much more extended region observed in soft X-rays \\citep{gal02}.   The FIR emission from $D_{{\\rm warm}}$ is probably dominated by the AGN, as pointed out in \\S\\ref{sec:continuum}. \\cite{dav04} estimated that, within a radius of 330 pc, the starburst luminosity accounts for $(0.8-1.3)\\times10^{12}$ \\Lsun; since the luminosity for the component $D_{{\\rm cold}}$ is $\\approx7.5\\times10^{11}$ \\Lsun, the starburst luminosity contribution from $D_{{\\rm warm}}$ is most probably $\\lesssim6\\times10^{11}$ \\Lsun. \\cite{dav04} also estimated that, within 100 pc of the AGN (the size of $D_{{\\rm warm}}$), the starburst contribution is $(3.2-4.7)\\times10^{11}$ \\Lsun, in agreement with the afore mentioned upper limit. On the other hand, the extended ($360\\times260$ pc$^2$) radio continuum 1.4 GHz emission detected by \\cite{car98} is compatible with both the whole FIR emission arising from the starburst \\citep[yielding an IR-to-radio flux density ratio of $Q\\sim2.5$,][]{car98}, and only the $D_{{\\rm cold}}$ FIR emission arising the starburst (yielding $Q=2.1$). If one now assumes that the starburst contribution to the FIR emission of $D_{{\\rm warm}}$ is $3\\times10^{11}$ \\Lsun, the resulting $Q$ is 2.3 \\citep[in agreement with the median value of $Q=2.3\\pm0.2$ in normal galaxies;][and references therein]{con92}. Also, the surface brightness and luminosity-to-mass ratio of $D_{{\\rm warm}}$ due to only the starburst component would then be $2\\times10^{12}$ \\Lsun\\ kpc$^{-2}$ and 550 \\Lsun/\\Msun, which are probably upper limits for a starburst. In summary, our preferred model $D$ favors that $\\sim2/3$ of the bolometric luminosity from Mkn 231 is due to the AGN.    Although the chemistry of the extended starburst $D_{{\\rm cold}}$ component, where the lower-lying OH lines and the bulk of the [C II] 158 $\\mu$m and [O I] 63 $\\mu$m lines are generated, appears to be dominated by PDRs, the dominant chemistry of the warm component may be of a different nature. Since the compactness of $D_{{\\rm warm}}$ suggests that its FIR emission is dominated by the AGN, and given the high X-ray intrinsic luminosity of the AGN ($\\sim10^{44}$ erg s$^{-1}$) in the 2-10 keV band \\citep{bra04}, we consider the possibility that the derived OH and \\hdo\\ column densities can be better explained in terms of XDR chemistry. Calculations by \\cite{mei05} show that, in dense XDRs with $N_{{\\rm H}}\\sim10^{24}$ cm$^{-2}$, the OH and \\hdo\\ column densities are expected to attain values of $\\sim10^{18}$ cm$^{-2}$, similar to the values that we have derived in the mixed case. Furthermore, the Compton-thick screen with $N_{{\\rm H}}\\sim2\\times10^{24}$ cm$^{-2}$ required to block the primary X-ray emission \\citep{bra04} could be partially identified with $D_{{\\rm warm}}$, with similar column density. Therefore an XDR may naturally explain the observations of high-lying OH and \\hdo\\ lines in Mkn 231. Nevertheless, both OH and \\hdo\\ are also tracers of starburst chemistry, and a number of PDRs and hot cores overlapping along the line of sight could in principle yield the column densities required to explain our observations at least in the screen case, and could then mimic an XDR. We conclude that further observations are still required to discern between XDR and starburst chemistry.    We suggest that the [C II] 158 $\\mu$m line deficit in Mkn 231, and also probably in Arp 220 (see Paper I), is primarily due to a component that dominates the FIR emission but emits weakly in the [C II] line, as previously suggested by \\cite{luh03}. However, we also suggest that this component may still be rich in C$^+$, but with an extreme luminosity-to-gas-mass ratio that limits the [C II] luminosity per unit of luminous power in the continuum. Our derived luminosity to H$_2$ mass ratio of 500 \\Lsun/\\Msun\\ is high; future studies will indicate whether it can be applied to other ULIRGs with similar [C II] deficits.    Given the high C$^+$ column densities (a few $\\times10^{19}$ \\cmd) required to account for the [C II] 158 $\\mu$m line emission, the H$_2$ columns derived toward the extended component are unlikely to be associated with non-overlapping star-forming regions surrounded by optically thick envelopes. A single PDR is not expected to have $N({\\rm C^+})$ in excess of ${\\rm a \\,\\,few}\\times10^{17}$ \\cmd. A scenario with a crowded population of PDRs overlapping along the line of sight is therefore more plausible. If a typical {\\it single} OB stellar cluster emits up to $10^{39}$ erg s$^{-1}$, the maximum value found by \\cite{sco01} in M51 and also the luminosity of W49, then the number of single clusters to produce the Mkn 231 starburst luminosity of $10^{12}$ \\Lsun\\ is $\\gtrsim4\\times10^6$. If these clusters are concentrated in a disk of radius 460 pc and thickness 25 pc (DS98), the mean distance between neighboor clusters is only $\\lesssim1.5$ pc \\citep[see also][]{ket92}. Merging of PDRs may be compatible with some area filling factor if, for example, some spiral structure is invoked within the disk.   In NGC 1068, \\cite{spi05} observed the OH 119 $\\mu$m line in emission against the continuum. On the basis of a possible XDR chemistry, and given the corresponding mass, density and temperature that characterize the nuclear region, \\cite{spi05} suggested that the OH 119 $\\mu$m emission line could be formed in that region. This option also relied on the fact that relatively weak FIR emission is expected to arise from the circumnuclear disk of the Seyfert 2 galaxy, so that the OH 119 $\\mu$m line can be excited through collisions and emit above the continuum in these dense and warm environments. The situation is different for Mkn 231 and Arp 220. In these ULIRGs, the extremely high luminosity arising from the nuclear region, together with the high concentrations of gas there, make the nuclear FIR emission component very bright. Any possible emission in the 119 $\\mu$m line will be obscured by the strong absorption, which in these objects is $\\sim30$ times stronger than the emission feature in NGC 1068. Also, the strong FIR radiation density pumps higher-lying OH and \\hdo\\ levels, thus producing absorptions in the higher-excitation lines. The OH and \\hdo\\ molecules are, therefore, potentially powerful tracers of circumnuclear regions around AGNs. Future Herschel observations of Mkn 231 and other sources will allow us to apply the models developed in this paper to these sources, refine them accordingly, and will certainly give new insights into the physical and chemical conditions of bright infrared galaxies."
    },
    "0711/0711.0621_arXiv.txt": {
        "abstract": "Nuclear stellar clusters are a common phenomenon in spirals and in starbursts galaxies, and they may be a natural consequence of the star formation processes in the central regions of galaxies. HST UV imaging of a few Seyfert 2 galaxies have resolved nuclear starbursts in Seyfert 2 revealing stellar clusters as the main building blocks of the extended emission. However, we do not know whether stellar clusters are always associated with all types of nuclear activity. We present NUV and optical images provided by HST to find out the role that stellar clusters play in different types of AGNs (Seyferts and LLAGNs). Also with these images, we study the circumnuclear dust morphology as a probe of the circumnuclear environment of AGNs. ",
        "introduction": "\\subsection{Starburst-AGN connection: The role of stellar clusters}  Nuclear stellar clusters are a common phenomenon in spirals, having been detected in 50-70\\% of these sources (Carollo et al.  1998, 2002; B{\\\" o}eker et al. 2002, 2004). Therefore stellar clusters are a natural consequence of the star formation processes in the central region of spirals. On the other hand, evidence has been accumulating during the past few years about the ubiquity of black hole (BH) in the nuclei of galaxies. Furthermore, the tight correlation of the BH mass and stellar velocity dispersion (Ferrarese \\& Merrit 2000; Gebhardt et al.\\ 2000) implies that the creation and evolution of a BH is intimately connected to that of the galaxy bulge. Recently, in a HST survey in the Virgo Cluster, C\\^ot\\'e et al. (2006) have detected compact sources in a comparable fraction of elliptical galaxies. These compact stellar clusters, referred to as nuclei by the authors (see also Ferrarese et al 2006a), have masses that scale directly to the galaxy mass, in the same way as do the BH masses in high luminosity galaxies (Ferrarese et al. 2006b). Therefore, a natural consequence of the physical processes that formed present-day galaxies should be the creation of a compact massive object in the nucleus, either a BH and/or a massive stellar cluster.   \\subsection{Starburst-AGN connection: Seyfert galaxies}  According to the unified scheme of AGNs, the main components of a Seyfert nucleus are: 1) A BH and its associated accretion disk. 2) A circumnuclear dusty torus that collimate the AGN radiation through its polar axis. So, a Seyfert 2 nucleus should be a Seyfert 1 that is viewed close to the equatorial plane.  3) A mirror of dust and warm electrons located along the polar axis of the torus, that reflects and polarizes the AGN radiation. Seyfert 2 nuclei exhibit a featureless continuum (FC) that comprises much of the near-UV. It was long-thought that this FC was light from the hidden Seyfert 1 nucleus. Cid Fernandes \\& Terlevich (1995) proposed that a heavily-reddened starburst provides this FC.  Because of the high sensitivity of UV wavelengths to the presence of massive stars, HST UV observations were obtained for a few Seyfert 2 galaxies to probe the role of starbursts in their nuclei and the origin of the FC (Heckman et al. 1997; Gonz\\'alez Delgado et al. 1998). HST high spatial resolution imaging shows that the UV continuum source is spatially extended ($\\sim$ 100 pc) and it is resolved in knots with sizes of a few parcsecs and properties similar to the stellar clusters detected in starburst galaxies. UV  spectra for four galaxies, corresponding to the central $\\sim$ 100 pc, were obtained.  The data provided direct evidence of the existence of a nuclear starburst that accounts for the FC. Absorption features formed in the photospheres and in the stellar winds of massive stars are detected. Their UV colors indicate that the starburst is quite reddened. Their bolometric luminosities are similar to the estimated luminosities of the hidden Seyfert 1 nuclei. Subsequently, near-UV and optical spectra of a large sample of Seyfert 2 were obtained proving unambiguous identification of circumnuclear starburts in $\\sim$40$\\%$ of nearby Seyfert 2 galaxies as well as their energetic significance (Gonz\\'alez Delgado et al. 2001; Cid Fernandes et al. 2001). These results have been latter confirmed  studying a large sample observed by SDSS (Heckman et al. 2004; Wild et al. 2007).  \\subsection{Starburst-AGN connection: LLAGNs}  Low-luminosity active galactic nuclei (LLAGNs) constitute a significant  fraction of the nearby AGN population. These include LINERs, and transition-type objects (TOs, also called weak--[OI] LINERs) whose properties are in between classical LINERs and H{\\sc ii} nuclei.  LLAGNs comprise $\\sim$ 30$\\%$ of all bright galaxies and are the most common type of AGN (Ho, Filippenko \\& Sargent 1997; hereafter HFS97).  What powers them and how they fit in the global picture of AGN has been at the forefront of AGN research for over two decades. Are they all truly ``dwarf Seyfert galaxies'' powered by accretion onto a nearly dormant super-massive black hole, or can some of them be explained at least partly in terms of stellar processes?  HST observations at UV wavelengths of a few LLAGNs have also proven that at least weak-[OI] LINERs could be powered by young massive stars (Maoz et al. 1998; Colina et al. 2002). Nuclear stellar clusters are detected in these objects through high spatial resolution UV imaging and spectra. NGC 4303 is probably the best example (Colina et al. 2002).  Additional evidence comes from optical studies, in which we have focused on the study of the stellar population in the nuclei and circumnuclear region of LLAGNs (Cid Fernandes et al.\\ 2004; Gonz\\'alez Delgado et al.\\ 2004). These studies identified a class of objects, called ``Young-TOs'', which are clearly separated from LINERs in terms of the properties and spatial distribution of the stellar populations. They have stronger stellar population gradients, a luminous intermediate age stellar population concentrated toward the nucleus ($\\sim$100~pc) and much larger amounts of extinction than LINERs. These objects, which underwent a powerful star formation event $\\sim$ 1 Gyr ago, could correspond to post-Starburst nuclei or to evolved counterparts of the Seyfert 2 with a composite nucleus, characterized too by harboring nuclear starbursts.  However, HST UV monitoring observations of 17 LLAGNs have detected variability with amplitudes from a few percent to 50$\\%$ (Maoz et al. 2005). The variability is more frequently detected in those LINERs that have a compact radio core, as expected from bona fide AGNs. Thus, it is not clear which fraction of LLAGNs have a nuclear stellar cluster and which are \"dwarf Seyfert galaxies\" like those that show UV variability.  \\subsection{Goals of the project}  The main goals of the project are to detect and characterize the properties of nuclear and circumnuclear  compact sources in different types of AGNs (type 1 and 2 Seyferts, LINERs and TOs), to  determine their nature and their evolution. The high angular resolution provided by HST is crucial to determine the morphological structure of the nuclear and circumnuclear region of these objects. Due to the steep increase of the surface brightness towards the center of early type galaxies, the near-UV is the ideal wavelength range to isolate the contribution of the nuclear and circumnuclear sources from the bulge component (see Figure 1). Although young stellar clusters are powered by massive stars (O and B), which emit most of their flux at the far-UV, the near-UV  can also trace equally well their distribution (see Figure 1).  In addition, near-UV is more useful than far-UV observations to detect for example clusters with intermediate ages (few 100 Myr), considering the weakness of these populations at shorter wavelengths. Further information at longer wavelengths (e.g. optical and NIR) is also very useful to get colors and to estimate the stellar clusters ages. Additionally high spatial resolution optical images are essential to trace the morphology of the circumnuclear dust which is a valuable probe of the presence of cold interstellar gas in galaxies, and it is very sensitive to the perturbations that drive the gas toward the center and feed the AGN. Thus, the circumnuclear dust can also provide relevant information about the origin of nuclear activity.  To reach these goals we are carrying out several projects with HST+ACS imaging at the near-UV wavelengths a sample of Seyferts (ID. 9379, PI. Schmitt) and LLAGNs (ID. 10548, PI.  Gonz\\'alez Delgado). These observations are complemented with WFPC2 optical data retrieved from the HST archive. Here we present a summary of the the first results obtained for the sample of Seyferts and LLAGN galaxies. Further explanations are in Mu\\~noz-Mar\\'\\i n et al (2007) and Gonz\\'alez Delgado et al (2007). ",
        "conclusions": "The main conclusions obtained for Seyfert sample are: 1) The UV morphology is very irregular, with clumpy and compact structures in most of the cases. 2) Seyfert  1 are very compact and completely PSF dominated; while Seyfert  2 have the nucleus resolved. 3) Circumnuclear stellar clusters  are more common in Seyfert  2 than in Seyfert 1; but the distribution of the flux in stellar clusters does not change much among different Seyfert types, when only galaxies with clusters detected are considered. 4) The contribution of the clusters to the total flux is more important in Seyfert  2 than in other Seyfert types.  The main results obtained for LLAGNs  are: 1) We have not found any correlation between the presence of nuclear compact sources and emission-line type.  Thus, nucleated LINERs are as frequent as nucleated TOs.  2) The nuclei of \"Young-TOs\" are brighter than the nuclei of \"Old-TOs\" and LINERs. These results confirm our previous results that Young-TOs are separated from other LLAGNs classes in terms of their central stellar population properties and brightness. 3) Circumnuclear dust is detected in 88$\\%$ of the LLAGNs, being almost ubiquitous in TOs.  4) The dust morphology is complex and varied, from nuclear spiral lanes to chaotic filaments and nuclear disk-like structures. Chaotic filaments are as frequent as dust spirals; but nuclear disks are mainly seen in LINERs.  These results suggest an evolutionary sequence of the dust in LLAGNs, LINERs being the more evolved systems and Young-TOs the youngest."
    },
    "0711/0711.2304_arXiv.txt": {
        "abstract": "We present a new technique for detecting scattered starlight from transiting, close-orbiting extrasolar giant planets (CEGPs) that has the virtues of simplicity, robustness, linearity, and model-independence. Given a series of stellar spectra obtained over various phases of the planetary orbit, the goal is to measure the strength of the component scattered by the planet relative to the component coming directly from the star. We use two complementary strategies, both of which rely on the predictable Doppler shifts of both components and on combining the results from many spectral lines and many exposures. In the first strategy, we identify segments of the stellar spectrum that are free of direct absorption lines and add them after Doppler-shifting into the planetary frame. In the second strategy, we compare the distribution of equivalent-width ratios of the scattered and direct components. Both strategies are calibrated with a ``null test'' in which scrambled Doppler shifts are applied to the spectral segments. As an illustrative test case, we apply our technique to spectra of HD 209458 taken when the planet was near opposition (with orbital phases ranging from 11 to 34$\\arcdeg$, where 0$\\arcdeg$ is at opposition), finding that the planet-to-star flux ratio is $(1.4\\pm 2.9)\\times10^{-4}$ in the wavelength range 554$-$681~nm. This corresponds to a geometric albedo of $0.8\\pm1.6$, assuming the phase function of a Lambert sphere. Although the result is not statistically significant, the achieved sensitivity and relatively small volume of data upon which it is based are very encouraging for future ground-based spectroscopic studies of scattered light from transiting CEGP systems. ",
        "introduction": "The discovery of the first exoplanet orbiting an ordinary main sequence star, 51 Peg, a dozen years ago \\citep{mayor95} also revealed the first example of a largely unexpected but common class of astronomical objects, namely close-orbiting extrasolar giant planets (CEGPs).  The discovery of the first transiting exoplanet system, HD 209458b also a CEGP \\citep{henry00,mazeh00,charbonneau00,brown01}, a few years later allowed study of these enigmatic objects to move beyond orbital parameters and minimum masses.  Although radial velocity, transit and microlensing techniques have made the study of exoplanets one of the most active and exciting areas of astrophysics in recent years and have produced a wealth of additional major discoveries and important theoretical problems \\citep{perryman00,chauvin07}, the effort to understand CEGPs remains central to the field.  The light of the primary star scattered from CEGPs reflects the nature (composition, clouds, energy balance, structure, motions {\\it et cetera}) of their atmospheres \\citep[e.g.,][]{seager98,marley99,seager00,sudarsky00,sudarsky03,burrows05}; its direct detection would thus test and discriminate among emerging models of CEGPs.  There have been previous attempts to detect scattered starlight from CEGPs against the much brighter direct light from the primary via high-resolution spectroscopic searches for the scattered, and hence Doppler-shifted, stellar absorption lines \\citep[][and references therein]{charbonneau99,leigh03}. The problem is very difficult because the expected contribution of the scattered light to the total light is of order $10^{-4}$ or less for known CEGPs. Due to both the very small amplitude of the expected signal and the unknown orbital inclination angle $i$ for non-transiting systems, the methods so far employed rely on model-dependent assumptions to estimate prior distributions for planetary radii, orbital inclinations and the form of the phase function \\citep{leigh03}.  The advantage of such an approach is that it can be applied to the brightest (but non-transiting) CEGP systems that thus have the least photon noise. In other words and as is usually the case, relying on a prior model allows one to maximize purely statistical signal-to-noise (S/N) but at the price of increased systematic uncertainty. Two previously published methods \\citep[e.g.][]{charbonneau99,cameron99} depend on a $\\chi^2$ minimization applied to the whole stellar spectrum, which is not an optimally sensitive method for this problem because it dilutes the detectability of the scattered lines (which carry the signal of interest) by considering the quality of the fit to the continuum regions.  The simplest example of the advantage of studying transiting CEGP systems is that measurement of the scattered light fraction allows direct determination of the albedo in isolation, and not merely in combination with the planet's radius squared (which must be estimated from a theoretical model for non-transiting systems). The major advantage of studying non-transiting systems is that some of them are nearly two orders of magnitude brighter than the brightest transiting system, thus yielding far smaller photon statistics noise.  Here we propose a largely model-independent technique to detect the scattered stellar absorption lines in transiting CEGP systems. In this method, no assumptions concerning the form of the phase function, the planet's radius, the orbital inclination {\\it et cetera} are needed to measure the planet-to-star flux ratio and albedo.  Rather, only physical parameters of the system measured via transit observations are required. Moreover, the technique maximizes sensitivity by concentrating only on those portions of the spectrum in which scattered stellar lines are known to be present (at some amplitude to be measured).  As a practical ``worked example'' we apply this technique to {\\it Subaru} HDS spectra of HD 209458. The required input physical parameters of the system and their relation to the expected scattered light signal are briefly reviewed in \\S 2. In \\S 3, we describe the technique in detail, and apply it to HD 209458b, using high precision optical spectra acquired on UT 2002 October 27, with the Subaru High Dispersion Spectrograph \\citep[HDS;][]{noguchi02}.  Results are summarized in \\S 4. Prospects for future studies are then examined based on these results. ",
        "conclusions": "We have developed a simple, robust, linear, and model-independent technique to detect scattered starlight from transiting CEGPs, by measuring scattered spectral features at predictable Doppler shifts from the stellar lines due to planets' orbital motions. There are two main variations of this technique. First, spectrum fragments have been aligned and combined over a large set of stellar absorption lines as well as over many exposures taken near the secondary eclipse. Second, the planet-to-star flux ratio has been determined by examining the distribution of the EW ratios between the scattered and stellar absorption lines. We have made a comparison sample by Doppler-shift correcting the exposures with multiple sets of scrambled, instead of the real, redshifts. This comparison sample has been passed through the same data pipeline, and provided the control comparisons, as well as the significance levels of any detections. Jackknife- and bootstrap- re-samplings have yielded consistent errors with those determined from the variations of the reference spectra. We have also produced spectra with simulated scattered signals, using the stellar absorption lines scaled by a factor of an assumed planet-to-star flux ratio. The simulated signals have been fully recovered through the data pipeline.  Physical parameters of the system are required to have been determined, in order to predict the planet's relative radial velocity at each orbital phase, and further localize its imprinted signals on the recorded spectra. These parameters can be determined from radial-velocity and photometric observations of transiting CEGPs with high precision and reliability.  As a case study, we have analyzed very high S/N optical spectra of HD 209458b, acquired with the Subaru HDS in UT 2002 October. Our results suggest an average planet-to-star flux ratio of (1.4$\\pm$2.9)$\\times10^{-4}$, in the wavelength from 554 nm to 681 nm, when the planet was during the orbital phases $11.0^{\\circ} < \\alpha < 33.9^{\\circ}$. Assuming a Lambert-sphere phase function \\citep{pollack86,charbonneau99}, the equivalent geometric albedo can be estimated as $0.8\\pm1.6$, according to Equation \\ref{equ:ratio}, where we take the phase of the middle exposure $\\phi \\approx 0.561$, corresponding to $\\alpha = 22.3^{\\circ}$, out of $11.0^{\\circ} < \\alpha < 33.9^{\\circ}$ for all the exposures we used. The best previous existing constraint on scattered light from HD 209458b is based on direct imaging with the $MOST$ satellite \\citep{rowe06}. \\citet{rowe06} found a 1-$\\sigma$ limit on the planet-to-star flux ratio of $4.88\\times10^{-5}$, corresponding to a geometric albedo upper limit in $MOST$ bandpass (400-700 nm) of 0.25. Our results provide a useful point of comparison for this difficult measurement, covering a  somewhat different wavelength range and based on different observational strategies and analysis techniques.  These results do not seem yet able to support a comparison to models of the HD 209458b atmosphere or to justify detailed physical interpretation, but they do not miss the mark by far and are thus quite encouraging.  Specifically, the Subaru HDS data analyzed here constitute well under 5 hours of total exposure time; it would be possible to obtain more data, and there is no indication in the present data of systematic problems (such as instrumental instabilities) that would prevent further integration from producing a greater sensitivity. Moreover, one significantly more favorable transiting CEGP system (HD 189733) is known, and others may be discovered. In combination with measurements of re-emitted thermal radiation \\citep[e.g.][]{charbonneau05}, scattered light studies offer the possibility of major qualitative advances in our understanding of CEGPs, a central riddle in the study of exoplanets generally."
    },
    "0711/0711.0247_arXiv.txt": {
        "abstract": "Since its discovery in 1996, the source of the bright \\Ha\\ emission (up to 750 mR\\footnote{1~Rayleigh (R) =  $10^6/4\\pi$ photons cm$^{-2}$ s$^{-1}$ sr$^{-1}$, equivalent to $5.7\\times 10^{-18}$ erg cm$^{-2}$ s$^{-1}$ arcsec$^{-2}$ at \\Ha.}) along the Magellanic Stream has remained a mystery. There is no evidence of ionising stars within the \\HI\\ stream, and the extended hot halo is far too tenuous to drive strong shocks into the clouds. We now present a hydrodynamical model that explains the known properties of the \\Ha\\  emission and provides new insights on the lifetime of the Stream clouds. The upstream clouds are gradually disrupted due to their interaction with the hot halo gas. The clouds that follow plough into gas ablated from the upstream clouds, leading to shock ionisation at the leading edges of the downstream clouds. Since the following clouds also experience ablation, and weaker \\Ha\\ (100$-$200 mR) is quite extensive, a disruptive cascade must be operating along much of the Stream. In our model, the clouds are evolving on timescales of 100$-$200~Myr, such that the Stream must be replenished by the Magellanic Clouds at a fairly constant rate. The ablated material falls onto the Galaxy as a warm drizzle which suggests that diffuse ionized gas at 10$^4$K may be an important constituent of galactic accretion. The observed \\Ha\\ emission provides a new constraint on the rate of disruption of the Stream and, consequently, the infall rate of metal-poor gas onto the Galaxy. When the ionized component of the Stream is fully accounted for, the rate of gas accretion is 0.4 M$_\\odot$ yr$^{-1}$, roughly twice the rate deduced from \\HI\\ observations alone. ",
        "introduction": "How do galaxies get their gas? This important question has never been fully answered, either through observation or through numerical simulation. \\HI\\ observations of the nearby universe suggest that galaxy mergers and collisions are an important aspect of this process (Hibbard \\& van Gorkom 1996), but tidal interactions do not guarantee that the gas settles to one or other galaxy. The most spectacular interaction phenomenon is the Magellanic \\HI\\ Stream that trails from the LMC-SMC system (10:1 mass ratio) in orbit about the Galaxy. Since its discovery (Wannier \\& Wrixon 1972; Matthewson et al 1979), there have been repeated attempts to explain the Stream in terms of tidal and/or viscous forces (q.v. Mastropietro et al. 2005; Connors, Kawata \\& Gibson 2005). Indeed, the Stream has become a benchmark against which to judge the credibility of N-body$+$gas codes in explaining gas processes in galaxies. A fully consistent model of the Stream continues to elude even the most sophisticated codes.  Here, we demonstrate that \\Ha\\ detections along the Stream (Weiner \\& Williams 1996; Putman et al 2003) are providing new insights on the present state and evolution of the \\HI\\ gas. At a distance of $D\\approx 55$~kpc, the expected \\Ha\\ signal excited by the cosmic and Galactic UV backgrounds are about 3~mR and 25~mR respectively (Bland-Hawthorn \\& Maloney 1999, 2002), significantly lower than the mean signal of 100$-$200 mR, and much lower than the few bright detections in the range $400-750$ mR (Weiner, Vogel \\& Williams 2002). This signal cannot have a stellar origin since repeated attempts to detect stars along the Stream have failed (e.g. Ostheimer, Majewski \\& Kunkel 1997).  Some of the Stream clouds exhibit compression fronts and head-tail morphologies (Br\\\"uns et al 2005) and this is suggestive of confinement by a tenuous external medium. But the cloud:halo density ratio ($\\eta = \\rho_c/\\rho_h$) necessary for confinement can be orders of magnitude {\\it larger} than that required to achieve shock-induced \\Ha\\ emission (e.g. Quilis \\& Moore 2001). Indeed,  the best estimates of the halo density at the distance of the Stream ($\\rho_h \\sim 10^{-4}$ cm$^{-3}$;  Bregman 2007) are far too tenuous to induce strong \\Ha\\ emission at a cloud face. It is therefore surprising to discover that the brightest \\Ha\\ detections lie at the leading edges of \\HI\\ clouds (Weiner\\etal\\  2002) and thus appear to indicate that shock processes are somehow involved.  We now present a model that goes a long way towards explaining the \\Ha\\ mystery. The basic premise is that a tenuous external medium not only confines clouds, but also disrupts them with the passage of time. The growth time for Kelvin-Helmholtz (KH) instabilities is given by $\\tau_{\\rm KH} \\approx \\lambda \\eta^{0.5}/v_h$ where $\\lambda$ is the wavelength of the growing mode, and $v_h$ is the apparent speed of the halo medium ($v_h\\approx 350$ km s$^{-1}$; see \\S 2). For cloud sizes of order a few kiloparsecs and $\\xi \\approx 10^4$, the KH timescale can be much less than an orbital time ($\\tau_{\\rm MS} \\approx 2\\pi D/v_h \\approx 1$ Gyr). Once an upstream cloud becomes disrupted, the fragments are slowed with respect to the LMC-SMC orbital speed and are subsequently ploughed into by the following clouds. In \\S 2, the new hydrodynamical models are described and the results are presented in \\S 3.  In \\S 4, we discuss the implications of our model and suggest avenues for future research. ",
        "conclusions": "We have seen that the brightest \\Ha\\ emission along the Stream can be understood in terms of shock ionization and heating in a transsonic (low Mach number) flow. For the first time, the Balmer emission (and associated emission lines) provides diagnostic information at any position along the Stream that is independent of the \\HI\\  observations. Slow Balmer-dominated shocks of this kind (e.g. Chevalier \\& Raymond 1978) produce partially ionized media where a significant fraction of the \\Ha\\ emission is due to collisional excitation. This can lead to Balmer decrements (\\Ha/\\Hb\\ ratio) in excess of 4, i.e. significantly enhanced over the pure recombination ratio of about 3, that will be fairly straightforward to verify in the brightest regions of the Stream.  The shock models predict a range of low-ionization emission lines (e.g. OI, SII), some of which will be detectable even though suppressed by the low gas-phase metallicity. There are likely to be EUV absorption-line diagnostics through the shock interfaces revealing more extreme kinematics (Fig. 3, inset), but these detections (e.g. OVI) are only possible towards fortuitous background sources (Sembach\\etal\\ 2001; Bregman 2007). The predicted EUV/x-ray emissivity from the post-shock regions is much too low to be detected in emission.  The characteristic timescale for large changes is roughly 100$-$200 Myr, and so the Stream needs to be replenished by the outer disk of the LMC at a fairly constant rate (e.g. Mastropietro\\etal\\ 2005). The timescale can be extended with larger $\\eta$ values (see equation (1)), but at the expense of substantially diminished \\Ha\\ surface brightness. In this respect, we consider $\\eta$ to be fairly well bounded by observation and theory.  What happens to the gas shedded from the dense clouds? Much of the diffuse gas will become mixed with the hot halo gas suggesting a warm accretion towards the inner Galactic halo. If most of the Stream gas enters the Galaxy via this process, the derived gas accretion rate is $\\sim 0.4$\\Msun\\ yr$^{-1}$. The higher value compared to \\HI\\ (e.g. Peek\\etal\\ 2007) is due to the gas already shredded, not seen by radio telescopes now. In our model, the HVCs observed today are unlikely to have been dislodged from the Stream by the process described here. These may have come from an earlier stage of the LMC-SMC interaction with the outer disk of the Galaxy.  The ``shock cascade'' interpretation for the Stream clears up a nagging uncertainty about the \\Ha\\ distance scale for high-velocity clouds.  Bland-Hawthorn et al (1998) first showed that distance limits to HVCs can be determined from their observed \\Ha\\ strength due to ionization by the Galactic radiation field, now confirmed by clouds with reliable distance brackets from the stellar absorption line technique (Wakker 2001).  HVCs have smaller kinetic energies compared to the Stream clouds, and their interactions with the halo gas are not expected to produce significant shock-induced or mixing layer \\Ha\\ emission, thereby supporting the use of \\Ha\\ as a crude distance indicator.  Here, we have not attempted to reproduce the \\HI\\ observations of the Stream in detail. This is left to a subsequent paper where we explore a larger parameter space and include a more detailed comparison with the \\HI\\ and \\Ha\\ power spectrum, inter alia. We introduce additional physics, in particular, the rotation of the hot halo, a range of Stream orbits through the halo gas, the gravitational field of the Galaxy, and so on.  If we are to arrive at a satisfactory understanding of the Stream interaction with the halo, future deep \\Ha\\ surveys will be essential. It is plausible that current \\Ha\\ observations are still missing a substantial amount of gas, in contrast to the deepest \\HI\\ observations. We can compare the particle column density inferred from \\HI\\ and \\Ha\\ imaging surveys. The limiting \\HI\\ column density is about  $N_H \\approx \\langle n_H \\rangle L \\approx 10^{18}$ cm$^{-2}$ where $\\langle n_H \\rangle$ is the mean atomic hydrogen density, and $L$ is the depth through the slab. By comparison, the \\Ha\\ surface brightness can be expressed as an equivalent emission measure, $E_m \\approx \\langle n_e^2 \\rangle L \\approx \\langle n_e\\rangle N_e$. Here $n_e$ and $N_e$ are the local and column electron density. The limiting value of $E_m$ in \\Ha\\ imaging is about 100 mR, and therefore $N_e \\approx 10^{18}/\\langle n_e \\rangle$ cm$^{-2}$. Whether the ionized and neutral gas are mixed or distinct, we can hide a lot more ionized gas below the imaging threshold for a fixed $L$, particularly if the gas is at low density ($\\langle n_e \\rangle \\ll 0.1$ cm$^{-3}$). A small or variable volume filling factor can complicate this picture but, in general, the ionized gas still wins out because of ionization of low density \\HI\\ by the cosmic UV background (Maloney 1993). In summary, even within the constraints of the cosmic microwave background (see Maloney \\& Bland-Hawthorn 1999), a substantial fraction of the gas can be missed if it occupies a large volume in the form of a low density plasma."
    },
    "0711/0711.4464_arXiv.txt": {
        "abstract": "Models of brown dwarf atmospheres suggest they exhibit complex physical behaviour.  Observations have shown that they are indeed dynamic, displaying small photometric variations over timescales of hours. Here I report results of infrared (0.95--1.64\\,\\micron) spectrophotometric monitoring of four field L and T dwarfs spanning timescales of 0.1--5.5\\,hrs, the goal being to learn more about the physical nature of this variability.  Spectra are analysed differentially with respect to a simultaneously observed reference source in order to remove Earth-atmospheric variations. The variability amplitude detected is typically 2--10\\%, depending on the source and wavelength.  I analyse the data for correlated variations between spectral indices.  This approach is more robust than single band or \\chisq\\ analyses, because it does not assume an amplitude for the (often uncertain) noise level (although the significance test still assumes a shape for the noise power spectrum).  Three of the four targets show significant evidence for correlated variability. Some of this can be associated with specific features including Fe, FeH, VO and \\k1, and there is good evidence for intrinsic variability in \\h2o\\ and possibly also \\ch4.  Yet some of this variability covers a broader spectral range which would be consistent with dust opacity variations.  The underlying common cause is plausibly localized temperature or composition fluctuations caused by convection. Looking at the high signal-to-noise ratio stacked spectra we see many previously identified spectral features of L and T dwarfs, such as \\k1, \\na1, FeH, \\h2o\\ and \\ch4. In particular we may have detected methane absorption at 1.3--1.4\\,\\micron\\ in the L5 dwarf SDSS\\,0539$-$0059. ",
        "introduction": "Numerous observational studies over recent years have revealed photometric variability in old brown dwarfs and very low mass stars with spectral types late M, L and T (collectively ultracool dwarfs, or UCDs) (e.g.\\ Gelino et al.\\ \\cite{gelino02}, Koen et al.\\ \\cite{koen04}). The variability is always of low amplitude (a few tens of millimagnitudes) and often nonperiodic, leading some authors to interpret this as a result of intrinsic atmospheric variability (Bailer-Jones \\& Mundt \\cite{bjm01}) rather than a rotational modulation. The idea here is that, while UCDs are rapid rotators (Mohanty \\& Basri \\cite{mohanty03}, Bailer-Jones \\cite{bj04}) with likely periods of 3--10\\,hrs, heterogeneous variations in opacity on a shorter timescale could account for the observed signature.  In an earlier review of the literature, Bailer-Jones \\cite{bj05} estimated that 40\\% of UCDs are variable, although this may be a lower limit due to limited observations, plus the fact that a given UCD is not always variable. For a more recent review and discussion see Goldman \\cite{goldman05}.  Among young (1--10\\,Myr) low mass stars and brown dwarfs (with mid and late M spectral types) variability is also common (e.g.\\ Joergens et al.~\\cite{joergens03}), with also about $\\sim 40$\\% of objects variable, but this is typically periodic and with larger amplitudes of up to a few tenths of a magnitude (reviewed in Bailer-Jones \\cite{bj05}). This is similar to that seen in (low mass) T Tauri stars, where the variability is related to cool spots.  For the older, cooler objects considered here the physical mechanism behind the variability is less clear, because the predicted neutrality of their atmospheres inhibits formation of magnetically-cooled regions.  Local opacity variations could arise from local temperature or composition variations, one or both of which could be caused from convective flows of matter to/from the deeper, opaque regions of the atmosphere.  To investigate this further, multiband monitoring is required. Bailer-Jones \\cite{bj02} (hereafter BJ02) used synthetic UCD spectra to predict the variability signature due to spots and dust clouds and used these to analyse spectral time series of an L1.5 dwarf. While there was no evidence for variability in a single channel, there was a suggestion of correlated variations which were consistent with dust. Nakajima et al.\\ \\cite{nakajima00} found a suggestion of water variability in a T dwarf, while Clarke et al. \\cite{clarke03} found no variability in the dust-sensitive band heads of TiO, CrH and FeH of the L2V Kelu-1.  Here I report on infrared spectrophotometric monitoring of four UCDs.  This is an extension of a pilot study in BJ02. The objective is to identify which (if any) parts of the spectrum are most variable and, moreover, which regions show significant correlated variability with one another. If the correlated variations can be traced to specific molecules (e.g.\\ water, methane, FeH) or dust-sensitive regions, then not only is this a more robust detection of intrinsic variability (as opposed to instrumental or telluric artefacts), but may provide some clues to the nature of the atmospheric variations. ",
        "conclusions": "I have monitored four UCDs with differential infrared spectrophotometry to look for evidence of correlated variability over timescales between 0.1 and 5.5\\,hrs.  Of the four targets monitored, three show significant evidence for positively correlated variations. In the L2 dwarf SSSPM\\,0828 the most significantly correlated bands lie between 1010 and 1070\\,nm, with the next most significant regions (\\modcor\\,$>$\\,0.8) coinciding with Fe, FeH, VO, \\k1\\ and water. SSSPM\\,0828 also shows similar broad band correlation patterns (and lack thereof) seen in the L1.5V 2M1439 reported by BJ02 (observed from a different site with a different instrument).  The other L2V monitored, 2MASS\\,0847, is devoid of any correlated variability. The T5 dwarf 2MASS\\,0559 shows well correlated (\\modcor\\,$>$\\,0.8) variability at 925--962\\,nm, 1347--1351\\,nm and 1434--1448\\,nm which is probably due to water, possibly also methane. The L5 dwarf SDSS\\,0539 is different from the other three sources in that it exhibits many more regions of correlated variability, larger correlation coefficients and some anticorrelated regions. The latter coincide with TiO and FeH features.  Generally, however, the correlations in SDSS\\,0539 are quite broad band and cannot always be associated with specific features. This would be consistent with variations in the dust opacity.  I have demonstrated that the observed correlated variability is not due to OH lines nor to telluric (water) absorption.  For example, much of the correlated variability in 2MASS\\,0559 and SDSS\\,0539 occurs at wavelength regions where there is little correlated variability in the direct (non-relative) reference star spectra.  In conclusion, this work provides good evidence for correlated variability in three of four ultracool dwarfs.  It is intrinsic to the sources and is not the result of telluric variability or colour-dependent extinction effects (Bailer-Jones \\& Lamm \\cite{bjl}) in the Earth's atmosphere (because the relative fluxes are formed over very narrow bands). The analysis does not require an estimate of the noise level (flux amplitude), although the estimation of confidence intervals does use an assumption on the shape of the noise distribution. This multiband correlation analysis is more robust than single band analyses, because correlated variations are more likely to be intrinsic. This is not to say, however, that only wavelength-correlated variability in UCDs exists. Some variability is seen in narrow UCD features, such as the \\k1\\ lines and FeH, CrH, TiO, but it is not limited to readily-identified spectral features. Much of the variability is broad band in nature, consistent with water and perhaps even methane, but also broader still, as would be expected by dust. Such a variety of elements and chemical phases in the variability signature suggests a common cause. A good candidate is local temperature and/or composition variations. The opacity is very sensitive to temperature, and local variations caused by rising convective columns of hot gas (for example), would change the relative ratios of molecular and gas species ."
    },
    "0711/0711.2148_arXiv.txt": {
        "abstract": "For the investigation of collisions among protoplanetesimal dust aggregates, we performed microgravity experiments in which the impacts of high-porosity mm-sized dust aggregates into 2.5 cm-sized high-porosity dust aggregates can be studied. The dust aggregates consisted either of monodisperse spherical, of quasi-monodisperse irregular or of polydisperse irregular micrometer-sized dust grains and were produced by random ballistic deposition with porosities between 85\\% and 93\\%. Impact velocities ranged from $\\sim 0.1 ~\\rm m~s^{-1}$ to $\\sim 3 ~\\rm m~s^{-1}$ and impact angles were almost randomly distributed. In addition to the smooth surfaces of the target aggregates formed in our experiments, we ``molded'' target aggregates such that the radii of the local surface curvatures corresponded to the projectile radii, decreasing the targets porosities to $80-85$\\%. The experiments showed that impacts into the highest-porosity targets almost always led to sticking, whereas for the less porous dust aggregates, consisting of monodisperse spherical dust grains, the collisions with intermediate velocities and high impact angles resulted in the bouncing of the projectile with a mass transfer from the target to the projectile aggregate. Sticking probabilities for the impacts into the ``molded`` target aggregates were considerably decreased. For the impacts into smooth targets, we measured the depth of intrusion and the crater volume and could derive some interesting dynamical properties which can help to derive a collision model for protoplanetesimal dust aggregates. Future models of the aggregate growth in protoplanetary disks should take into account non-central impacts, impact compression, the influence of the local radius of curvature on the collisional outcome and the possible mass transfer between target and projectile agglomerates in non-sticking collisions. ",
        "introduction": "Young stars are surrounded by accretion disks. It is now widely accepted that the formation of planetesimals, the km-sized precursors of the terrestrial planets and of the cores of the gas planets, is initiated by the process of agglomeration of protoplanetary dust in these disks. For the young Solar System, the circumstellar disk is also known by the name of ''solar nebula``. Due to the decreasing rate of accretion, the gas disk cools down and fine dust particles condense. Typical initial grain sizes, as evidenced by astronomical observations of circumstellar accretion disks \\citep{kesetal06,przetal03}, primitive meteorites from our own Solar System \\citep{ker93}, and recent measurements from comets \\citep{haretal05,keletal05,hoeetal06} are in the (sub-)micrometer size range.  The velocities of these dust grains relative to the gas disk are caused by Brownian motion, gravity-induced drift motion and gas turbulence and result in mutual collisions among the grains \\citep{weicuz93} which can lead to the growth of dust agglomerates. When the collision velocities are sufficiently low, the dust particles show a hit-and-stick behavior \\citep{popetal00,bluwur00,domtie97} and the dust agglomerates are bound by weak van der Waals forces \\citep{heietal99}. As a consequence, fractal dust agglomerates are formed whose fractal dimensions $D_{\\rm f}$, defined by the relation between agglomerate mass $m$ and size $s$, $m \\propto s^{D_{\\rm f}}$, range from $D_{\\rm f} \\approx 1.4 \\ldots 1.5$ \\citep{bluetal00,krablu04,pasdom06} for Brownian motion-dominated growth to $D_{\\rm f} \\approx 1.8 \\ldots 1.9$ for growth caused by drift motions \\citep{bluetal98} or gas turbulence \\citep{wurblu98}. Due to the increasing agglomerate mass and the consequentially increasing collision velocities for protoplanetary dust \\citep{weicuz93}, collisions will eventually lead to the compaction of the agglomerates \\citep{domtie97,bluwur00} so that the fractal dimensions will increase to $D_{\\rm f} = 3$ \\citep{weicuz93,blu04}. Such agglomerates are, however, not compact but can be quite porous. Simulation experiments suggest that the volume filling factors, given by \\begin{equation}\\label{vff} \\phi = 1 - p =\\frac{\\rho}{\\rho_0}~, \\end{equation} with $p$, $\\rho$ and $\\rho_0$ being the porosity, the mass density of the aggregate and of the solid grain material, respectively, should be as low as $\\phi \\le 0.20 \\ldots 0.33$ even if the collision velocities are as high as 50 $\\rm m~s^{-1}$ (\\citet{blusch04}; \\citet{bluetal06}, hereafter referred to as paper I). As long as the collision velocities are $\\lesssim 1 \\rm m~s^{-1}$, the maximum volume filling factor for loose particle ensembles is $\\phi = 0.07 \\ldots 0.15$ (\\citet{blusch04}; paper I). Very low collision velocities could lead to even lower volume filling factors of $\\phi \\approx 0.01$ \\citep{ormetal07}.  For fractal dust agglomerates, concurring theoretical \\citep{domtie97} and experimental \\citep{bluwur00} results show that above a threshold velocity of $\\sim 1 \\rm \\, m~s^{-1}$, dust agglomerates do no longer stick together but bounce off and fragment upon collision. For non-fractal dust aggregates, impact experiments by \\citet{wuretal05a} and \\citet{wuretal05b}, who used cm-sized porous projectiles and decimeter-sized porous targets, showed that for high-porosity targets ($\\phi = 0.12 \\ldots 0.26$) no sticking occurs in the velocity range $16.5 \\ldots 37.5 \\, \\rm m~s^{-1}$ and a crater is formed on the target which leads to a mass loss from the target. For compacted targets ($\\phi = 0.34$), a mass gain of the target of $\\sim 50$ percent of the projectile mass could be found for impact velocities $\\ge 13 \\, \\rm m~s^{-1}$.  Not much is known about the collision behavior of high-porosity dust agglomerate in the velocity regime around $1 \\, \\rm m~s^{-1}$. Earlier experiments by \\citet{blumue93}, who used mm-sized dust agglomerates with $\\phi = 0.26$ and collision velocities in the range $0.15 \\ldots 3.8 \\, \\rm m~s^{-1}$, showed that none of the collisions between (almost) equal-sized agglomerates led to sticking, while at the highest collision velocities fragmentation dominated. With our new technology for the formation of high-porosity macroscopic dust agglomerates (\\citet{blusch04}; paper I), realistic macroscopic dust agglomerates with volume filling factors in the range $\\phi = 0.07 \\ldots 0.15$ are now available with which the collision behavior of protoplanetary dust agglomerates can be investigated.  Sect. \\ref{sectobj} gives an overview of the scientific objectives of this work, Sect. \\ref{sectsetup} presents the experimental technology and the properties of the used dust agglomerates, Sect. \\ref{sectexpres} gives the experimental results on sticking efficiencies, mass transfer, coefficients of restitution, energy loss, tensile strength, crater formation and impact fragmentation. Sect. \\ref{sectdisc} discusses the low-velocity impact dynamics of high-porosity dust agglomerates, Sect. \\ref{sectsummary} gives a summary of our experimental findings, and in Sect. \\ref{sectconclusion}, we draw conclusions from our work for the formation of larger objects in the protoplanetary nebula. ",
        "conclusions": "Conclusions} In paper I, we have shown that macroscopic protoplanetary dust aggregates are expected to be very porous. Depending on the collisional history and the size of the aggregates, we expect the dusty objects to have volume filling factors between $\\phi \\approx 0.1$ (for all objects whose collision velocities never exceeded $\\sim 1~\\rm m s^{-1}$, i.e. for sizes $\\lesssim$ cm) and $\\phi \\approx 0.3$ (for all objects with sizes $\\gtrsim 1$ m). In this paper, we confirmed that collisions among fluffy protoplanetary dust aggregates with velocities $\\sim 1~\\rm m s^{-1}$ lead to impact pressures of $\\lesssim 2,000$ Pa and, thus, to a moderate increase in volume filling factor. Collisions between (sub-)mm-sized dust aggregates and cm- to dm-sized fluffy objects are very abundant in the solar nebula and the impact velocities range around $1~\\rm m s^{-1}$ \\citep{weicuz93}. Thus, our microgravity experiments match the solar-nebula conditions very closely and are directly applicable to growth models.  We find that the sub-mm sized dust aggregates always stick to much larger target aggregates, independent of impact velocity and impact angle. The larger, mm-sized aggregates, however, behave differently. Sticking is restricted to the higher velocities and to the smaller impact angles. This trend was found for dust aggregates consisting of monodisperse spherical particles, for quasi-monodisperse irregular particles as well as for irregular monomer particles with a wide size distribution, although quantitative differences in the sticking probabilities exist. Thus, for random impacts in the solar nebula, the sticking probability in collisions between mm-sized dust aggregates and cm-dm sized dusty bodies is below unity. Moreover, non-sticking (and even fragmentation) is favored when the local radius of curvature at the point of impact of the larger body is similar to the projectile radius. This can be the case for collisions between similar-size dust aggregates or between projectiles and targets with irregular, non-flat surface textures. Another interesting feature we found in our experiments is that non-sticking impacts (i.e. those with higher impact angles) lead to a mass transfer from the larger to the smaller body.  As a consequence, future growth models for protoplanetary dust should take into account that the outcome of a single collision between dust aggregate A and dust aggregate B is not only dependent on the size of A and B and the mutual collision velocity, but is also influenced by the collisional history (e.g. the distribution of local radii of curvature on the surfaces of A and B; compaction) and the (random) impact angle of the particular collision. In addition to that, due to the possible occurrence of mass transfer from a larger body A to a smaller body B (in bouncing collisions) or fragmentation of the smaller object B (e.g. when it hits a surface part of A with a small local radius of curvature), the numerical description of the collisional outcome (and its use in e.g. Smoluchowski's growth equation) becomes rather complex.  From our experimental findings, we consider it rather unlikely that protoplanetary bodies can grow beyond dm-sizes in a direct and simple hit-and-stick manner. Although our experiments suggest that small projectiles stick at higher velocities than large projectiles, this trend does not imply that km-sized bodies can form by the accumulation of very small dust aggregates or single particles. Recent experiments by Schr\\\"apler \\& Blum (unpublished) show that at impact velocities $\\gtrsim 15~\\rm m~s^{-1}$ dust aggregates cannot grow by the accumulation of single grains (and, therefore, also not by the accumulation of small dust aggregates). However, m-sized bodies in protoplanetary disks possess relative velocities with respect to small grains of the order of $50~\\rm m~s^{-1}$ \\citep{weicuz93}. Moreover, experiments by \\citet{wuretal05a} suggest that impacts between dusty projectiles into fluffy dusty targets around $\\sim 10~\\rm m s^{-1}$ never lead to an accumulation of mass on the larger body but lead to strong fragmentation of the projectile and cratering (i.e. mass loss) of the target. A possible way out of this dilemma could be the indirect effect of projectile or fragment capturing by aerodynamic \\citep{wuretal01a,wuretal01b} or electrostatic \\citep{blu04} forces or by gravitational collapse in locally overdense regions in the midplane of protoplanetary disks \\citep{johetal07}."
    },
    "0711/0711.2979_arXiv.txt": {
        "abstract": "We present Giant Meterwave Radio Telescope (GMRT) and Westerbork Synthesis Radio Telescope (WSRT) observations of the recently discovered Local Group dwarf galaxy, Leo~T. The peak \\HI\\ column density is measured to be $7\\times10^{20}$ \\cm, and the total \\HI\\ mass is $2.8\\times10^5$ \\msun, based on a distance of 420 kpc. Leo~T has both cold ($\\sim500$ K) and warm ($\\sim 6000$ K) \\HI\\ at its core, with a global velocity dispersion of 6.9 \\kms, from which we derive a dynamical mass within the \\HI\\ radius of $3.3\\times10^6$ \\msun, and a mass-to-light ratio of greater than 50. We calculate the Jeans mass from the radial profiles of the \\HI\\ column density and velocity dispersion, and predict that the gas should be globally stable against star formation. This finding is inconsistent with the half light radius of Leo~T, which extends to 170 pc, and indicates that local conditions must determine where star formation takes place. Leo~T is not only the lowest luminosity galaxy with on-going star formation discovered to date, it is also the most dark matter dominated, gas-rich dwarf in the Local Group. ",
        "introduction": "\\label{sec:intro}  Leo~T is an impressively small, yet complex dwarf galaxy. A member of the Local Group, Leo~T is the lowest luminosity galaxy discovered to date with on-going star formation (Irwin et al. \\citeyear{Irwin07}). Its colour-magnitude diagram reveals both a red giant branch and young blue stars, $\\sim6-8$ Gyrs and $\\sim200$ Myr in age, respectively. Although its stellar morphology and intermediate-aged red stars are similar to the dwarf spheroidal (dSph) galaxies, many of which have been discovered recently (Belokurov et al. \\citeyear{Belokurov07}, and references within), the presence of a younger, blue stellar population is more typical of a dwarf irregular (dIrr) galaxy. This duality has lead to the 'transition' label, hence the name Leo~T. In addition, dSph galaxies are usually found within 250 kpc of Milky Way, whereas Leo~T is located at a distance of 420 kpc, similar to that of other transitional dwarfs, such as Phoenix, and dIrr galaxies \\citep{Grebel00}. The presence of cool gas is another trademark of transitional dwarfs and dIrrs. As reported in Irwin et al. (\\citeyear{Irwin07}), Leo~T has a spatially coincident detection of \\HI\\ in the Northern \\HI\\ Parkes All Sky Survey (HIPASS, Wong et al. \\citeyear{Wong06}). Recent optical spectroscopy has confirmed that the stellar recessional velocity matches the \\HI\\ radial velocity measurement \\citep{Simon07}.  The study of the smallest dwarf galaxies provides insight into how the least massive dark matter haloes retain cool gas and form stars. The number of dark haloes in the Local Group predicted by cosmological simulations is typically of the order of hundreds \\citep{Klypin99, Moore99}. Only those dark haloes that can maintain sufficient cool gas, allowing star formation to proceed, produce luminous dwarf galaxies. The various processes that are thought to suppress star formation in low mass haloes include a change in the Jeans mass due to global reionization \\citep[e.g.,][]{Efstathiou92, Benson02, Cooray05}, the heating and removal of gas via supernovae and stellar wind feedback \\citep[e.g.][]{Dekel03, Ricotti05}, ram pressure and tidal stripping \\citep[e.g.][]{Blitz00}, or simply a temperature floor in the interstellar medium \\citep{Kaufmann07}. Mass-to-light measurements of Local Group dwarfs suggest that each galaxy is embedded in a dark matter halo with a mass of about $10^7$ \\msun\\ \\citep{Mateo98}. The idea that there is a minimum dark matter halo mass able to form stars is supported by analytic calculations \\citep{Taylor05}, observational data \\citep{Gilmore07} and numerical simulations \\citep{Ricotti05, Read06}. Thus observations of the presence, morphology and kinematics of cool gas in dwarf galaxies provide important constraints on these predictions and hold ramifications for our understanding of galaxy formation and evolution. Leo~T is a crucial piece of evidence in this line of inquiry, as it is the faintest dwarf galaxy detected to date with on-going star formation.   In this letter we present Giant Meterwave Radio Telescope (GMRT) and Westerbork Synthesis Radio Telescope (WSRT) observations that confirm the presence of \\HI\\ in Leo~T at higher spatial and velocity resolution than the initial detection in HIPASS. In section~3 we give the \\HI\\ parameters of Leo~T, and calculate the Jeans mass profile. A discussion of how star formation has proceeded in Leo~T and a comparison to other Local Group dwarfs is given in section~4. ",
        "conclusions": "Given its stellar mass of $\\sim1.2\\times10^{5}$\\msun\\ and age of $\\sim 6-8$ Gyrs, on average, Leo~T has been forming stars at the slow rate of $1.5-2\\times10^{-5}$\\msun\\ per year. Evidently this gentle star formation rate has neither heated nor blown out all the gas in Leo~T, allowing the stellar mass to build. A reservoir of approximately $10^5$\\msun\\ of cool gas ($\\sigma_v=2$ \\kms, $T\\sim 500 K$) remains, which would take about $5\\times10^{9}$ years to exhaust at the average past star formation rate. Since Leo~T appears to be on such a gentle trajectory with respect to the Galactic Standard of Rest (-58 \\kms) and Local Group (-97 \\kms), it is possible that its gas has avoided being stripped or heated by an interaction. Such a low central velocity dispersion ($\\sigma_v=2$ \\kms) may be the reason why this galaxy has been able to form stars despite its peak \\HI\\ column density ($7\\times10^{20}$\\cm) being lower than those of typical low surface brightness dwarf galaxies \\citep{vanzee97}.  The Plummer law fit gives a stellar radius of 1.4\\arcmin\\ (Irwin et al. \\citeyear{Irwin07}), corresponding to a physical length of 170 pc. An excess of stars above the background is detected out to a radius of 5\\arcmin. Visual inspection of the image in Figures~\\ref{fig:GMRTolay} and \\ref{fig:WSRTolay} shows that the over density of blue stars lies within 1\\arcmin, or 120 pc of the centre. The Jeans mass analysis suggests that the gas is stable against collapse, contrary to what is observed. This prediction is based the assumption of a spherical distribution and radially averaged \\HI\\ column density and velocity dispersion. Flattening the density distribution causes the Jeans mass decrease. Since Leo~T appears to be stable against star formation globally, it must be local processes that govern the star formation. The tendency for locations with narrow velocity components to coincide with \\HI\\ column densities above $3\\times 10^{20}$ \\cm\\ supports this idea. Furthermore, a small level of star formation may increase the turbulence, inhibiting the global collapse, but promoting local instabilities via density fluctuations \\citep{Schaye04}.  It would be interesting to repeat the radial Jeans mass analysis on compact high velocity clouds (HVCs) to see whether their \\HI\\ properties suggest that star formation is possible. The analysis would be straightforward since column density and velocity dispersion measurement are independent of distance, an issue that has plagued the interpretation of HVCs. Although Leo~T is more compact than any of the compact HVCs for which high resolution imaging is available \\citep{Braun00}, their peak column densities ($\\sim4\\times10^{20}$\\cm) and minimum velocity dispersions ($<1$ \\kms) are similar to those of Leo~T. The initial detection of an \\HI\\ cloud at the location of Leo~T in the HIPASS data was indistinguishable from an HVC. It is only the high resolution \\HI\\ imaging presented here, together with a recessional velocity measurement of the stars in Leo~T that has lead to the confirmation that the \\HI\\ detected is of Leo~T itself. Given the similarity in \\HI\\ characteristics, it is quite possible that compact HVCs represent a population of 'failed' galaxies that contain gas and no stars, with masses akin to that of Leo~T. Confirming this suggestion, however, will be quite difficult, as distance measurements rely on the chance alignment with both a foreground and background object, and the sky area covered by compact HVCs is small.  Compared with other dwarf galaxies in the Local Group, Leo~T contains a significant component of dark matter given its large gas fraction. In Figure~\\ref{fig:mass}, dIrr galaxies tend to populate the top left of the diagram, characterized by high gas fractions (\\mhi/$L_B>0.1$) and typical total mass-to-light ratios less than 10. On the other hand, the lower half of the diagram is occupied by dSph galaxies, which have low gas fractions (mostly upper limits as no \\HI\\ is detected in many cases) and total mass-to-light ratios ranging from 1 to 100. The three other galaxies classified as transitional dwarfs have properties that lie in between the dIrr and dSph galaxies in Figure~\\ref{fig:mass}. Leo~T appears to stand alone in the top right of diagram, it is significantly more gas rich and also contains more dark matter than other transitional dwarfs.  The number and nature of Local Group dwarf galaxies provide important constraints on models of galaxy formation and evolution. How many other extremely low-luminosity, gas-rich galaxies are yet to be found in the Local Group?  The SDSS (Sloan Digital Sky Survey) has surveyed one fifth of the sky, and has uncovered just one such gas-rich dwarf, namely Leo~T. The completeness for similar objects over the SDSS survey area is close to one (Koposov et al. \\citeyear{Koposov07}), thus only a handful of other discoveries of comparable luminosity are expected in other directions, assuming an isotropic distribution. The other avenue to explore is detection via the \\HI\\ emission line. Compared with HIPASS, current surveys, such as the Arecibo Galactic \\HI\\ Survey (Stanimirovic et al. \\citeyear{Stanimirovic06}) will provide four-times better spatial resolution with 0.2 \\kms\\ velocity resolution, and the Galactic All Sky Survey \\citep[GASS,][]{McClure-Griffiths06} will provide a 20-fold improvement in velocity resolution. Disentangling extragalactic \\HI\\ (with a chaotic velocity field) from emission associated with the Galaxy at low Galactocentric velocities will always remain an issue; the only definitive test is a matching velocity for the stellar component of the object.    We have confirmed the presence of $2.8\\times10^5$ \\msun\\ of \\HI\\ in the Local Group dwarf galaxy, Leo~T. The gas is essentially centred on the stellar emission with a radial extent of 300 pc. The \\HI\\ consists of both a cold and warm neutral medium, revealed by a double Gaussian fit with $\\sigma_v=2,7$ \\kms\\ to pixels in the central region; the velocity dispersion of the global profile is 6.9 \\kms. The comparatively low stellar mass gives Leo~T a high gas fraction of 80 per cent. The \\HI\\ extent and velocity dispersion of gas have been used to estimate a total dynamical mass for Leo~T, M$_{\\rm dynamical}\\gtrsim 3.3\\times10^6$ \\msun. This value for the total mass is just lower than the fiducial $10^7$\\msun, below which galaxies are expected to be dark \\citep[e.g.,][]{Read06}, due to supernova feedback and reionization.  Although the simulations predict that most gas-rich dwarf galaxies form in haloes with M $>10^8$\\msun, there are some rare examples of galaxies in simulations that have a dark matter halo mass of $\\sim10^7$\\msun and a small baryon fraction that do form at least some stars, albeit inefficiently \\citep{Ricotti05} -- a description which fits Leo~T. It is interesting to note that Leo~T has a significant fraction of dark matter; given that most gas-rich dwarfs (i.e. dIrr) have M$_{\\rm dyn}$/L$_V$ values between 1 and 10, the total mass of Leo~T would be expected be about $10^5$\\msun, yet a mass in excess of $10^{6}$\\msun\\ has been measured, thus upholding the idea of a minimum dark matter halo mass for dwarf galaxies \\citep{Mateo98, Gilmore07}. We have compared the cumulative dynamical and Jeans masses to determine that the \\HI\\ in Leo~T is {\\em{globally}} stable against star formation, this is inconsistent with the observed presence of young stars. From this we conclude that local rather than global processes must be responsible for Leo~T's stellar component. The very low past average star formation rate may be the reason why such cold gas is able to reside in Leo~T."
    },
    "0711/0711.0467_arXiv.txt": {
        "abstract": "Massive high-redshift galaxies form in over-dense regions where the probability of forming other galaxies is also strongly enhanced.  Given an observed flux of a galaxy, the inferred mass of its host halo tends to be larger as its inferred redshift increases.  As the mass and redshift of a galaxy halo increase, the expected clustering of other galaxies around it gets stronger.  It is therefore possible to verify the high-redshift identity of a galaxy (prior to an unambiguous spectral identification) from the clustering of other galaxies around it.  We illustrate this method for the massive galaxy suggested by \\citet{Mobasher05} to be at redshift $z \\! \\sim \\!6.5$.  If this galaxy were to exist at $z \\! \\sim \\! 6.5$, there should have been a mean of $\\sim \\! 10$ galaxies larger than a hundredth of its mass and having z-band magnitudes less than $\\sim \\!25$ detected as i-dropouts in the HUDF.  We calculate an approximate probability distribution for neighbor galaxies and determine that there is less than a $\\sim \\! 0.3\\%$ chance of detecting no massive neighbor galaxies.  The lack of other massive $z \\! \\sim \\! 6.5$ galaxies in the HUDF image argues that the \\citet{Mobasher05} galaxy is instead a low redshift interloper.  We generalize our results to other galaxy masses and redshifts. ",
        "introduction": "The purported detection of a very massive galaxy, HUDF-JD2, at a redshift of $z \\! \\sim \\! 6.5$ by \\citet{Mobasher05} provides a critical test for the standard paradigm of galaxy formation in a $\\Lambda$CDM cosmology.  However, the lack of an unambiguous spectral identification allows this galaxy to be a $z \\!\\sim \\! 2$ interloper in which obscuration by dust mimics the Lyman break \\citep{Dunlop07,Chary07}.  If the galaxy is at $z \\!\\sim\\! 6.5$, the estimated mass of the halo containing this galaxy is $M_{JD2}\\!\\sim\\!2\\times10^{13}\\,M_{\\odot}$ for a reasonable star formation efficiency, $f_{\\star}$, of $10\\%$.  In the concordance cosmological model, the expected number of such galaxies in the Hubble Ultra Deep Field (HUDF) field-of-view is less than $10^{-6}$ \\citep{BL06}.  Since the expected probability of finding such a galaxy in the HUDF is extremely low, its redshift identification is of great importance for testing the standard cosmological model.  Here we point out that, despite the low average abundance of galaxies like JD2 (due to its large mass and high redshift), such a galaxy cannot exist alone.  The large-scale over-density implied by its existence naturally results in neighboring halos over and above what would be expected from random fluctuations in the average galaxy population.  The surrounding over-dense region behaves as if it is part of a closed universe, in which the formation of all galaxies occurs earlier.  We approach this problem analytically in this {\\it Paper}.  In addition to the deeper fundamental understanding gained from such a treatment, the extreme rarity of objects like JD2 make a statistical analysis using numerical simulations difficult. In \\S \\ref{Met}, we show how the excursion set formalism can be used to calculate an approximate probability distribution for the number of neighbors around massive galaxies.  We then calculate, in particular, the expected clustering of bright galaxies around the \\citet{Mobasher05} galaxy in \\S\\ref{JD2}, and explore the dependence of our results on the star formation efficiency, duty cycle, and power-spectrum normalization in \\S\\ref{param}.  In \\S\\ref{future}, we generalize these results to other halo masses and redshifts.  Finally, \\S\\ref{discussion} summarizes our main conclusions.  Unless otherwise noted, we assume a flat, $\\Lambda$CDM model for the universe with the {\\it WMAP3} cosmological parameters \\citep{Spergel07}. ",
        "conclusions": "\\label{discussion}  We have presented a clear signature for the existence of a massive galaxy at high redshift.  Due to large correlations among massive halos at high redshifts, once such an object is found, it is unlikely to be alone.  On the contrary, the more massive the halo and the higher its redshift, the greater the number of neighbors it has.  This signature is helpful in situations where only photometric data is available or when the spectroscopic redshift identification is ambiguous. For a galaxy of a given observed flux, the inferred halo mass increases as its suggested redshift increases.  It therefore becomes easier to rule out a higher-redshift hypothesis, using dropout techniques to locate the neighboring LBGs that are expected at a similar redshift.  We derived a probability distribution for the number of neighbors around a massive galaxy at high redshift, which led to calculations of the mean number of such neighbors and the probability of detecting no such halos.  In the case of JD2, the number of neighbor galaxies observed in this way can distinguish between the $z\\! \\sim \\!6.5$ interpretation of \\citet{Mobasher05} and that of a galaxy at $z\\! \\sim \\!2$ \\citep{Dunlop07, Chary07}.  If at $z=6.5$, the mean number of excess i-dropouts in the HUDF with z-band magnitudes less than $\\sim\\!25$ is predicted to be $\\sim\\! 10$.  The lack of such bright i-dropouts in the HUDF \\citep{Bouwens06, Bouwens07} implies that JD2 cannot be at such high redshift with $99.7\\%$ confidence.  Thus, future detections of massive galaxies at high redshift could be verified by looking for bright, massive neighbors.  Additional research into exactly how LBGs form in dark matter halos and how biased they are compared to these halos would firm-up the predictions for the probability distribution of neighbors.  Such work would undoubtedly include various feedback processes and environmental effects on LBG formation.  Previous study indicates that LBGs are more clustered than dark matter halos \\citep{ST03}, and our results for the number of neighbors are, thus, lower limits on the true values."
    },
    "0711/0711.4558_arXiv.txt": {
        "abstract": "We have studied extreme mass-ratio inspirals (EMRIs) in spacetimes containing a rotating black hole and a non self-gravitating torus with a constant distribution of specific angular momentum. We have found that the dissipative effect of the hydrodynamic drag exerted by the torus on the satellite is much smaller than the corresponding one due to radiation reaction, for systems such as those generically expected in Active Galactic Nuclei and at distances from the central supermassive black hole (SMBH) which can be probed with LISA. However, given the uncertainty on the parameters of these systems, namely on the masses of the SMBH and of the torus, as well as on its size, there exist configurations in which the effect of the hydrodynamic drag on the orbital evolution can be comparable to the radiation-reaction one in phases of the inspiral which are detectable by LISA.  This is the case, for instance, for a $10^6M_{\\odot}$ SMBH surrounded by a corotating torus of comparable mass and with radius of $10^3-10^4$ gravitational radii, or for a $10^5 M_{\\odot}$ SMBH surrounded by a corotating $10^4 M_{\\odot}$ torus with radius of $10^5$ gravitational radii. Should these conditions be met in astrophysical systems, EMRI-gravitational waves could provide a characteristic signature of the presence of the torus. In fact, while radiation reaction always increases the inclination of the orbit with respect to the equatorial plane (\\textit{i.e.,} orbits evolve towards the equatorial retrograde configuration), the hydrodynamic drag from a torus corotating with the SMBH always decreases it (\\textit{i.e.,} orbits evolve towards the equatorial prograde configuration). However, even when initially dominating over radiation reaction, the influence of the hydrodynamic drag decays very rapidly as the satellite moves into the very strong-field region of the SMBH (\\textit{i.e.,} $p\\lesssim5M$), thus allowing one to use pure-Kerr templates for the last part of the inspiral. Although our results have been obtained for a specific class of tori, we argue that they will be qualitatively valid also for more generic distributions of the specific angular momentum. ",
        "introduction": "Introduction}  One of the most exciting prospects opened up by the scheduled launch of the space-based gravitational-wave detector LISA~\\cite{LISA} will be the possibility of mapping accurately the spacetime of the supermassive black holes (SMBHs) which are believed to reside in the center of galaxies~\\cite{SMBH}. Among the best candidate sources for this detector there are Extreme Mass Ratio Inspirals (EMRIs), \\textit{i.e.} stellar-mass black holes ($m\\approx1-10M_\\odot$) or compact objects orbiting around the SMBH and slowly inspiraling due to the loss of energy and angular momentum via gravitational waves (radiation reaction).  In order for the signal to fall within the sensitivity band of LISA, the SMBH must have a mass $M\\approx10^5-10^7M_\\odot$, \\textit{i.e.}, the low end of the SMBH mass function.  It is currently expected that a number of such events ranging from tens to perhaps one thousand could be measured every year \\cite{gair_event_rates}, but since they will have small signal-to-noise ratios, their detection and subsequent parameter extraction will require the use of matched-filtering techniques.  These basically consist of cross-correlating the incoming gravitational-wave signal with a bank of theoretical templates representing the expected signal as a function of the parameters of the source.  This will not only allow one to detect the source, but also to extract its properties. For instance, the accurate modelling of the motion of a satellite in a Kerr spacetime will allow one to measure the spin and the mass of the SMBH.  Although producing these pure-Kerr templates has proved to be a formidable task, particularly because of the difficulty of treating rigorously the effect of radiation reaction (see Ref.~\\cite{Poisson_review} for a detailed review), considerable effort has gone into trying to include the effects of a deviation from the Kerr geometry. These attempts are motivated by the fact that possible ``exotic'' alternatives to SMBHs have been proposed (\\textit{e.g.,} boson stars~\\cite{monica}, fermion balls \\cite{fermion_balls} and gravastars~\\cite{gravastars}), although the presence of these objects would require to modify radically the mechanism with which galaxies are expected to form. On the other hand, non-pure Kerr templates might allow one to really map the spacetimes of SMBHs and to test, \\textit{experimentally}, the Kerr solution.  Different approaches to this problem have been considered in the literature. EMRIs in a spacetime having arbitrary gravitational multipoles should be considered in order to maintain full generality \\cite{ryan}, but this method does not work very well in practice and would only apply to vacuum spacetimes.  For this reason, alternative approaches have been proposed and range from EMRIs around non-rotating boson stars~\\cite{gair_boson}, to EMRIs in bumpy black-hole spacetimes \\cite{bumpyBH} (\\textit{i.e.}, spacetimes which are \\textit{almost} Schwarzschild and require naked singularities or exotic matter) or in quasi-Kerr spacetimes~\\cite{quasi_kerr} (\\textit{i.e.}, spacetimes, consisting of Kerr plus a small quadrupole moment).  Interestingly, none of these methods is suitable for taking into account the effect of the matter which is certainly present in galactic centers. SMBHs can indeed be surrounded by stellar disks (as in the case of the Galactic center~\\cite{stellar_disks}) or, as in the case of Active Galactic Nuclei (AGNs)~\\cite{agn}, in which we are most interested, by accretion disks of gas and dust which can be even as massive as the SMBH~\\cite{Hure}. While the gravitational attraction of a disk can have important effects on EMRIs if this disk is very massive and close to the SMBH~\\cite{BHTorus}, an astrophysically realistic accretion disk can influence an EMRI only if the satellite crosses it, thus experiencing a ``hydrodynamic'' drag force.  This drag consists of two parts. The first one is due to the accretion of matter onto the satellite black hole (this was studied analytically by Bondi \\& Hoyle~\\cite{bondi} and subsequently confirmed through numerical calculations~\\cite{petrich,font1,font2}). This transfers energy and momentum from the disk to the satellite, giving rise to a short-range interaction.  The second one is instead due the gravitational deflection of the material which is not accreted, which is therefore far from the satellite, but which can nevertheless transfer momentum to it. This long-range interaction can also be thought of as arising from the gravitational pull of the satellite by its own gravitationally-induced wake (\\textit{i.e.,} the density perturbations that the satellite excites, by gravitational interaction, in the medium), and is often referred to as ``dynamical friction''. This effect was first studied in a collisionless medium by Chandrasekar~\\cite{chandra}, but acts also for a satellite moving in a collisional fluid \\cite{rephaeli,petrich,ostriker,sanchez,kim,my_dyn_frict}.  The effect of this disk-satellite interaction on EMRIs has been studied by different authors for a number of disk models. More specifically, in a series of papers Karas, Subr and Vokrouhlicky considered the interaction between stellar satellites and thin disks \\cite{karas_subr0,karas_subr1,karas_subr2}. In Ref.~\\cite{karas_subr2}, in particular, Subr and Karas found that the effect of the star-disk interaction on EMRIs dominates over radiation reaction for thin disks, both for non-equatorial orbits crossing the disk only twice per revolution and for equatorial orbits embedded in the disk. The only exceptions to these conclusions come if the satellite is very compact (a neutron star or a black hole) or the disk has a low density (\\textit{e.g.,} in the region close to the SMBH if the flow becomes advection-dominated). These results agree with those found by Narayan~\\cite{narayan}, who focused on Advection Dominated Accretions Flows (ADAFs), which were believed to describe accretion onto ``normal'' galactic nuclei (\\textit{i.e.,} ones much dimmer than AGNs)\\footnote{Accretion onto ``normal''  galactic nuclei is now believed to be better described by Advection Dominated Inflow Outflow Solutions (ADIOS)~\\cite{ADIOS}. However, this is not expected to change significantly Narayan's results since ADIOS's, like ADAFs, have very low densities in the vicinity of the SMBH.}. Overall, he found that for compact objects and white dwarfs the effect of the hydrodynamic drag is negligible with respect to radiation reaction, whereas it is not negligible for main sequence and giant stars. More recently, Levin~\\cite{levin} has proposed a scenario in which massive stars form in a thin accretion disk in an AGN, ultimately producing stellar-mass black holes embedded in the disk. The small black holes are then dragged towards the (non-rotating) SMBH, but if this is accreting at a rate comparable to the Eddington limit, the drag from the accreting gas will not affect the final part of the inspiral (\\textit{i.e.,} at radii smaller than $10$ Schwarzschild radii) significantly. Finally, Chakrabarti~\\cite{chakra}, studied instead the orbital evolution of a satellite black hole on a circular equatorial orbit embedded in a disk with a non-Keplerian distribution of angular momentum, and found that the exchange of angular momentum between the disk and the satellite can lead to significant orbital modifications.  All of these studies have been been carried out within a Newtonian or pseudo-Newtonian description of gravity (with the partial exception of Ref.~\\cite{karas_subr0}, in which the orbits are Kerr geodesics, but the disk model and the hydrodynamic drag is not relativistic). In this paper, instead, we provide a first relativistic treatment for satellite  black holes moving on generic orbits around a rotating SMBH surrounded by a thick disk (\\textit{i.e.,} a torus). We consider the torus to have constant specific angular momentum and neglect its self-gravity (\\textit{i.e.,} we consider the metric to be pure Kerr). Under these assumptions, an analytical solution exists for this system~\\cite{constant_l_disks1,constant_l_disks2}. This configuration can be proved to be marginally stable with respect to axisymmetric perturbations~\\cite{marginally_stable} (\\textit{i.e.,} if perturbed, such a torus can accrete onto the SMBH), and is expected to be a good approximation at least for the inner parts of the accretion flow~\\cite{constant_l_disks1,constant_l_disks2}.  We have found that for a system composed of a SMBH with mass $M=10^6 M_\\odot$ and a torus with mass $M_{\\rm t}\\lesssim M$ and outer radius $r_{\\rm out}=10^5 M$, the effect of the hydrodynamic drag on the motion of the satellite black hole is much smaller than radiation reaction at those distances from the SMBH which can be probed with LISA (\\textit{i.e.,} $\\sim 10 M$ for $M=10^6 M_\\odot$). Although these values for $M$, $M_{\\rm t}$ and $r_{\\rm out}$ are plausible for AGNs, an overall uncertainty is still present and has motivated an investigation also for different masses and sizes of the torus. In this way we have found that the effect of the torus can be important in the early part of the inspiral and that it could leave an observable imprint in the gravitational waveforms detected by LISA, if the radius of the torus is decreased to $r_{\\rm out}=10^3-10^4 M$ or, even for $r_{\\rm out}=10^5 M$ and $M_{\\rm t}\\lesssim M$, if $M=10^5 M_\\odot$. In this latter case, in fact, LISA could detect an EMRI event at distances as large as $r\\sim 45 M$ from the SMBH, although the event needs to be sufficiently close to us because the amplitude of the gravitational-wave signal decreases as $M/r$.  In addition, if non-negligible, the effect of the hydrodynamic drag would have a distinctive signature on the waveforms. Radiation reaction, in fact, always increases the inclination of the orbit with respect to the equatorial plane (\\textit{i.e.,} orbits evolve towards the equatorial retrograde configuration)~\\cite{NHSO}. The hydrodynamic drag from a torus corotating with the SMBH, on the other hand, always decreases this angle (\\textit{i.e.,} orbits evolve towards the equatorial prograde configuration). Should such a behavior be observed in the data, it would provide a strong qualitative signature of the presence of the torus. However, it is important to point out that even for those configurations in which the hydrodynamic drag plays a major role, this is restricted to the initial part of the inspiral detectable by LISA, whereas its effect rapidly vanishes in the very strong-field region of the SMBH (\\textit{i.e.,} $p\\lesssim5M$). As a result, the pure-Kerr templates would provide a faithful description of the last part of the inspiral even in these cases.  The rest of the paper is organized as follows. In Sec. \\ref{sec:torus} we review the equilibrium solutions that we used for the orbiting torus. In Sec. \\ref{sec:interaction} we present the equations governing the interaction between the satellite black hole and the torus, while in Sect.~\\ref{sec:adiabatic} we apply the adiabatic approximation to the hydrodynamic drag. Results are then discussed in Sec. \\ref{sec:eq_circ} for equatorial circular orbits and in Sec. \\ref{sec:generic} for generic (inclined and eccentric) orbits. Finally, the conclusions are drawn in Sec.~\\ref{sec:conclusion}. Throughout this paper we use units in which $G=c=1$. ",
        "conclusions": "Conclusions}  SMBHs are expected to be surrounded by matter, either in the form of stellar disks, as in the case of ``normal'' galactic centers, or in the form of accretion disks of gas and dust, as in the case of AGNs. In order to assess under what conditions and to what extent the interaction with matter could modify the gravitational-wave signal from EMRIs in AGNs, we have studied EMRIs in spacetimes containing a SMBH surrounded by a non self-gravitating torus. For simplicity, and in order to handle the equilibrium solution analytically, we have considered a torus with a constant distribution of specific angular momentum, using as reference dimensions and masses those for the accretion disks expected in AGNs, but bearing in mind that these also come with rather large uncertainties. We have extrapolated our results also to cases in which the mass of the torus is comparable with that of the SMBH, although we stress that in this limit our test-fluid approximation for the torus is no longer valid.  Overall, we have found that the effect of the hydrodynamic drag exerted by the torus on the satellite black hole can have important effects sufficiently far from the central object, and that these effects are qualitatively different from those of radiation reaction. In particular, if the torus is corotating with the SMBH, the hydrodynamic drag always \\textit{decreases} the inclination of the orbits with respect to the equatorial plane (\\textit{i.e.,} orbits evolve towards the equatorial prograde configuration), whereas radiation reaction always \\textit{increases} the inclination (\\textit{i.e.,} orbits evolve towards the equatorial retrograde configuration).  In the case of a system composed of a SMBH with mass $M=10^6 M_\\odot$ and a corotating torus with mass $M_{\\rm t}\\lesssim M$, the effect of the torus will be marginally observable by LISA only if the radius of the torus is as small as $r_{\\rm out}\\approx 10^3-10^4 M$. However, if the SMBH has a lower mass, EMRIs will be detectable by LISA at larger distances from the SMBH, and the effects of a torus will be more evident.  For instance, for a SMBH with $M=10^5 M_\\odot$ and a corotating torus with outer radius $r_{\\rm out}=10^5 M$ and mass $M_{\\rm t}=0.1 M$ ($M_{\\rm t}=M$), the inclination with respect to the equatorial plane will decrease, due to the hydrodynamic drag, for orbits with semi-latus rectum $p\\gtrsim35M$ ($p\\gtrsim25 M$), while the EMRI signal will start being detectable by LISA already at a distance of $\\approx 50 M$ from the SMBH.  Note, however, that unless one considers as the satellite an intermediate-mass black hole with $m\\sim 100 M_\\odot$ around a $10^5 M_\\odot$ SMBH (a configuration which may be possible but about which too little is yet known), considering EMRIs at such large distances from the SMBH has the obvious drawback that the amplitude of the gravitational-wave signal will be proportionally smaller. This will considerably reduce the detection volume, although the decrease in the event rate could be mitigated by the fact that weak-field EMRIs are probably more numerous than strong-field EMRIs, which are the ones accounted for in standard calculations of event rates.  In general, we expect measurements of the evolution of the inclination angle in the early phases of EMRIs to be a potential source of important information about the presence of thick tori which could not be detected by other techniques. Moreover, because for any astrophysically plausible torus configurations the effect of the hydrodynamic drag becomes rapidly negligible in the very strong-field region of the SMBH (\\textit{i.e.,} $p\\lesssim5M$), the presence of a torus would not prevent high-precision tests of the Kerr nature of the SMBH being performed.  Although obtained with a simple model for the torus (\\textit{i.e.}, with a constant specific angular momentum), the important feature that distinguishes the hydrodynamic drag from radiation reaction, namely the decrease of the inclination angle, cannot be affected by a change of the specific angular momentum distribution (we recall that $\\ell$ must be increasing with radius for stability). Such a feature, in fact, is simply due to the conservation of the momentum of the satellite in the $\\theta$-direction during accretion and to the dynamical friction of the fluid: both effects force the satellite to smaller inclinations by reducing its $\\theta$-velocity. However, the calculation of the magnitude of the hydrodynamic drag and how it compares with radiation reaction for more general disk models is not straightforward.  Tori built with increasing distributions of specific angular momentum, in fact, would have two substantial differences with respect to those considered here. Firstly, the separation between the specific angular momentum of the torus and the Keplerian specific angular momentum will generally decrease for orbits corotating with the torus, thus reducing the relative motion between the satellite and the fluid and consequently the hydrodynamic drag, whereas it will increase for orbits counter-rotating relative to the torus, thus enhancing the hydrodynamic drag. The magnitude of this effect depends on the precise angular momentum distribution considered and rough estimates can be made assuming a power-law for the specific angular momentum, \\textit{i.e.}, $\\ell/M\\sim (r/M)^{\\alpha}$, with $\\alpha<0.5$ for the torus to have an outer radius and a cusp~\\cite{font_daigne,zanotti_etal:05}. Using the general formulas reported in Sec.~\\ref{sec:torus}, it is easy to check that for $r$ between $20 M$ and $50 M$ and prograde orbits the relative motion decreases by $\\sim 20-30\\%$ for $\\alpha=0.1$ and by at least $95\\%$ for $\\alpha=0.4$ (this significant decrease is due to the fact that for $\\alpha=0.4$ the center moves to a radius $r_{\\rm center}\\sim 27 M$, just in the middle of the radial interval which we are considering). Conversely, in the same radial range the increase for counter-rotating orbits is of about $8-10\\%$ ($30-45\\%$), for $\\alpha =0.1$ ($0.4$). Secondly, the density in the inner parts of the torus will generally decrease. Using again the expressions in Sec.~\\ref{sec:torus}, it is easy to check that the density decreases by about $18\\%$ ($97\\%$) at $r\\sim20M$ and by about $11\\%$ ($90\\%$) at $r\\sim50M$, for $\\alpha =0.1$ ($0.4$).  Overall, therefore, the decrease of the inclination angle due to the hydrodynamic drag could be detectable by LISA also for non-constant $\\ell$ tori, especially if $\\ell$ varies slowly with the radius or, if $\\ell$ varies rapidly with the radius, if the EMRI is counter-rotating relative to the torus.  Finally, let us comment on two further effects that can in principle occur in the systems considered in this paper. First, the motion of the satellite will be influenced by the gravitational attraction exterted by the torus. This is clearly a conservative effect, and cannot therefore influence the infall of the satellite towards the SMBH, which is instead regulated by the dissipative forces (radiation reaction and hydrodynamic drag). However, this effect can in principle cause the periastron to advance, thus introducing a phase-shift in the gravitational waveforms. (Note that a similar advance is caused by the conservative part of the gravitational self-force~\\cite{pound}.) To calculate the order of magnitude of this effect, let us consider for simplicity a thin disk of outer radius $r_{\\rm out}$, mass $M_{\\rm D}$ and constant surface density, and a satellite of mass $m$ located on the equatorial plane at a distance $d\\ll r_{\\rm out}$ from the central SMBH, the mass of which we denote by $M$. The potential energy of the satellite due to the gravitational field of the disk can be easily calculated to be, up to a constant and to leading order, \\eq U\\approx\\frac{m M_{\\rm D}\\, d^2}{2 r_{\\rm out}^3}\\,. \\eeq This potential energy can be used to compute the Newtonian periastron precession of the satellite's orbit during a revolution [use for instance eq.~(1) of Ref.~\\cite{landau}, chapter 3, exercise number 3]: \\eq \\delta \\phi\\approx- \\frac{3 M_{\\rm D} \\pi  d^3}{M r_{\\rm out}^3} \\eeq for almost circular orbits. Using this equation and the well-known Netwonian formula for the revolution period, it is easy to check that, for orbits relevant for LISA, the total phase-shift accumulated in $1$ year is $\\ll 2 \\pi$ as long as $r_{\\rm out} \\gtrsim 10^4 M$ and $M=10^5-10^6 M_\\odot$. Because LISA is not expected to detect phase lags less than 1 cycle over its lifetime\\footnote{This corresponds indeed to a dephasing time of the order of LISA's lifetime.}, this periastron advance and the consequent phase-shift cannot be observed. On the other hand, for disks or tori with $r_{\\rm out} \\sim 10^3 M$, this effect could in principle be marginally visible by LISA (especially if $M=10^5M_\\odot$).  A second effect which could in principle affect EMRIs in the presence of a torus is the spin of the satellite black hole, which increases due to accretion of the torus material. The satellite's spin couples with the orbital angular momentum as well as with the spin of the SMBH, but its effect on the motion is negligible over a timescale of $1$ year~\\cite{leor_spin}, unless it is close to its maximal value (in which case it might be marginally observable)~\\cite{leor_spin}."
    },
    "0711/0711.4591_arXiv.txt": {
        "abstract": "Sparticle landscapes in mSUGRA, in SUGRA models with nonuniversalities (NUSUGRA), and in  D-brane models are analyzed. The analysis exhibits the existence of Higgs Mass Patterns (HPs) (for $\\mu>0$) where the CP odd Higgs could be the next heavier particle beyond the LSP and sometimes even lighter than the LSP. It is shown that the Higgs production cross sections from the  HPs are typically  the largest enhancing the  prospects for their  detection at the  LHC. Indeed it is seen that the recent Higgs  production limits from CDF/D\\O\\ are beginning to put constraints on the HPs. It is also seen that the $B_s\\to \\mu^+\\mu^-$ limits  constrain the HPs more stringently. Predictions of the Higgs  production cross sections for these patterns at the LHC are  made.   We compute the neutralino-proton cross  sections  $\\sigma (\\tilde\\chi_1^0 p)$  for dark matter experiments and show that the largest $\\sigma (\\tilde\\chi_1^0 p)$ also arise from the HPs and further that  the HPs and some of the other patterns are beginning to be constrained by the most recent data from CDMS and from Xenon10 experiments. Finally,  it is shown that the prospects are bright for the discovery of dark matter with $\\sigma(\\tilde\\chi_1^0 p)$ in the range $10^{-44\\pm .5}$cm$^2$ due to a ``Wall\" consisting of a copious number of parameter points in the Chargino Patterns (CPs) where the chargino is the NLSP.  The Wall, which appears in all models considered (mSUGRA, NUSUGRA and D-branes) and runs up to about a TeV in LSP mass,   significantly enhances the chances for the observation of dark matter by SuperCDMS,  ZEPLIN-MAX, or LUX  experiments  which are expected to achieve  a sensitivity of $10^{-45}$ cm$^2$ or more. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.4244_arXiv.txt": {
        "abstract": "In order to develop a pipeline for automated classification of stars to be observed by the TAUVEX ultraviolet space Telescope, we employ an artificial neural network (ANN) technique for classifying stars by using synthetic spectra in the UV region from 1250\\AA~ to 3220\\AA~ as the training set and International Ultraviolet Explorer (IUE) low resolution spectra as the test set. Both the data sets have been pre-processed to mimic the observations of the TAUVEX ultraviolet imager. We have successfully classified 229 stars from the IUE low resolution catalog to within 3-4 spectral sub-class using two different simulated training spectra, the TAUVEX spectra of 286 spectral types and UVBLUE spectra of 277 spectral types. Further, we have also been able to obtain the colour excess (i.e. E(B-V) in magnitude units) or the interstellar reddening for those IUE spectra which have known reddening to an accuracy of better than 0.1 magnitudes. It has been shown that even with the limitation of data from just photometric bands, ANNs have not only classified the stars, but also provided satisfactory estimates for interstellar extinction. The ANN based classification scheme has been successfully tested on the simulated TAUVEX data pipeline. It is expected that the same technique can be employed for data  validation in the ultraviolet from the virtual observatories. Finally, the interstellar  extinction estimated by applying the ANNs on the TAUVEX data base would provide an extensive extinction map for our galaxy and which could in turn be modeled for the dust distribution in the galaxy. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.1657_arXiv.txt": {
        "abstract": "{ The event-by-event analysis of multiparticle production in high energy hadron and nuclei collisions can be performed using the discrete wavelet transformation. The ring-like and jet-like structures in two-dimensional angular histograms are well extracted  by wavelet analysis. For the first time the method is applied to the jet-like events with background simulated by event generators, which are developed to describe nucleus-nucleus collisions at LHC energies. The jet positions are located quite well by the discrete wavelet transformation of angular particle distribution even in presence of strong background. \\PACS{ {21.65.Qr}{}   \\and {25.75.Gz}{}   \\and {02.30.-f}{} } % } % \\authorrunning{Dremin} \\titlerunning{wavelet analysis of event topology} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.4905_arXiv.txt": {
        "abstract": "New observational techniques and sophisticated modelling methods has led to dramatic breakthroughs in our understanding of the interplay between the surface magnetism, atomic diffusion and atmospheric dynamics in chemically peculiar stars. Magnetic Doppler images, constructed using spectropolarimetric observations of Ap stars in all four Stokes parameters, reveal the presence of small-scale field topologies. Abundance Doppler mapping has been perfected to the level where distributions of many different chemical elements can be deduced self-consistently for one star. The inferred chemical spot structures are diverse and do not always trace underlying magnetic field geometry. Moreover, horizontal chemical inhomogeneities are discovered in non-magnetic CP stars and evolving chemical spots are observed for the first time in the bright mercury-manganese star \\aand. These results show that in addition to magnetic fields, another important non-magnetic structure formation mechanism acts in CP stars. ",
        "introduction": "Chemically peculiar stars are privileged astrophysical laboratories for investigation into the structure formation processes taking place in the surface layers of stars. Unlike the turbulent atmospheres of cool solar-type stars or mass loss-dominated envelopes of massive stars, atmospheres of slowly rotating main sequence A and B stars are very stable, and this in turn allows several remarkable processes to operate in these stars. The most prominent characteristic of many CP stars is the presence of strong fossil magnetic fields at their surfaces and in interiors. Moreover, due to the atmospheric stability, further enhanced by the magnetic field, chemical segregation processes, determined by the competition between gravitational settling and radiative levitation, can operate efficiently. These processes are responsible for conspicuous non-solar chemical abundance  patterns observed in CP stars and can also lead to strong vertical abundance gradients in the line-forming regions. The balance between gravitation and radiation pressure is affected by the magnetic field and hence prominent horizontal chemical inhomogeneities emerge at the surfaces of CP stars. Rapid pulsational variability in cool Ap stars, also related to non-solar chemistry and the presence of strong magnetic field, adds dynamical aspect to the intricate picture of the CP-star phenomenon.  Various manifestations of this unique and poorly understood combination of complex surface properties of CP stars are readily observed in the form of periodic light and spectrum variation, anomalous line profile shapes, significant Zeeman broadening, circular and linear polarization in spectral lines. With some effort these observables can be interpreted in terms of local magnetic field strength and orientation, horizontal and vertical abundance maps, and time-dependent pulsation velocity field. The resulting extensive information on the surface structures and atmospheric dynamics  is often utterly unique and is usually unaccessible for other kinds of stars, making CP stars the only source of crucial direct evidence and powerful observational constraints on the physical processes governing chemical diffusion, emergence and evolution of large-scale magnetic field.  At the same time, the complexity and remarkable diversity of the phenomena observed in the atmospheres of CP stars drive the development of new astronomical instrumentation, modelling techniques and computer codes. On this background continuing efforts of many CP-star specialists, with enthusiastic support from the rest of the stellar community, has led to the revival of this research field in the last decade. New discoveries have overturned some of the old paradigms, generally resulting in a more complex, yet also more interesting picture of chemically peculiar stars.  In this review I highlight recent progress achieved in the studies of magnetic fields and horizontal chemical inhomogeneities in CP stars. The closely related topics of chemical stratification and magnetoacoustic pulsations in roAp stars are dealt with in other contributions (Ryabchikova, Sachkov et al., this meeting). ",
        "conclusions": ""
    },
    "0711/0711.4134_arXiv.txt": {
        "abstract": "Recent work to include R-Matrix data within a larger model comprised mostly of distorted-wave and plane-wave-Born data has resulted in the General Spectral Modeling (GSM) code.  It employs a quasi-static approximation, a standard, low-density methodology that assumes the ionization balance is separable from a determination of the excited-state populations that give rise to the spectra.  GSM further allows for some states to be treated statistically as contributions to effective rates, instead of being included explicitly in the kinetics model.  While these two approximations are known to be valid at low densities, this work investigates using such methods to model high-density, non-LTE emission spectra and determines at what point the approximations break down by comparing to spectra produced by the Los Alamos National Laboratory code ATOMIC which makes no such approximations.  As both approximations are used by other astrophysical and low-density modeling codes, the results should be of broad interest.  He-like K$\\alpha$ emission spectra are presented for three elements, Ni, Fe, and Ca, in order to gauge the effect of both the statistical methods and the ground-state-only, quasi-static approximation employed in GSM.  This work confirms that at and above the temperature of maximum abundance of the He-like ionization stage, the range of validity for both approximations is sufficient for modeling the low- and moderate-density regimes one typically finds in astrophysical and magnetically confined fusion plasmas.  However, a breakdown does occur for sufficiently high densities; we obtain quantitative limits that are significantly higher than previous works.  Additionally, this work demonstrates that, while the range of validity for both approximations is sufficient to accurately predict the density-dependent quenching of the z line, the approximations begin to break down at higher densities.  Thus these approximations should be used with greater care when modeling high-density plasmas such as those found in laser-driven inertial confinement fusion and electromagnetic pinch devices. ",
        "introduction": "The subject of K$\\alpha$ emission lines arising from helium-like ions has been a topic in the literature for the better part of a century \\cite{Edlen-Tyren:1939}.  The work that the present effort builds on begins with Gabriel and Jordan's work on the Sun \\cite{Gabriel-Jordan:1969-nature,Gabriel-Jordan:1969-helines}.  There are four K$\\alpha$ lines arising from the helium-like ionization species, each involving decay from an excited state to the ground state. The `w' line represents a dipole allowed transition arising from the decay of the $1s2p$ $^1P^{o}_1$ state.  Because of mixing, the  $1s2p$ $^3P^{o}_1$ state can also decay via a dipole allowed transition, giving rise to the intercombination line `y'.  The $1s2s$~$^3S_1$ state can decay via a relativistic magnetic-dipole transition, giving rise to the forbidden line, `z'.  Lastly an additional forbidden line, `x', is formed from the $1s2p$~$^3P^{o}_2$ state decaying via a magnetic-quadrupole transition.  Due to the sometimes small energy separation between the $1s2p$ $^3P^{o}_1$ and $1s2p$~$^3P^{o}_2$ states, the x and y lines can not be independently resolved for all elements or plasma conditions.  The $1s2s$ $^1S_0$ state must also be included explicitly in the determination of the excited-state populations as it can decay via a two-photon transition which has an important effect on the model, as does the presence of the $1s2p$~$^3P^{o}_0$ state \\cite{Gabriel-Jordan:1969-helines,Mewe-Schrijver:1978stat}.  From these four lines, line ratios have been developed in order to diagnose certain conditions about the corresponding plasma. One of those line ratios, R, is defined as \\cite{Gabriel-Jordan:1969-helines} \\begin{equation} R=\\frac{I(\\mathrm{z})}{I(\\mathrm{x})+I(\\mathrm{y})}\\;. \\label{Rformula} \\end{equation} The intensities, $I$, of the non-dipole allowed lines (x and z) and the intercombination line (y) vary in a similar manner as a function of temperature.  However, electron collisions can transfer population from the $1s2s$~$^3S_1$ state to the $1s2p$~$^3P^{o}_{0,1,2}$ states, thereby quenching the z line and increasing the intensity of the x and y lines.  Therefore the ratio $R$ can be used as a diagnostic of electron density.  Building on this earlier work, the study of K$\\alpha$ emission from He-like ions was extended to heavier elements by including the effect of dielectronic satellite lines \\cite{Gabriel:1972-sats}, and to a broader range of elements and transient conditions in the seminal works of Mewe and Schrijver \\cite{Mewe-Schrijver:1978stat,Mewe-Schrijver:1978nonstat}.  Improved data and improved models continued to be brought to bear as they became available \\cite{Pradhan-Shull:1981-helines,Pradhan:1982-helines,Pradhan:1985-cascade,Lee-etal:1985,Lee-etal:1986}.  Photoionized media have also been investigated \\cite{Porquet-Dubau:2000}.  Lastly, more recent Breit-Pauli R-Matrix \\cite{BPRM,Nahar-Pradhan:1994-urr} data were used to refine the quantities of interest even further \\cite{Oelgoetz-Pradhan:2001,Oelgoetz-Pradhan:2004}.  Additionally the K$\\alpha$ lines are also observed in laboratory plasmas \\cite{Edlen-Tyren:1939,Negus-Peacock:1979,Peacock-Burgess:1981,Bitter-etal:1993,Bertschinger-etal:1999,Marchuk-etal:2004,vonHellermann-etal:2005,Safronova-etal:2006}, and these same methods have been successfully used to model and interpret the spectra of magnetically confined fusion plasmas \\cite{Bitter-etal:1993,Bertschinger-etal:1999,Marchuk-etal:2004,vonHellermann-etal:2005}.  While the above narrative and citations are only a small portion of the work done in this field since these lines were initially observed, they are the platform upon which the General Spectral Modeling (GSM) code \\cite{Oelgoetz-Dissertation}, one of the codes used in the present work, is based.  It should be noted that much of the previous work explicitly assumed that, for the purposes of finding the populations of the He-like excited states, the populations of the adjacent ionization stages are entirely in the ground state.  While this particular type of quasi-static approximation (the ground-state-only, quasi-static approximation) is clearly valid at low densities, it breaks down at higher densities as electron-impact ionization and recombination out of excited states become important mechanisms for populating the excited states.  One could envision extending this treatment by including not only the ground state in the set of states with which the excited states are in equilibrium, but also select metastable states.  The effect of expanding the quasi-static treatment employed in this way is not explored in the present work, as the vast majority of papers on astrophysical He-like K$\\alpha$ emission make these same approximations.  Additionally many of the above works treat many of the excited states as cascade corrections (see, for example, \\cite{Pradhan:1985-cascade}) to direct rate coefficients.  The effect of such a statistical approximation is explored in the present work.  From a more general perspective, expanding beyond the specific case of He-like ions described above, the trend over time has been to incorporate more accurate atomic data in plasma kinetics and spectral modeling.  Often, these better atomic data are difficult or impossible to calculate for all processes between all levels.  Furthermore, if the atomic data are calculated using a method that accounts for resonance structure, such as the R-Matrix method, it is much more time consuming to calculate rate coefficients from the more complex cross sections than from smooth cross sections.  Reducing the computational requirements makes the quasi-static and statistical approximations appealing.  A separate, but related, consideration is that as the plasma transitions from non-local-thermodynamic-equilibrium (non-LTE) conditions to the high density regime, an accurate modeling approach should produce the appropriate LTE limit.  However, each of the two approximations mentioned above typically preclude a transition to the correct LTE population distribution in the high-density limit, regardless of the quality of the atomic data that is used.  The flexibility of GSM allows for an investigation of the interplay between these various considerations when attempting to model plasmas over a broad range of conditions.  The primary interest of this work is benchmarking and examining the range of validity for the ground-state-only, quasi-static approximation as well as the statistical treatment of excited states in the context of He-like ions, as these approximations are employed by most of the works referenced above.  Many of the connections between excited states are neglected by these two approximations; thus the constraints on the system are not sufficient to produce the correct high-density, local thermodynamic equilibrium (LTE) limit.   Using these approximations results in a significant speed up however, and allows one to include the effects of many more states with the same computing power.  Thus the applicability of these approximations to high-density, non-LTE plasmas such as those found in vacuum sparks \\cite{Negus-Peacock:1979}, electromagnetic pinch type devices \\cite{Bakshaev-etal:2001,Safronova-etal:2006} and laser produced, inertial confinement fusion plasmas \\cite{Peacock-Burgess:1981} is of great interest.  To gauge the effect of these approximations, the spectra produced by GSM are compared to the results from another code, ATOMIC \\cite{ATOMIC}.  ATOMIC does not employ either the quasi-static approximation or statistical methods for treating excited states; instead every fine-structure level in the model is treated identically, such that all included processes between all levels are considered explicitly.  It should be noted that the present work is not the first comparison of a code based on the quasi-static approximation to ATOMIC; comparisons between the code ADAS \\cite{ADAS} and ATOMIC have been previously considered \\cite{Loch-etal.2004}.  However, the primary aim of the present work is to investigate the breakdown of the approximations involved and to provide quantitative limits for the breakdown, while the earlier work considered the effects of differing atomic data sets on plasma kinetics modeling.  The secondary aim of this work is to confirm that using a statistical treatment and the ground-state-only, quasi-static approximation are adequate for modeling emission spectra in the low- and moderate-density regimes encountered in magnetic confinement fusion devices and most astrophysical plasmas.  As the line ratio $R$ is useful for determining the density of the emitting plasma, and many previous calculations of $R$ employ both a statistical treatment of excited states and the ground-state-only, quasi-static approximation, we also confirm that the methods used for such calculations are indeed adequate.  The test cases presented in this work are for Ni, Fe, and Ca as they are directly relevant to astrophysical and laboratory plasmas. ",
        "conclusions": "The ground-state-only, quasi-static approximation and statistical treatments for excited states are often-used methodologies to simplify the spectral modeling of plasmas at low densities.  This work confirms their low-density validity, based on detailed atomic models, and provides numerical values for the electron densities at which these approximations break down for Ni, Fe, and Ca at various temperatures near and above the temperature of maximum abundance for the He-like ionization stage.  At the temperatures considered, the quasi-static approximation using only the ground states of the H-like and Li-like ionization stages appears to be valid for densities that are significantly higher than those predicted by previous works \\cite{Gabriel:1972-sats,Mewe-Schrijver:1978stat}.  The ground-state-only, quasi-static approximation is also valid up to densities that are orders of magnitude higher than the density limits calculated for the statistical treatment employed in the GSM7 models in all cases considered.  This work also confirms that the ground-state-only, quasi-static approximation as well as the methods employed to treat excited states as contributions to effective rates via a statistical treatment are valid for predicting the quenching of the z line, and thus the drop in the $R$ ratio, as a function of electron density in the temperature range considered.  However, this study further indicates that these approximations should be used with care at densities where the $R$ ratio would be expected to be close to zero, such as those found in laboratory produced sparks \\cite{Negus-Peacock:1979}, high-density electromagnetic pinch devices \\cite{Bakshaev-etal:2001,Safronova-etal:2006}, and inertial confinement fusion experiments \\cite{Peacock-Burgess:1981}.  Further study is required to determine if the  quasi-static approximation using only the ground states of the H-like and Li-like ionization stages has a smaller range of validity with respect to density than the statistical treatment at temperatures significantly below the temperature of He-like maximum abundance, or for conditions found in transient plasmas.  We do expect the range of validity of the ground-state-only, quasi-static approximation to be narrower in density at temperatures lower than those considered in this work because the population of the Li-like ionization stage should be greater, and thus the effect of the Li-like ionization stage on K$\\alpha$ spectral formation via collisional ionization would play a more significant role.  Another avenue of future research concerns the improvement of both the statistical treatment and the quasi-static approximation used in this work without abandoning either approach.  The models that employ statistical methods could be improved by treating more states explicitly and fewer statistically.  Similarly, the quasi-static approximation employed could be improved by allowing for more than just the ground state to be populated in the ionization-balance portion of the calculation.  Investigating both improvements as a function of temperature and density is left to future work."
    },
    "0711/0711.3008_arXiv.txt": {
        "abstract": "The ratio of the luminosity of diffuse 511 keV positron annihilation radiation, measured by INTEGRAL in its four years, from a Galactic ``positron bulge\" ($<$ 1.5 kpc) compared to that of the disk is $\\sim$ 1.4. This ratio is roughly 4 times larger than that expected simply from the stellar bulge to disk ratio of $\\sim$ 0.33 of the Galactic supernovae, which are thought to be the principal source of the annihilating positrons through the decay of radionuclei made by explosive nucleosynthesis in the supernovae. This large discrepancy has prompted a search for new sources.  Here, however, we show that the measured 511 keV luminosity ratio can be fully understood in the context of a Galactic supernova origin when the differential propagation of these $\\sim$ MeV positrons in the various phases of the interstellar medium is taken into consideration, since these relativistic positrons must first slow down to energies $\\leq$ 10 eV before they can annihilate. Moreover, without propagation, none of the proposed positron sources, new or old, can explain the two basic properties on the Galactic annihilation radiation: the fraction of the annihilation that occurs through positronium formation and the ratio of the broad/narrow components of the 511 keV line.  In particular, we show that in the neutral phases of the interstellar medium, which fill most of the disk ($>$ 3.5 kpc), the cascade of the magnetic turbulence, which scatters the positrons, is damped by ion-neutral friction, allowing positrons to stream along magnetic flux tubes. We find that nearly 1/2 of the positrons produced in the disk escape from it into the halo. On the other hand, we show that within the extended, or interstellar, bulge ($<$ 3.5 kpc), essentially all of the positrons are born in the hot plasmas which fill that volume. We find that the diffusion mean free path is long enough that only a negligible fraction annihilate there and $\\sim$ 80\\% of them escape down into the HII and HI envelopes of molecular clouds that lie within 1.5 kpc before they slow down and annihilate, while the remaining $\\sim$ 20\\% escape out into the halo and the disk beyond. This propagation accounts for the low observed annihilation radiation luminosity of the disk compared to the bulge.  In addition, we show that the primary annihilation sites of the propagating positrons in both the bulge and the disk are in the warm ionized phases of the interstellar medium. Such annihilation can also account for those two basic properties of the emission, the fraction ($\\sim 93\\pm7$\\%) of annihilation via positronium and the ratio ($\\sim$ 0.5) of broad ($\\sim$ 5.4 keV) to narrow ($\\sim$ 1.3 keV) components of the bulge 511 keV line emission. Moreover, we expect that the bulk of this broad line emission comes from the tilted disk region ($0.5 < R < 1.5$ kpc) with a very large broad/narrow flux ratio of $\\sim$ 6, while much of the narrow line emission comes from the inner bulge ($R < 0.5$ kpc) with a negligible broad/narrow flux ratio. Separate spectral analyses of the 511 keV line emission from these two region should be able to test this prediction, and further probe the structure of the interstellar medium.  Lastly, we show that the asymmetry in the inner disk annihilation line flux, which has been suggested as added evidence for new sources, can also be fully understood from positron propagation and the asymmetry in the inner spiral arms as viewed from our solar perspective without any additional sources. ",
        "introduction": "The discovery (Johnson et al.\\ 1972; Leventhal et al.\\ 1979) of the diffuse Galactic positron annihilation radiation at 511 keV from the inner Galaxy has led to extensive studies of the possible origin of the annihilating positrons. These studies have shown (e.g.\\ Ramaty \\& Lingenfelter 1979; Kn\\\"odlseder et al.\\ 2005) that, of all the various potential sources, the positrons from the decay of radioactive nuclei produced by explosive nucleosynthesis in supernovae are the largest and, hence, most likely source.  Recent measurements, however, have raised new questions about the source. Analyses of the measurements by the gamma-ray spectrometer (SPI) on ESA's INTEGRAL spacecraft of the distribution of the diffuse Galactic positron annihilation radiation have shown that the luminosity of a Galactic positron bulge within 1.5 kpc of the Galactic Center is $\\sim$ 1.4 times that of the Galactic disk (Kn\\\"odlseder et al. 2005; Weidenspointner et al. 2007, 2008a). Their best fit models are either concentric shell sources, or Gaussian spheroidal sources, in a bulge extending out to $\\sim$ 1.5 kpc combined with a thick disk with a half thickness of $z \\sim$ 0.4 kpc, that extends from the Galactic Center to well beyond the Solar distance of 8 kpc. They also find a large additional contribution from a spherical halo extending beyond the bulge.  In order to compare with these observations, we use these model fitting definitions of the positron bulge, thick disk, and halo, following Kn\\\"odlseder et al.\\ (2005) and Weidenspointner et al. (2007, 2008a) in all of the discussion that follows, although both have been defined in various other ways in the literature.  \\subsection{ 511 keV Line Bulge/Disk Luminosity Ratio }  The best-fit analysis of the SPI/INTEGRAL first year's data (Kn\\\"odlseder et al.\\ 2005, Shells+D1) gave a 511 keV line luminosity of $L_B \\sim (0.78\\pm0.04)\\times10^{43}$ photons s$^{-1}$ in a spherical bulge from 0 to 1.5 kpc and an old, thick disk luminosity of $L_D \\sim (0.26\\pm0.07)\\times10^{43}$ photons s$^{-1}$, scaled to a Solar distance of $R_o \\sim$ 8 kpc. This gave a bulge to disk 511 keV luminosity ratio of $L_B/L_D \\sim 3.0\\pm0.8$. The subsequent best-fit analyses of the first two years' data, which included more coverage of the disk (Weidenspointner et al.\\ 2007, for 4 Nested Shells and Old Disk), gave a bulge flux of $\\sim (0.79\\pm0.07)\\times10^{-3}$ photons cm$^{-2}$ s$^{-1}$ in spherical shells out to 1.5 kpc, a thick disk flux of $\\sim (1.41\\pm0.35)\\times10^{-3}$ photons cm$^{-2}$ s$^{-1}$, and a marginally detected spherical halo flux of $\\sim (0.86\\pm0.59)\\times10^{-3}$ photons cm$^{-2}$ s$^{-1}$ from 1.5 to 8 kpc. To infer the corresponding 511 keV line luminosities, we use the flux to luminosity ratios from their earlier modeling (Kn\\\"odlseder et al.\\ 2005 Table 3), which are essentially just $4\\pi R^2$ where $R \\sim R_o$ for the bulge, about 0.57$R_o$ for the thick disk, and about 0.93$R_o$ for the halo. From the best-fit fluxes, we infer a 511 keV positron bulge luminosity $L_B \\sim (0.60\\pm0.05)\\times10^{43}$ photons s$^{-1}$ in the spherical shells out to 1.5 kpc, and a thick disk luminosity of $L_D \\sim (0.36\\pm0.09)\\times10^{43}$ photons s$^{-1}$ for a $L_B/L_D \\sim 1.7\\pm0.5$.  Most recently from more than four years of observation, Weidenspointner et al.\\ (2008a, model BD) fit the 511 keV fluxes and luminosities to a two component narrow (FWHM of 3$^o$, or 0.42 kpc) and wide (FWHM of 11$^o$, or 1.5 kpc) spheroidal Gaussian distribution and thick disk from more extensive observations that increased the effective sampling distance to 0.75$R_o$ for the disk. They found best-fit fluxes of $\\sim (0.75\\pm0.09)\\times10^{-3}$ photons cm$^{-2}$ s$^{-1}$ from the bulge and $\\sim (0.94\\pm0.16)\\times10^{-3}$ photons cm$^{-2}$ s$^{-1}$ from the disk. Scaled to $R_o \\sim$ 8 kpc, these fluxes give a best-fit 511 keV luminosity $L_B \\sim (0.57\\pm0.08) \\times10^{43}$ photons s$^{-1}$ in the spherical bulge $<$ 1.5 kpc, and a thick disk luminosity of $L_D \\sim (0.40\\pm0.06) \\times10^{43}$ photons s$^{-1}$. This gives a bulge to disk 511 keV luminosity ratio of $L_B/L_D \\sim 1.4\\pm0.3$. They also found a best-fit combined halo and bulge luminosity of $\\sim (1.55\\pm0.1)\\times10^{43}$ photons s$^{-1}$ for similar scaling, and subtracting the bulge luminosity, suggests a spherical halo luminosity of $L_H \\sim 1\\times10^{43}$ photons s$^{-1}$ beyond the bulge, for a total Galactic 511 keV luminosity $L_G \\sim (2.0\\pm0.2)\\times10^{43}$ photons s$^{-1}$.  Best-fit SPI/INTEGRAL spectral analyses by Weidenspointner et al.\\ (2008b) also show an asymmetry in the positron annihilation line flux from two opposing inner disk components at 0$^o$ to 50$^o$ to either side of the Galactic bulge. These show an $80\\pm40$\\% excess in the 511 keV line flux at the negative longitudes compared to the positive.  Assuming simply that the positron annihilation and production rates are in local equilibrium, these observations of the luminosities have been taken to imply a similar bulge/disk ratio and asymmetric spatial distribution for Galactic positron production. This suggested production ratio and distribution has challenged supernova source models, since it is much larger than that expected from the distribution and mean bulge/disk ratio of Galactic supernovae (e.g. Kn\\\"odlseder et al.\\ 2005; Weidenspointner et al.\\ 2008b). As we show, however, these spatial properties can be fully explained by positron propagation.  \\subsection{511 keV Line Width \\& Positronium Fraction }  We also consider the other two fundamental spectral properties of the Galactic positron annihilation: the ratio of the broad to narrow 511 keV line emission and the fraction ($f_{Ps}$) of the annihilation that occurs via positronium (Ps) formation. As shown by Jean et al. (2006) neither of these can be explained by any of the proposed positron sources, either old or new, without extensive positron propagation, since all of these sources are expected to produce most of the positrons in the hot tenuous phases of the interstellar medium, which cannot be their primary annihilation site. For this reason alone a detailed treatment is required of the production, propagation and annihilation in each phase and region of the Galaxy.  Studies (e.g. Guessoum, et al.\\ 1991; Guessoum, et al.\\ 2005) show that positron annihilation in different phases of the interstellar medium occurs in differing ratios of direct annihilation on both free and bound electrons, to indirect annihilation via positronium formation, depending on the ionization fraction and temperature. Direct annihilation produces two photons at 511 keV, while positronium annihilation produces either two 511 keV photons or a three photon continuum depending on the spin state. Positronium is formed 25\\% of the time in the singlet state, parapositronium, which annihilates with a mean life of $\\sim 1.25\\times10^{-10}$ s into two 511 keV photons, while 75\\% of the time it forms in the triplet state, orthopositronium, which annihilates with a mean life of $\\sim 1.4\\times10^{-7}$ s into a three photon continuum between 0 and 511 keV (e.g.\\ Guessoum et al.\\ 1991).  The SPI/INTEGRAL measurements of the ratio of the Galactic 2$\\gamma$ 511 line flux to that of the 3$\\gamma$ continuum in the bulge and disk shows that the bulk of the positrons annihilate via positronium (Ps) with $f_{Ps} \\sim 94\\pm6$\\% (Churazov et al.\\ 2005), $\\sim 92\\pm9$\\% (Weidenspointner et al.\\ 2006), and $\\sim 95\\pm3$\\% (Jean et al. 2006). The weighted mean positronium fraction of $\\sim 94\\pm4$\\% also allows us to determine the total Galactic positron annihilation rates, $A$, in the bulge and disk, since $A = (e^+/\\gamma_{511})L$. From the above we see that $(e^+/\\gamma_{511}) = [2(f_{Ps}/4) + 2(1 - f_{Ps})]^{-1}$, which equals 1.69$\\pm$0.17. Thus the inferred best-fit positron bulge and Galactic disk annihilation rates are $A_B \\sim (0.96\\pm0.17)\\times10^{43}$ e$^+$ s$^{-1}$ and $A_D \\sim (0.68\\pm0.12)\\times10^{43}$ e$^+$ s$^{-1}$ using the previously derived L$_{B}$ and L$_{D}$. In the hot tenuous plasma of the halo, however, the positronium fraction depends on the refractory grain abundance, ranging from only about 18\\% with narrow (FWHM $\\sim$ 2 keV) 511 keV line emission, if all the refractories are in grains, to 42\\% with broad (FWHM $\\sim$ 11 keV) if the grains all disintegrated (Jean et al. 2006). Thus we assume a possible range of $(e^+/\\gamma_{511}) \\sim 0.65\\pm0.07$ there. From the best-fit halo flux, we thus infer a halo annihilation rate of $A_H \\sim (0.65\\pm0.11)\\times 10^{43}$ e$^+$ s$^{-1}$, implying a total Galactic positron production rate of $\\sim (2.3\\pm0.2)\\times10^{43}$ e$^+$ s$^{-1}$.  The width of the 511 keV line also depends on the temperature of the medium and its state of ionization which effects the fraction of positronium formed by charge exchange with H, H$_2$, and He by superthermal, $\\sim$10 eV, positrons as they slow down (Guessoum et al.\\ 1991; Guessoum et al.\\ 2005). The prompt annihilation of this fast moving positronium produces a characteristic broad (FWHM $\\sim$ 6 keV) component of the 511 keV line, while the subsequent thermal annihilation, either directly on free and bound electrons or via positronium formed by radiative combination or charge exchange give much narrower ($<$ 2 keV) lines. Even the annihilation in the hot ($T \\sim 10^6$K), tenuous ionized medium which would produce a broader ($\\sim$10 keV) line from the hot free elections, can instead produce a narrow line if most of the positrons annihilate on refractory dust grains (Guessoum et al.\\ 1991; Guessoum et al.\\ 2005). The dust, as it does with interstellar molecule formation, provides a 2-dimensional regime with much higher interaction rates for surface chemistry.  The width of the Galactic bulge 511 keV line measured by SPI/INTEGRAL from the first years' data has been recently fitted (Churazov et al.\\ 2005; Jean et al.\\ 2006) by two components with about 67$\\pm$10\\% of the emission in a narrow line with a width of 1.3$\\pm$0.4 keV and the remaining 33$\\pm$10\\% in a broad line with a width of 5.4$\\pm$1.2 keV by positronium formation in flight, giving a broad/narrow 511 keV line flux ratio of $\\sim$ 0.5. The combined positronium fraction and relative fluxes of the broad and narrow 511 keV line emission strongly constrain the bulge annihilation processes, as we discuss in detail below.  \\subsection{Positron Production \\& Propagation }  The large difference between the ratio of the 511 keV annihilation line bulge/disk luminosity and that of the expected positron production by supernovae has led to suggestions there is some new, unrecognized source of positrons in the Galactic bulge. This seems unlikely, however, since various recent reviews (e.g.\\ Dermer \\& Murphy 2001; Kn\\\"odlseder et al.\\ 2005; Guessoum, Jean \\& Prantzos 2006) of the potential Galactic positron sources all conclude that supernovae are still the most plausible source and that other suggested sources, including cosmic-ray interactions, novae, and various exotic processes, all seem to be weaker and much less certain.  However, as we have shown (e.g.\\ Guessoum, Ramaty \\& Lingenfelter 1991), the 511 keV annihilation emission only illuminates the annihilation sites, not the sources, of the positrons, since these positrons must first slow down to energies $\\leq$ 10 eV before they can annihilate. Thus, we examine in detail the propagation and annihilation of the $\\beta^+$-decay positrons in the bulge, disk, and halo, and we show here that the expected positron propagation and annihilation are very different in the bulge and disk, and the cloud and inter-cloud environments. When these differences are taken into consideration, we find that the spatial distribution, the 511 keV line widths and the positronium fraction can all in fact be clearly understood in the context of a Galactic supernova origin.  Although positron propagation has previously been discussed, the arguments range from the simple suggestion (Prantzos 2006) that most of the disk positrons might diffuse along dipolar field lines and annihilate in the bulge, to Monte Carlo transport simulations (Gillard et al.\\ 2007) suggesting that they travel less than a few hundred parsecs in all but the hot phase. We therefore also carefully examine both the physics of MeV electron propagation and the observational evidence of such propagation from extensive measurements of Solar flare and Jovian electrons in the heliosphere.  Here, we examine in detail the mechanisms of the relativistic positron propagation in the various phases of the interstellar medium.  Overall, we assume that the Galactic stellar and interstellar distributions are defined by two superimposed systems: a stellar bulge-disk-and-halo system of stars and an interstellar bulge-disk-and-halo system of gas and plasma, as shown in Figure 1. The stellar bulge and disk populations essentially determine the distribution of supernovae from the decay of whose radionuclei the positrons are born. Following Ferri\\'ere et al. (2007), the interstellar bulge, disk and halo define the distribution of gas and plasma where the positrons die by annihilation.  As we discuss in more detail later, the stellar bulge is confined within the interstellar bulge, which is blown out by a bulge wind to about 3.5 kpc in the Galactic plane and above the disk it feeds and merges into the halo. The stellar disk extends not only throughout the interstellar disk out to at least about 15 kpc, but also all the way into the interstellar bulge. The interstellar disk, however, extends in only to the outer edge of the interstellar bulge, and starts in the so-called ``molecular ring\" at around 3.5 kpc, which is defined primarily by the two opposing innermost spiral arms.  Threading the combined stellar and interstellar systems, is the Galactic magnetic field, along whose flux tubes the positrons travel from their birth to their death. We assume that the flux tubes are nearly vertical (e.g. Beck 2001), perhaps dipolar, in the inner part of the interstellar bulge and are blown out by the bulge wind (e.g. Bregman 1980; Blitz et al. 1993) roughly radially in the outer bulge, out to the molecular ring, where they begin to be tightly wound into the spiral arms of the interstellar disk (e.g. Beck 2001). There, in giant star formation regions along these spiral arms, hot, massive OB star associations are formed, whose collective supernovae generate superbubbles, that blow out plasma as well as flux tubes into the overlying halo (e.g. Parker 1979).  Within this framework we model the various aspects of the production, propagation and annihilation of Galactic positrons.  We show that roughly half of the Galactic positrons are born in the extended interstellar bulge (e.g. Ferri\\'ere, Gillard \\& Jean 2008) in the inner $\\sim$ 3.5 kpc of the Galaxy, and nearly all result from SNIa occurring in the hot, tenuous high-pressure plasma fills nearly all of that volume. We show that $\\sim$ MeV positrons created in these SNIa supernovae are expected to escape from the remnants into the surrounding tenuous plasma. There they propagate along magnetic flux tubes by 1-dimensional diffusion, resonantly scattered in pitch angle by turbulently generated magnetohydrodynamic (MHD) waves that cascade down to the thermal electron gyroradius. In these plasmas the diffusion mean free path is long enough that only a negligible fraction annihilate there and $\\sim$ 80\\% of the positrons escape down into the dense HII and HI shells of neighboring molecular clouds within 1.5 kpc, where they quickly slow down and annihilate, while the remaining $\\sim$ 20\\% escape out into the halo and disk beyond.  In the Galactic disk beyond 3.5 kpc, the filling factor of hot ($\\sim 10^6$ K) plasma is more modest ($\\sim$ 20\\%). Here the plasma is mostly in superbubbles, created by the correlated core-collapse supernovae of massive stars. Consequently, the bulk of positron production via $^{26}$Al decay, synthesized by these stars, occurs preferentially in the superbubbles. The larger bubbles blow out into the halo, sweeping the magnetic flux tubes up with them. Again the diffusion mean free paths of the positrons are such that the positrons escape along the flux tubes, either up into the halo, or down into the warm ionized envelopes, surrounding the base of the superbubble generated chimneys, where they slow down and annihilate.  In the warm ($\\sim 10^4$ K), but essentially neutral HI gas, which fills the largest fraction of the disk $>$ 3.5 kpc, we show that the cascade of magnetic turbulence, created at parsec scales, is damped by friction between the ions and neutral gas long before the cascading turbulence reaches the positron resonant gyroradii. Thus the positrons suffer much less scattering, and stream along the flux tubes with an isotropic pitch angle distribution. Therefore the positrons also mostly stream out into adjoining phases, either into the halo or into those warm ionized envelopes, separating the ubiquitous warm HI gas from the hot superbubbles, where they slow down and annihilate.  Therefore we find that roughly 1/2 of the positrons, produced anywhere in the disk $>$3.5 kpc, slow down and annihilate there, primarily in the warm ionized gas, while the rest escape from the disk into the halo. This escape thus accounts for the low observed annihilation radiation luminosity of the disk and explains the observed ratio of annihilation rates in the bulge and disk.  Moreover, the efficient diffusion of positrons out of the hot tenuous plasmas in both the bulge and disk, where nearly 70\\% are born, and into the warm ionized envelopes and cold neutral interiors of molecular clouds, where they annihilate, also explains both the observed positronium fraction and broad to narrow 511 keV line flux ratio.  The positrons escaping into the halo from the bulge and disk also lead to the annihilation of roughly 1/3 of all the positrons in the halo, even though there in no significant production there.  \\subsection{Outline }  Although the processes of Galactic positron production, propagation and annihilation are rather straightforward, all the details still make the determination of the expected 511 keV line emission both lengthly and complicated.  We consider positron production from the decay of three separate long-lived radioisotopes, each synthesized by distinct classes of stars: the $^{56}$Ni and $^{44}$Ti, respectively, by SNIa and SNIp from thermonuclear explosions of $\\sim$ Gyr-old accreting white dwarves in close binary systems, and the $^{26}$Al by young ($<$ 40 Myr), massive ($>$ 8 M$_{\\odot}$) stars, most likely through Wolf Rayet (WR) winds and core collapse SNII and SNIbc. The supernovae, SNIa and SNIp, from older stellar populations are known to occur in galactic bulges and disks, while WR stars, SNII and SNIbc, formed in extremely young stellar populations, occur in active, or recent, galactic star-formation sites, such as spiral arms, molecular-cloud complexes, and galactic nuclei. Thus in section 2 we briefly re-examine the production of $\\beta^+$-decay positrons by Galactic supernovae and their spatial and temporal distribution in the light of recent observations and theory.  In section 3 we then explore in detail the expected energy loss, scattering, and propagation processes of the relativistic positrons from these supernovae in the different phases of the interstellar medium. We consider collisionless scattering of these positrons by small-scale fluctuations generated by magnetohydrodynamic turbulence in the ionized interstellar phases, based on extensive studies and observations of relativistic electrons in the analgous turbulence in the interplanetary medium. Within this context we adopt a self-consistent transport model tied to field-aligned turbulence dependent on the properties of the ambient interstellar phases.   In sections 4, 5, and 6 we determine the expected spatial distribution of positron production, propagation, slowing down and annihilation in the various phases of the positron bulge, disk and halo. In each of these we review the fundamental properties of the interstellar medium, which constrain positron transport as well as leave signatures in the positron annihilation line emission features. Here we explore the very large differences between the interstellar properties, and the resulting positron propagation and annihilation, in the well-studied local interstellar medium compared to those in the inner $\\sim$ 1.5 kpc of our Galaxy, where the diffuse 511 keV emission is most intense.  Finally, in sections 7 \\& 8 we give a summary and conclusions, comparing the expected 511 keV line emission with the SPI/INTEGRAL observations. We show that the observed spatial and spectral features of the Galactic positron annihilation radiation can all be fully understood in the context of a supernova radionuclide origin, when the differential propagation of these $\\sim$ MeV positrons in the various phases of the interstellar medium is taken into consideration. ",
        "conclusions": "In conclusion, we find that roughly half of the Galactic positrons from supernova generated radionuclei are produced in he bulge and inner disk within the stellar bulge of $R < 3$ kpc, and half in the stellar disk beyond that. Essentially all of the bulge positrons are produced in the very hot plasma, through which they diffuse with a calculated mean free path of only 1\\% of the size of the region. Therefore, before they can either slow down and annihilate in the plasma or escape beyond, most of them are stopped and annihilate in the very dense ($\\sim$100-1000 H cm$^{-1}$) warm outer shells of the molecular clouds that lie within $R <$ 1.5 kpc. The disk positrons, on the other hand, are produced about equally in the warm HI gas and the hot superbubbles, which blow out into the halo. We calculate that roughly half of the disk positrons, which either stream through the low density ($\\sim$0.1 H cm$^{-1}$) neutral HI at close to $c$, or diffuse through the hot superbubble plasma, escape into the overlying halo before they can stop and annihilate in the low density (0.3-3 H cm$^{-1}$) photoionized outer shells of clouds and superbubbles in the disk. Thus, positron propagation easily explains the larger annihilation flux from the bulge. Moreover, it explains the high observed postronium fraction and narrow observed 511 keV line width, that both require annihilation predominantly in the warm ionized gas and not in the hot plasma in which nearly all of the positrons are produced, either in this or in other suggested sources.  We expect these conclusions to be quite robust, since they are based on 1) extensive observations of galactic supernova rates and distributions, 2) observationally confirmed models of supernova yields of radionuclei and their decay-positron survival fractions, 3) wide ranging observations of the properties and distributions of the various phases of the interstellar medium and magnetic fields, and 4) the widely used photodissociation region (PDR) models of molecular cloud shells. From these we have calculated the expected propagation of positrons in the various phases, assuming streaming in the neutral phases and diffusion in the ionized gas and plasma phases, using diffusion mean free paths in the latter phases calculated from a two-component, phenomenological model of anisotropic turbulence, which is consistent with observations of particle propagation in our best natural laboratory, the interplanetary medium.  Thus, as we show in Table 7, when positron propagation is considered, the positrons from the $\\beta^+$-decay chains of the radioactive nuclei, $^{56}$Ni, $^{44}$Ti, and $^{26}$Al, produced in Galactic supernovae, can fully account for {\\em all} of the features of the diffuse Galactic 511 keV and 3-$\\gamma$ continuum annihilation radiation observed by SPI/INTEGRAL. {\\em We have also predicted additional measurable features that can not only further test such an origin of the positrons but provide new information on the nature of the interstellar medium.}   This work was supported by NASA's {\\em International Gamma-Ray Astrophysics Laboratory} Science Program, grant NNG05GE70G.  \\newpage"
    },
    "0711/0711.1019_arXiv.txt": {
        "abstract": "Observations of the optical polarization of NGC~4151 in 1997--2003 show variations of an order of magnitude in the polarized flux while the polarization position angle remains constant.  The amplitude of variability of the polarized flux is comparable to the amplitude of variability of the total U-band flux, except that the polarized flux follows the total flux with a lag of $8 \\pm 3$ days. The time lag and the constancy of the position angle strongly favor a scattering origin for the variable polarization rather than a non-thermal synchrotron origin.  The orientation of the position angle of the polarized flux (parallel to the radio axis) and the size of the lag imply that the polarization arises from electron scattering in a flattened region within the low-ionization component of the broad-line-region. Polarization from dust scattering in the equatorial torus is ruled out as the source of the lag in polarized flux because it would produce a larger lag and polarization perpendicular to the radio axis.  We note a long-term change in the percentage polarization at similar total flux levels and we attribute this to a change in the number of scatterers on a timescale of years. ",
        "introduction": "\\citet{shklovskii53} pointed out that synchrotron radiation should be highly polarized, and the detection in 1953 of polarization of the Crab Nebula demonstrated the synchrotron nature of its continuum radiation \\citep{dombrovskii58}.  It was therefore natural to search for polarization of the optical emission of Seyfert galaxies \\citep{dibai_shakhovskoy66}.  Optical polarization of NGC~4151 was discovered by \\citet{babadzhanyants_hagen-thorn69} who also reported that the polarization was variable.  The polarization was studied in more detail by \\citet{kruszewski71} who found that the polarization in the $V$ band varied by a factor of four over a nine-month period. The detection of variable polarization was taken as evidence of a synchrotron origin of the continuum.  The discovery of polarization of emission lines \\citep{angel76} in NGC~1068 showed that the polarization of type-2 AGNs was due to scattering and not synchrotron emission.  The origin of the polarization in type-1 AGNs, however, has remained ambiguous. \\citet{stockman79} discovered that the optical polarization of radio-loud AGNs was parallel to the radio jet axis. They considered this to be either the result of intrinsic polarization of optical synchrotron emission, or the result of scattering from a flattened distribution of scatters. Subsequently \\citet{antonucci82, antonucci83} showed that, in general, polarization tends to be parallel to the radio axis in type-1 AGNs and parallel in type-2 AGNs.  Since the polarization of broad lines is roughly similar to the polarization of the continuum \\citep{goodrich_miller94} this showed that this polarization was due to scattering.  Nonetheless because of the variability of optical polarization it has often been assumed, by analogy with blazars (where there is no doubt that the optical polarization variations have a non-thermal origin) that synchrotron radiation also contributes to optical polarization in non-blazar AGNs (e.g., Giannuzzo \\& Salvati 1993).  In this paper we attempt to resolve the question of the nature of the optical polarization by investigating the relationship between the variability of the polarized flux and the total flux. ",
        "conclusions": "The polarized flux of NGC~4151 appears to lag the unpolarized flux by $\\sim 8$ d which is comparable to the light-crossing time of the broad-line region.  We interpret this as the result of the extra light-travel time of the scattered photons. Dust in the torus is ruled out as the source of polarization by the angle of polarization and because the observed lag is too short. Both the direction of polarization and the lag can naturally be explained instead by electron scattering in a flattened region of similar size to the low-ionization BLR.  We have also detected a possible change in the percentage polarization of NGC~4151 at similar flux levels on a timescale of years, and we attribute this to long-term changes in the number of scattering electrons."
    },
    "0711/0711.1196_arXiv.txt": {
        "abstract": "Six glitches have been recently observed in the rotational frequency of the young pulsar PSR B1737$-$30 (J1740$-$3015) using the 25-m Nanshan telescope of Urumqi Observatory. With a total of 20 glitches in 20\\,years, it is one of the most frequently glitching pulsars of the $\\sim$1750 known pulsars. Glitch amplitudes are very variable with fractional increases in rotation rate ranging from $10^{-9}$ to $10^{-6}$. Inter-glitch intervals are also very variable, but no relationship is observed between interval and the size of the preceding glitch. There is a persistent increase in $|\\dot\\nu|$, opposite in sign to that expected from slowdown with a positive braking index, which may result from changes in the effective magnetic dipole moment of the star during the glitch. ",
        "introduction": "Glitches in pulsars are rare and extraordinary events. The study of such glitches and their post-glitch recoveries can give an insight into the interior of neutron stars, the physics of ultradense matter and provide limits on the equation of state \\citep{pin91,accp93,fle00}. Pulsar glitches are usually detected from long-term and frequent timing observations. The number of observed glitches has increased dramatically in the past few years, which probably is due to both the increased continuous observations of active pulsars and a large number of new young pulsars detected by the pulsar surveys \\citep[e.g., the Parkes surveys, see][]{kbm+03} over the last decade.  Glitches are characterized as sudden increases in the rotation frequency $\\nu=1/P$ (where $P$ is the pulsar period), often followed by an interval of approximately exponential recovery or relaxation back towards the pre-glitch frequency. The post-glitch relaxation can have timescale of days to years. Frequency jumps with magnitudes $\\Delta\\nu/\\nu \\sim 10^{-6}$ are recognized as ``giant'' glitches, and have been observed mostly in pulsars with characteristic ages $\\tau_c = P/(2\\dot{P}) \\sim 10^4$~yr, such as PSRs~B0833$-$45 (the Vela pulsar), B1046$-$58 and B1338$-$62 \\citep{wmp+00}. Giant glitches have not been observed in the youngest radio pulsars, such as PSR~B0531+21 (the Crab pulsar) and PSR~B1509$-$58. Detected glitches are generally of magnitude $10^{-9}<\\Delta\\nu/\\nu<10^{-6}$, and the relative increment in slow-down rate $\\Delta\\dot\\nu/\\dot\\nu$ is in the range $10^{-3}$ to $10^{-2}$. Very small glitches with $\\Delta\\nu/\\nu<10^{-9}$ are difficult to distinguish from timing noise in young pulsars \\citep{hlk06}. Of the 1750 known pulsars, about 170 glitches in 54 pulsars have been observed so far\\footnote{See http://www.atnf.csiro.au/research/pulsar/psrcat/ and \\citet{mhth05}}.  Half of the published glitches have fractional amplitudes $\\Delta\\nu/\\nu\\sim 10^{-8}$. Most glitches detected in the Vela pulsar are giant glitches, and the largest glitch, with $\\Delta\\nu/\\nu\\approx16\\times10^{-6}$, was observed in PSR J1806$-$2125 \\citep{hlj+02}. About 60\\% of the 54 glitching pulsars have glitched only once, and 16\\% have glitched twice.  Observations with Rossi X-ray Timing Explorer (RXTE) have revealed that PSR J0537$-$6910, the 16-ms pulsar associated with the supernova remnant N157B in the Large Magellanic Cloud, is the most frequently glitching known pulsar with 23 glitches detected in seven years of monitoring \\citep{mmw+06}. Most of the glitches have $\\Delta\\nu/\\nu$ of a few$\\times 10^{-7}$ and, interestingly, a strong correlation is observed between the amplitude of a glitch and the time to the next one. Furthermore, despite the usual post-glitch recovery (decrease) in $|\\dot\\nu|$, a persistent long-term increase in this parameter is observed, corresponding to a negative braking index $n=\\nu\\ddot\\nu/{\\dot\\nu}^2 \\sim -1.5$. This increase is most probably related to changes on or in the neutron star that result from the glitch activity.  \\begin{figure*} \\centerline{\\psfig{file=glt1.ps,width=60mm,angle=-90} \\psfig{file=glt2.ps,width=60mm,angle=-90} \\psfig{file=glt3.ps,width=60mm,angle=-90}} \\end{figure*} \\begin{figure*} \\centerline{\\psfig{file=glt4.ps,width=60mm,angle=-90} \\psfig{file=glt5.ps,width=60mm,angle=-90} \\psfig{file=glt6.ps,width=60mm,angle=-90}} \\caption{Timing residuals relative to the pre-glitch timing model for six glitches in the rotation history of PSR B1737$-$30. Residuals are defined in the sense of observed arrival time minus predicted arrival time. }\\label{fg:gl_phs} \\end{figure*}   In this paper we report on six recent glitches detected in the period of PSR~B1737$-$30 and discuss the implications of these results. PSR~B1737$-$30 is a young radio pulsar with a characteristic age of $2.06\\times10^{4}$~yr which glitches frequently, with 20 glitches of all sizes ranging from ``small'' to ``giant'' ($\\Delta\\nu/\\nu \\sim 10^{-9}$ to $10^{-6}$) observed since its discovery in 1986 \\citep[][this paper]{ml90,sl96,klgj03}. As for PSR J0537$-$6910, there appears to be a cumulative shift in the spin-down rate of PSR B1737$-$30 resulting from its frequent glitches, which most likely accounts for its long-term braking index of $n = -4\\pm2$ \\citep{ura02}. ",
        "conclusions": "Our observations of PSR B1737$-$30 show six glitches in PSR B1737$-$30 over a period of seven years. In total, 20 glitches have been observed over a 20-year period.  Glitch sizes cover a wide range from just a few parts in $10^9$ to large glitches comparable to those seen in the Vela pulsar. The inter-glitch intervals are highly variable, ranging from 3 weeks to more than 3 years. Unlike in PSR J0537$-$6910, there is no clear relationship of glitch interval with size of the preceding glitch. Although the age of PSR B1737$-$30 is comparable to that of the Vela pulsar, their glitch behaviours are different, with glitch sizes and intervals being much more variable in PSR B1737$-$30. Exponential and quasi-linear relaxations in $\\dot\\nu$ are observed in both pulsars, but with quite different parameters. It is remarkable that the three most highly glitching pulsars have such different glitch and post-glitch properties, yet another example of the diversity and complexity of magnetised neutron stars."
    },
    "0711/0711.0190_arXiv.txt": {
        "abstract": "We report on our \\spz\\  observations of the anomalous X-ray pulsar \\axp, made following a large X-ray burst that occurred on 2007 February 7. To search for mid-infrared flux variations, four imaging observations were carried out at 4.5 and 8.0 $\\mu$m with the Infrared Array Camera from February 14 to 21. No significant flux variations were detected, and the average fluxes were 32.1$\\pm$2.0 $\\mu$Jy at 4.5 $\\mu$m and 59.8$\\pm8.5$ $\\mu$Jy at 8.0 $\\mu$m, consistent with those obtained in 2005.  The non-detection of variability is interesting in light of reported rapid variability from this source in the near-infrared, but consistent with the fact that the source already went back to its quiescent state before our observations began, as indicated by contemporaneous X-ray  observations. In order to understand the origin of the near-infrared variability, frequent, simultaneous multi-wavelength observations are needed. ",
        "introduction": "The Anomalous X-ray Pulsars (AXPs) are a small group of isolated, young neutron stars having X-ray luminosities greater than their rotational energy loss rates. They are believed to be magnetars, neutron stars possessing $\\sim$10$^{14}$ G magnetic fields, with X-ray emission powered by the decay of their superstrong magnetic fields \\citep{td96}. The magnetar nature of AXPs has been strongly supported by their high spin-down rates and similar radiative properties to soft gamma repeaters, including short X-ray bursts (\\citealt{wt06}; \\citealt{kas07}). The bursting activity likely reflects structural changes in the surface magnetic field of a magnetar.  As the brightest and nearest among the currently known AXPs, \\axp\\ has been well studied over wavelength ranges from the X-ray to mid-infrared (MIR).  Its X-ray emission in the 0.5--10 keV range can be well described by a blackbody ($kT$= 0.5 keV) plus power law (photon index $\\Gamma$ = 3.4; e.g., \\citealt{pat03}), which is believed to arise from the surface and magnetosphere of the star (e.g., \\citealt*{tlk02}).  Its optical emission is pulsed at the 8.7-s spin period (\\citealt{km02, dhi+05}) and appears to have a power-law--like spectrum \\citep*{hvk00}, likely also originating from the magnetosphere (although no apparent connection can be established between the optical and X-ray spectra). A surprise came from the detection of the source in the MIR (4.5 and 8.0 $\\mu$m), which revealed a rising spectrum from the near-IR (NIR) to MIR \\citep*{wck06}. The detection can be interpreted as a surrounding debris disk \\citep{wck06}. The putative disk, presumably formed from fallback material after the supernova explosion, is irradiated by the X-rays from the central pulsar and emits mainly in the MIR. The existence and appearance of such a disk has been predicted (e.g., \\citealt*{lwb91, phn00}).  Recently, optical/NIR flux variability from \\axp\\ has been reported \\citep{dv06}. In a total of only 9 observations, significant flux variations were detected, suggesting such variability is very common. This variability is intriguing, since no such flux changes have been detected in X-rays, even though X-ray observations of the source have been made much more frequently than the optical/NIR observations. An obvious question is whether there is similar variability in the MIR. Rapid variability in the MIR would be very hard to understand in the disk model in the absence of X-ray variability.  On 2007 February 7, a large fast-rise X-ray burst was detected during the {\\em Rossi X-Ray Timing Explorer} ({\\em RXTE}) monitoring observations of \\axp\\ \\citep{gav+07}. The peak flux of the burst was more than two orders of magnitude larger than that in the source's quiescent state. This event provided an opportunity to test the fallback disk model, since a MIR flux increase is expected from an X-ray irradiated disk when the input X-ray flux is increased (see Figure \\ref{fig:disk}). For the purposes of searching for variability and testing the fallback disk model, we observed \\axp\\  with the {\\em Spitzer Space Telescope}. In this paper, we report on the results of the observations. ",
        "conclusions": "\\label{sec:disc}  There are a total of 11 reported $K$-band observations of \\axp\\ in the past 9 years, and three of them were with simultaneous $JH$ observations \\citep{dv06,gon+07}. In the observations, the AXP exhibited large and rapid flux variations in the $K$ band, and the $JH$ fluxes correlated with the $K$ but varied less. Assuming that the $K$-band variations are sufficiently sampled, the probability of seeing a $>$15\\% flux change would be 60\\%. In one extreme case, a 40\\% $K$ flux decrease occured within one day. Our \\spz\\  observations were planned so that such strong variability would have been detected. Correlated flux changes between the MIR and $K$ bands are expected in the proposed dust disk model for \\axp\\ (Figure~\\ref{fig:disk}; \\citealt{wck06}), if X-ray flux changes are considered as the primary cause of the variability.  We did not detect the $K$-like flux variations in our MIR observations. The \\spz/IRAC fluxes were consistent with being a constant within the uncertainties.  Given that the upper limit on the MIR 4.5 $\\mu$m flux variations was $\\lesssim 13$\\%, the MIR emission variability was similar to the NIR $JH$ bands and less variable than in the $K$ band. We note that in a $K$-band observation made one day prior to our first \\spz\\ observation, the source had $K=19.9\\pm0.1$ mag \\citep{gon+07}, approximately the average $K$ magnitude of \\axp. Because the time scale of the $K$ flux variations is not determined, this could suggest that the source was in a stable state in the NIR during our observations, which could be a reason for the absence of MIR variability. To further constrain the relation between the MIR and NIR emission, frequent, simultaneous observations are needed.  If the MIR emission is truly less variable, the restricted $K$-band variability would demand explanation. Interestingly, similar NIR variability is seen in a few dust disk systems around young stars (e.g., \\citealt{eir+02}). In those systems, the variability has been suggested to be caused by structural changes of the inner disk. If this is the case for \\axp, the properties in the source's IR variability might be explainable. Since emission in the $K$ band would mainly arise from the inner edge of the disk and thus be more sensitive to inner disk changes, it would vary more strongly than emission in the $JH$ and \\spz\\ MIR bands. It is not clear what would be the cause of the structural changes in the disk in \\axp.  One possibility is variable X-ray heating, such as due to burst-related X-ray brightening. However, no significant short-term flux variations have been found in the X-ray observations of the source \\citep{dv06}, and thus far only a few bursts have been found in bi-weekly \\textit{RXTE} X-ray monitoring observations \\citep{dkgw06,gav+07}, indicating infrequency of the burst events. Thus in the absence of X-ray variability, the $K$-band variability challenges the dust disk model \\citep{dv06}.  In addition, we did not detect any significant MIR flux changes from \\axp\\  in our \\spz\\ IRAC observations following the large X-ray burst. The 4.5/8.0 $\\mu$m fluxes were nearly the same as those obtained in 2005. As demonstrated by X-ray observations of the source that were made several days prior to our first \\spz\\ observation (Gonzalez et al. 2007), the AXP unfortunately had already returned to quiescence during our observations. Also as shown in the long-term X-ray monitoring observations, the quiescent X-ray flux of \\axp\\  has been nearly constant and stable over the past decade (\\citealt*{dkg07}; Gonzalez et al. 2007). Therefore, our non-detection is consistent with the source's state at X-ray energies at the time."
    },
    "0711/0711.2689_arXiv.txt": {
        "abstract": "We study cosmological perturbations in the context of an interacting dark energy model, in which the cosmological term decays linearly with the Hubble parameter, with concomitant matter production. A previous joint analysis of the redshift-distance relation for type Ia supernovas, barionic acoustic oscillations, and the position of the first peak in the anisotropy spectrum of the cosmic microwave background has led to acceptable values for the cosmological parameters. Here we present our analysis of small perturbations, under the assumption that the cosmological term, and therefore the matter production, are strictly homogeneous. Such a homogeneous production tends to dilute the matter contrast, leading to a late-time suppression in the power spectrum. Nevertheless, an excellent agreement with the observational data can be achieved by using a higher matter density as compared to the concordance value previously obtained. This may indicate that our hypothesis of homogeneous matter production must be relaxed by allowing perturbations in the interacting cosmological term. ",
        "introduction": "The cosmological constant problem has acquired a renewed importance since several independent observations have been pointing to the presence of a negative pressure component in the cosmic fluid \\cite{Pad}. From the point of view of quantum field theories, the natural candidate for such a dark energy is the quantum vacuum. Since, at the macroscopic level, it has the symmetry of the background, its energy-momentum tensor has the form $T_{\\mu}^{\\nu}= \\Lambda g_{\\mu}^{\\nu}$, where $\\Lambda$ is a scalar function of coordinates. This leads, in the case of an isotropic and homogeneous space-time and co-moving observers, to the equation of state $p_{\\Lambda} = - \\rho_{\\Lambda} = - \\Lambda$, where $\\Lambda$ may be, in general, a function of time. In the case of a constant $\\Lambda$, the vacuum contribution plays the role of a cosmological constant in Einstein's equations.  However, the estimation of the vacuum energy density by quantum field theories in the flat space-time leads, after some regularization procedure, to a very huge result when compared to the observed value. A possible way out of this difficult is to argue that such a result is valid only in a flat background, in which the very Einstein equations predict a null total energy-momentum tensor. Therefore, the huge vacuum density should be canceled by a bare cosmological constant, like in a renormalization process. Now, if we could obtain the vacuum density in the FLRW space-time, after the subtraction of the Minkowskian result it would remain an effective time-dependent $\\Lambda$ term, which decreases with the expansion.  The idea of a time-dependent cosmological term has found different phenomenological implementations \\cite{Ozer}, being a subject of renewed interest in recent years \\cite{Schutzhold,Horvat,Fabris}. A general feature of all those approaches is the production of matter, concomitant with the vacuum decay in order to assure the covariant conservation of the total energy \\cite{Barrow}. Indeed, in the FLRW space-time, the Bianchi identities lead to the conservation equation \\begin{equation} \\dot{\\rho}_T+3H(\\rho_T+p_T)=0, \\label{continuidade} \\end{equation} where $\\rho_T$ and $p_T$ stand for the total energy density and pressure, respectively, and $H = \\dot{a}/a$ is the Hubble parameter. By writing $\\rho_T = \\rho_m + \\Lambda$ and $p_T = p_m - \\Lambda$ (where $\\rho_m$ and $p_m$ are the energy density and pressure of matter), the above equation reduces to \\begin{equation} \\dot{\\rho}_m + 3H(\\rho_m + p_m)=-\\dot{\\Lambda}, \\label{continuidade3} \\end{equation} which shows that, in the case of a varying $\\Lambda$, matter is not independently conserved\\footnote{Properly speaking, we should also consider the pressure and energy associated to the very process of matter production, that is, the energy-momentum tensor of the interaction between matter and vacuum. In this sense, decaying vacuum models do not differ essentially from interacting dark energy models \\cite{Zimdahl}, with the scalar function $\\Lambda$ replaced by a scalar field interacting with matter. Nevertheless, if the vacuum decays into non-relativistic particles, as we will consider here, the interaction term can be neglected, and the above decomposition may be considered a good approximation.}.  An important point to be clarified in this kind of model is the homogeneity of matter production. Of course, in a strictly homogeneous space-time the production is homogeneous, since $\\rho_m$ and $\\Lambda$ depends only on time. But, in the presence of density perturbations, is the new matter produced homogeneously, or just where matter already exists \\cite{Dirac}? In the case of a homogeneous production, the new matter tends to dilute the density perturbations, leading to a suppression of the density contrast. In some models, this suppression is strong enough to impose very restrictive observational limits to them \\cite{Opher}.  In this paper we will analyze the evolution of density perturbations in a particular, spatially flat, cosmological model with vacuum decay \\cite{Borges,Jailson}. It can be based on a phenomenological prescription for the variation of $\\Lambda$ with time \\cite{SC}, given by $\\Lambda \\approx (H + m)^4 - m^4$, where $m$ is a characteristic energy that can be identified with the scale of the QCD vacuum condensation, the latest cosmological vacuum transition. Although it can be corroborated by holographic arguments \\cite{SC,SC2}, based on the thermodynamics of de Sitter space-times, here we will take it just as a phenomenological ansatz. In the limit of very early times, we have $\\Lambda \\approx H^4$, which provides a non-singular inflationary solution \\cite{SC}.  In the opposite limit of large times we have $\\Lambda = \\sigma H$, with $\\sigma \\approx m^3$. This scaling law for the vacuum density was also suggested in \\cite{Schutzhold}, on the basis of different arguments. It leads to a cosmological scenario in qualitative agreement with the standard one \\cite{Borges}, with an initial radiation era followed by a long phase dominated by dust. This dust phase tends asymptotically to a de Sitter universe, with the deceleration/acceleration transition occurring some time before the present epoch. On the other hand, a quantitative analysis has shown a good accordance with supernova observations, leading to age and matter density parameters inside the limits imposed by other independent observations \\cite{Jailson}.  Since the radiation phase we obtain is indistinguishable from the standard one, our analysis will be initially focused on the evolution of density perturbations of non-relativistic matter in the dust-dominated phase, considering wavelengths inside the horizon\\footnote{As already commented, we will assume that the vacuum is decaying into non-relativistic particles, in order to avoid any conflict with CMB observations and with the observed coldness of dark matter. We will also suppose that only dark matter is produced, since the baryon content is well constrained by nucleosynthesis. Evidently, these assumptions cannot be verified without a microscopic theory of the vacuum-matter interaction.}. In this way, it will be possible to make use of a generalization of the Newtonian linear treatment of the problem, which includes the effects of matter production \\cite{Waga}. We will show that, even in the case of a homogeneous vacuum decay, the contrast suppression is important only for late times, not affecting the process of galaxy formation. On the other hand, it dominates for future times, and we will discuss how this behavior can possibly alleviate another problem related to the cosmological term: the cosmic coincidence problem.  Subsequently, a relativistic analysis will be performed, in order to construct the matter power spectrum. Again, the hypothesis of homogeneous matter production will be used, leading as well to a consequent power suppression. A second interesting difference as compared to the $\\Lambda$CDM model is a shift of the spectrum turnover to the left, that is, to smaller wavenumbers. The late-time suppression is not very sensitive to the value used for the matter density, a feature that can already be noted from the Newtonian analysis. On the other hand, the correction of the turnover position, by taking a higher matter density, displaces all the spectrum to the right, compensating the late-time power suppression. In this way we can obtain an excellent fit of data, but with a higher matter density in comparison with the standard case.  The article is organized as follows. In next section we review the main features of our interacting model. In Section III we perform the Newtonian analysis of evolution of density perturbations in the matter era. In Section IV the matter power spectrum is constructed, on the basis of a simplified relativistic calculation. In Section V the reader can find our concluding remarks. ",
        "conclusions": "In spite of the physical plausibility of a time dependent cosmological term, a complete theoretical development of this idea, including the microscopic details of the vacuum-matter interaction, is still lacking. On the other hand, macroscopic approaches depend on some phenomenological hypothesis, leading some times to diverse prescriptions for the vacuum decay\\footnote{For example, the linear dependence of $\\Lambda$ on the Hubble parameter we use here contrasts with the quadratic dependence used in reference \\cite{Fabris}. In that work, the quadratic dependence is due to the computation of quantum effects of matter field in a cosmological background, which leads to a running cosmological term. The authors also used the matter power spectrum data to constrain the fundamental parameters of the quantum model.}. For this reason, a careful comparison with current observations is very important, playing the role of corroborating or ruling out the different models.  We have already analysed the supernova observations \\cite{Jailson}, obtaining good fits and cosmological parameters in accordance with other independent tests, as the age of globular clusters and dynamical limits to the matter density \\cite{dynamical}. Other precise tests, as the position of the first acoustic peak of the cosmic microwave background and the baryonic acoustic oscillations have also been performed \\cite{Jailson2}, showing a good concordance when jointed to the supernova analysis.  In the present paper we have studied the evolution of matter density perturbations, in particular the contrast suppression associated to the process of matter production. We have shown that, even in the case of a homogeneous production, the evolution of the contrast is the same as in the standard recipe along the entire era of galaxy formation, diverging from the later only for $z < 5$.  On the other hand, the suppression would be dominant for future times, and this may have an interesting relation with another problem related to the cosmological term, namely the approximate coincidence between the present densities of matter and dark energy. Indeed, we can see from figures 1 and 2 that the matter contrast has its maximum just before today, when matter and vacuum give similar contributions to the total density. The largest structures formed until now tend to disaggregate in the future, and their existence then coincides with the time of approximate equality between the matter and vacuum densities.  This could alleviate the cosmic coincidence problem, if galaxies also follow such a process. However, we should remember that galaxies have left the linear regime of growth a long time ago, and that now their evolution is non-linear, driven essentially by their self-gravitation. Therefore, an explanation of the cosmic coincidence in the terms above will depend on a non-linear study of density perturbations in the context of the present model. Only such an investigation would tell us whether the contrast suppression described here can affect smaller structures like galaxies.  Its also important to have in mind that the homogeneity of the matter production, implicit in the derivation of solution (\\ref{solution}) and in our simplified relativistic treatment, is just an {\\it ad hoc} hypothesis, to be verified from both the theoretical and observational viewpoints. For a constant, non-interacting vacuum term it is certainly true, but not necessarily in the present case. Any inhomogeneity of the vacuum density around matter distributions may lead to an inhomogeneous production, reducing in this way the contrast suppression. This would allow us to fit the observed power spectrum with a smaller matter density, closer to the concordance value obtained in \\cite{Jailson2}. Whether the matter contrast will still have a maximum around the present time, with the discussed implications for the coincidence problem, is a matter of investigation. A relativistic study of this case, that is, with $\\delta \\Lambda \\neq 0$, is already in progress."
    },
    "0711/0711.2530_arXiv.txt": {
        "abstract": "{The $\\beta$~Cephei stars have been studied for over a hundred years. Despite this, many interesting problems related to this class of variable stars remain unsolved.  Fortunately, these stars seem to be well-suited to asteroseismology. Hence, the results of seismic analysis of $\\beta$~Cephei stars should help us to better understand pulsations and the main sequence evolution of massive stars, particularly the effect of rotation on mode excitation and internal structure. It is therefore extremely important to increase the sample of known $\\beta$~Cephei stars and select targets that are useful for asteroseismology.} {We analysed ASAS-3 photometry of bright early-type stars with the goal of finding new $\\beta$~Cephei stars. We were particularly interested in $\\beta$~Cephei stars that would be good for seismic analysis, i.e., stars that (i) have a large number of excited modes, (ii) show rotationally split modes, (iii) are components of eclipsing binary systems, (iv) have low-frequency modes, that is, are hybrid $\\beta$~Cephei/SPB stars.} {Our study was made with a homogeneous sample of over 4100 stars having MK spectral type B5 or earlier. For these stars, the ASAS-3 photometry was analysed by means of a Fourier periodogram.} {We have discovered 103 $\\beta$~Cephei stars, nearly doubling the number of previously known stars of this type. Among these stars, four are components of eclipsing binaries, seven have modes equidistant or nearly equidistant in frequency. In addition, we found five $\\beta$~Cephei stars that show low-frequency periodic variations, very likely due to pulsations. We therefore regard them as candidate hybrid $\\beta$~Cephei/SPB pulsators. All these stars are potentially very useful for seismic modeling. Moreover, we found $\\beta$~Cephei-type pulsations in three late O-type stars and fast period changes in one, HD\\,168050.} {} ",
        "introduction": "This is the second paper in the series presenting results of a search for $\\beta$~Cephei-type pulsating variables in the $V$-filter photometry in the third phase of the All Sky Automated Survey, ASAS-3 \\citep{pojm97,pojm00,pojm05}. In the first paper of the series \\citep[][hereafter Paper I]{pipo07}, we published the results of an analysis of the ASAS-3 photometry for 41 of about 112 presently known $\\beta$~Cephei-type stars, that is, all that had sufficiently good photometry in the ASAS-3 database. The analysis allowed us to detect amplitude changes in three $\\beta$~Cephei stars, BW~Cru, V836~Cen and V348~Nor, and period changes in KK~Vel and V836~Cen. In addition, new modes were found or ambiguities in the frequencies of previously detected modes were removed for some stars.  The published ASAS catalogues, containing photometry of variable star candidates, have already been used to discover $\\beta$~Cephei stars. In total, 19 new members of this group have been found in the ASAS data \\citep{pigu05,hand05}. However, owing to the fact that the method of selection of variable stars was biased towards large-amplitude variables, all these 19 stars had relatively large amplitudes. Therefore, it was obvious that once a detailed analysis of the whole ASAS-3 database has been done, many new stars of this type, with smaller amplitudes, will be found. Such an analysis will be carried out in the future, but because the whole ASAS database contains millions of stars and the period(s) alone are not sufficient to classify a given star as a $\\beta$~Cephei variable, we decided to proceed in a different way. Using the information in the existing stellar catalogues, we selected stars that are known to be either O-type or early B-type stars and analysed ASAS-3 photometry for these stars only. The selection criteria are explained in Sect.~2. As a result, nearly 300 $\\beta$~Cephei stars were discovered of which the first part is discussed in this paper.  This will allow us to define the properties of the whole class much better and, in view of the increasing interest in the study of these stars by means of asteroseismology, to select appropriate targets for subsequent work. In the present paper we describe the results of the search for $\\beta$~Cephei-type stars among objects selected from the five volumes of the Michigan Catalogue for HD stars.  The ASAS-3 photometry for stars selected using other catalogues will be discussed in the next paper. ",
        "conclusions": "According to theory \\citep[see, e.g.,][]{pamy99}, for the zero-age main-sequence models of massive stars, acoustic (p) and gravity (g) modes separate very well in frequency. This separation is not so good for evolved stars because in the course of the main-sequence evolution frequencies of g modes increase and eventually replace those of p modes through avoided crossing \\citep{dzpa93}. Such modes have a `mixed' character: of a p mode in the envelope and of a g mode in the stellar interior. When only unstable modes are considered, however, the separation in frequency of p and g modes is, for a given model, preserved. Pulsations in p modes and mixed modes are attributed to $\\beta$~Cephei-type variability, whereas those in g modes, to the SPB type. As $\\beta$~Cephei stars populate mainly the later stages of main-sequence evolution, mixed modes are expected to be excited in these stars.  Theoretical calculations show also \\citep{pamy99} that $\\beta$~Cephei and SPB instability strips partially overlap, which means that hybrid stars, showing pulsations of both types, might exist. This possibility was already indicated by \\citet{dzie93} for a region around spectral type B2-B3 and masses equal to 7--8~$M_\\odot$. Newer calculations \\citep{pamy99} confirmed this and showed in addition that g-mode instability extends towards higher masses. The latter instability is confined, however, to the very late main-sequence stages.  This picture of p and g mode instability domains was calculated neglecting the effects of rotation on pulsations. The theory that includes fast rotation is still not fully satisfactory, but it is already clear that rotation changes stability of modes. A good example of the theoretical calculations that include the effect of rotation on the stability of g modes was published by \\citet{town05}, who showed that for fast-rotating stars the SPB instability strip widens towards higher masses. The region of possible hybrid behaviour becomes therefore much larger.  In addition, the frequencies of modes, especially of non-axisymmetric modes, change considerably. In particular, frequencies of retrograde sectoral mixed modes are shifted to much lower frequencies than they would have in non-rotating stars, i.e., will occur in the range where normally the g modes are observed. On the other hand, frequencies of prograde sectoral g modes shift to the region occupied normally by p modes. It must be also remembered that in the presence of fast rotation modes can no longer be classified in a simple scheme described by three quantum numbers because the geometry of such modes differs from spherical harmonics.  The effects related to fast rotation are impontant not only because some $\\beta$~Cephei stars are known to rotate fast. The region of $\\beta$~Cephei and SPB stars in the H-R diagram is populated by Be stars which, by definition, rotate very fast. These stars were known to show both short-period variability and line-profile variations. The short-period photometric variability called, after the prototype, the $\\lambda$~Eri-type \\citep{balo90} was usually attributed to rotationally-induced effects because no clear evidence for multiperiodicity was found \\citep[see, e.g.,][]{baja02}. On the other hand, there were many convincing examples of line-profile variations in Be stars that were interpreted in terms of non-radial pulsations \\citep{baad84,yang90,stef95,floq00}, including multiperiodic behaviour \\citep{rivi98}. Despite this, it was not widely accepted that non-radial pulsations are present in Be stars mainly because there was no clear relation between periods seen in photometry and spectroscopy. An excellent review of variabilities observed in Be stars was recently given by \\citet{pori03}.  A strong argument in favour of non-radial pulsations in Be stars, as seen in photometry, came from the observations carried out by the MOST satellite, in particular, those of an early Be star, HD\\,163868, in which \\citet{walk05-Be} found over 60 periodicities in three distinct groups. A theoretical interpretation of this variability \\citep{walk05-Be,dzie07,savo07} includes pulsations in g modes. It is therefore reasonable to use the term `SPBe stars' for Be stars with detected g-mode pulsations, as proposed by \\citet{walk05-Be}, although it has to be remembered that many previous investigators already indicated a common cause of periodic variability in SPB and Be stars.  The two grouping of modes seen in HD\\,163868 around 1.7 and 3.4 d$^{-1}$ might be a clue to understanding a grouping of modes in $\\beta$~Cephei and SPB candidates in the LMC and SMC, found by \\citet{kola04,kola06}. In many of these stars, the frequencies of modes grouped around two values whose ratio was close to 2. Whether fast rotation is responsible for this could be verified by future spectroscopic observations. For the SMC we know that many of these stars are fast-rotating because they show emission in Balmer lines.  Since inclusion of fast rotation enhances the possibility of getting unstable low-frequency modes among early and mid-B stars, we should see them observationally. It seems that we get more and more examples that confirm these predictions. First, low-frequency modes were detected in two classical $\\beta$~Cephei stars, $\\nu$~Eri \\citep{jerz05} and 12~Lac \\citep{hand06}. In the ASAS-3 data (Sect.~3.4), we found five $\\beta$~Cephei stars in which low-frequency modes were detected. In addition, another $\\beta$~Cephei star with low-frequency modes, HD\\,133823, was found in Paper I (see also Fig.~\\ref{fdf}). On the other hand, periodic variations with very short periods were found in some mid and late B-type stars like V2104\\,Cyg \\citep{uytt07} and stars in NGC\\,3293 \\citep{hand07}. A preliminary analysis of the data obtained for the other young open cluster, h Persei (NGC\\,869) also reveals the presence of mid-B stars, of which some are known as Be stars, showing relatively short periods, even below 0.3~d (Majewska-\\'Swierzbinowicz et al., in preparation).  The complexity of mode behaviour in the presence of fast rotation may pose a problem as far as the classification of hot pulsators is concerned. This is because period(s) alone may not be sufficient to recognize $\\beta$~Cephei and SPB star.  Mode identification becomes a neccessity if we want to classify stars pulsating in p modes as $\\beta$~Cephei stars, while those pulsating in g modes, as SPB stars.  The other important results that come from the discovery of 103 $\\beta$~Cephei stars can be summarized as follows: \\begin{enumerate} \\item The number of known $\\beta$~Cephei stars has grown to over 200 and thus some statistical studies become possible. The sample of new stars is particularly well suited for this purpose because the stars were selected and observed in a homogeneous way. \\item From the new sample, a selection of excellent targets for asteroseismology can be done. These would include stars with large amplitudes (Sect.~3.2), large number of modes, equidistant multiplets (Sect.~3.3) and components of eclipsing binaries (Sect.~3.7). \\item We have found three O-type $\\beta$~Cephei stars (Sect.~3.6); this type of pulsation was also found in $\\zeta$~Oph from the MOST data \\citep{walk05}.  Thus, we may conclude that $\\beta$~Cephei-type instability extends towards O-type stars, but seems to be confined to very late subtypes (O8-09). This still does not agree well with the theory because calculations do not predict an upper luminosity limit for p-mode instability \\citep{pamy99}. However, these calculations have not included the effect of stellar wind which may stabilize pulsations for hotter O-type stars. \\item The very recent stability calculations using new opacities \\citep{migl07} indicate a much wider region in the H-R diagram where hybrid, $\\beta$~Cephei/SPB-type behaviour, should be observed. The inclusion of fast rotation seems to have a similar effect \\citep{town05}.  Observationally, this fact is confirmed by the growing number of detections of hybrid stars, including five $\\beta$~Cephei stars with long periods we found (Sect.~3.4) as candidates. In order to verify their hybridity, mode identification is needed. A detailed study of these stars may allow to testing the physics used in calculations of the evolutionary models of massive stars but also in the theoretical predictions of their pulsational stability. \\end{enumerate} Finally, it is interesting to note that four stars from our sample, HD\\,48553, 171305, 173006, and 186610 are located within the `CoRoT eyes' \\citep{bagl02} and, if selected for observations by this satellite, could be studied in unprecedented detail."
    },
    "0711/0711.0703_arXiv.txt": {
        "abstract": "{The fast classification of new variable stars is an important step in making them available for further research. Selection of science targets from large databases is much more efficient if they have been classified first. Defining the classes in terms of physical parameters is also important to get an unbiased statistical view on the variability mechanisms and the borders of instability strips.} {Our goal is twofold: provide an overview of the stellar variability classes that are presently known, in terms of some relevant stellar parameters; use the class descriptions obtained as the basis for an automated `supervised classification' of large databases. Such automated classification will compare and assign new objects to a set of pre-defined variability training classes.} {For every variability class, a literature search was performed to find as many well-known member stars as possible, or a considerable subset if too many were present. Next, we searched on-line and private databases for their light curves in the visible band and performed period analysis and harmonic fitting. The derived light curve parameters are used to describe the classes and define the training classifiers.} {We compared the performance of different classifiers in terms of percentage of correct identification, of confusion among classes and of computation time. We describe how well the classes can be separated using the proposed set of parameters and how future improvements can be made, based on new large databases such as the light curves to be assembled by the CoRoT and Kepler space missions.} {The derived classifiers' performances are so good in terms of success rate and computational speed that we will evaluate them in practice from the application of our methodology to a large subset of variable stars in the OGLE database and from comparison of the results with published OGLE variable star classifications based on human intervention. These results will be published in a subsequent paper.} ",
        "introduction": "The current rapid progress in astronomical instrumentation provides us with a torrent of new data. For example, the large scale photometric monitoring of stars with ground-based automated telescopes and space telescopes delivers us large numbers of high quality light curves. The HIPPARCOS space mission is an example of this and led to a large number of new variable stars discovered in the huge set of light curves. In the near future, new space missions will deliver even larger numbers of light curves of much higher quality (in terms of sampling and photometric precision).  The CoRoT mission (Convection Rotation and planetary Transits, launched on $27$ December $2006$) has two main scientific goals: asteroseismology and the search for exoplanets using the transit method. The latter purpose requires the photometric monitoring of a large number of stars with high precision. As a consequence, this mission will produce excellent time resolved light curves for up to $60000$ stars with a sampling rate better than $10$ minutes during $5$ months. Even higher numbers of stars ($>100000$) will be measured for similar purposes and with comparable sampling rate by NASA's Kepler mission (launch end $2008$, duration $4$ years). The ESA Gaia mission (launch foreseen in $2011$) will map our Galaxy in three dimensions. About one billion stars will be monitored for this purpose, with about $80$ measurements over $5$ years for each star.  Among these large samples, many new variable stars of known and unknown type will be present. Extracting them, and making their characteristics and data available to the scientific community within a reasonable timescale, will make these catalogues really useful. It is clear that automated methods have to be used here. Mining techniques for large databases are more and more frequently used in astronomy. Although we are far from reproducing capabilities of the human brain, a lot of work can be done efficiently using intelligent computer codes.  In this paper, we present automated supervised classification methods for variable stars. Special attention is paid to computational speed and robustness, with the intention to apply the methods to the huge datasets expected to come from the CoRoT, Kepler and Gaia satellite missions. We tackle this problem with two parallel strategies. In the first, we construct a Gaussian mixture model. Here, the main goals are to optimize speed, simplicity and interpretability of the model rather than optimizing the classifiers' performance. In the second approach, a battery of state-of-the-art pattern recognition techniques is applied to the same training set in order to select the best performing algorithm by minimizing the misclassification rate. The latter methods are more sophisticated and will be discussed in more detail in a subsequent paper (Sarro et al., in preparation).  For a supervised classification scheme, we need to predefine the classes. Every new object in a database to be classified will then be assigned to one of those classes (called definition or training classes) with a certain probability. The construction of the definition classes for stellar variability is, therefore, an important part of this paper. Not only are these classes necessary for this type of classification method, they also provide us with physical parameters describing the different variability types. They allow us to attain a good view on the separation and overlap of the classes in parameter space. For every variability class, we derive descriptive parameters using the light curves of their known member stars. We use exclusively light curve information for the basic methodology we present here, because additional information is not always available and we want to see how well the classes can be described (and separated) using only this minimal amount of information. This way, the method is broadly applicable. It is easy to adapt the methods when more information such as colors, radial velocities, etc. is available.  The first part of this paper is devoted to the description of the stellar variability classes and the parameter derivation. The classes are visualized in parameter space. In the second part, a supervised classifier based on multivariate statistics is presented in detail. We also summarize the results of a detailed statistical study on Machine Learning methods such as Bayesian Neural Networks. Our variability classes are used to train the classifiers and the performance is discussed. In a subsequent paper, the methods will be applied to a large selection of OGLE (Optical Gravitational Lensing Experiment) light curves, while we plan to update the training classes from the CoRoT exoplanet light curves in the coming two years. ",
        "conclusions": "We presented a uniform description of the most important stellar variability classes currently known. Our description is based on light curve parameters from well-known member stars for which high quality data are available. The parameters are derived using Fourier analysis and harmonic fitting methods, and can be calculated on short timescales. The class descriptions obtained form the basis for a supervised classification method which produces class probabilities given a set of time series attributes. It is shown that our class descriptions are accurate enough to separate the monoperiodic variables, some of the multiperiodic variables, and eclipsing binaries. An obvious improvement to these capabilities will come from the addition of color information to the class definitions. This will be discussed in a subsequent paper, where our methodology will be applied to the OGLE database.  We obtained overall good classification results. The Gaussian mixture method is relatively simple and robust, and allows for an easy astrophysical interpretation. The machine learning algorithms, on the other hand, achieve a lower rate of misclassifications at the expense of longer training times, reduced interpretability and a higher level of statistical knowledge of the user. The following extensions/improvements are planned for the future: \\begin{itemize} \\item Extending the number of variability classes and the number of member stars when more and better quality light curves become available, e.g.\\ those from the exoplanet field data of the CoRoT mission. In particular, we will add the classes of stars with transits due to exoplanets, main-sequence stars with solar-like oscillations and stars with magnetic activity. \\item Improve the description of light curves by using methods other than Fourier analysis. Wavelet analysis, e.g., may be more suitable for describing non-periodic variables. Also, additional information derived from the shape of the power spectrum will be considered. \\item Adapt our codes with the intention to apply them to future large scale databases, such as those to be assembled by CoRoT, Kepler and Gaia. In particular, we are now preparing the classification of light curves which will be measured in the framework of the CoRoT Exoplanet programme. \\item Introduce cost matrices to generate specialized classifiers. When the goal of a classifier is to generate clean samples of a given class, it is generally preferred that the number of false positives be diminished even at the expense of missing some of the less clear candidates. In these cases, the introduction of cost matrices in machine learning algorithms allows for the differential weighting of errors in the training process, resulting in classifiers specialized in particular tasks. \\end{itemize} The results presented here are a brief summary of all the experiments and analyses that were applied to the data. For example, only summaries of the performance measures of some of the approaches taken were included. The full statistical analysis and detailed research to characterize the confusion regions of the classifiers will be published in a specialized journal (Sarro et al., in preparation). In particular, questions about the general properties of the subsamples of class $i$ misclassified as $j$, their class probability distributions and if they constitute a separable subset of class $i$ were not discussed here, to avoid entering into a highly technical discussion. Given the large number of classes that constitute this classification problem, it is difficult to include the answers to these questions for the more than 400 possible combinations of $i$ and $j$ but, at the same time, they are of paramount importance for the correct interpretation of classifier results. This analysis is available from the authors upon request.  It is important to realize that the methodology presented here should be evaluated in a statistical sense, i.e., one can never be sure that each individual star is correctly and unambiguously classified. Our method was specifically designed for databases that are so large that individual inspection of all the stars is impossible. Of course, such inspection can and should be done after a first classification with our methods has been achieved, for the specific class of interest to the user. We also note that our basic methods described here assume the simplest possible input: single-band photometric time series. Any additional independent information (color indices or time series, radial velocity time series, spectral type or $\\log g$, etc.) will imply an improved performance as will be shown in our application to the OGLE database for which such additional attributes were retrieved through the Virtual Observatory.  We will judge the performance of the different classifiers presented here in a future paper by comparing our classifications for the OGLE stars with published results obtained with extractor-type methods requiring human interaction."
    },
    "0711/0711.3120_arXiv.txt": {
        "abstract": "We present observations between 14.2 and 17.9\\,GHz of sixteen Galactic {\\sc Hii} regions made with the Arcminute Microkelvin Imager (AMI). In conjunction with data from the literature at lower radio frequencies we investigate the possibility of a spinning dust component in the spectra of these objects. We conclude that there is no significant evidence for spinning dust towards these sources and measure an average spectral index of $\\alpha = 0.15\\pm0.07$ (where $S\\propto\\nu^{-\\alpha}$) between 1.4 and 17.9\\,GHz for the sample. ",
        "introduction": "Recent observations of Galactic targets (Finkbeiner et~al. 2002; 2004; Casassus et~al. 2004;2006; Watson et~al. 2005; Scaife et~al. 2006; Dickinson et~al. 2007) have provided some evidence for the anomalous microwave emission commonly ascribed to spinning dust (Drain \\& Lazarian 1998a,b). This emission was first seen as a large scale phenomenon in CMB observations and represented a problem as it emits in the frequency range ~ 10--60\\,GHz (Kogut et~al. 1996; Leitch et~al. 1997; de Oliviera-Costa et~al. 2002; 2004) close to the minimum of the combined synchrotron, Bremsstrahlung and thermal dust emissions at 70\\,GHz. Indeed the models of Draine \\& Lazarian predict a spectrum for spinning dust which is strongly peaked between 20 and 40\\,GHz. Arising as a consequence of rapidly rotating small dust grains the emission has been suggested to occur in a number of distinct astronomical objects and to be highly correlated with thermal dust emission; this has been supported by some of the pointed observations referred to earlier and by recent evidence of diffuse emission at medium to high galactic latitudes from correlations made with WMAP data (Davies et~al. 2006).  Previous observations of {\\sc Hii} regions in the microwave region of the spectrum have shown evidence both for (Watson et~al. 2005; Dickinson et~al. 2007) and against (Dickinson et~al. 2006; Scaife et~al. 2007) the presence of a spinning dust component in the emission of these objects. Since the behaviour of {\\sc Hii} regions at radio frequencies is relatively well understood they provide an excellent testing ground for examining this phenomenon. At frequencies below $\\sim$100\\,GHz {\\sc Hii} regions are expected to be dominated by free-free emission, or thermal Bremsstrahlung. This mechanism progresses from the optically thick regime to the optically thin at around 1\\,GHz, and possesses a characteristically shallow spectrum ($\\alpha=0.1$; where $S\\propto\\nu^{-\\alpha}$) at frequencies above this turn over.  Here we present observations of sixteen Galactic {\\sc Hii} regions selected from the VLA survey of optically visible Galactic {\\sc Hii} regions (Fich 1993). Using spectral data from the AMI in conjunction with measurements from the literature we model the spectrum of free-free radiation and compare it with our own measurements in order to place limits on possible excess emission at 15\\,GHz, which may arise from spinning dust. We use the correlated FIR (100, 60, 25 and 12\\,$\\mu$m) emission to further constrain this excess. ",
        "conclusions": "The observations of fifteen bright {\\sc Hii} regions and one planetary nebula reported here show no evidence for anomalous emission due to spinning dust at six frequencies between 14 and 18\\,GHz. This result confirms the dominance of free-free emission in these objects with a spectral index consistent with the canonical value of 0.1. No significant evidence for spinning dust emission has been found."
    },
    "0711/0711.1495.txt": {
        "abstract": "Following the outburst of the unusual variable star V838 Monocerotis in 2002, a spectacular light echo appeared. A light echo provides the possibility of direct geometric distance determination, because it should contain a ring of highly linearly polarized light at a linear radius of $ct$, where $t$ is the time since the outburst. We present imaging polarimetry of the V838~Mon light echo, obtained in 2002 and 2005 with the Advanced Camera for Surveys onboard the {\\it Hubble Space Telescope}, which confirms the presence of the highly polarized ring. Based on detailed modeling that takes into account the outburst light curve, the paraboloidal echo geometry, and the physics of dust scattering and polarization, we find a distance of $6.1\\pm0.6$~kpc. The error is dominated by the systematic uncertainty in the scattering angle of maximum linear polarization, taken to be $\\theta_{\\rm max}=90^\\circ \\pm \\, 5^\\circ$. The polarimetric distance agrees remarkably well with a distance of $6.2 \\pm 1.5$~kpc obtained from the entirely independent method of main-sequence fitting to a sparse star cluster associated with V838~Mon. At this distance, V838~Mon at maximum light had $M_V\\simeq-9.8$, making it temporarily one of the most luminous stars in the Local Group. Our validation of the polarimetric method offers promise for measurement of extragalactic distances using supernova light echoes. ",
        "introduction": "In early 2002, the previously unknown variable star V838~Monocerotis underwent an unusual outburst. Its eruption light curve showed an initial brightening to 10th magnitude, followed about a month later by a dramatic rise to a sharp blue peak at 6th magnitude. Over the ensuing nearly three months there were several further episodes of re-brightening, after which it faded at optical wavelengths to near its pre-outburst brightness. Unlike a classical nova, V838~Mon became progressively redder, eventually becoming the coolest known luminous star (e.g., Bond et al.\\ 2003; Evans et al.\\ 2003; and papers in the proceedings of a conference on the object, edited by Corradi \\& Munari 2007).  Early in the outburst, a light echo surrounding V838~Mon was discovered by Henden, Munari, \\& Schwartz (2002). The echo phenomenon occurs when light from an eruptive variable is scattered by dust in its vicinity, reaching the Earth at progressively later times as the wave of illumination propagates outward.  Galactic light echoes are extremely rare.  The only other known example of extent similar to that of V838~Mon was the echo produced by Nova GK Persei 1901 (Kapteyn 1902; Perrine 1902; Ritchey 1902). Following early misunderstandings, light-echo geometry was properly described by Couderc (1939), and more recent discussions are given by many authors, including Chevalier (1986), Felten (1991), Sparks (1994), Sugerman (2003), and references therein. For an instantaneous light flash, the echo surface at a given time is an ellipsoid with the source at one focus and the Earth at the other. For astronomical distances, this ellipsoid is well approximated by the paraboloid given by $z=x^2/2ct-ct/2$, where $x$ is the projected distance from the star in the plane of the sky, $z$ is the distance from this plane along the line of sight (LOS) toward the Earth, $c$ is the speed of light, and $t$ is the time since the outburst.  The appearance of a light echo is governed by the time-dependent brightness of the illuminating source, the density and scattering properties of the  dust, and the distance to the star. Sparks (1994, 1996) showed that {\\it polarimetric\\/} observations of a light echo may be used to derive a {\\it purely geometric\\/} distance to the source. While the intensity distribution of an echo can be highly complex since it depends largely on the density of the scattering medium, the linear-polarization distribution should be simple, smooth, and dependent only on projected linear radius from the star and the time since the outburst. This is true because the degree of linear polarization depends almost exclusively on the scattering angle, which in turn is a monotonically decreasing function of projected radius due to the parabolic geometry of a light echo.  Furthermore, linear polarization due to scattering typically maximizes at a scattering angle of $90^\\circ$, i.e., for dust located in the plane of the sky. As the above equation shows, the  $90^\\circ$ scattering, occurring for dust at $z=0$, takes place at a linear distance of $x=ct$ from the star. Hence there should be a ring of linearly polarized light surrounding the source, whose radius expands linearly with time, and whose angular radius yields the distance geometrically. See Sparks (1994) for a detailed description of this novel method for geometric determination of astronomical distances.  The spectacular light echo of V838~Mon thus offers a unique opportunity to test this technique, and also to contribute to a physical understanding of this extraordinary object by determining its distance and luminosity. ",
        "conclusions": "We have applied the method of light-echo polarimetric imaging described by Sparks (1994) to the derivation of a geometric distance to V838~Mon. Following a careful reduction of polarimetric images of the light echo, obtained with the Advanced Camera for Surveys onboard the {\\it Hubble Space Telescope}, we confirmed the presence of an apparently expanding ring of highly polarized light surrounding V838~Mon. Setting the radius of this ring to a linear size of $ct$ yielded initial distance estimates of 5.4--6.2~kpc. We then developed modeling and fitting procedures of increasing complexity, which yielded consistent results.  Our final best value for the distance is $6.1\\pm0.6$~kpc, where the error is dominated by the systematic uncertainty in the scattering angle of maximum linear polarization. This distance agrees extremely well with a value of $6.2\\pm1.2$~kpc, obtained by an entirely independent method of main-sequence fitting, applied to a sparse stellar cluster to which V838~Mon belongs.  At this distance, the outburst of V838~Mon in early 2002 is shown to have been extremely luminous---it was brighter than all but the most luminous classical novae, and temporarily one of the brightest stars in the entire Local Group. The mechanism that produced such an outburst remains uncertain, but scenarios involving stellar mergers may be the most plausible.  Our results strongly verify the polarimetric method for determining geometric distances. The technique shows great promise for application to distance determinations for extragalactic supernovae."
    },
    "0711/0711.4646_arXiv.txt": {
        "abstract": "Sterile neutrino with mass of several keV can be the cosmological dark matter, can explain the observed velocities of pulsars, and can play an important role in the formation of the first stars.   We describe the production of sterile neutrinos in a model with an extended Higgs sector, in which the Majorana mass term is generated by the vacuum expectation value of a gauge-singlet Higgs boson.  In this model the relic abundance of sterile neutrinos does not necessarily depend on their mixing angles, the free-streaming length can be much smaller than in the case of warm dark matter produced by neutrino oscillations, and, therefore, some of the previously quoted bounds do not apply.  The presence of the gauge singlet in the Higgs sector has important implications for the electroweak phase transition, baryogenesis, and the upcoming experiments at the Large Hadron Collider and a Linear Collider. ",
        "introduction": "The discovery of the neutrino masses can  easily be incorporated into the Standard Model (SM) by adding two or more SU(3)$\\times$SU(2)$\\times$U(1) singlet fermions, often called right-handed neutrinos, which are allowed to have the Yukawa couplings to the Higgs boson and the standard,  left-handed neutrinos.  The Yukawa couplings generate the Dirac mass terms for the neutrinos after the spontaneous symmetry breaking.   In addition, the singlet fermions can have some Majorana masses. The interplay between the Dirac mass and the Majorana mass, known as the seesaw mechanism~\\cite{seesaw}, can accommodate the observed neutrino masses for a variety of Majorana masses.   If the Majorana mass terms are large, the particles associated with the singlet fields are very heavy. However, if one or more Majorana masses are below the electroweak scale, the so called sterile neutrinos appear among the low-energy degrees of freedom.  These new particles can be the cosmological dark matter~\\cite{dw,production_oscillations,shi_fuller,nuMSM,shaposhnikov_tkachev, Kusenko:2006rh,Kadota:2007mv}, their production in a supernova can explain the pulsar kicks~\\cite{pulsars} and can affect the supernova explosion in a variety of ways~\\cite{supernova_misc}; the same particles can play an important role in the formation of the first stars~\\cite{reion} and some other astrophysical phenomena~\\cite{Biermann:2007ap}.  The properties of the sterile dark matter, and, in particular, how warm or cold it is for a given mass, depend on the production mechanism.  One mechanism, which generates a population of relic sterile neutrinos at the sub-GeV temperature was proposed by Dodelson and Widrow (DW)~\\cite{dw}.  If the lepton asymmetry is negligible, this scenario appears to be in conflict with a combination of the X-ray bounds~\\cite{x-rays} and the Lyman-$\\alpha$ bounds~\\cite{viel,silk}.  This conclusion is based on the state-of-the-art calculations of the sterile neutrino production in neutrino oscillations~\\cite{production_oscillations}.  It is possible to evade this constraint if the lepton asymmetry of the universe is greater than $O (10^{-3})$~\\cite{shi_fuller}.  On the other hand, some astronomical observations~\\cite{dSphs,cdm-wdm} point to a non-negligible free-streaming length for dark matter, which favors warm dark matter.  Moreover, warm dark matter can cause filamentary structure on small scales~\\cite{Gao:2007yk}, in contrast with cold dark matter.  It is also possible that the sterile neutrinos make up only a fraction of dark matter~\\cite{silk}, in which case they can still be responsible for the observed velocities of pulsars~\\cite{pulsars,Kusenko:2006rh}.  Dodelson--Widrow mechanism is not the only mechanism by which sterile dark matter could be produced.  The relic population of sterile neutrinos could be generated in a variety of ways, for example, from their coupling to the inflaton~\\cite{shaposhnikov_tkachev}, the electroweak-singlet Higgs boson~\\cite{Kusenko:2006rh}, or the radion~\\cite{Kadota:2007mv}. Whatever the production history of sterile neutrinos might be at the high temperature, there is always some additional amount produced in neutrino oscillations at some sub-GeV temperatures~\\cite{dw,production_oscillations}. The two components can have very different momentum distributions. Therefore, generically this form of dark matter is a mixed two-component dark matter, which can have some very interesting observable consequences~\\cite{boyan_mixed}.  In this paper we concentrate on the possibility that the relatively light Majorana mass could arise via the Higgs mechanism in a model with an SU(2)$\\times$U(1)-singlet Higgs boson coupled to the Standard Model Higgs boson~\\cite{Chikashige:1980ht}, and that the sterile neutrinos could be produced from the Higgs decays at a temperature as high as 100~GeV~\\cite{Kusenko:2006rh}.  We will explore various scenarios for such production and the implications for the electroweak phase transition.  In particular, we will address the cooling and the red-shifting of dark matter, which have important implications for dark matter profiles in halos~\\cite{cdm-wdm,dSphs}, the small-scale structure inferred from Lyman-$\\alpha$ observations~\\cite{viel}, and the velocity dispersion in dwarf spheroids~\\cite{dSphs}. ",
        "conclusions": "The inclusion of singlet fermions (right-handed neutrinos) in the Standard Model is the usual way to generate the observed neutrino masses~\\cite{seesaw}. In contrast with many other models, we assume that both Dirac and Majorana neutrino masses are generated via the Higgs mechanism~\\cite{Chikashige:1980ht}. The immediate advantage of this model is the possibility to produce dark matter, in the form of sterile neutrinos, which is cold enough to satisfy the bounds on the small-scale structure and the bounds from X-ray observations, while explaining the pulsar kicks at the same time~\\cite{Kusenko:2006rh}. The same sterile neutrinos  can play an important role in the formation of the first stars~\\cite{reion}. We have considered different ways in which the dark-matter sterile neutrinos can be produced from the Higgs decays in the early universe.  If the production from the Higgs decays dominates over the production by neutrino oscillations, the resulting dark matter population is colder than in the Dodelson-Widrow case for the same mass. The Higgs structure of the model has important implications for the collider physics an can be probed at the at the Large Hadron Collider and a Linear Collider~\\cite{Profumo:2007wc,singlet_higgs_LHC}.  This work was supported in part by the DOE grant DE-FG03-91ER40662 and by the NASA ATP grant NAG~5-13399.  A.K. appreciates hospitality of the Aspen Center for Physics."
    },
    "0711/0711.4470_arXiv.txt": {
        "abstract": "In this paper we report on the real nature of the star HD\\,145792, classified as He weak in {\\it ``The General Catalogue of Ap and Am stars''}. By means of FEROS@ESO1.52m high resolution spectroscopic data, we refined the atmospheric parameters of the star, obtaining: T$_{\\rm eff}$\\,=\\,14400~$\\pm$~400~K, $\\log g$\\,=\\,4.06~$\\pm$~0.08 and $\\xi$\\,=\\,0 $^{+0.6}$ km s$^{-1}$. These values resulted always lower than those derived by different authors with pure photometric approaches.  Using our values we undertook an abundance analysis with the aim to derive, for the first time, the chemical pattern of the star's atmosphere. For metals a pure LTE synthesis (ATLAS9 and SYNTHE) has been used, while for helium a hybrid approach has been preferred (ATLAS9 and SYNSPEC). The principal result of our study is that HD\\,145792 belongs to He strong class contrary to the previous classification. Moreover, helium seems to be vertically stratified in the atmosphere, decreasing toward deepest layers.  For what that concerns metals abundances, we found the following: overabundance of oxygen, neon, silicon, phosphorus, sulfur and calcium; carbon, nitrogen, magnesium, aluminum, titanium, chromium and nickel are normal, being the discrepancies from the solar values within the experimental errors; iron resulted to be slightly underabundant. ",
        "introduction": "HD\\,145792 (=HR\\,6042\\,=V\\,1051~Sco) is a chemically peculiar star belonging to the Sco-Cen OB association (\\citet{morris61}; \\citet{eggen98}) and classified as Helium weak star in the {\\it General Catalogue of Ap and Am stars} by \\citet{renson91}. Hipparcos photometry has been analyzed by \\citet{paunzen98}, they discovered light variability with a period of 0.84780 days. A few magnetic field measurements are present in the catalogue by \\citet{bychkov03} with an average value of H$_{\\rm eff}$\\,=\\,285~$\\pm$~190 gauss.  Regarding atmospheric parameters, there are in the literature several inconsistent determinations of $T_{\\rm eff}$ and $\\log g$. For instance, \\citet{nissen74} found $T_{\\rm eff}$\\,=\\,16000~K and $\\log g$\\,=\\,4.34, \\citet{geus89} found $T_{\\rm eff}$\\,=\\,15135~K and $\\log g$\\,=\\,4.38 and \\citet{gulati89} found effective temperatures ranging between 18450~K and 18920~K. All these studies are based on photometric approach. No previous abundances determination have been reported in the literature.  The principal aim of this study was to determine for the first time the chemical abundances in the HD\\,145792's atmosphere, with particular attention to helium behavior. With this goal in mind, we obtained high resolution spectrum of the star with FEROS@ESO1.52~m telescope. Atmospheric parameters like $T_{\\rm eff}$, $\\log g$\\ and microturbulence have been derived by us with spectroscopic methods. In the following sections we present our results, giving particular emphasis to the actual peculiarity of the star: according to our results HD\\,145792 is a Helium strong star contrarily to the previous classification reported in \\citet{renson91}.  \\begin{figure*} \\includegraphics[width=12cm]{balmer.ps} \\caption[]{Comparison between the observed and computed hydrogen line profiles. From top to bottom: Balmer series from H$_{(n=11 \\rightarrow 2)}$ to H$_{\\alpha}$ and Paschen series from P$_{(n=17 \\rightarrow 3)}$ to P$_{(n=10 \\rightarrow 3)}$. The readers should note the strong telluric lines contamination around this line.} \\label{balmer} \\end{figure*} ",
        "conclusions": "In this paper we report on the first abundances analysis ever presented in the literature on the star HD\\,145792. Principal results of our paper can be summarized in this two important points:  \\begin{itemize}  \\item the object resulted to be a He strong star, contrarily to the classification reported in \\citet{renson91}. The erroneous classification could be linked to effective temperatures and gravities reported in the literature that, as we discussed in Sect.~\\ref{atmos}, are always greater than ours ($T_{\\rm eff}$\\,=\\,14400~$\\pm$~400~K and $\\log g$\\,=\\,4.06~$\\pm$~0.08).  \\item helium seems to be not homogeneously distributed along the atmosphere. It shows approximatively a constant and almost solar abundance from the inner layers up to $\\log \\tau_{5000}$\\,=\\,-0.5 and a moderate increase toward the outer atmospheric layers (see Fig.~\\ref{helium_str})  \\end{itemize}  For what that concern the chemical pattern of metals, we found, with respect the solar case, a strong overabundance of neon; a slight overabundance of oxygen, silicon, phosphorus, sulfur and calcium; iron appears to be slightly underabundant; the others elements are normal."
    },
    "0711/0711.0065_arXiv.txt": {
        "abstract": "When thin accretion disks around black holes are perturbed, the main restoring force is gravity.  If gas pressure, magnetic stresses, and radiation pressure are neglected, the disk remains thin as long as orbits do not intersect.  Intersections would result in pressure forces which limit the growth of perturbations.  We find that a discrete set of perturbations is possible for which orbits remain non-intersecting for arbitrarily long times.  These modes define a discrete set of frequencies.  We classify all long-lived perturbations for arbitrary potentials and show how their mode frequencies are related to pattern speeds computed from the azimuthal and epicyclic frequencies.  We show that modes are concentrated near radii where the pattern speed has vanishing radial derivative.  We explore these modes around Kerr black holes as a possible explanation for the high-frequency quasi-periodic oscillations of black hole binaries such as GRO J1655-40.  The long-lived modes are shown to coincide with diskoseismic waves in the limit of small sound speed.  While the waves have long lifetime, they have the wrong frequencies to explain the pairs of high-frequency quasi-periodic oscillations observed in black hole binaries. ",
        "introduction": "X-ray emission from accretion disks around compact objects often shows high-frequency flickering in the form of quasi-periodic oscillations (QPO).  Several galactic binaries containing black holes show pairs of QPO of fixed, relatively high frequencies having ratios very close to 3:2 \\citep{Remillard2006}.  These include GRO J1655-40 with 300 and 450 Hz QPO \\citep{Strohmayer2001}; XTE J1550-564 with 184 and 276 Hz \\citep{Miller2001}; GRS 1915+105 with 113 and 168 Hz \\citep{Remillard2004}; and H1743-322 with 165 and 240 Hz \\citep{Homan2005}.  The origin of these quasi-periodic oscillations is still unclear. One possibility is that orbiting blobs or hot spots follow non-circular orbits whose epicyclic frequencies are related to the observed frequencies (Stella, Vietri \\& Morsink 1999).  This model has been investigated in detail by \\cite{Schnittman2004} and \\cite{Schnittman2005}, who use a ray-tracing code to model the power spectrum of light curve variations.  The model has two serious shortcomings.  First, special orbits must be chosen by hand (e.g., the radii of the nearly circular orbits) to produce the correct frequencies, so this model lacks predictive power.  Moreover, the blobs can persist only a few rotation periods before being sheared apart by differential rotation.  Another model that can reproduce a 3:2 frequency ratio is the relativistic torus model of \\cite{Rezzolla2003}.  To produce the 3:2 ratio, the model requires a black hole spin close to maximal. The actual value of frequency depends on the radial extent of the torus. This is a free parameter of the model, which therefore also lacks predictive power.  The fixed frequencies and 3:2 ratio of these QPO pairs suggests that they may represent nonlinearly driven resonant frequencies in the curved spacetime around the black hole (Abramowicz \\& Klu{\\'z}niak 2001; Rebusco 2004; Lee, Abramowicz \\& Klu{\\'z}niak 2004).  This idea has been studied using a mathematical model for parametric resonance but has not yet been shown to arise automatically from nonlinear dynamics in general relativity.  Another possible explanation for the fixed frequencies of high-frequency QPO is diskoseismic waves similar to helioseismic $g$-modes \\citep{Wagoner2001}.  This model has the advantage of predicting a discrete set of frequencies.  However, a frequency ratio of 3:2 only obtains for a few particular values of the black hole spin that are unlikely to obtain for all observed systems. Thus, similarly to the hot spot model, this model requires tuning of parameters.  The models can be classified as to whether the oscillating structures are particle-like (blobs, with or without resonant interactions) or wave-like (oscillating tori, diskoseismic waves). Blob models are in many respects the simplest, as the only important physics is gravity.  It is natural to ask whether a similarly simple description is possible for waves.  In this paper we show that waves in an accretion disk can also be described using simple kinematics in curved spacetime. We investigate nearly kinematic density waves, which are waves of compression and rarefaction of particles freely orbiting in a disk. These density waves are very similar to the density waves which underlie the Lin-Shu hypothesis of galaxy spiral structure, apart from several important differences. Our model neglects self-gravity, pressure, magnetic stresses, and all forces except the gravity of the central compact object. A kinematic model of perturbations in thin accretion disks around Kerr black holes was considered in \\cite{Amin2006}. However, they allow for fluid streamlines to intersect forming shocks, as we will describe below.  Early models for galactic spiral structure (Lindblad 1957, Kalnajs 1973) failed to predict sufficiently long-lived patterns because differential rotation relatively quickly winds up material arms; moreover, these models neglected self-gravity which is invalid for galactic disks (Toomre 1977; Binney \\& Tremaine 1987).  However, accretion disks around black holes contain relatively much less mass than their central ``bulge'', so self-gravity is much less important.  If other forces are unimportant --- as might be expected for a thin disk --- then the remaining challenge is to prevent the waves from becoming tightly wound in a short time.  This is exactly the problem considered in the galactic context by \\cite{Kalnajs1973}.  Our model is a generalization of his kinematic density waves to arbitrary potentials and to the strongly curved spacetime around a Kerr black hole.  We choose to investigate a thin disk, since it is analytically tractable. In this case pressure forces are not important, otherwise the disk would not remain very thin.  For consistency, shocks must not form in the gas.  Thus the perturbations must evolve so that fluid streamlines avoid intersections.  The boundary between regions where streamlines cross and where they do not is a caustic surface. Our model attempts to find long-lived kinematic density waves free from caustics. ",
        "conclusions": "Starting with a simple kinematical consideration --- fluid streamlines should not intersect --- we were able to classify all long-lived linear perturbations in pressureless thin accretion disks in generic potentials including the Kerr metric.  The result is a set of patterns (planar spiral arms, vertical warps, and linear combinations) whose mode functions are concentrated at discrete radii satisfying a simple condition: the pattern frequency for non--self-intersecting orbits must be approximately constant over a range of radii, i.e.\\ its derivative with respect to radius must vanish. This condition automatically picks out discrete set of frequencies for modes with a long lifetime against winding up.  The lifetimes are $O(r/\\delta r)^2$ oscillation periods where $\\delta r$ is the radial extent of the mode.  Although our treatment neglected pressure, the results are in excellent agreement with the diskoseismic $g$-modes investigated by \\cite{Perez1997}, when the sound speed is much less than the orbital speed.  Our treatment, based on the epicyclic approximation, is much simpler than the solution of the relativistic fluid equations.  One of our goals was to investigate the possibility of explaining the $3:2$ ratio in high-frequency quasi-periodic oscillations observed in accreting black hole binaries. Our model gives several mode pairs with a frequency ratio close to 1.5 for a wide range of black hole spin. This result is a direct consequence of the discrete radii at which the spiral patterns exist.  However, we cannot recover naturally the correct frequencies unless the black hole mass estimates are seriously in error.  Planar spiral density waves have frequencies too high, while the beat frequency between a blob and these waves is too low, compared with the observed frequencies. Moreover, the best matches follow for modes with $|m|=2$ and $|m|=3$, with no explanation for the failure to observe the mode with $|m|=1$.  Because our analysis classified all long-lived perturbations of pressureless disks, the failure to reproduce the observed high-frequency quasi-periodic oscillations implies that pressure or other non-gravitational forces (e.g.\\ magnetic stresses or radiation pressure) must be important in producing the QPOs.  Based on the diskoseismic results, we would expect that small non-gravitational forces do not significantly change the mode frequencies.  Large non-gravitational forces would thicken the disk and might destroy the coherence and longevity of modes, although this remains to be seen. Even so, some mechanism is needed to pick out discrete frequencies of the high-frequency QPOs.  These results make more attractive the possibility that nonlinear coupling between modes, e.g.\\ parametric resonance, is responsible for selecting the frequency ratios (Abramowicz \\& Klu{\\'z}niak 2001; Rebusco 2004).  One approach to investigating this process is to carry the perturbation methods of the current paper to one higher order, to include a coupling between radial and vertical oscillations."
    },
    "0711/0711.0779_arXiv.txt": {
        "abstract": "There has been growing interest in the possibility of testing more precisely the assumption of statistical isotropy of primordial density perturbations.  If it is to be tested with galaxy surveys at distance scales $\\alt 10~{\\rm Mpc}$, then nonlinear evolution of anisotropic power must be understood.  To this end, we calculate the angular dependence of the power spectrum to third order in perturbation theory for a primordial power spectrum with a quadrupole dependence on the wavevector direction.  Our results suggest that primordial power anisotropies will be suppressed by $\\alt 7$\\% in the quasilinear regime.  We also show that the skewness in the statistically anisotropic theory differs by no more than 1\\% from that in the isotropic theory. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.4047_arXiv.txt": {
        "abstract": "% I review the evidence for the importance of feedback from massive stars at small and large scales. The feedback mechanisms include accretion luminosity, ionizing radiation, collimated outflows, and stellar winds. The good news is that feedback doesn't entirely prevent the formation of massive stars, while the bad news is that we don't know what does limit their masses. Feedback from massive stars also influences their surroundings.  I argue that this does not produce a triggering efficiency above unity, nor does it prevent lots of prompt star formation in GMCs, though it may preserve massive remnants of the clouds for many dynamical times. ",
        "introduction": "Massive stars influence their surroundings through a wide variety of feedback mechanisms.  Their ionizing radiation \\citep*[e.g.,][]{vacca96} and line-driven winds \\citep*{cak} have been known for many years, while only more recently have collimated jets and wide-angle outflows been observed from young, massive stars \\citep[e.g.,][]{bs05}.  The initial mass function for stars above a solar mass or so follows a power law with slope -2.35 \\citep{salpeter} up to an upper cutoff mass of $m_{up} \\sim 100$~M$_{\\odot}$ \\citep{f05}. The two most favored ideas currently for what determines the slope of the IMF are turbulent fragmentation of the parent cloud \\citep{p07} or competitive accretion \\citep{bonnell}. What determines the value of the cutoff mass $m_{up}$ remains unknown.  Limitation of accretion by increasingly strong feedback in increasingly massive stars remains a viable candidate mechanism \\citep{zy07}.  In fact, the limitation of accretion by radiative feedback had long been a severe problem in theories of massive star formation, because in models of spherically symmetric accretion, radiation pressure on standard interstellar dust prevents accretion on to stars with mass $M > 30$~M$_{\\odot}$ \\citep{wolfire-cassinelli87}.  However, recent models have called this conclusion into question because of better treatment of dust properties, non-spherical accretion, and the high accretion rates expected in high pressure cores \\citep[e.g.,][]{mt03} such as those that have been identified by the {\\em Midcourse Space Experiment} and {\\em Spitzer Space Telescope} \\citep{shirley03}.  \\citet{ys02} performed two-dimensional, nested-grid, numerical models of the collapse of slowly rotating, unmagnetized, molecular cores, using a wavelength-dependent treatment of dust grain properties for a mixture of small amorphous carbon grains, astrophysical silicates, and, in colder regions, water and ammonia ice coatings.  They found that a combination of decreased opacity, disk accretion, and focusing of radiative losses through the polar cavity (the ``flashlight effect'' first discussed by \\citealt{yb99}) allowed formation of 35--45~M$_{\\odot}$ stars.  \\citet*{kmk05} used computations of radiative transfer to demonstrate that the flashlight effect can be further enhanced if a collimated gas outflow clears a cavity in the envelope.  The efficacy of radiation in preventing accretion is further reduced by two other effects.  First, \\citet*{kkm05} showed that in three dimensions, the radiation-driven bubbles that form when the radiation pressure begins to repel gas are subject to Rayleigh-Taylor instability.  Essentially, the radiation acts as a light fluid that is trying to support the heavy gas.  As a result, the bubbles tend to collapse, allowing additional accretion.  Second, in a magnetized core, the photon bubble instability \\citep{a92} can also act to allow radiation to escape \\citep*{tqy07}.  This occurs when compressive MHD waves are amplified by radiation pressure, forming low-density regions into which the radiation streams.  Although radiation does not prevent accretion entirely, it does heat the surrounding core once a protostellar object has formed in the center. The luminosity from accretion onto even a sub-solar mass protostar can heat the surrounding envelope enough to prevent fragmentation during the collapse of a centrally condensed core \\citep{k06}.  Adaptive mesh refinement simulations including radiative transfer using a temperature dependent opacity show that fragmentation is strongly suppressed when compared to models with an isothermal equation of state \\citep*{kkm07}.  At the time of publication, these models had only reached 10~M$_{\\odot}$, but more recent conference reports have not yet given an upper mass limit.  Although this work calls into question the competitive accretion scenario \\citep*{kmk05b}, it remains to be seen whether the two scenarios converge to a common description of collapse when beginning from a turbulent cloud rather than a centrally condensed core (see Clark or Bonnell in this volume).  The good news here is that we seem to understand why feedback does not entirely prevent massive star formation.  The bad news is that we still don't know if it does or does not ultimately determine the upper mass limit of around 100--150~M$_{\\odot}$ \\citep[e.g.,][]{f05}. Several alternate explanations can be imagined.  For example, if fragmentation cannot be prevented for cores above the upper mass limit (perhaps because of insufficient central radiative heating), then the upper mass limit would be determined by the largest mass reservoirs available.  Alternatively, disk fragmentation could cut off accretion around too massive stars.  Another suggestion is that massive stars grow by direct collision, so that the stellar density is the limiting factor.  If feedback does act to cut off accretion, which mechanism is dominant?  Jets and winds appear to evolve during the formation of massive stars.  \\citet{bs05} suggest an evolutionary sequence in which the outflow type evolves as the central star accretes additional mass. In this sequence, high mass protostellar objects with central stars appearing as mid-to-early B stars have strongly collimated jets, early~B to late~O stars that have already formed hypercompact H~{\\sc ii} regions start to show a line-driven wind in addition to the jet, while early~O stars that have accreted most or all of their mass and have ultracompact H~{\\sc ii} regions (UCHRs) lose the jet and retain only the wind.  This sequence could be symptomatic of disk destruction and accretion cutoff in the final stages, a hypothesis supported by the lack of observation of O~stars with collimated jets \\citep*{shepherd03,sollins+04,arce+07} or disks \\citep{cesaroni+07}. Whether the winds themselves can destroy the disks remains an open question, however.  Jets have been suggested by \\citet{nl07} as a way of supporting cluster-forming cores.  However, the relatively low-resolution numerical models they used do not yet resolve the question of whether the turbulence driven in the cluster can actually cut off accretion. Furthermore, their outflow-driven turbulence model predicts a prominent break in the kinetic energy power spectrum at a scale close to the outflow length.  Such a break has not yet been observed in a molecular cloud. For example, \\citet{om02} found no break in the Polaris Flare---admittedly not a cluster forming core.  \\citet{b07} examine the velocity distribution produced by a single Mach 5 jet and find very little supersonic material outside the head of the jet, suggesting that jets may have difficulty driving the observed supersonic turbulence in clouds.  On the other hand, \\citet{matzner07} offers an analytic treatment that argues that jets can indeed drive the observed turbulence, consistent with \\citet{nl07}.  Ionizing radiation from the growing central star offers another mechanism for cutting off accretion.  \\citet{keto02b} made the important point that the formation of an H~{\\sc ii} region around a massive star does not cut off accretion until the Str\\\"omgren radius grows to beyond the Bondi-Parker radius where the sound speed in the ionized gas exceeds the local escape velocity.  \\citet{keto07} semi-analytically predicted the evolutionary sequence for stars forming in rotating, collapsing, cores, using the similarity solution derived by \\citet{tsc84}.  A gravitationally confined, quasi-spherical, hypercompact H~{\\sc ii} region forms first, which then transforms to a bipolar UCHR as accretion increases the central star's mass and luminosity.  As the ionizing luminosity increases further, only a remnant disk remains neutral, and finally that evaporates as well.  The basic prediction of this model that hypercompact H~{\\sc ii} regions should show evidence of accreting ionized gas has been borne out by observations both by \\citet{keto02a,kw06} and by others \\citep{beltran06}.  If ionizing radiation is to cut off accretion, the final stage of accretion will be through a photoevaporating disk.  This configuration was first analytically described by \\citet{h94} in the context of solar mass star formation, while in the context of massive stars an analytic treatment of the structure of the photoevaporating wind was done by \\citet*{lugo04}.  The best numerical model to date was done by \\citet{ry97} using a two-dimensional radiation gas dynamics code that included an explicit treatment of dust scattering.  Dust scattering leads to a factor of 3--4 increase in the photoevaporation mass loss rate, which they show to be proportional to the stellar luminosity $S_*^{0.58}$, close to the analytic prediction of \\citet{h94}, though flattening somewhat at photon luminosities $S_* > 10^{47}$~s$^{-1}$. They also explicitly included line-driven stellar winds in some models, demonstrating that they were collimated by the disk and photoevaporating wind.  If UCHRs expand at the sound speed of ionized gas, $c_i \\sim 10$~km~s$^{-1}$, they should have lifetimes of roughly $10^4$~yr. Less than 1\\% of an OB star's lifetime of a few megayears should therefore be spent within an UCHR, so the same fraction of OB stars should now lie within UCHRs.  However, \\citet{wc89} surveyed UCHRs and found numbers in our Galaxy consistent with over 10\\% of OB stars being surrounded by them, or equivalently, lifetimes $>10^5$~yr.  A number of explanations have been proposed for this lifetime problem, including thermal pressure confinement in cloud cores, ram pressure confinement by infall or bow shocks, champagne flows, disk evaporation, and mass-loaded stellar winds (see \\citealt{c99}). Several of these explanations have basic problems that suggest they likely cannot explain the lifetime problem.  Confinement by thermal pressure of the surrounding molecular gas requires pressures of $P/k \\sim 10^8$--$10^9$~cm$^{-3}$~K \\citep*{drg95,gf96}. At typical molecular cloud temperatures of 10--100~K, this implies densities $n > 10^6$~cm$^{-3}$.  However, the \\citet{j1902} mass \\begin{equation} \\label{mjeans} M_J = \\left(\\frac{4 \\pi \\rho}{3}\\right)^{-1/2} \\left(\\frac{5 k T}{2 G \\mu}\\right)^{3/2} = 6.1 \\mbox{ M}_{\\odot}  \\left(\\frac{n}{10^6 \\mbox{ cm}^{-3}}\\right)^{-1/2} \\left(\\frac{T}{100 \\mbox{ K}}\\right)^{3/2}, \\end{equation} where we have assumed a mean mass per particle $\\mu = 3.87 \\times 10^{-24}$~cm$^{-3}$ appropriate for fully molecular gas with one helium atom for every ten hydrogen nuclei.  Therefore cores massive enough to form OB stars contain multiple Jeans masses and are thus very likely to be freely collapsing (e.g.\\ \\citealt{mk04}). The free-fall time \\begin{equation} \\label{tff} t_{\\rm ff} = \\left(\\frac{3 \\pi}{32 G \\rho}\\right)^{1/2} = (3.4 \\times 10^4 \\mbox{ yr}) \\left(\\frac{n}{10^6 \\mbox{ cm}^{-3}}\\right)^{-1/2}. \\end{equation} Typical lifetimes of $>10^5$~yr would thus require massive cores to last $>3 t_{\\rm ff}$ at the hypothesized densities, rather than dynamically collapsing.  Although these high pressures are indeed observed, they are unlikely to occur in objects with lifetimes long enough to solve the problem.  \\citet{x96} instead proposed a variation on this theme: confinement by turbulent rather than thermal pressure. However, turbulent motions decay quickly, with a characteristic timescale of less than a free-fall time under molecular cloud conditions \\citep*{sog98,m99}. Turbulent pressure would thus have to be continuously replenished to maintain confinement for multiple free-fall times, which would be difficult at such high densities and small scales.  Another option is ram pressure confinement of UCHRs by infall of surrounding gas.  However, this is unstable for two different physical reasons. First, the density and photon flux will follow different power laws in a gravitationally infalling region ionized from within, so they can never balance each other in stable equilibrium \\citep{y86,h94}.  Either the ionized region will expand, or the infall will smother the ionizing source. Second, the situation is Rayleigh-Taylor unstable, as this option requires the rarefied, ionized gas of the UCHR to support the infalling, dense gas.  Bow shock models \\citep{v90,m91,ah06} require high values of ram pressure $P_{ram} \\propto n v_*^2$.  At $n=10^5$~cm$^{-3}$, a velocity of $\\sim 10$~km~s$^{-1}$ is required \\citep{v90}.  A star moving at such a high velocity in a straight line would travel a parsec over the supposed UCHR lifetime of $10^5$~yr, requiring a uniform-density region of mass $>5 \\times 10^3$~M$_{\\odot}$ for confinement.  As collapse would occur on the same timescale, more mass would actually be required to solve the lifetime problem.  Expansion of an H~{\\sc ii} region in a density gradient can drive supersonic champagne flows down steep enough gradients \\citep{t79}. Two-dimensional models first studied expansion across sharp density discontinuities \\citep*{bty79}, but then examined other configurations such as freely-collapsing cloud cores \\citep*{ybt82}, clouds with power-law density gradients in spherical \\citep*{ftb90}, and cylindrical \\citep{gf96} configurations, and exponential density gradients \\citep{ah06}. Stellar wind combined with a champagne flow down a continuous gradient was also modeled by \\citet{ah06}.  This model does fit well the observed velocity structure in some cometary UCHRs such as G29.96$-0.02$. However, these models face the same timescale problem as thermal pressure confinement models: regions dense enough to explain the observations are gravitationally unstable, and collapse on short timescales.  Another class of models relies on mass-loaded stellar winds to reproduce the observed properties of UCHRs \\citep*{dwr95,rwd96,wdr96,l96}. In these models, an expanding stellar wind entrains a distribution of small, self-gravitating, pressure-confined clumps that take substantial time to evaporate. These models can reproduce many of the basic features of the observations including some line profiles \\citep{dwr95}, core-halo, shell \\citep{rwd96}, and cometary and bipolar shapes \\citep*{rwd98}, but at the cost of requiring an arbitrary distribution of pre-existing clumps that cannot be self-consistently predicted.  Finally, \\citet{m07} examines the expansion of H~{\\sc ii} regions into turbulent, self-gravitating gas.  At densities high enough for massive stars to form, the expanding shell driven by newly ionized gas quickly becomes gravitationally unstable \\citep{v88,mn93}, collapsing even more promptly than the surrounding gas.  These regions of secondary collapse in the shell may be externally ionized to form objects with the properties of UCHRs.  As the shell expands to larger sizes, new regions can form, extending the lifetime during which UCHRs remain visible well beyond the expansion time of the original H~{\\sc ii} region.  Some combination of secondary collapse and disk evaporation appears to me most likely to ultimately be the solution to the UCHR lifetime problem. ",
        "conclusions": ""
    },
    "0711/0711.3667_arXiv.txt": {
        "abstract": "New, very long (25\\arcmin), cuts of spatially resolved profiles of the \\ha\\ and \\nii\\ optical emission lines have been obtained over the face of the Helix planetary nebula, NGC 7293. These directions were chosen to supplement previous similar, though shorter, cuts as well as crossing interesting phenomena in this nebular envelope. In particular, one new cut crosses the extremes of the proposed CO J=2-1 emitting outer `torus' shown by Huggins and his co--workers to be nearly orthogonal to its inner counterpart. The second new cut crosses the  extensive outer filamentary arcs on either side of the bright nebular core. It is shown that NGC~7293 is composed of multiple  bipolar outflows along different axes. Hubble-type outflows over a dynamical timescale of 11,000 yr are shown to be occurring for all of the phenomena from the smallest He{\\sc ii} emitting core out to the largest outer filamentary structure. All must then have been ejected over a short timescale but with a range of ejection velocities. ",
        "introduction": "The Helix planetary nebula (NGC~7293) continues to attract both observational and theoretical interest simply because it is one of the closest (213 pc distant -- Harris et al. 1997) bright, evolved planetary nebulae (PNe) and hence open to investigation on a wide range of spatial scales. In both O'Dell, McCullough \\& Meixner (2004) and Meaburn et al. (2005b), and references therein, much of the previous work is summarised up to those dates. More recently, Hora et al. (2006) have made Spitzer Space Telescope infra-red observations where the knotty nature of the whole PN envelope is spectacularly apparent in their images. Ultra-violet (UV) images of NGC~7293 released in 2005 and taken with the Galaxy Evolution Explorer (GALEX) satellite reveal particularly well many of the halo features including what was thought to be the large filamentary bipolar lobe as well as a possible `jet' and the counter bow-shaped feature previously noted in \\niil\\ emission by Meaburn et al. (2005b). (The GALEX FUV (135-175 nm) image combined with the NUV (175-280 nm) one can be examined in photojournal.jpl.nasa.gov/jpeg/PLA07902.jpg). Theoretically, Dyson et al (2006) have predicted the creation of the unambiguously accelerating tails of the inner toroidal system of cometary knots as their dense core are overrun by the mildly supersonic AGB `superwind'. Garcia-Segura et al (2006) have also considered the effects on the PN envelope as the highly supersonic `fast' wind switched off for this is no longer observed in the spectrum of the central star (Cerruti--Sola \\& Perinotto 1985).  The global structure of the complex PN envelope of NGC~7293 is becoming clarified for the key to its understanding is knowledge of its large--scale kinematical behaviour. Initially, Meaburn \\& White (1982) from wide-field Fabry-Perot line profiles and imagery using the \\niil\\ and \\oiiil\\ emission lines showed that the helical appearance of the bright core of NGC~7293 is the manifestation of a bipolar shell expanding at $\\approx$ 25 \\kms\\ viewed at a 37\\degree\\ to its axis. This central bipolar structure was also shown to contain an inner \\oiiil\\ emitting spherical shell now (Meaburn et al 2005b) known to be expanding at only 12 \\kms. Healey \\& Huggins (1990) and Young et al (1999) with complete coverage of CO J=2--1 profiles over the bright, apparently helical structure refined this viewpoint and furthermore suggested that the bipolar lobes  emanated from a central, clearly identifiable, torus expanding at 29~\\kms. Incidentally, this central structure, with two diametrically opposite, elongated, expanding shells emanating from a central expanding torus will be referred to throughout this paper as `bipolar' in keeping with its previous description (see fig.11 in Meaburn et al. 2005). Overall it could equally be described as a quasi-ellipsoidal expanding structure with a dense, toroidal waist. All features of this model for the bright helical structure were consolidated by the morphological/kinematical modelling of the optical observations in Meaburn et al. (1998 \\& 2005b). The spatially resolved \\niil\\ profiles with the first Manchester Echelle Spectrometer (MES -- Meaburn et al 1984) on the Anglo-Australian 3.9--m telescope over 17\\arcmin\\ diameter east--west and north--south cuts (6 \\& 7 in Fig. 1) (Meaburn et al 1996 \\& 1998) proved decisive in this analysis.  These previous long cuts of optical line profiles have now been supplemented by even longer cuts at intermediate and strategically desirable position angles over the nebular surface. Of particular interest are the anomalous velocity features in the CO J=2--1 profiles discovered by Healey \\& Huggins (1990) and mapped completely by Young et al (1999) which suggests that there is a separate neutral torus, not seen hitherto in the imagery or line profiles at optical wavelengths but which is both outside the nebular core and expanding nearly orthogonally to it. One of the new extra-long cuts of optical line profiles, to be reported here, crosses both the `point symmetric' CO features (see fig. 3 of Young et al 1999) that lead to this orthogonal interpretation, as well as being orientated nearly down the axis of the proposed `inner' bipolar lobe. The second new extra-long cut crosses the main `helical' structure but extends continuously out to the faint halo filaments on either side of this and along the inner part of a possible jet. Only limited observations of the optical line profiles from the north-eastern halo filaments had been made previously (Walsh \\& Meaburn 1987; Meaburn et al 2005b though note the calibration correction made for the latter in the present paper).  The combination of these new observations with the previous long cuts of optical line profiles now provides unprecedented kinematical coverage of the principal features of NGC~7293. Consequently, the propensity in PNe for the ejection of multiple bipolar lobes, each perpendicular to its own expanding torus themselves at various tilt angles, is clearly revealed in NGC~7293. ",
        "conclusions": "\\subsection{The morphology of the large--scale features.}  Many of the large--scale optical phenomena are sketched in Fig. 6 as dark solid lines for the brighter features and dark dashed lines for the fainter ones. These can be seen in the deep \\ha\\ + \\niiab\\ image presented in fig. 1a-c of Meaburn et al (2005b) and spectacularly in the near UV (175--280 nm) Galaxy Evolution Explorer (GALEX) image (NASA/JPL--Caltech/SSC circulated in 2005). The faint nebulosity visible in the latter image may be a consequence of the dominant C{\\sc iii}] 1906 \\& 1909 \\AA\\ emission lines. The bright \\niil\\ emitting  central inner `ring' in Fig. 6 is on the inside edge of a corresponding CO J=2-1 emitting ring (Huggins \\& Healy 1986; Healy \\& Huggins 1990 and Young et al 1999) expanding at 29 \\kms\\ yet surrounds the system of slower moving cometary knots whose expansion is only 14.5 \\kms\\ (Meaburn et al 1996). Bipolar lobes project (Meaburn et al 1996\\& 2005b) along the common axis of these ionized/neutral tori. The lobes of this central bipolar system manifest themselves as the bright helical filaments in Fig. 6.  Young et al (1999) reveal clearly the presence of an outer CO expanding ring whose curve of radial velocities (their fig. 5) shows that it is orientated nearly orthogonally to its inner companion. The sides of this outer ring with its highest approaching and receding radial velocities are depicted in Fig. 6 by the southeastern and northwestern hashed `CO' circles respectively. It is notable that the axis of this outer CO torus has approximately the same position angle (PA = 140\\degree) as the faint extensions marked L1 and L2 in Fig. 6.  The jet--like feature and its possible bow--shaped filamentary counterpart are also very clear in the UV GALEX image. The possible jet extends from inside the image of the filaments of the SW outer arc to the edge of the outer envelope and because of the tentative identification of its nature is marked jet? in Fig. 6.  \\subsection{The kinematics of the large--scale features}  \\subsubsection{Inner bipolar structure along PA = 125\\degree.}  The detailed kinematical--morphological modelling in Meaburn et al (2005b) of the pv arrays of \\niil\\ profiles along the previous cuts 6 \\& 7 in Fig.1 showed that the bright helical filaments and inner ring in Fig. 6 are the manifestations of an inner bipolar nebula with a central expanding torus. In this model the axis of the central torus is taken to be along PA $\\approx$ 125\\degree, orientated at 37\\degree\\ to the sight line  and expanding at 14.5\\kms\\ i.e. the expansion velocity of the system of cometary knots. The lobes are expanding with Hubble--type velocities and reach expansion velocities of 24 \\kms. The southeastern and northwestern lobes are tilted towards and away from the observer respectively. The tilt of the bipolar axis and that of the central torus with respect to the sight line was determined simply by measuring the apparent dimensions of the ring (see Fig. 6) in optical images and by assuming the torus is circular.  This interpretation is consolidated in detail by the new profiles over this core along cuts 4 \\& 5 in Figs. 2 \\& 3 respectively. The line profiles of the inner \\niil\\ emitting ring of the central torus is revealed particularly well in Fig. 2, along cut 4, which is aligned nearly with the bipolar axis at PA = 125\\degree. Here its emission manifests itself as the two brightness maxima at 4\\arcmin\\ and -4\\arcmin\\ from the central star at radial velocities of \\vhel\\ = -36.7 $\\pm$ 1 and -11.2 $\\pm$ 1 \\kms\\ respectively. These measurements were made by Gaussian fitting the single line profiles extracted from those parts of the pv arrays containing the maxima. The expansion velocity of the \\niil\\ emitting torus can now be refined compared with the 14.5 \\kms\\ (the expansion velocity of the system of cometary knots) used in the model in Meaburn et al (2005b). With this direct measurement this becomes 12.75 (sin 37\\degree)$^{-1}$ = 21.2 \\kms. A systematic change in expansion velocity throughout the inner neutral/ionized torus is then occurring with the system of cometary knots nearest the central star expanding at 14.5 \\kms, the larger diameter intermediate \\niil\\ part of the torus (the `ring' in Fig. 6) at 21.2 \\kms\\ and the outermost CO component of this central torus (Young et al 1999) at 29 \\kms.  The change of expansion velocity of the central \\niil\\ emitting torus from 14.5 to 21.2 \\kms\\ in the morphological/kinematical model in Meaburn et al (2005b) makes only minor cosmetic differences to the predictions of the model for comparison with the previous observations along cuts 6 \\& 7 in Fig. 1. In fact, the detailed comparison is improved quantitatively but will not be repeated here. The tilt to more positive velocities from the southeastern to the northwestern (bottom to top) is clear in Fig. 2 as predicted by the same bipolar/torus model.  \\subsubsection{Outer CO torus}  Firstly, the outer CO ring appears to have been detected at optical wavelengths for the first time in the new pv array along cut 4. The very deep, negative, greyscale image of this is presented in Fig. 4 and two sets of velocity `spikes' (A and B in Fig. 4) can be seen over the approaching and receding maxima depicted in Fig. 6. These have the same radial velocity ranges as the corresponding CO features of Young et al (1999). The higher angular resolution of the present \\niil\\ data compared with the CO maps (30\\arcsec) appears to have resolved this outer neutral/ionized torus into a double ring structure with its axis along PA = 140\\degree\\ compared with PA = 125\\degree\\ for the inner torus (Sect. 3.2.1). There is no easy way to estimate directly the tilt of the axis of this outer torus to the sight line for it does not appear on optical images and most of its extent gets confused in the CO maps as it crosses the central nebula, consequently an estimation of only a lower limit of its expansion velocity can be made i.e. $\\geq$ 27 \\kms. The possible bipolar lobes L1 and L2 in Fig. 6 (Sect. 3.1.1) associated with this outer torus would have receding and approaching radial velocities respectively. Neither have yet been measured. Also the axis of the outer CO torus and the common axis of these possibily related bipolar lobes would expected to be the same.  \\subsubsection{Outer bipolar structure along PA = 50\\degree.}  Evidence for some type of outer bipolar structure, as suggested on morphological grounds in Sect. 3.1, enveloping the inner one, is also present in this new kinematical data. As the pv array crosses from the bright helical structure to the NE outer arc in Fig. 3 along cut 5 the profiles stay on \\vhel\\ = -50 \\kms until they reach -95 \\kms\\ to then come back to \\vsys\\ = -27 \\kms\\ at the bright filamentary edge. Similar behaviour is shown in Fig. 5 for the previous pv array along slit 2. Some sort of three dimensional expansion at $\\geq$ 68 \\kms\\ must be occurring with the bright NE outer arc filaments being viewed tangentially through the edge of an expanding shell. However, this partial (?) shell cannot have a simple, radially expanding, quasi-spherical structure for no receding velocities are detected within its circumference.  The pv arrays in Fig. 3 for the southwestern end of cut 5 cover the filamentary feature that could be an inner extension of the `jet' in Fig. 6. Within this complication, radial velocities are in receding directions as the SW outer arc is crossed which is consistent with the NE and SW outer arcs in Fig. 6 being part of the same, albeit not simple, bipolar structure.  \\subsubsection{Jet and bow--shock}  Only very limited  information is in the present kinematical data concerning the nature of the possible jet and bow--shock whose morphology is discussed in Sect. 3.1. The most telling behaviour is the Hubble--type increase in radial velocity along the possible inner part of the jet as  shown along the southwestern end of the pv array in Fig. 3. The radial velocities in the knots in the `jet' change systematically from \\vhel\\ = -10 \\kms\\ at about -8\\arcmin\\ from the central star out to \\vhel\\ = 10 \\kms\\ at the bottom end of the array. Unfortunately the jet-like feature nearest the outer envelope was not covered in the present observations.  No direct kinematical information has yet been obtained over the bow--shaped filament in the northwestern quadrant (Fig. 6). However, the UV GALEX image reveals this to be the edge of extensive complex filamentary structure up to the central bright helical filaments and the outer arc may only be the edge of a three dimensional structure typical of bow-shocks. In this case it could be possible that the extreme velocity feature out to \\vhel\\ = -110 \\kms\\ at 6\\arcmin\\ from the central star along cut 5 in Fig. 3 is from this more extended structure. More observations are needed to investigate this possibility conclusively."
    },
    "0711/0711.1381_arXiv.txt": {
        "abstract": "Detection of paired kilohertz quasi-periodic oscillations (kHz QPOs) in the X-ray emission of a compact object is compelling evidence that the object is an accreting neutron star. In many neutron stars, the stellar spin rate is equal or roughly equal to $\\Delta\\nu$, the frequency separation of the QPO pair, or to $2\\Delta\\nu$. Hence, if the mechanism that produces the kilohertz QPOs is similar in all stars, measurement of $\\Delta\\nu$ can provide an estimate of the star's spin rate. The involvement of the stellar spin in producing $\\Delta\\nu$ indicates that the magnetic fields of these stars are dynamically important.  We focus here on the implications of the paired kHz QPOs recently discovered in the low-mass X-ray binary (LMXB) system Cir X-1 \\citep{cirx1}. The kHz QPOs discovered in Cir X-1 are generally similar to those seen in other stars, establishing that the compact object in the Cir X-1 system is a neutron star. However, the frequency $\\nu_u$ of its upper kHz QPO is up to a factor of three smaller than is typical, and $\\Delta\\nu$ varies by about a factor 2 (167 Hz, the largest variation so far observed). Periodic oscillations have not yet been detected from Cir X-1, so its spin rate has not yet been measured directly. The low values of $\\nu_u$ and the large variation of $\\Delta\\nu$ challenge current models of the generation of kHz QPOs. Improving our understanding of Cir~X-1 will improve our knowledge of the spin rates and magnetic fields of all neutron stars. ",
        "introduction": "\\label{sec:1}  Power density spectra of the X-ray brightness of LMXBs reveal various types of QPOs. Some two dozen of the 187 known LMXBs \\citep{lmxbs} have been observed to produce simultaneously two QPOs with frequencies $\\sim100-1300$ Hz that increase and decrease together by hundreds of Hz \\citep{pbk}. Some of the key observational features of these kHz QPOs and their connection to our current theoretical understanding can be summarized as follows \\citep{calgary-review}: \\begin{enumerate} \\item The frequency of one of the two kHz QPOs probably reflects the azimuthal frequency of gas in a streamline at a particular radius in the inner disk. The frequencies of the kHz QPOs are similar to orbital frequencies near a neutron star and vary by hundreds of Hertz on time scales as short as minutes. Such large, rapid variations are possible if the frequency is related to the orbital frequency at a radius that varies. This can be used to constrain the equation of state of neutron stars \\citep{mlp98}. \\item The star's spin appears to be involved in producing the frequency separation $\\Delta\\nu$ of the two kilohertz QPOs in a pair. This is evident in the millisecond pulsars (MSPs) XTE~J1807$-$294, where $\\Delta\\nu\\approx \\nu_s$ \\citep{linares-05}, and SAX~J1808.4$-$3658, where $\\Delta\\nu\\approx \\nu_s/2$ \\citep{Wijnands-03}. It is strongly indicated in the other kilohertz QPO sources, because in all cases where both $\\Delta\\nu$ and $\\nu_s$ have been measured, the largest value of $\\Delta\\nu$ is consistent or approximately consistent with either $\\nu_s$ or $\\nu_s/2$ \\citep{vdKlis06}. Thus, if paired kHz QPOs are detected but accretion- and nuclear-powered X-ray oscillations are not, a rough estimate of the star's spin rate can be made using $\\Delta\\nu$. \\item The spin rate may be communicated directly by the star's magnetic field or indirectly, e.g., via an anisotropic radiation pattern that rotates with the star. In eihter case the star's magnetic field must be dynamically important. \\item In order to explain the oscillation amplitudes observed, the X-rays must be coming from the surface of the neutron star, because the energy available at larger radii is far too small. A viable model should explain (a) what determines the frequency (radius) at which the principal QPO is generated; (b) why the width of a kHz QPO peak (range of orbital radii) is so small; (c) how the frequency of the principal QPO changes by factors $\\sim2-3$; and (d) how the spin frequency of the neutron star is involved. So far, the only model that may explain these key features of the kHz QPOs and their amplitudes, energy-dependence, and dependence on the mass accretion rate is the sonic-point spin-resonance model \\citep{spsr2}. \\end{enumerate}  \\begin{figure}[t] \\includegraphics[height=.25\\textheight]{figure1.ps} \\caption{Spectral energy density of variations in the X-ray brightness of Cir X-1, from an archived RXTE observation. Right: 136 and 404 Hz ``kHz'' QPOs; middle: low-frequency QPO; left: very low frequency noise \\citep{cirx1}.}\\label{twins} \\end{figure} ",
        "conclusions": ""
    },
    "0711/0711.2457_arXiv.txt": {
        "abstract": "We present an analysis of the faint M star population seen as foreground contaminants in deep extragalactic surveys. We use space-based data to separate such stars from high redshift galaxies in a publically-available dataset, and consider the photometric properties of the resulting sample in the optical and infrared. The inferred distances place these stars well beyond the scale height of the thick disk. We find strong similarities between this faint sample (reaching $i'_{AB}=25$) and the brighter disk M dwarf population studied by other authors. The optical-infrared properties of the bulk of our sources spanning 6000\\AA-4.5$\\mu$m are consistent with those 5-10 magnitudes brighter.  We also present deep spectroscopy of faint M dwarf stars reaching continuum limits of $i'_{AB}\\approx26$, and measure absorption line strengths in the CaH2 and TiO5 bands. Both photometrically and spectroscopically, our sources are consistent with metallicities as low as a tenth solar: metal-rich compared with halo stars at similar heliocentric distances. We comment on the possible MACHO identification of M stars at faint magnitudes. ",
        "introduction": "\\label{sec:intro}  The advent in recent years of extremely deep imaging surveys, designed to probe distant galaxies through multiwavelength imaging extending deep into the infrared \\citep[e.g. GOODS,][]{2004ApJ...600L..93G}, has made possible studies of extremely faint, and often extremely red, sources. In order to target high redshift galaxies, colour cuts have been developed to exclude the majority of Galactic stars - including the low metallicity dwarfs expected in the halo of our own galaxy. Nonetheless, it has become clear that an unexpected number of extremely red stars remain at faint magnitudes. The spectroscopic component of high-redshift galaxy searches have confirmed the nature of such stars, in several cases obtaining deep spectroscopy as well as photometry \\citep[e.g.][]{2004ApJ...607..704S,2005A&A...434...53V}, but the overall population has not been systematically characterised.  Low mass stars of class M or later comprise more than 80 per cent of the stellar population of our galaxy. Faint stars of late-M, L and T classes probe the tail of the stellar mass function, bridging the divide between the main sequence and low mass brown dwarfs. Such stars have unusual colours, arising from the presence of deep molecular absorption bands in their cool atmospheres.  The faintness of these stars has made their investigation challenging, particularly with increasing distance from our sun. Efforts to determine the properties of such stars in the halo and thick disk have concentrated on nearby stars with high proper motions indicative of a halo origin \\citep[e.g.][]{1999AJ....117..508G}. While deep surveys of the M star population in order to identify distant examples have been undertaken with WFPC2 on the {\\em Hubble Space Telescope} ({\\em HST}), reaching $I_{AB}$=23.2 \\citep{2001ApJ...555..393Z}, these have been limited to wavebands shortwards of 9000\\AA\\ where the coolest stars have very little flux.  Meanwhile, the more local population - late M stars at disk metallicities within 1.5\\,kpc of the sun - have been the subject of detailed study in recent years. The Sloan Digital Sky Survey \\citep[SDSS,][]{2000AJ....120.1579Y} is relatively shallow in the optical bands, reaching an $I$ band limit of $i'\\approx20$. However, the addition of a redder $z'$ band in the optical to similar depths has enabled the photometric selection of cooler, class L and T stars as sources with large flux decrements between adjacent passbands, caused by sharply-defined, strong absorption and transmission features. The $z'$-band imaging also allows the characterisation of late M stars by their $i'-z'$ colour. Such surveys have enabled templates based on thousands of local (mainly disk) dwarf stars to be constructed \\citep{2002AJ....123.3409H,2007AJ....133..531B}.  At the same time, advancements in infrared detector technology have made it increasingly easy to study these faint, cool stars at wavelengths longwards of the optical.  Their near and mid-infrared bands are dominated by molecular absorption, leading to extreme colours in infrared bands \\citep{2006ApJ...651..502P}. Again, such studies have been limited to local stars, in order to attain the maximum possible signal to noise.  However, given the lower typical metallicity of stars in the halo when compared to the galactic disc, the characteristics of local disk-dominated M star samples may well be unrepresentative of fainter sources at larger distances and the colour selection methods applied locally might be expected to detect nothing at halo distances.  In this paper we present an analysis of faint low mass stars, selected for their extreme colours between two adjacent bands in high spatial resolution, deep optical imaging. Such a colour selection is expected to select either relatively-nearby late M stars at sub-solar metallicities or near-solar metallicity early M stars at the large distances more normally associated with the Galactic halo, if such a population exists. These two alternatives cannot be distinguished on the basis of a single red colour, but can be explored using spectroscopy and photometry probing well into the infrared.  In section \\ref{sec:photometry} we present the data used in this analysis and we discuss the selection of stellar candidates in section \\ref{sec:phot-sel}. In sections \\ref{sec:halo-pop} and \\ref{sec:spitzer-properties} we discuss the optical and infrared properties of our sample, and in section \\ref{sec:cooler} we examine the photometry of cooler stars in the same fields. In section \\ref{sec:spect} we discuss existing spectroscopy of such faint stars, including a spectroscopic sample presented here for the first time, and spectroscopic abundance indicators that may hint at a typical metallicity for this population.  Finally, in section \\ref{sec:halo-pop} we discuss the interpretation of the overall properties of the faint, cool star population including their metallicity and distance distribution.  All magnitudes in this paper (optical and infrared) are quoted in the AB system \\citep{1983ApJ...266..713O}. ",
        "conclusions": "\\label{sec:conclusions}  We have investigated the properties of a faint population of M stars selected using a photometric drop method similar to that used to identify high redshift galaxies.  The M-dwarfs in this study are up to five magnitudes fainter than those considered by previous authors, and hence lie at distances up to ten times further than those investigated before.  This population extends well above the galactic plane, with little evidence for a decline in number at faint magnitudes/large distances. This is difficult to reconcile with models of known disk structure.  The M-stars selected by this photometric method show similar colours to those within a kiloparsec of the sun, both in the optical and the infrared.  Together with the evidence from line indices, this suggests that there are significant numbers of M-stars with moderate to high metallicities extending well above the plane of the galaxy.  Our sample of stars with $v-i' > 1.7$ should be complete for mid-M dwarfs of near-solar metallicity but is incomplete for low-metallicity M dwarfs many of which have bluer colours or are more distant. Determining the metallicity of individual sources is difficult without exceptionally deep spectroscopy. If one or more of our sources are at low metallicity, they may be even cooler - and hence somewhat closer - than we believe, although our spectroscopic results indicate that fewer than 6\\% of M stars are likely to fall into this category."
    },
    "0711/0711.0722_arXiv.txt": {
        "abstract": "{} { We study a peculiar object with a projected position close to the nucleus of M51. It is unusually large for a star cluster in M51 and we therefore investigate the three most likely options to explain this object: (a) a background galaxy, (b) a cluster in the disk of M51 and (c) a cluster in M51, but in front of the disk.} { {\\rm We use broad-band images of the Advanced Camera for Surveys and the Near Infrared Camera and Multi-Object Spectrometer}, both on board the \\emph{Hubble Space Telescope}, to study the properties of this object. Assuming the object is a star cluster, we fit the metallicity, age, mass and extinction using simple stellar population models. Assuming the object is a background galaxy, we estimate the extinction from the colour of the background around the object. We study the structural parameters of the object by fitting the spatial profile with analytical models.} { We find de-reddened colours of the object which are bluer than expected for a typical elliptical galaxy, and the central surface brightness is brighter than the typical surface brightness of a disc galaxy. It is therefore not likely that the object is a background galaxy. Assuming the object is a star cluster in the disc of M51, we estimate an age and mass of $0.7^{+0.1}_{-0.1}$~Gyr and $2.2^{+0.3}_{-0.3}\\times 10^{5}~\\msun$, respectively (with the extinction fixed to $\\ebv = 0.2$). Considering the large size of the object, we argue that in this scenario we observe the cluster just prior to final dissolution. If we fit for the extinction as a free parameter, a younger age is allowed and the object is not close to final dissolution. Alternatively, the object could be a star cluster in M51, but in front of the disc, with an age of $1.4^{+0.5}_{-0.2}$~Gyr, mass $M = 1.7^{+0.8}_{-0.3}\\times 10^{5}~\\msun$. Its effective radius is between $\\sim$12--25~pc. This makes the object a ``fuzzy star cluster'', raising the issue of how an object of this age would end up outside the disc.} {} ",
        "introduction": "\\label{sec:Introduction}  Recently, \\citet{scheepmaker07} studied the effective (half-light) radii of 1284 star clusters in the disc of M51. The largest cluster candidate in their sample, carrying their ID number ``212995'', is remarkable in several ways.  If the object is a cluster in the disk of M51, it is exceptionally large, with an estimated effective radius of 21.6~pc, and located at a galactocentric distance of 1~kpc in a region with many dust lanes.  A cluster with this size and galactocentric distance is not expected to survive for long in the spiral disc, where the disruptive effects of the tidal field and other external perturbations (spiral arms, GMCs) are large \\citep{gieles05, gieles06c}. If the cluster candidate is indeed in the disc, it should therefore be young and highly reddened to account for its observed colours of $\\bv = 0.75$ and $\\vi = 1.05$. Using only 2 colours, however, \\citet{scheepmaker07} could not accurately determine an age and extinction.  Observations of HII regions in the disc of M51 show that the metallicity of the gas is approximately solar to twice solar \\citep{diaz91, hill97}. If the object is in the disc of M51, we would therefore expect its metallicity to be around solar. Outside of the disc, the metallicity could be lower.  There is also the possibility that the object is a background galaxy. The observed colours of the object are similar to the colours of typical background galaxies, but this suggests that the reddening through the disc of M51 is negligible. This would be quite remarkable at a distance of 1~kpc from the nucleus, where the extinction is expected to be high, as suggested by the presence of numerous dust lanes. The hypothesis of a background galaxy therefore suggests a chance alignment, such that the object is observed through a ``hole'' in a region of the disc with many dust lanes.  In this work we present a more detailed study of the nature of the object by extending the data with archival \\hst/\\nicmos\\ observations in \\jband\\ and \\hband, and by estimating the extinction of the object. Having better constraints on the intrinsic spectral energy distribution of the object, we compare the two hypotheses, of the object being a star cluster or a background galaxy, in a quantitative way. ",
        "conclusions": "\\label{sec:Discussion and conclusions}  We have studied the photometry, extinction and structural parameters of a peculiar object in M51 using \\hst/\\acs\\ and \\hst/\\nicmos\\ data. We have considered different possible scenarios for the object, from which we conclude the following: \\begin{enumerate} \\item If we assume that the object is a background galaxy, we estimate that the extinction through the disc of M51 is at least $\\ebv = 0.2$. Correcting for this \\emph{minimum} extinction, we find intrinsic colours of $\\bvzero = 0.43$, $\\vizero = 0.68$ and $\\jhzero = 0.76$. These colours suggest that the object is not a background galaxy, because \\bvzero\\ and \\vizero\\ are $\\sim$0.6~mag bluer than expected for a typical elliptical galaxy (Fig.~\\ref{fig:colour-colour_galaxies}) and the surface brightness is $\\sim2~\\mathrm{mag}~\\mathrm{arcsec}^{-2}$ brighter than a typical disc galaxy. If the extinction would be larger than $\\ebv = 0.2$, these differences in colour and surface brightness would be even larger.  \\item The surface brightness distribution can be well approximated by typical cluster profiles (i.e.\\ EFF or King profiles) with a FWHM of $\\sim$9 pixels. A de Vaucouleurs $R^{1/4}$ law, typical for elliptical galaxies, does not provide a good fit to the inner part of the observed profile (Fig.~\\ref{fig:profile}). Because the exact concentration (or power-law) index of the object can not be determined, the effective radius of the object depends on the choice of the profile. Assuming the object is at the distance of M51, our best estimates for the effective radius are between $\\sim$12~pc and $\\sim$25~pc (Table~\\ref{tab:radius}).  \\item If we assume that the object is a star cluster in the disc of M51, we estimate an age of $0.7^{+0.1}_{-0.1}$~Gyr and a mass $M = 2.2^{+0.3}_{-0.3}\\times 10^{5}~\\msun$, assuming $\\ebv = 0.2$ and solar metallicity. Considering the large size of the object, however, we estimate that in this scenario the cluster must initially have been very massive ($>4.5\\times 10^{5} \\msun$), and that we observe the cluster just prior to final dissolution. Although not impossible, this scenario is considered to be unlikely.  \\item We can not fully exclude the possibility that our best-fitting cluster ages and extinctions are over- and underestimated, respectively, due to the lack of U-band photometry of our object. If we consider the \\emph{maximum}-fitted extinction of $\\ebv = 0.7$ (model 6 in Table~\\ref{tab:galev}), instead of the \\emph{best}-fitted extinction, the cluster has an age of only $\\sim$4~Myr and is not close to final dissolution. With our current data, however, the likelihood of this fit is very low.  \\item If we assume that the object is a star cluster well outside the disc of M51 and seen in projection against it, we estimate an age of $1.4^{+0.5}_{-0.2}$~Gyr and a mass of $M = 1.7^{+0.8}_{-0.3}\\times 10^{5}~\\msun$. This scenario naturally explains why the best-fitting extinction is low, $\\ebv = 0.1^{+0.5}_{-0.1}$, although we observe the cluster in between two dust lanes. Furthermore, this scenario implies that we are not observing the cluster in a special phase of its lifetime.  \\item We conclude that, without additional U-band data or spectroscopy, both the ``disc'' and ``projection'' scenarios are possible within the uncertainties. With the current data, however, the cluster models favour the scenario in which the object is a star cluster seen in projection against the disc. In this scenario, the cluster just qualifies as being a ``faint fuzzy'' star cluster, namely having $\\bvzero > 0.6$, $\\vizero > 1.0$ and $\\reff > 7$~pc \\citep{larsen00a}. Seen in projection against the spiral disc, the (projected) galactocentric distance of 1~kpc is not unexpected. However, with an estimated age of $1.4^{+0.5}_{-0.2}$~Gyr, the formation mechanism of such a cluster, in combination with its extended shape and its current location outside of the spiral disc, are interesting issues that leave room for further investigation. \\end{enumerate}  \\noindent \\emph{Note added in proof:} After the paper was accepted we became aware of the study of background galaxies seen though the disk of M~51 by \\citet{holwerda05b}. These authors found that the average opacity values for M~51 are the highest for any disk ($A_I = 1.5$), making it very unlikely that the object is a background galaxy."
    },
    "0711/0711.4517_arXiv.txt": {
        "abstract": "\\noindent We discuss flavor-mixing probabilities and flavor ratios of high energy astrophysical neutrinos. In the first part of this paper, we expand the neutrino flavor-fluxes in terms of the small parameters $U_{e3}$ and $\\pi/4 - \\theta_{23}$, and show that there are universal first and second order corrections. The second order term can exceed the first order term, and so should be included in any analytic study. We also investigate the probabilities and ratios after a further expansion around the tribimaximal value of $\\sin^2 \\theta_{12} = 1/3$. In the second part of the paper, we discuss implications of deviations of initial flavor ratios from the usually assumed, idealized flavor compositions for pion, muon-damped, and neutron beam sources, viz., $(\\nue : \\numu : \\nutau)=(1 : 2 : 0)$, $(0 : 1 : 0)$, and $(1 : 0 : 0)$, respectively. We show that even small deviations have significant consequences for the observed flavor ratios at Earth. If initial flavor deviations are not taken into account in analyses, then false inferences for the values in the PMNS matrix elements (angles and phase) may result. ",
        "introduction": "Introduction} There has been recently much discussion on the flavor mixing of high energy astrophysical neutrinos \\cite{LP,fluxes,GR,Michael,Pasquale,xing2,WW,xing,KT,old,WR,MO,CP,kachneu,new}. Neutrino mixing modifies the initial flavor distribution of fluxes $\\Phi_e^0 : \\Phi_\\mu^0 : \\Phi_\\tau^0$ in calculable ways. In terrestrial neutrino telescopes such as IceCube \\cite{ice} or KM3Net \\cite{km3}, one can measure neutrino flavor ratios\\cite{ratio1} and thereby obtain information on the neutrino parameters and/or the sources\\footnote {In principle, more exotic neutrino properties such as neutrino decay, Pseudo-Dirac structure, magnetic moments, interaction with dark energy, or breakdown of fundamental symmetries could also be studied \\cite{decay,Pseudo,magn,HungPaes,CPT,MinSmi}.}. There are two essential ingredients to such analyses.  One is the initial flavor composition which depends on the nature of the production process at the source.  The other is the neutrino mixing scheme, in particular the values of the parameters governing neutrino mixing. The latter plays the role of altering the flavor mix from the original due to non-trivial lepton mixing. We will discuss in this paper precision issues for both these aspects.  The current best-fit values as well as the allowed $1\\sigma$ and $3\\sigma$ ranges of the oscillation parameters are \\cite{data}: \\begin{eqnarray} \\label{eq:data} \\sin^2 \\theta_{12} &=& 0.32^{+0.02\\,, \\,0.08}_{-0.02\\,, \\,0.06} ~, \\nonumber \\\\ \\sin^2\\theta_{23} &=& 0.45^{+0.07\\,, \\,0.20}_{-0.07\\,, \\,0.13} ~,\\\\ |U_{e3}|^2 &<& 0.02~(0.04)~.\\nonumber \\end{eqnarray} The CP phase is completely unknown and has a range between zero and $2\\pi$. The current information on the mixing parameters \\cite{APS} therefore suggests that $\\theta_{13}$ and the deviation from maximal atmospheric neutrino mixing are small parameters and therefore to expand the formulae in terms of them. This is quite useful in order to obtain an analytical understanding of basically all phenomenological problems of interest. Furthermore, the deviation from the value $\\sin^2 \\theta_{12} = \\frac 13$ is also small and thereby a third small expansion parameter is introduced. An idealized description for the leptonic mixing, or Pontecorvo-Maki-Nakagawa-Sakata (PMNS) matrix $U$, which is nevertheless perfectly compatible with all experimental information is tribimaximal mixing \\cite{tbm} \\be \\label{eq:tbm} U \\simeq U_{\\rm TBM} = \\frac{1}{\\sqrt{6}} \\left( \\barrr \\D 2 & \\D \\sqrt{2}  & 0 \\\\[0.cm] \\D -1 & \\D \\sqrt{2}  & \\D -\\sqrt{3}\\\\[0.cm] \\D -1 & \\D \\sqrt{2}  & \\D \\sqrt{3} \\ea \\right)~. \\ee The transition probabilities of astrophysical neutrinos are functions of the elements of $U$, whose usual PDG parameterization is given in the Appendix. Expanding these probabilities in terms of the small parameters motivated by the observed oscillation phenomenology is one of the purposes of this paper. Previously, expansions in terms of $|U_{e3}|$ and $\\pi/4-\\theta_{23}$ have been discussed. In this case, we will show here first that in addition to the known universal first order correction \\cite{xing,WR} there is a second order universal correction which can exceed the first order one and has to be included in analytical studies. Furthermore, by expanding the probabilities in terms of the deviation from $\\sin^2 \\theta_{12} = \\frac 13$ we find extremely compact expressions.   In what concerns the initial flavor mix at the source, there are three simple and -- as we will make clear in this paper -- idealized possibilities: \\begin{itemize} \\item The most conventional one is  neutrino emission from purely hadronic processes, such as $ p + p \\rightarrow \\pi^\\mp \\rightarrow \\mu^\\mp + \\stackrel{(-)}{\\nu_\\mu} \\rightarrow e^\\mp +  \\stackrel{(-)}{\\nu_e} + \\nu_\\mu + \\overline{\\nu}_\\mu$, in which the initial neutrino flavor mix is identical to the one in atmospheric neutrinos: \\be \\label{eq:120} \\Phi_e^0 : \\Phi_\\mu^0 : \\Phi_\\tau^0 = 1 : 2 : 0~. \\ee If the dominant decay process is $p + \\gamma \\rightarrow \\pi^+ + X$ then the flavor mix is still $1 : 2 : 0$, however the initial $\\overline{\\nu}_e$ flux is absent. This can in principle be checked by taking advantage of the Glashow resonance reaction $\\nuebar+e^-\\rightarrow W^-$~\\cite{LP,GR1,GR}, which occurs at an incident $\\nuebar$~energy of $6.3$ PeV.   \\item The muons may lose energy so that the $\\stackrel{(-)}{\\nu_e}$ flux is depleted at the energies of interest. This can happen in a variety of ways \\cite{010c,010mod,010b,010ba,010a,lipari,kachneu}; e.g., muons can lose energy in strong magnetic fields, or get absorbed in matter. In these cases the initial flavor composition is simply \\be \\label{eq:010} \\Phi_e^0 : \\Phi_\\mu^0 : \\Phi_\\tau^0 = 0 : 1 : 0~ \\ee without any electron or tau neutrinos. Specific models have been discussed. In general, one expects that any pion source will have a transition to a muon-damped composition, with the transition energy depending on the source properties. Qualitatively, muon damping occurs when the muon's energy-loss length $\\sim\\frac{E}{dE/dx}$ is shorter than its decay length $(E/m_\\mu)\\, c  \\tau_0$, i.e., when $dE/dx \\gs c \\tau_0/m_\\mu$.   \\item The third case is that of sources which emit dominantly neutrons. These neutrons originate in the photo-dissociation of heavy nuclei. The decays of the neutrons give rise to an initial pure ``$\\beta$-beam'' of $\\overline{\\nu}_e$~\\cite{100a}, i.e., \\be \\label{eq:100} \\Phi_e^0 : \\Phi_\\mu^0 : \\Phi_\\tau^0 = 1 : 0 : 0~, \\ee with no $\\nu_e$ or any muon or tau neutrinos. \\end{itemize}  One other expected source of both pion-decay neutrinos and neutron-decay neutrinos, in separated energy regions~\\cite{Engel:2001hd}, is the GZK nucleonic reaction $p_{\\rm CR} + \\gamma_{\\rm CMB} \\rightarrow \\Delta^+ \\rightarrow n + \\pi^+$. The subsequent pion and muon decays produce neutrinos with energies about 20 times (i.e., $m_\\pi/m_N\\times 1/4$) below $E_{\\rm GZK}\\sim 5\\times 10^{19}$~eV, while neutron decay produces $\\nuebar$'s with energies about a thousand times (i.e., $\\beta$-decay $Q$-value/$m_n$) below $E_{\\rm GZK}$.  We will emphasize in the present work that the above three flavor mixes in Eqs.~(\\ref{eq:120}, \\ref{eq:010}, \\ref{eq:100}) are idealizations. Realistically, one should expect deviations from these simple flux compositions. These deviations should be taken into account in analyses to avoid incorrect conclusions about the violation/conservation of CP, the octant of $\\theta_{23}$, or the magnitude of $|U_{e3}|$. We shall outline the need for care with several examples.  This paper is built up as follows. In Section~\\ref{sec:formalism} we discuss an expansion of the mixing probabilities and flux ratios up to second order (i) in the small parameters $\\sin^2 \\theta_{23} - \\frac 12$ and $|U_{e3}|$, and eventually, (ii) in the small parameter $\\sin^2 \\theta_{12} - \\frac 13$. In Section \\ref{sec:uncertain} we assess the validity of the idealized and often used initial flavor compositions given in Eqs.~(\\ref{eq:120}, \\ref{eq:010}, \\ref{eq:100}). We illustrate the effects of impure neutrino mixes with various examples. We sum up and conclude in Section \\ref{sec:concl}. ",
        "conclusions": ""
    },
    "0711/0711.1332_arXiv.txt": {
        "abstract": "We propose a new concept of pupil motion sensor for astronomical adaptive optics systems and present experimental results obtained during the first laboratory validation of this concept. Pupil motion is an important issue in the case of extreme adaptive optics, high contrast systems, such as the proposed Planet Finder instruments for the ESO and Gemini 8-meter telescopes. Such high contrast imaging instruments will definitively require pupil stabilization to minimize the effect of quasi-static aberrations. The concept for pupil stabilization we propose uses the flux information from the AO system wave-front sensor to drive in closed loop a pupil tip-tilt mirror located in a focal plane. A laboratory experiment validates this concept and demonstrates its interest for high contrast imaging instrument. ",
        "introduction": "Nowadays, the field of high contrast imaging in astronomy is in great progress. Indeed, most of the large telescopes in operation are equipped with adaptive optics systems (NACO on the European Southern Observatory (ESO) Very Large Telescope (VLT) \\cite{rousset2003}, Altair on Gemini North \\cite{stoesz2004}, Keck-AO on the Keck II telescope \\cite{wizinowich2006} for example. Stellar coronagraphy is also being used in combination with some of these systems \\cite{boccaletti2004}.  Some astrophysical programmes are still out of reach with the existing instruments, due to the relatively limited performance of these systems in terms of the achieved contrast. More specifically, extra-solar planets cannot be directly detected with the current systems, except for a few recently confirmed planetary mass companions to young brown dwarfs \\cite{chauvin2005a, chauvin2005b, neuhauser2005}. But these very favorable cases, thanks to the young age and relatively low mass of the host stars not very demanding in terms of contrast requirements, are not representative of the vast majority of the envisioned targets.  The very high contrast required for direct detection of extra-solar planets around most of the target stars, is an extremely challenging goal, for which new and dedicated instruments should be developped.  In particular, such a new instrument, called SPHERE \\cite{beuzit2006, dohlen2006} is currently being developped to equip the ESO VLT by mid of 2010. SPHERE will combine eXtreme Adaptive Optics (XAO) \\cite{fusco2006}, high performance coronagraphy \\cite{boccaletti2004} and spectroscopic and polarimetric differential imaging techniques \\cite{racine1999, berton2006, schmid2005}.  For very high contrast observations, SPHERE system performance analysis has shown that the system pupil stability is one of the main limitations. The new generation of instruments will therefore need a way to very accurately control the stability of the pupil, from the telescope mirrors to the science detector, especially up to the coronagraph.  We present here our study on pupil stabilization for XAO systems, performed as part of the SPHERE phase A study. We first discuss the impact of pupil stability for extra-solar planets detection in section~\\ref{sec:planets}. In section~\\ref{sec:design} we present the principle of the pupil motion control and in section~\\ref{sec:error} we discuss the possible sources of errors in the measurement of the pupil shift. Finally, in section~\\ref{sec:experiment} we present the laboratory experiment which allowed to fully validate the proposed concept. ",
        "conclusions": ""
    },
    "0711/0711.3047_arXiv.txt": {
        "abstract": "Multi-wavelength imaging polarimetry at far-infrared wavelengths has proven to be an excellent tool for studying the physical properties of dust, molecular clouds, and magnetic fields in the interstellar medium.  Although these wavelengths are only observable from airborne or space-based platforms, no first-generation instrument for the Stratospheric Observatory for Infrared Astronomy (SOFIA) is presently designed with polarimetric capabilities.  We study several options for upgrading the High-resolution Airborne Wideband Camera (HAWC) to a sensitive FIR polarimeter.  HAWC is a $ 12 \\times 32$ pixel bolometer camera designed to cover the 53 -- 215 $\\micron$ spectral range in 4 colors, all at diffraction-limited resolution (5 -- 21 arcsec). Upgrade options include: (1) an external set of optics which modulates the polarization state of the incoming radiation before entering the cryostat window; (2) internal polarizing optics; and (3) a replacement of the current detector array with two state-of-the-art superconducting bolometer arrays, an upgrade of the HAWC camera as well as polarimeter.  We discuss a range of science studies which will be possible with these upgrades including magnetic fields in star-forming regions and galaxies and the wavelength-dependence of polarization. ",
        "introduction": "Mapping polarimetry at far-infrared/submillimeter wavelengths began on the Kuiper Airborne Observatory (KAO) more than 20 years ago. Over the course of the KAO's lifetime polarimetric instruments evolved from single pixel devices\\cite{rhh84} to larger arrays\\cite{stokes} and eventually moved to ground-based observatories\\cite{hertz2,scubapol}. Far-infrared (FIR) and submillimeter observations of dense molecular clouds, protostellar disks and envelopes, the Galactic center, and one external galaxy have shown that dust emission from these objects is polarized (at levels of $\\sim 2$ -- 10\\%) at almost every point\\cite{archivestokes,archivehertz,li06}. The alignment of dust grains with the local magnetic field\\cite{dg51} allows one to infer the direction of the magnetic field projected onto the plane of the sky. This allows studies of magnetic fields' interactions with the local interstellar medium and their role in star and galaxy formation\\cite{dasth,chuss03,mwf01,m82scuba} as well as studies of dust and cloud properties.\\cite{pspec,lgm97}  As successor to the KAO, the Stratospheric Observatory for Infrared Astronomy (SOFIA) offers a new frontier for observations in the infrared-to-submillimeter spectral regime. Unfortunately, no first-light SOFIA instruments\\cite{krabbe02} have existing capability for accurate polarimetry. Both a far-infrared (\\emph{Hale})\\cite{hale} and mid-infrared\\cite{packham} polarimeter have been proposed as future SOFIA instruments.  Here we discuss three options for upgrading the High-resolution Airborne Wideband Camera (HAWC)\\cite{hawc} to a sensitive FIR polarimeter.  HAWC is a first-light instrument for SOFIA\\@. It consists of a $12 \\times 32$ pixel bolometer camera designed to cover the 53 -- 215 $\\micron$ spectral range in 4 colors, all at diffraction-limited resolution ($5\\arcsec$ -- 21\\arcsec). A complete upgrade of the HAWC camera into a polarimeter (which we call \\emph{Hale}) will include replacing the existing 384 pixel array of silicon bolometers with two superconducting bolometer arrays.  Each such array would consist of over 5000 pixels and detect orthogonal polarization components selected by wire grids close to the detectors. A cold (4\\,K) crystal half-wave plate (HWP) upstream from the grids modulates the incoming radiation. In this way we create a sensitive, diffraction-limited, 4-color polarimeter.  Additionally, the dual-polarization measurement and the detector upgrade yield a photometer with a much larger field-of-view than the original HAWC instrument.  We consider two other options (called HAWC-pol) as ``pathfinder'' programs working towards the more complete \\emph{Hale} instrument. These options are attractive as it is feasible that they may be incorporated into HAWC soon after its first light on SOFIA\\@. In both HAWC-pol options we retain the existing $12\\times 32$ silicon bolometer array and measure only a single component of polarization. Option (1) places one or more rotating HWPs at the cold HAWC pupil stop, followed by a polarizing grid. In option (2) we place only a polarizing grid at the pupil and modulate the polarization outside the cryostat using a variable-delay polarization modulator, or polarization ``switch''. Details of the upgrade components are discussed in Sect.\\ \\ref{sect:instrument}; their locations in the HAWC cryostat are shown in Fig.~\\ref{fig:cryostat}.  The technical specifications for both \\emph{Hale} and HAWC-pol are shown in Table \\ref{tbl:specs}.  \\begin{figure}[tb] \\begin{center} \\includegraphics[width=15cm]{hale4} \\end{center} \\caption[example] {The HAWC cryostat showing the locations of several proposed polarimetric upgrades: dual-polarization detector arrays, wave-plates (HWP), and the variable-delay modulator. \\label{fig:cryostat} } \\end{figure}  {\\input{spectab}} ",
        "conclusions": ""
    },
    "0711/0711.4451_arXiv.txt": {
        "abstract": "{ Numerous investigations on the fundamental properties of low-mass stars using eclipsing binaries indicate a strong discrepancy between theory and observations that is still awaiting explanation. Current models seem to predict radii for stars less massive than the Sun that are some 10\\% smaller than observed, while their effective temperatures are some 5\\% larger. Here we discuss recent new observational data that are relevant to this issue and review the progress made in understanding the origin of the important differences with theoretical calculations. Notably, we provide evidence that stellar activity may be responsible for the mismatch between observations and theory through two different channels: inhibition of convection or effects of a significant starspot coverage. The activity hypothesis is put to a test with observational diagnostics and some of the consequences of the large starspot coverage are evaluated. We conclude that stellar activity likely plays a key role in defining the properties of active low-mass stars and that this should be properly taken into account when investigating young, active stars in clusters or star-forming regions. ",
        "introduction": "Evidence collected over the past years convincingly shows that the components of low-mass eclipsing binary stars are not adequately reproduced by current evolution models. The observations yield stars that are systematically larger and cooler than theoretical calculations by about 10\\% and 5\\%, respectively. But, in contrast, these binary components are found to have luminosities that are in good agreement with those of single stars and also with model predictions. This is the reason why studies focused solely on the mass-luminosity relationship did not reveal any mismatch between observations and theory and, further, obtained very tight mass and luminosity correspondence when using infrared photometry (e.g., \\citealt{Delfosse2000}).  \\begin{figure*}[t] \\centering \\includegraphics[scale=0.4]{fig1a.eps} \\includegraphics[scale=0.4]{fig1b.eps} \\caption{\\footnotesize Left: $M$--$R$ plot for low-mass eclipsing binary stars with empirical determinations. Right: Same but for stars with uncertainties below 3\\%. The solid line represents a theoretical isochrone of 1 Gyr calculated with the \\citet{Baraffe1998} models.} \\label{fig:fig_EBs} \\end{figure*}  A revision of the available observational evidence and some hints about the possible resolutions of the discrepancies were provided by \\citet{Ribas2006}. We refer to that study for background information and detailed references. Here we shall only discuss new results appeared since 2006 and also the progress made in understanding the origin of the difference between model predictions and observations of fundamental properties of low-mass stars. ",
        "conclusions": "Without excluding at this point any other scenario, it seems clear from the tests discussed above that stellar activity plays a key role in defining the structure and radiative properties of low-mass stars. This could be very important in the context of young, and therefore active, low-mass stars, for example in clusters or star-forming regions. Failure to account for the effects of stellar activity can lead to strong potential biases in the determination of the ages of these objects and their ensembles. As discussed by \\citet{Morales2007}, a simple calculation using the temperature differentials found for single active stars indicates that ages of young clusters determined from (active) low-mass stars could be systematically underestimated by about 40\\%. Interestingly, this difference agrees well with the discrepancy between color-magnitude diagram ages and Li depletion boundary ages in young clusters (e.g., \\citealt{BarradoyNavascues2004})."
    },
    "0711/0711.1274_arXiv.txt": {
        "abstract": "From a uniform analysis of a large (8.5 Ms) Rossi X-ray Timing Explorer data set of Low Mass X-ray Binaries, we present a complete identification of all the variability components in the power spectra of black holes in their canonical states. It is based on gradual frequency shifts of the components observed between states, and uses a previous identification in the black hole low hard state as a starting point. It is supported by correlations between the frequencies in agreement with those previously found to hold for black hole and neutron stars. Similar variability components are observed in neutron stars and black holes (only the component observed at the highest frequencies is different) which therefore cannot depend on source-specific characteristics such as the magnetic field or surface of the neutron star or spin of the black hole. As the same variability components are also observed across the jet-line the X-ray variability cannot originate from the outer-jet but is most likely produced in either the disk or the corona. We use the identification to directly compare the difference in strength of the black hole and neutron star variability and find these can be attributed to differences in frequency and strength of high frequency features, and do not require the \\emph{absence} of any components. Black holes attain their highest frequencies (in the hard-intermediate and very-high states) at a level a factor $\\sim$6 below the highest frequencies attained by the corresponding neutron star components, which can be related to the mass difference between the compact objects in these systems. ",
        "introduction": "\\label{sec:introbhns}  Matter orbiting a compact object in an accretion disk is expected to be relatively unaffected by the nature of the compact object down to a few Schwarzschild radii from the center. Therefore, similar phenomena may occur in the process of accretion onto black holes and low magnetic field neutron stars in low mass X-ray binaries (LMXBs), despite some of the differences between these two types of compact objects, such as the presence or absence of a solid surface or an intrinsic magnetic field. Indeed, many similarities are seen in X-ray observations of neutron stars and black holes, and in fact, traditionally, distinguishing between neutron star and black holes based on other characteristics than X-ray bursts, pulsations and mass estimates has proven to be difficult. However, in both X-ray spectroscopy \\citep[e.g.][]{nar97,rut00,barret01,gar01,nar02,dogier03,barret04,titshap05} and X-ray timing, differences between black holes and neutron stars have begun to emerge. In particular with respect to timing, neutron stars are characterized by strong, highly tunable twin kHz quasi periodic oscillations (QPOs) at frequencies up to $\\sim1300$ Hz \\citep[for a review see][]{klis00,klis06} whose phenomenology is quite different from that of the high-frequency QPOs seen in black holes \\cite[e.g.][]{homan03a,homan05a,rem99b,rem02c,stroh01a,klis06}. Also, \\cite{sun00} have suggested that in some states neutron stars show considerably more power above several 100 Hz as compared black holes in any state.  \\begin{figure}[!htb] \\epsscale{1.0} \\plotone{f1.ps} \\caption[]{The canonical black hole states and their location in the hardness-intensity diagram. The track shown is an example of those that are followed by black hole transients in outburst, but persistent sources too can follow parts of such as track. Note that as depicted flares can occur in the HS during which the energy spectrum get harder and the power spectra are very similar to those seen in the hard or soft VHS. In this paper we use the names in bold face, see text. The grey area indicates the IMS region, which we further distinguish in the hard- and soft VHS and the hard IMS. In this figure we also show the jet-line \\citep{fender04,kording06} which can be found parallel to the line indicating the transition to and from the HS. The jet-line marks a transition between radio loud to a radio quiet states. Note that during a hard flare in the HS, the radio can be switched on again.\\\\ } \\label{fig:hidstatesBH} \\end{figure} The observation of recurrent, correlated spectral and timing behavior led to the introduction of the ''canonical'' states for the black hole and neutron star systems. For the black holes the definition of these states has undergone some evolution over the years. Initially luminosity was seen as the main determinant of state while in later work spectral hardness and two-dimensional classification schemes dominated, and the terminology reflects this historical process \\citep[see for instance][]{klis95,tan95,homan01,bell05b,klis06}. We use the state definitions and terminology that resulted from this process, and in particular classify source state primarily based on the relative location in the hardness-intensity diagram (HID) and power spectral characteristics. We distinguish the following states, see Fig.~\\ref{fig:hidstatesBH} for a schematic HID: the \\emph{low-hard state} (LS), characterized by a hard power-law dominated X-ray spectrum and strong, low-frequency band-limited noise in the power spectrum, the \\emph{high-soft state} (HS) with a soft, disk black-body dominated X-ray spectrum and little rapid variability, and an \\emph{intermediate state} (IMS), with both power law and disk black-body contributing appreciably to the spectrum and exhibiting the most complex variability characteristics, including most of the QPOs. All three states occur over a wide range in luminosity, although towards the lowest luminosity levels sources consistently enter the LS. Following \\cite{bell05b} within the IMS we further distinguish a \\emph{hard intermediate state}, which shows somewhat harder X-ray spectra (although not as hard as in the LS) and power spectra with moderately strong band limited noise (BLN; somewhat weaker and at higher frequency than that in the LS), and the \\emph{soft intermediate state} with softer X-ray spectra (not as soft as in the HS) and power spectra lacking this BLN component but showing power-law noise. The hard IMS has characteristics in common with the LS and the soft IMS with the HS; the hard--soft IMS transition may be the most physically abrupt one between what may considered to be just two basic black hole states: hard and soft. The IMS was first observed at high luminosity (in the peak of black hole transient outbursts) and for that reason was originally called very high state \\citep[VHS;][]{miya91}, and the rapid transitions between hard and soft IMS at high luminosity were already a defining characteristic of the VHS when it was first named. At lower luminosity the IMS usually shows up as hard IMS. To distinguish between high and low luminosity occurrences of the IMS, we shall indicate the high-luminosity realizations of hard and soft IMS as \\emph{hard} and \\emph{soft VHS}, respectively, while the low flux instances of the IMS will simply be referred to as hard (and if it occurs, soft) IMS. We consider hard IMS and hard VHS as basically the same state but at different luminosity levels, as their power spectra very similar (see for instance Sect.~\\ref{sec:lshvhshims}).  For completeness we note that \\cite{clintock04} proposed alternative definitions and names for black hole states, predominantly based on the energy spectral behavior and therefore not in one-to-one correspondence with the states distinguished here. They define a 'hard state' which is closest to the LS and a 'thermal-dominant state' closest to the HS, and further distinguish a 'steep power law state' that is closest to the soft IMS. Contrary to the current work, that classification does not aim to provide coverage of all observable states, but rather to exemplify a number of specific combinations of X-ray spectra and rapid X-ray variability in order to provide a focus for physical modeling. As a consequence of this, many observations not fitting the \\cite{clintock04} definitions occur. These are described as 'intermediates' between these three states; in particular the hard IMS in their description is intermediate between hard and steep power law states. In the work presented here we will use the canonical state definitions as summarized above, in which the hard and soft VHS and hard IMS are all different realizations of the IMS found in between the LS and HS, as shown in Fig.~\\ref{fig:hidstatesBH}. In that figure we also indicated the possibility of soft IMS occurrences during the low luminosity intervals of the IMS, similar to the soft VHS at high luminosities, as suggested for instance by \\cite{bell05b}. Up to date no unambiguous detection of such a soft IMS has been reported, perhaps due to a combination of the fact that the source generally moves rapidly through this part of the HID, and the lower statistics at these low luminosities. We finally note that the results presented here are not sensitive to the precise naming scheme for the black hole states.  In Fig.~\\ref{fig:hidstatesBH} we also indicate the position of the jet-line which marks the transition between radio-loud and -quiet states \\citep{fender04,kording06}. The jet-line in Fig.~\\ref{fig:hidstatesBH} is a generic one and in individual cases the transition between radio-loud and -quiet states is not always a sharp one. Radio emission in the radio-quiet (also known as \"quenched\") states is not ruled out \\citep[][and references therein]{kording06} but occurs at flux levels of up to a factor 50 lower than compared to the radio-loud states \\citep[][]{corbel01}. Remnant emission from large ejection events that occur close to the right-hand side of the jet-line may still be observed to the left of it \\citep[in the radio-quiet states;][]{kording06}. In any case, in individual cases the jet-line is observed to be diagonal (Fig.~\\ref{fig:hidstatesBH}), i.e. in a typical transient outburst the transition back to radio-loud at low X-ray luminosity occurs at a spectrally harder position than the quenching of the jet at high X-ray luminosity, but its exact position is still a matter of debate (Fender R., 2007, private communications).  The neutron star LMXBs are subdivided into Z sources and atoll sources based on differences in their correlated X-ray spectral/timing behavior. The names are derived from the tracks these sources produce in a color-color diagram \\citep[CD;][]{has89}. Below we will only discuss the atoll sources, which have the following states (Fig.~\\ref{fig:hidstatesNS}): the extreme island state (EIS), the island state (IS) and the banana branch subdivided into lower left banana (LLB), lower banana (LB) and upper banana (UB) \\citep[for a review see for instance][]{klis06}. The precise topology of these tracks at low X-ray luminosity (EIS) is a matter of some debate \\citep{gier02,muno02,barret02,straaten03,reig04} but that is of no concern to us here.  \\begin{figure} \\epsscale{1.25} \\plotone{f2.ps} \\caption[]{The canonical atoll neutron star states and their location in the hardness-intensity diagram. Note that we use here the same, but for neutron stars unconventional, black hole HID representation as in Fig.~\\ref{fig:hidstatesBH}. The inset shows the same track using the conventional neutron star CD. The track shown is a typical one followed for neutron star atoll sources (persistent or transient); some sources show transitions between all states, while others are only found in few of them.} \\label{fig:hidstatesNS} \\end{figure}  The behavior of the black holes and the neutron stars in their different states shows similarities. Based on the spectral and timing characteristics, \\cite{klis94a} already suggested that the spectrally hard black hole state (LS) is similar to the spectrally hard neutron star atoll state (EIS; although at that time the distinction between IS and EIS was not yet commonly made), and that some of the softer black hole and neutron star states might also be identified with one another. Another way to associate neutron star and black hole phenomenology is by using the strong correlations that exist between the characteristic frequencies of most of their power spectral components. Some of these frequency-frequency relations look the same for black holes and (atoll) neutron stars, suggesting that the same physical processes underly these power spectral components, or, at least, the correlations between them. The frequency-frequency relations first identified were the Wijnands-van der Klis  relation \\citep[WK;][]{wvdk99} and the Psaltis-Belloni-van der Klis relation \\citep*[PBK;][]{pbk99}. The WK relation relates the break frequency of the band-limited noise and the centroid frequency of a low-frequency QPO above this break, and is the same for atoll sources and black holes. Z sources follow a parallel but slightly offset relation. The PBK relation relates the frequency of the 1-10 Hz high-frequency bump and that of a 0.1--1 Hz low-frequency QPO in the black hole and weak neutron star sources, and coincides with an extrapolation of the relation between the 100-1000 Hz lower kilohertz QPO frequency and the 10--100 Hz low-frequency (LF) QPO frequency found in both atoll and Z sources. More recent studies confirm these results, although it is also clear that the entire phenomenology cannot be caught in just these two correlations \\citep[e.g.][]{nowak00,bpk02,pott03,kalemci03,bell05b,stratos06}. \\cite{straaten02,straaten03} for the atoll sources extended this set of relations into a proposed 'universal scheme' of correlations which encompass the WK and PBK relations among several others.  \\cite{war02} and \\cite{mauche02} show that the PBK relation may even extend to white dwarf systems, which implies that even while some of the neutron star and black hole frequencies may require strong field gravity for their generation, at least the mechanism producing the {\\it correlations} between the frequencies must be generic to a broad class of accretion flows.  A complication regarding the frequency-frequency relations is the proper identification of the different types of variability phenomena, especially in the black hole power spectra. For the (neutron star) atoll sources the identification of the different components in the power spectra now seems relatively well established \\citep{disalvo01,straaten02,straaten03,straaten05,diego05,manu05}, even across different states, as components are observed to gradually change across state transitions. For the black holes no complete set of characteristic variability components observed across all canonical states exists. Only in the case of the black hole LS an identification, similar to that of the neutron star EIS, has been suggested \\citep{wvdk99,pbk99,nowak00,bpk02}.  In this paper we attempt for the first time to perform a complete identification of all black hole variability components across all canonical states. For this purpose we use the identification of the variability components in the black hole LS as presented by \\cite{bpk02}, and the observed amplitude changes and shifts in frequency of these components during transitions to the other states as observed in a number of black hole transients. Using these new identifications, we study the frequency-frequency relations and the variations in amplitude with frequency in a large number of black hole observations. We find that the frequency-frequency relations confirm our identifications, and are able to identify which components are responsible for the differences between the black hole and neutron star power spectra.  The analysis presented here is part of a much larger program in which we compare all of the available public \\emph{Rossi X-ray Timing Explorer (RXTE) Proportional Counter Array (PCA)} data for the eight neutron star sources Aquila~X-1, 4U~0614+091, Terzan 2, 4U~1636--536, 4U~1728--34, 4U~1608--52, GX~3+1, 4U~1820--30 and the nine black hole sources XTE~J1118+480, GX~339--4, XTE~J1650--500, GRO~J1655--40,XTE~J1748--288, GRS~1758--258, XTE~J1859+266, XTE~J1550--564 and 4U~1543--47. Our current analysis covers total exposure of $\\sim8500$ ks up to AO8. ",
        "conclusions": ""
    },
    "0711/0711.4384_arXiv.txt": {
        "abstract": "Young cavities in the X-ray emitting hot gas in galaxy clusters are often filled with radio synchrotron emission from cosmic rays. However, in the M87/Virgo cluster, where cavities are less prominent, X-ray observations show a 30-kpc long nearly radial filament of relatively cooler gas that projects from the cluster core into a large (40 kpc) radio lobe. We describe the dynamical relationship between these two very dissimilar observations with gas dynamical calculations that include the dynamical effects and spatial diffusion of cosmic rays. After cosmic rays inflate the cavity, they diffuse through the cavity walls, forming a much larger lobe. The cavities, which are most visible just after they have formed (in about 20 Myr), require a total cosmic ray energy that is more than 10 times larger than that usually assumed, $E = 4PV$. During the relatively brief cavity lifetime, a jet-like, low-entropy thermal filament is formed in the buoyant flow and moves at high subsonic velocities through the cavity center and beyond. After 100 million years, long after the cavity has disappeared, the relatively dense filament extends to 20-30 kpc and the cosmic rays have diffused into a quasi-spherical lobe 40 kpc in diameter. These computed X-ray and radio features agree well with those observed in M87/Virgo and resemble those in other ``relic'' cluster radio sources such as Abell 13 and Abell 133. Eventually, the filament falls back and shocks at the center of the cluster -- perhaps stimulating the famous non-thermal M87 jet -- and only the aging radio lobe remains. ",
        "introduction": "The means by which accretion energy is communicated from massive black holes in cluster-centered galaxies to the surrounding hot gas are commonly thought to be fundamentally non-thermal. In galaxy clusters non-thermal energy is evidently transported and deposited in distant cluster gas via jets visible at radio frequencies. This notion is supported by numerous observations of concentrated non-thermal radio emission in young kpc-sized cavities in the hot cluster gas observed with X-rays (e.g. Bohringer et al. 1993; Fabian et al. 2000; Birzan et al. 2006; Clarke et al. 2006). A full understanding of these remarkable energy outbursts can only be achieved with quantitative observational and theoretical studies of both cosmic rays and hot thermal gas in galaxy clusters.  In Mathews \\& Brighenti (2007) we showed how a localized source of cosmic rays, as for example the working surface of a jet, carves out cavities in an initially uniform hot gas. When cosmic rays diffuse rapidly through the thermal gas, the cavities produced with a given total cosmic ray energy are small and tend to disappear rapidly after the cosmic ray source turns off. If the same total energy is imparted to cosmic rays that diffuse more slowly through the gas, the size of the cavities and their longevity both increase. In this case the non-thermal energy is efficiently trapped within the cavity (Churazov et al. 2001). The case of most interest probably lies between these extremes. Non-thermal radio emission from cosmic ray electrons is largely confined within cavities (or jets) at early times, but at later times the cosmic rays diffuse into the surrounding hot gas to form radio lobes with sizes much larger than the cavities.  In the exploratory calculations described here we extend our previous calculations to study the 2D evolution of cavities formed with cosmic rays in hot cluster gas confined by gravity. Because of the diffusive nature of cosmic rays and the requirement that they eventually evolve into the large radio lobes observed, the total non-thermal energy required to create typical X-ray cavities of radius $r_{cav}$ can be much larger than the purely hydrodynamic work $\\sim P 4 \\pi r_{cav}^3/3$ required to displace uniform hot gas with pressure $P$.  Another, more specific motivation for the calculations described here are the {\\it Chandra} X-ray observations of M87/Virgo by Young et al. (2002) and Forman et al. (2005; 2007) that reveal a $\\sim25$-kpc long, nearly radial gaseous filament extending from the central galaxy M87 toward the SW. (We adopt a distance of 16 Mpc from Tonry et al. 2001.) This thermal feature, beautifully shown in Fig. 2 of Forman et al. (2007), is remarkable because it lies just along the radius of the large SW quasi-spherical ``outer'' radio lobe of radius $\\sim25$ kpc observed at 90 cm with the VLA (Owen et al. 2000). This geometrical relationship between the thermal filament and radio lobe can be seen in Fig. 5 of Young et al. (2002). By including the dynamical and diffusive effects of cosmic rays, we show how X-ray cavities evolve into radial filaments surrounded by large (radio) lobes. Our calculation described here is also relevant the evolution of some ``relic'' radio sources observed in other clusters. According to Giovannini \\& Feretti (2004) these extended, steep spectrum radio sources are usually located in the outer parts of the hot cluster gas, unrelated to any non-central galaxy. However, cosmic rays that diffuse from X-ray cavities can only produce relic radio lobes that are relatively close to the cluster core, such as those in Virgo, A13 or A133. Most relic sources are too large (400 - 1500 kpc) or too distant ($\\gta 1000$ kpc) to be explained in this way.  Previous computational studies of the evolution of X-ray cavities, in which cavities filled with super-heated, non-relativistic buoyant gas are formed or placed in the hot gas cluster atmospheres, have shown that the cold gas flowing around and beneath the buoyant cavity eventually projects up through the cavity, resulting in a radially outflowing thermal feature that is cooler and more metal-rich than the surrounding gas (e.g. Reynolds et al. 2005; Gardini 2007; Ruszkowski et al. 2007; Roediger et al. 2007). We therefore envisage a cavity-lobe evolution with two jets. The first jet is an AGN jet, perhaps similar to the famous 2 kpc non-thermal jet in M87, that injects cosmic ray energy into expanding cavities. While it is commonly (even universally) assumed that X-ray cavities are formed by jets from the galactic nucleus, there is very little observational evidence of these jets. The second jet is the thermal jet that forms from the vortex flow around the buoyant cavity which, as we show here, can project to great radial distances. We shall refer to this as a ``cavity jet'' because of its origin and resemblance to the long radial filament visible in {\\it Chandra} images. We show below that the cavity jet can be visible long after the associated X-ray cavity has disappeared from view and the cosmic rays that initially created it have diffused into a large (radio) lobe surrounding the cavity jet.  In our discussion of the cavity-lobe evolution that follows we have elected not to compute the radio frequency surface brightness in detail although this will be essential to fully confirm the accuracy of our proposed evolution. We impose this limitation to avoid introducing new uncertain parameters -- such as the radial variation of magnetic field strength in the cluster gas, the ratio of of cosmic ray electrons to protons, synchrotron losses, etc. -- that would further complicate the dynamical presentation which is sufficiently detailed already. Consequently, we shall regard our model for the radio lobe a success if the cosmic rays occupy approximately the same volume as the outer radio lobe in M87/Virgo and have enough energy to produce the radio luminosity observed. Finally, our intention is to study the physical nature of the cavity-radio lobe evolution and not attempt to solve the global cooling flow problem although the two are obviously related. ",
        "conclusions": "Radial jet-like thermal filaments and associated large radio lobes evolve from X-ray cavities that are inflated with cosmic rays that eventually diffuse through the cavity walls. The radio features created from cavity CRs are unlike either FRI or FRII morphologies, where non-thermal energy flows directly from the central black hole. In general cavities and cavity jets are not likely to coexist. X-ray cavities containing cosmic rays that later diffuse into larger radio lobes are likely to be observed early-on as they are being filled with cosmic rays or shortly thereafter. By the time the cosmic rays have diffused from an X-ray cavity to fill a large radio-emitting volume in the cluster gas, the initial cavity will have disappeared. The radio lobes will be accompanied for $\\gta 2 \\times 10^8$ yrs by nearly radial, low temperature, low entropy, metal rich thermal filaments similar to the M87/Virgo filament seen in {\\it Chandra} images.  Since X-ray cavities are observed in about 20 percent of all clusters (Birzan et al. 2004), it may be surprising that the fraction of clusters like M87/Virgo, A13 and A133 with cavity jets is much smaller. Moreover, the double outer radio lobes in M87/Virgo appear to be nearly identical radio sources symmetric about the center of M87, but only one lobe has a cavity jet. If the filaments last about 4-5 times longer than X-ray cavities, as in our model, the number of cluster filaments should be larger by the same factor, but this has not been observed to date.  The initial position of the cavity does not seem to have a large effect on the development or X-ray visibility of cooler cavity jets over times $\\sim10^8$ yrs. To explore this we repeated the calculation described in Figure 7 but with the initial cavity formed at $z = 20$ kpc instead of 10 kpc. In this case the cavity jet was broader and remained visible for slightly longer times (but with a somewhat lower amplitude). The cavity jet is visible from well within 20 kpc from the center of M87/Virgo to 30 kpc and beyond. The CR energy density $e_c(r,z)$ at time $10^8$ yrs resembles that in Figure 7 but with contours shifted out by $\\sim 10$ kpc along the $z$-axis.  Evidently many cluster filaments have been missed because most X-ray clusters are much fainter than M87/Virgo and have not been observed with deep, $\\sim$1000 ksec {\\it Chandra} exposures. If the CR diffusion coefficient $\\kappa$ is larger than our values for M87/Virgo, the cavities will be smaller for given $E_{ctot}$ and their lifetimes may be too short to form cavity jets. Nevertheless, we hope that our calculations will inspire a more dedicated search for additional cavity jets in X-ray cooling core clusters having large diffuse radio lobes. Combined observations of thermal filaments and diffuse radio lobes can be used to date the explosive events from which they both evolved and to study in detail the propagation of cosmic rays, the magnetic structure in the hot gas and the total cosmic ray energy involved. If M87/Virgo is typical, the total energy of cosmic rays released $E_{ctot}$ can exceed the work necessary to simply form and fill the cavities, $E_{c,est} = 4 P V$, by more than an order of magnitude.  The hot gas in the Virgo cluster is currently experiencing a sequence of strong -- and very different -- energy releases from the central black hole. In addition to the cavity jet event discussed here, the more famous (and more recent) 2 kpc jet is moving in a NW direction unrelated to the orientation of the cavity jet and its large radio lobes. In addition, a strongly buoyant mushroom-shaped flow directly to the East of M87 is seen in both X-rays and radio images (Churazov et al. 2001). Ruszkowski et al. (2007) propose that an intensified magnetic field in the buoyant stem of this flow has caused it to be prominent in the 90 cm radio image of Owen et al. (2000). However, the cavity jet in the {\\it Chandra} image of M87/Virgo does not appear in this radio image. This difference in the magnetic field strength may result because the gas in the Eastern plume originated deeper within the cluster gas than that in the cavity jet, assuming that the field strength increases (with gas density) toward the cluster core.  In addition to these three energy outflows associated with strong radio emission, a number of less energetic outflows from the core of M87 are apparent in optical emission. Fig. 2 of Young et al. (2002) shows the central 7 $\\times$ 7 kpc region of M87/Virgo at three very different frequencies: {\\it Chandra} X-ray, 6 cm radio and optical emission in the H$\\alpha$ + [NII] lines. The optical line emission is particularly interesting because about 5 plumes of warm ($T \\sim 10^4$ K) gas can be seen emanating from the M87 core, all in different directions. The brightest of these is a H$\\alpha$ + [NII] feature that lies just North of the bright 2 kpc non-thermal jet, but at a slightly different angle. We see a parallel here with the dusty plume we have discovered and discussed in the X-ray bright galaxy group NGC 5044 (Temi, Brighenti \\& Mathews 2007). The NGC 5044 feature can be understood as dusty buoyant gas heated near the central black hole that is rapidly cooled to $\\sim10^4$ K by electron-dust cooling. It seems entirely plausible that the warm gas H$\\alpha$ + [NII] features in M87 are created in the same manner, adding to the already wide diversity of energy releases in this remarkable galaxy/cluster.  We expect that the numerous intermittent energy events seen in M87/Virgo have created subsonic velocities throughout the hot gaseous atmosphere. By contrast, our cavity jet was computed starting with stationary cluster gas exactly in hydrodynamic equilibrium and this results in a perfectly axisymmetric flow that may differ in detail from that observed in M87/Virgo. For example, the slight curvature of the cavity jet visible in Fig. 2 of Forman et al. (2007) could arise from either from a small velocity perpendicular to the accelerating jet or from an initial velocity asymmetry near the base of the pre-filament cavity that slightly altered the direction of the cavity jet as it formed. Similarly, the ``braided'' nature of this filament seen in the {\\it Chandra} image (Fig. 3 in Forman et al. 2007) may result from small random transverse velocities present when the X-ray cavity was formed.  \\vskip0.1in Studies of the evolution of hot gas in elliptical galaxies at UC Santa Cruz are supported by NASA grants NAG 5-8409 \\& ATP02-0122-0079 and NSF grant AST-0098351 for which we are very grateful."
    },
    "0711/0711.2277_arXiv.txt": {
        "abstract": "It is standard in theoretical neutrino astrophysics to use a broken power law approximation, based on the Band function, to describe the average photon flux of the prompt emission of Gamma-Ray Bursts. We will show that this approximation overestimates the contribution of high energy $\\gamma$-rays (and underestimates low energy $\\gamma$-rays). As a consequence models that rely on this approximation overestimate neutrino event rate by a factor of $\\approx$~2 depending on Earth's column density in the direction of the GRB. Furthermore the characteristic energy of neutrinos that trigger a km$^3$ detector is typically 10$^{16}$~eV, higher than previously predicted. We also provide a new broken power law approximation to the Band function and show that it properly represents the photon spectra. ",
        "introduction": "It is commonly believed that prompt emission by GRBs is due to synchrotron radiation by electrons accelerated in internal shocks associated with relativistic jets (with a bulk Lorentz boost $\\Gamma$ of 100-1000). A review of the theoretical and observational status of GRBs is beyond the scope of this paper. The reader is referred to \\citet{2006RPPh...69.2259M}. The average prompt photon emission is often described by fitting a Band function \\citep{1993ApJ...413..281B}. High energy neutrino emission by GRBs in coincidence with the prompt $\\gamma$-ray photons has been proposed as a consequence of GRBs being a candidate to produce (UHE) ultra high energy (up to $\\approx$~10$^{20}$~eV) cosmic rays \\citep{1995PhRvL..75..386W,1995ApJ...453..883V}.  \\citet{1997PhRvL..78.2292W} have calculated the diffuse flux due to GRBs. Their work has been further improved by calculating neutrino emission for individual bursts \\citep{2004APh....20..429G} and by performing detailed GEANT4 simulations of proton-photon interactions in the internal shocks \\citep{2006PhRvD..73f3002M}. This body of work supports the hypothesis that neutrino detection by km$^3$ Cherenkov detectors will probably be in coincidence with a handful of bright GRBs.  Previous studies have correlated photon emission with neutrino emission by approximating the photon spectra with a broken power law. We will show that this approximation overestimates the contribution of high energy photons (and conversely underestimates the contribution from low energy photons). This has consequences in the number of events expected by neutrino telescopes such as IceCube \\citep{2007..IceCube..ICRC} and KM3NET \\citep{2007..KM3NET..Faro} and in the characteristic energy of the neutrinos detected.  In section \\S \\ref{sec:band} we describe the Band function. In section \\S \\ref{sec:photopion} we follow the standard calculation of HE neutrinos emission from GRBs. In section \\S \\ref{sec:bandapprox} we compare the Band function with its broken power law approximation. In section \\S \\ref{sec:eventrates} we calculate neutrino event rates. And in section \\S \\ref{sec:discussion} we discuss the consequences of the calculation shown here. ",
        "conclusions": "\\label{sec:discussion}  We have shown that the usual choice to describe the photon spectra in the calculation of neutrino fluxes from GRBs overestimates the contribution of high energy photons (and therefore contribution of low energy neutrinos is overestimated). This results in a higher exepected event rate by a factor of $\\approx$~2 for all models that make this assumption. The actual value of the overestimation of the expected number of events depends on the matter column depth that neutrinos must cross through Earth in the direction of the GRB. Also we have shown that the typical neutrino energy is $\\approx$~10$^{15}$~eV. The characteristic energy of neutrinos is a factor of $e$ larger than the values obtained by \\citet{2004APh....20..429G} and a factor of 10 than the average value used by \\citet{1997PhRvL..78.2292W}. For back of the envelope calculations we provide a new approximation to the Band function that is adequate for the computation of neutrino spectra.  \\citet{2005PhRvL..95..181101K} have discussed the effects of muon and pion energy losses leading to neutrino flavor flux ratios (at Earth) different than 1:1:1. Previous AMANDA searches for neutrinos \\citep{2007ApJ...664..397A,2007ApJ...InPress}, based on Waxman-Bahcall-like assumptions have argued that this effect is not important experimentally because the typical energy of the neutrinos detected is close to the neutrino break energy. Because, as we have shown, the \\textit{effective} neutrino break energy is close to the typical syncrhotron energy break, the effects described by \\citet{2005PhRvL..95..181101K} must indeed be relevant."
    },
    "0711/0711.2041_arXiv.txt": {
        "abstract": "We compute rates of period change (\\pdots) for the 215s mode in G117-B15A and the 213s mode in R548, first for models without axions, and then for models with axions of increasing mass. We use the asteroseismological models for G117-B15A and R548 we derived in an earlier publication . For G117-B15A, we consider two families of solutions, one with relatively thick hydrogen layers and one with thin hydrogen layers. Given the region of parameter space occupied by our models, we estimate error bars on the calculated \\pdots \\ using Monte Carlo simulations. Together with the observed \\pdot \\ for G117-B15A, our analysis yields strong limits on the DFSZ axion mass. Our thin hydrogen solutions place an upper limit of 13.5 meV on the axion, while our thick hydrogen solutions relaxes that limit to 26.5 meV. ",
        "introduction": "G117-B15A and R548 are pulsating white dwarfs with atmospheres dominated by hydrogen. These stars are called DAVs or ZZ Ceti stars. Their effective temperatures are close to 12,000 K. They are non-radial, g-mode pulsators (the restoring force is buoyancy). G117-B15A is the most stable optical clock known. The rate of change in period (\\pdot = dP/dt) of its most stable mode (215s) is $3.57 \\pm 0.82 \\times 10^{-15}$ s/s \\citep{keplerpdot}. That is, G117-B15A loses one ``tick'' every 1.7 billion years. We can use this well-measured \\pdot \\ to constrain among other things the mass of the axion, the best motivated dark matter candidate and in general, the emission rate of weakly interacting particles. While measuring a \\pdot \\ for R548's most stable mode (near 213s) is less trivial because it is a doublet, the current upper limit, $5.5 \\pm 1.9 \\times 10^{-15}$ s/s \\citep{Anjum0}, also indicates that in the near future R548 may become useful to constrain the axion mass as well.  A \\pdot ~ of order $10^{-15}$ for DAVs is consistent with a slow increase in period due to the cooling of the interior. It is an evolutionary timescale $-$ a measure of how fast the star is cooling. Knowing \\pdot, we can determine how much energy a star like G117-B15A is losing every second. This energy loss includes of course the radiation of photons (Mestel cooling, \\citealp{mestel}), but in principle could also include the emission of weakly interacting particles. While neutrinos immediately come to mind, the Standard Model of Particle Physics predicts that neutrino emission should be negligible in G117-B15A and R548 \\citep{winget04,myphdthesis}. One possibility includes axions and we apply our method to the latter in the present study.  Axions arise from an elegant solution to a problem with the Standard Model of particle physics, the strong CP problem (e.g. \\citealp{axions,myphdthesis}). Along with supersymmetric particles, axions are currently favored candidates for dark matter. But they have not been discovered (neither have supersymmetric particles) and the theory of axions fails to place any constraint on their mass. The possible contribution of axions to dark matter depends on their mass. The mass of axions determines how strongly they interact with the matter we know, with more massive axions interacting more strongly. In turn, the stronger the interaction of axions with matter or light, the larger their emission rate. With pulsating white dwarfs, we can constrain the axion emission rates and therefore their mass.  We determine limits on the emission of weakly interacting particles by comparing their cooling effect on a white dwarf model with the observed cooling rates of that star (given by \\pdot). The strength of those limits depends on the uncertainties involved in the measurement of the \\pdot \\ and in the modeling. This method has been used before by \\citet{isern} and \\citet{corsico01} to derive an upper mass limit for axions. While the limits on the axion obtained by those authors were stronger than those obtained through other methods (see \\citealp{myphdthesis} for a review across fields of study), we now have the tools to greatly strengthen them. We present here the first study that takes into account the modeling uncertainties in a systematic way, allowing us to derive a strong upper mass limit for the axion. We present our results in a way that facilitates their application to other weakly interacting particles. We also publish for the first time the \\pdot \\ we would expect to observe in R548 as a function of the emission rate of weakly interacting particles. Together with a future measurement of \\pdot \\ for R548, these results can be used to reinforce the limits given by G117-B15A.  In \\citet{paper1}, thereafter paper I, we used a systematic, fine grid search to find best fit models to G117-B15A and R548. We were able to define the regions of parameter space the fits occupied. In this work, we use Monte Carlo simulations to calculate \\pdots \\ that include theoretical uncertainties, for varying axion mass. We begin with a brief introduction to axions in section 2. In section 3, we give an overview of the work done by \\citet{isern} and \\citet{corsico01}. We detail our method in section 4, discuss our results in section 5, and conclude in section 6. ",
        "conclusions": "We computed \\pdots \\  for the 215s mode in G117-B15A and the 213s mode in R548 first for models without axions, and then for models with axions of increasing mass. The thin hydrogen, zero axion models for G117-B15A have a \\pdot \\ of $2.98 \\pm 0.17 \\times 10^{-15} \\rm s/s$, consistent with the observed value of $3.57 \\pm 0.82 \\times 10^{-15} \\rm s/s$. The zero axion, thick hydrogen models have a \\pdot \\ of $1.92 \\pm 0.26 \\times 10^{-15} \\rm s/s$, 2 sigma below the observed \\pdot \\ and require additional cooling through axions to become consistent with that observed value. The zero axion models for R548 have a \\pdot \\ similar to the thin hydrogen \\pdot \\ for G117-B15A ($\\sim 3 \\times 10^{-15} \\rm s/s$), consistent with the best current limit on \\pdot \\  for R548's 213s mode.  If we believe that G117-B15A must be similar to R548, then we find an upper limit on the axion luminosity of $1.1 \\times 10^{-31}$ erg/s, corresponding to an axion-electron coupling constant \\gae \\ $=3.78 \\times 10^{-13}$ , or an axion mass of 13.5 meV. A second family of fits with thicker hydrogen layers relaxes that limit to $L_{a} \\leq 4.4 \\times 10^{-31}$ erg/s, \\gae \\ $\\leq 7.4 \\times 10^{-13}$ or $\\rm m_a \\leq $ 26.5 meV. What makes this limit unique is not the fact that it is lower than previous limits. What makes it unique is its strength. It is based on an unbiased exploration of the parameter space of the models. Compared to typical axion searches (see \\citealp{myphdthesis} for a review), the present analysis involved well-known physics and relatively small, well-identified uncertainties, considered systematically. Improved asteroseismological models will help us choose between the two families of best fits for G117-B15A and improve this limit.  An even longer data baseline for G117-B15A will reduce the size of the error bars on the measured \\pdot \\ and further narrow the allowed range for the axion mass. We can soon expect a useful \\pdot \\ for R548 as well, which will reinforce the results of this study. As a part of his planet search around pulsating white dwarfs,  Mullally has determined upper limits of $10^{-13}$ s/s for an additional 9 white dwarfs \\citep{fergalphdthesis} and has identified more blue edge DAVs with stable modes. With Argos \\citep{argos}, the high-speed photometry data collected at McDonald observatory as a part of the on-going observation of hot DAVs should begin to yield measurements of \\pdots \\ for those stars within the next 3 to 5 years. It is merely a question of time until we can place an even tighter constraint on the emission of weakly interacting particles in these stars.  In the present study, we assumed that any unseen source of energy loss was due to axions. The results we present can readily be used to constrain the emission of other weakly interacting particles possibly produced inside white dwarfs, such as supersymmetric particles."
    },
    "0711/0711.2331_arXiv.txt": {
        "abstract": "A shape of lateral distribution for charged particles in events with energy above $10^{19}$~eV is considered. Two methods were used for individual LDF parametrization. In the first approach, the index of power was determined for generalized Greisen-Linsley approximation. In second, mean square radius of the shower was determined for approximation proposed by Lagutin et al. Comparison of resulted parameters is presented for individual events arrived from different celestial regions~--- Galactic planes and the region with increased flux of particles with $E_{0} \\ge 10^{19}$~eV (according to Yakutsk array): $1.7\\text{h} - 3.7\\text{h}$ right ascension ; $45\\dgr - 60\\dgr$ declination. ",
        "introduction": "The knowledge of the lateral distribution function (LDF) of charged particles from extensive air shower (EAS) is vital for experiments in the field of ultra-high energy cosmic ray (UHECR) studying. It is LDF that defines main shower parameters such as $\\So$~ (charged particle density at the distance $600$~m from the core) and thus~--- primary energy.  In this paper we consider parameters of individual LDFs resulted from revision of high energy events registered at the Yakutsk EAS array. The aim of this work is to trace possible correlation between parameters of individual showers and their arrival directions on the sky, especially for Galactic planes and for the region with significantly increased UHECR flux, detected by Yakutsk group~\\cite{bib:Ivanov}. ",
        "conclusions": ""
    },
    "0711/0711.2936_arXiv.txt": {
        "abstract": "We give a gravitational upper limit for the mass of static degenerate fermionic dark matter objects. The treatment we use includes fully relativistic equations for describing the static solutions of these objects. We study the influence of the annihilation of the particles on this mass limit. We give the change of its value over the age of the Universe with annihilation cross sections relevant for such fermions constituting the dark matter. Our calculations take into account the possibility of Dirac as well Majorana spinors. ",
        "introduction": "Existence of a considerable amount of cold dark matter in the present Universe has been regarded as a robust ingredient in recent astrophysics and cosmology. Several hypothetical particles were suggested in order to solve the problem of the cold dark matter. Some of them are bosons: axions, Q-balls\\ldots The others are fermions: neutralinos, axinos, gravitinos, quintessino\\ldots (One can read a good summary about the non-barionic dark matter candidates in \\cite{Fornengo}.)  We study some general features of the degenerate fermionic matter here. Degeneracy means that every energy states are occupied, that is such a system is spherical symmetric and does not rotate. Our aim is to determine the maximal possible mass of {\\it static} degenerate fermion balls.  First we make a naive non-relativistic estimation because the non-relativistic system has been studied in the literature for decades. Several papers were devoted to the non-relativistic degenerate fermionic dark matter with or wihtout barionic component (see for example \\cite{ArbolinoRuffini,Viollier1,Bilic,Messer,HertelThirring1,HertelThirring2}). Papers whose topic is some aspect of the relativistic behaviour of these systems are rare (\\cite{Kupi}). We devote the second part of this paper to giving  an exact relativistic mass limit for these systems. In the third part of the paper we study the influence of the annihilation of the particles on this mass limit.  In non-relativistic case we use polytropic equation of state, in relativistic case we do not use such approximation but the exact equation which one can get for degenerate fermions.   If we assume that the dark matter consists of neutral, weakly interacting fermions, their spatial distribution is driven by gravity and Fermi distribution. $kT \\ll m_0 c^2$ is valid for such kind of particles today, that is if the density is sufficiently high, the degenerate ($T=0$ K) Fermi distribution approximates the real distribution well. $k$ is the Boltzmann constant, $T$ is their temperature, $m_0$ is their rest mass and $c$ is the speed of light. One can write that \\begin{equation} \\label{fermi} z\\pi {p_{F}}^3/3 h^3=\\varrho/m_0, \\end{equation} where $\\varrho$ is the mass density of particles, $p_{F}$ is their Fermi momentum and $h$ is the Planck constant. If the particles are Majorana spinors, the $z$ factor is 8. If Dirac spinors, then 16. ",
        "conclusions": "We studied self-gravitating systems of weakly interacting fermions. We determined the limit mass above which systems undergo gravitational collapse. We could see that this limit mass is a function of the rest mass of the particles: $M=constant/m^2$. However, this relation was known before we give the exact value of this $constant$ which depends on the type of fermions, namely whether they are Dirac or Majorana spinors. In Majorana case the central density cannot be as high as in Dirac case at same total mass. One can conclude from this that the value of this $constant$ is higher in Majorana case than in Dirac case. It is an upper or absolute masslimit of such objects.  We calculated the lower limit of the limit mass used the recent upper bound on annihilation cross section of dark matter particles. We could see that this limit is not strict but we can not state more on the base of the recent knowledge.  Further investigations are needed to answer the question what is the probability that such objects gather mass above the mass limit and undergo gravitational collapse or how much dark matter they could  \"burn\" if they could not collapse."
    },
    "0711/0711.2113_arXiv.txt": {
        "abstract": "{Theory and observations of the dominant thermal continuum emission in AGNs are examined.  After correction for reddening, the steady state AGN optical--UV spectral energy distributions (SEDs) are very similar. The SEDs are dominated energetically by the ``big blue bump'' (BBB), but this bump never shows the $\\nu^{+1/3}$ spectrum predicted for a standard thin accretion disk with a $r^{-0.75}$ radial temperature gradient.  Instead, the observed optical--UV SED implies a temperature gradient of $r^{-0.57}$ independent of the thickness of the disk. This means that there is some flow of heat outwards in the disk.  The disk is large and the region emitting the optical continuum is as large as the inner broad-line region (BLR). Because optical variability is seen in all AGNs on the light-crossing time of the BLR, variations must propagate at close to the speed of light, rather than on dynamical timescales. This argues that the energy-generation mechanism is electromagnetic rather that hydrodynamic. Since the velocities are near the speed of light, there can be significant local anisotropy in the emission. The large rapid variations of the BBB imply that the magnetohydrodynamic energy generation is fundamentally unstable. Because of the inevitable radial temperature gradient in the accreting material, {\\it different spectral regions come predominantly from different radii}, and variations in different spectral regions correspond to variability at different radii. This explains the frequently observed independence of X-ray and optical variations, cases of variability at lower energies leading variability at higher energies, and rapid changes in emission-line reverberation lags. Some observational tests of the local variability hypothesis are proposed.}  \\resumen{% }   \\addkeyword{accretion, accretion disks }  \\addkeyword{black hole physics}  \\addkeyword{galaxies: active}  \\addkeyword{quasars: general}  \\addkeyword{galaxies: Seyfert}   \\begin{document} ",
        "introduction": "A symposium in honor of Deborah Dultzin's 50th birthday\\footnote{The ``60th birthday'' in the symposium title is clearly a mistake. Deborah can't really have turned 60 because that would have implications for my own age that I am not sure I am quite ready to face yet!} is a good time to look at AGN variability because it is a major area of research in which Deborah has worked (see, for example, Dultzin-Hacyan et al.\\@ 1992, Crenshaw et al.\\@ 1996, Edelson et al.\\@ 1996, de Diego et al.\\@ 1998, and Ram{\\'{\\i}}rez et al.\\@ 2003). There is much discussion elsewhere in these proceedings of variability of the {\\it blazar} component of AGN emission (an area of research in which Deborah and her collaborators have been particularly heavily involved), but in this paper I am just going to focus on what I think are some of the fundamental implications of {\\it non}-blazar AGN continuum and continuum variability observations when confronted with simple accretion theory. For a more comprehensive overview of the state of our knowledge of AGN variability across the electromagnetic spectrum the reader see papers in Gaskell et al.\\@ (2006). ",
        "conclusions": "Variability rules out the standard popular geometrically-thin ``$\\alpha$-disk'' paradigm of Pringle \\& Rees (1971) and Shakura \\& Sunyaev (1972). Variability in the UV and optical is simply too fast and much too simultaneous to be due to fluctuations in the accretion rate in a standard thin disk.  The rapidity and simultaneity instead require that the variability mechanism propagates close to the speed of light.  Since the amplitude of AGN variability is enormous the variability mechanism is the main energy generating mechanism of AGNs.  Both theory and observations strongly point to variability being {\\it localized} in the disk, and there are tests that can be made of this hypothesis."
    },
    "0711/0711.2055_arXiv.txt": {
        "abstract": "We present detections of the ground-state $1_{10}\\rightarrow1_{01}$ transition of ortho-\\h2o\\ at 557 GHz in 18 molecular outflows based on data from the {\\it Submillimeter Wave Astronomy Satellite} (SWAS).  These results are combined with ground-based observations of the J=1-0 transitions of $^{12}$CO and $^{13}$CO obtained at the {\\it Five College Radio Astronomy Observatory} (FCRAO).  Data from {\\it Infrared Space Observatory} (ISO) for a subset of the outflows are also discussed.  Assuming the SWAS water line emission originates from the same gas traced by CO emission, we find that the outflowing gas in most outflows has an ortho-\\h2o\\ abundance relative to H$_2$ of between about 10$^{-7}$ and 10$^{-6}$. Analysis of the water abundance as a function of outflow velocity reveals a strong dependence.  The abundance of ortho-\\h2o\\ increases with velocity and at the highest outflow velocities some of the outflows have relative ortho-\\h2o\\ abundances of order 10$^{-4}$.  However the mass of very high velocity gas with such elevated \\h2o\\ abundances represents less that 1\\% of the total outflow gas mass. The ISO LWS observations of high-J rotational lines of CO and the 179.5 $\\mu$m transition of ortho-\\h2o\\ provide evidence for a warmer outflow component than required to produce either the SWAS or FCRAO lines. The ISO line flux ratios can be reproduced with C-shock models with shock velocities of order 25 km s$^{-1}$ and preshock densities of order 10$^5$ cm$^{-3}$; these C-shocks have post-shock relative water abundances greater than 10$^{-4}$. The mass associated with the ISO emission is also quite small compared with the total outflow mass, and is similar to that responsible for the highest velocity  water emission detected by SWAS.  Although the gas responsible for the ISO emission has elevated levels of water, the bulk of the outflowing gas has an abundance of ortho-\\h2o\\ well below what would be expected if the gas has passed through a C-shock with shock velocities greater than 10 km s$^{-1}$.  Gas-phase water can be depleted in the post-shock gas due to freeze-out onto grain mantles, however the rate of freeze-out is too slow to explain our results.  Therefore we believe that only a small fraction of the outflowing molecular gas has passed through shocks strong enough to fully convert the gas-phase oxygen to water.  This result has implications for the acceleration mechanism of the molecular gas in these outflows. ",
        "introduction": "Stellar winds are believed to play a pivotal role in the process of star formation. During the accretion phase of star formation, stellar winds carry away excess angular momentum that would otherwise prevent further collapse.  The winds interact with and accelerate the surrounding medium, and as a result, drive molecular outflows with masses often much greater than that of the young star itself.  However, the exact mechanism by which these winds interact with and accelerate the surrounding medium is still a matter of debate. A number of different models have been proposed including jets with bow shocks, jets with turbulent entrainment, and wide angle winds \\citep{arce02}.  The various models result in different spatial and velocity distributions of the entrained material, and therefore one approach to distinguish between the proposed models is to map the spatial and velocity structure of the outflowing molecular gas \\citep{arce06}.  Alternatively, the various acceleration mechanisms heat the gas in different ways, resulting in changes in the chemistry of the outflow as ices are sublimated and endothermic reactions pathways are opened.  Thus, determinations of the chemical abundances of the outflowing molecular gas may permit us to distinguish between the various mechanisms.  The review by \\citet{arce06} suggested that molecules such as SiO, CH$_3$OH, H$_2$O and sulfur-bearing molecules have their abundances affected by outflow activity.  For example, \\citet{bach01} study of L1157 revealed significant abundance variations and suggested that the chemistry of the outflows may be useful in understanding the time evolution of outflows.  The abundance of water can be strongly affected by shocks \\citep{kauf96,berg98} such as those predicted to exist in outflow regions.  Moderate velocity shocks are capable of producing large abundances of water by releasing frozen water from dust grains and by driving all free oxygen into water through a series of gas phase chemical reactions \\citep{kauf96}.  Thus, water should be a good tracer of shocks.  The chemical reactions that produce water have been thoroughly modeled and it has been shown that the production mechanism and resulting abundance is very sensitive to temperature \\citep{elit78a,elit78b}. At gas temperatures below $\\sim$300 K water is formed most readily through a series of ion molecule reactions that lead to the formation of H$_3$O$^+$. The dissociative recombination of H$_3$O$^+$ has several possible outcomes, but has been measured in the lab to produce water in approximately 25\\% of the interactions \\citep{jens00}.  This process of water production is relatively slow  and modelling has shown that in a quiescent medium of density 10$^5$ cm$^{-3}$ and temperature 30 K, a water abundance of roughly 10$^{-7}$ relative to molecular hydrogen is achieved after 10$^5$ years \\citep{berg98}.  Once the temperature of the gas rises above $\\sim$300 K a series of endothermic neutral-neutral reactions are activated.  At elevated temperatures, these reactions are very rapid compared to the ion molecule reactions, and all free oxygen is quickly driven into water \\citep{elit78a,elit78b}.  \\citet{kauf96} show that the passage of a moderate C-shock with a velocity greater than $\\sim$10 {\\kms} is able to elevate temperatures above 300 K long enough to produce greatly enhanced water abundances on the order of $\\sim10^{-4}$ relative to H$_2$.  More importantly, \\citet{berg98} showed that the enhanced water abundance persists in the post-shock gas after the gas has cooled and to significantly reduce the water abundance required  $\\sim10^5$ years even at densities as high as 10$^5$ cm$^{-3}$. Since outflow ages are typically estimated to be of order 10$^5$ years or less, any of the outflowing molecular gas that has passed through a moderate shock should be imprinted with this greatly elevated water abundance.  After $\\sim10^5$ years the abundance of water will drop as the water chemistry comes into equilibrium at the lower gas temperature and water freezes onto the cold dust grains.  Evidence for this enhancement is quite apparent in the water observations of some outflows such as in Orion-BN/KL \\citep{harw98,wrig00,meln00b}.  The ideal nature of water as a tracer of shocked regions is countered by the difficulty of detection from ground based observatories.  This problem was overcome by the successful launch of three satellites, the {\\it Infrared Space Observatory} (ISO), the {\\it Submillimeter Wave Astronomy Satellite} (SWAS) and {\\it Odin}. ISO was capable of detecting a number of water transitions excited in gas warmer than about 80 K \\citep{cleg96} while SWAS \\citep{meln00a} and {\\it Odin} \\citep{nord03, hjal03} can observe the fundamental ortho-{\\h2o} $1_{10}-1_{01}$ transition at 538.3 $\\mu$m. More recently, the {\\it Spitzer Space Telescope} has detected several mid-infrared water transitions toward NGC 2071 \\citep{meln07}. The fundamental transition of ortho-\\h2o\\ can be readily excited at temperatures greater than about 20 K, permitting SWAS and {\\it Odin} to trace out much cooler post-shock gas than ISO or {\\it Spitzer}.  SWAS has detected ortho-{\\h2o} emission associated with the quiescent dense molecular gas in numerous cloud cores.  The relative abundance of ortho-{\\h2o} to H$_2$ in these regions was found to be on the order of $\\sim$10$^{-9}$ to 10$^{-8}$ \\citep{snel00b}, roughly 2 orders of magnitude less abundant than predicted by quiescient gas phase chemistry. The discrepency is likely the result of water ice mantle formation on the dust grains \\citep{berg00}.  In this paper, we present SWAS detections of {\\h2o} emission from 18 well studied outflow regions.  We selected nearby outflows (within 1 Kpc) with significant SWAS observations that we were able to obtain additional mapping data from the {\\it Five College Radio Astronomy Observatory} (FCRAO). These outflows are driven by a mix of high and low luminosity Young Stellar Objects (YSOs). Eight of these outflows (L1448-mm, NGC 1333-SVS13, HH25mm, $\\rho$ Oph A (VLA1623), L1689N, Ser SMM1, L1157, L1228) are driven by relatively low luminosity YSOs that are in most cases individual low mass stars.  The other ten sources (GL490, Orion KL, OMC 2, NGC2071, MonR2, NGC2264 D, NGC2264 C, IC 1396N, S140, Ceph A HW2) are driven by much higher luminosity YSOs that are associated with newly forming groups and clusters. We combine the SWAS {\\h2o} detections with maps of {\\co} and {\\13co} emission obtained at the  FCRAO 14-m telescope to determine the water abundance in the outflowing gas.  Additionally, we use published ISO results for six of these outflows to further expand our understanding of the SWAS results. ",
        "conclusions": "The 1$_{10} - 1_{01}$ transition of ortho-\\h2o\\ at 538 $\\mu$m has been detected by SWAS in 18 molecular outflows.  The \\h2o\\ line profiles are similar to the line profiles observed for the J=1-0 transition of \\co\\ and suggest that the emission seen in both species may be produced by the same gas. If we assume that the SWAS \\h2o\\ emission arises in the same gas that makes up the bulk of the molecular outflow, we find that the outflowing gas has an ortho-\\h2o\\ abundance relative to H$_2$ typically between 10$^{-6}$ and 10$^{-7}$.  However, there are a few exceptions: most notably Orion KL and L1157 have anomalously high relative ortho-\\h2o\\ abundances of about 10$^{-5}$, and Mon R2 has an anomalously low relative abundance of about 10$^{-8}$.  The relative ortho-\\h2o\\ abundances in Table 2 have substantial uncertainty that arise almost solely due to the sensitivity of the derived water abundances to the assumed temperature and density of the outflowing gas, both of which are poorly known.  The {\\13co} J=5-4 observations by SWAS suggest that our physical model cannot be too much in error. and thus we believe that the abundance of water in most of these outflows is elevated relative to that measured in quiescent cloud gas.   We also derive the CI abundance in the outflowing gas, and find values that are similar to quiescent cloud material, and thus unlike water, the abundance of atomic carbon appears to be unaffected by the outflow activity.  We also analyzed the velocity dependence of \\h2o, and find that the abundance of ortho-\\h2o\\  varies significantly with velocity.  In nearly all of the outflows, we find a steady increase in the ortho-\\h2o\\ abundance with increasing radial velocity of the outflowing gas.  In the most striking examples (NGC 1333 SVS13 and HH25mm), the abundance of water increases by two orders of magnitude from the lowest velocity to the highest velocity outflowing gas.  Nevertheless, the mass of outflowing gas with greatly elevated \\h2o\\ abundance is very small, representing at most only  $\\sim$1\\% of the total outflow gas mass.  The ISO observations provide evidence for the presence of a much warmer outflow component in those outflows observed.  With the exception of Orion KL, our C-shock modeling of the ISO emission suggest that this warm gas constitutes less than 1 \\% of the total outflow mass and arises from an extremely small fraction of the molecular outflow solid angle. However, this gas is expected to have a relative abundance of water in excess of 10$^{-4}$. The presence of this warmer component adds an additional complexity to abundance determinations in the cooler gas.  Our shock modeling, although simple, suggests that a significant fraction of the water line flux observed by SWAS could arise in this warmer gas component.  Unfortunately the ISO observations did not cover the full extent of the SWAS beam nor did it have velocity resolution, so it is extremely uncertain what fraction of the SWAS line flux is associated with this warm gas probed by ISO.  We estimate that for the outflows with both ISO and SWAS observations, as much as 50 \\%\\ of the SWAS flux could originate in the warmer gas component. If any significant fraction of the SWAS \\h2o\\ line flux arises from the warm gas, then we have overestimated the line flux from the cooler gas traced by the low-J transitions of CO, and consequently overestimated the ortho-\\h2o\\ abundance in this cool gas.  Therefore our estimates of the water abundance presented in Table 2 are likely to be too large.  The extent to which the abundance is overestimated is very uncertain, but if only half of the SWAS line flux is due to the warm gas this only reduces the \\h2o\\ abundance by a factor of two, which is well below other uncertainties in our abundance determination. Thus, unless nearly all of the SWAS line flux is due to this small mass of warm, shocked gas, the ortho-\\h2o\\ abundance in the low velocity outflowing gas is enhanced relative to that measured for quiescent cloud gas (relative abundance of only 10$^{-9}$ to 10$^{-8}$ according to \\citet{snel00b}).  The studies by \\citet{kauf96} and \\citet{berg98} demonstrated that even a mild C-shock with shock velocities greater than 10 km s$^{-1}$ will elevate the gas temperature to permit the efficient conversion of oxygen into water and to vaporize a large fraction of the water ice on grains. Thus, most of the oxygen, not in CO, will be rapidly converted into water resulting in relative water abundances in excess of 10$^{-4}$. \\citet{berg98} suggested that these elevated levels of water will persist long after the gas has cooled and any enhancement of the water abundance may be present over the entire lifetime of these outflows.  The relative abundance of \\h2o\\ is thus a good tracer of the dynamical history of the outflowing molecular gas.  Our analysis of the SWAS data rules out the possibility that most of the outflow gas has passed through a shock in excess of 10 km s$^{-1}$.  The ISO and high velocity SWAS results for some outflows do indicate highly elevated water abundances in a small fraction of the outflowing gas mass that is consistent with moderate velocity  C-shocks. The current observations however do not have sufficient spatial and/or velocity resolution to determine exactly how these emissions might be related. It is suggestive that since the mass associated with the ISO and high velocity SWAS \\h2o\\ emission are very similar that these may arise from the same C-shock. For these outflows it would be extremely interesting to determine the spatial and kinematic relation between the ISO emission, the high velocity SWAS emission and well-established tracers of C-shocks, such as the 2.12 $\\mu$m H$_2$ emission seen in many of these outflows such as L1448-mm \\citep{davi95} and NGC1333-SVS13 \\citep{moli00}.   Future observations with {\\it Herschel} may help out in this regards.  Mechanisms for suppressing the \\h2o\\ abundance in the postshock gas were explored by \\citet{snel05}.  The only mechanism that is probably relevant for these outflows is the freeze-out of gas-phase water onto grain mantles. This was the primary mechanism in the model of \\citet{berg98} for the reduction of gas-phase water in post-shock gas.  The timescale for significant depletion of gas-phase water depends on the post shock density of the gas and the total grain cross-section.  In the model of \\citet{berg98}, for gas denser than 10$^6$ cm$^{-3}$, the timescale is shorter than 10$^4$ years, however at densities of 10$^4$ cm$^{-3}$ the time scale is longer than 3$\\times10^5$ years.  However since outflow lifetimes are extremely poorly understood, it is possible that the freeze-out of water is an important factor.  Based on the timescale for water freeze-out, \\citet{snel05} provided a simple expression for the hydrogen column density in the post-shock gas to the point where water will be frozen onto grain mantles. This calculation uses a slightly larger grain cross-section than the numerical model of  \\citet{berg98}.  Assuming a shock velocity of 10 km s$^{-1}$ and a shock compression factor of 10, the molecular hydrogen column would be of order 10$^{21}$ cm$^{-2}$.  If the relative abundance of ortho-\\h2o\\ in this column was as  $>$10$^{-4}$, as predicted by the C-shock models, the emission from this gas would greatly exceed what SWAS observed.  Unless a more efficient mechanism is found for depleting the gas-phase water in the post-shock gas, we believe that the \\h2o\\ abundance in outflows provides strong evidence that the bulk of the outflowing molecular gas was never subjected to shocks greater than 10 km s$^{-1}$ that would convert all of gas phase oxygen to water. This result places severe restrictions on the mechanism by which the molecular gas is accelerated in these outflows and favor mechanisms, such as turbulent entrainment, that can accelerate the gas gently.  However, some enhancement of the water abundance in the bulk of the outflowing gas over that measured in quiescent cloud gas may be necessary.  Slower shocks ($<$ 10 km s$^{-1}$) could heat the gas and dust to temperatures above the water-ice evaporation temperature, but below the temperature in which the rapid neutral-neutral reactions are activited, a suggestion made by \\citet{cern06} for the water abundance measured in Orion KL..  Finally we note that strong [OI] emission at a wavelength of 63 $\\mu$m  was observed by ISO in all of the outflows.  Many of the papers presenting ISO outflow results suggest that the origin of the [OI] emission is in J-shocks that dissociate the molecular gas.  Alternatively, the [OI] emission could arise from the same weak shocks that accelerate the bulk of the molecular gas. Future observations with {\\it Herschel}, which has better angular and spectral resolution, may help determine the relationship between the \\h2o\\ and [OI] emissions and other shock tracers in these outflows and provide a better understanding of the evolution of the \\h2o\\ abundance in these outflows.   We acknowledge the support by NASA contract NAS5-30702 that funded SWAS and NSF grant AST 01-00793 that funds the Five College Radio Astronomy Observatory. J.F acknowledges support by the NASA Goddard Center for Astrobiology through Cooporative Agreement NNG04G155A.  R.L.S., M.J.K., G.J.M., D.J.H, and E.A.B. acknowledge support of grant NNG06GB30G from NASA's LTSA prgrams.  D.A.N. acknowledge the support of grant NA65-13114 from NASA's LTSA program."
    },
    "0711/0711.0266_arXiv.txt": {
        "abstract": " ",
        "introduction": "The correlations between galaxy and central supermassive black hole properties found in the local Universe \\citep{1998AJ....115.2285M,2000ApJ...539L...9F} have led to notions that the growth of both must be intimately related. A fraction of the gas fuelling star formation is believed to reach the black hole, which will grow by accretion during an obscured quasar phase \\citep{1999MNRAS.308L..39F}, and is expected to develop winds or jets that eventually become powerful enough to expel the gas, preventing any further star formation \\citep{1998A&A...331L...1S}. This feedback process is needed to avoid forming more stars than are seen around local dormant black holes \\citep{2003ApJ...599...38B}, but few compelling examples of this phase have been reported to date.  Quasars are believed to be viewed as obscured or unobscured depending on the orientation of the line of sight with respect to the axis of symmetry and the dusty ``torus'' surrounding the accretion disk \\citep{1993ARA&A..31..473A}, yet the population of powerful obscured quasars had remained elusive.  Since most of the growth of supermassive black holes is obscured by dust, optical surveys miss this phase completely, and X-ray surveys miss a significant fraction due to associated gas column. Recently, mid-infrared surveys have been able to find obscured active galactic nuclei (AGN) in large numbers, but care must be taken to separate AGN from star-forming galaxies.  This can be achieved by making use of the particular mid-infrared colours of AGN \\citep[e.g.][]{2004ApJS..154..166L,2006ApJ...640..167A} or using a combination of mid-infrared and radio data \\citep[e.g.][]{2005Natur.436..666M,2005ApJ...634..169D}.  Indeed, from these surveys it seems that a large fraction of obscured AGN are undetected in X-rays \\citep[e.g.][]{2005ApJ...634..169D, 2006ApJ...640..167A,2007AJ....133..186L,2007MNRAS.379L...6M}, and that mid-infrared selected samples give a more complete picture than X-ray selected samples.  Some obscured quasars, however, do not show even the narrow emission lines expected if the central regions are obscured at certain angles by a torus \\citep[see][]{2005Natur.436..666M, 2006MNRAS.370.1479M}, some have evidence for dust in front of the narrow-line region \\citep{2007ApJ...663..204B}, and radio observations suggest that in some cases the radio jet is closely aligned with the observer's line of sight \\citep[][and Kl\\\"ockner et al., in prep.]  {2006MNRAS.373L..80M, 2007ApJ...667L..17S}, observations which are incompatible with the scenario of all obscuration being caused by the torus around the accretion disk.  The existence of objects which are not easily explained in the torus model has led to the suggestion that, in many cases, dust on scales of kpc is responsible for the obscuration.  Since young stars are born in dusty regions, it is possible that in these sources the host galaxy is undergoing an episode of star formation.  Mid-infrared spectroscopy gives a clean method of looking for star formation in the form of emission from complex molecules like polycyclic aromatic hydrocarbons \\citep [PAHs, e.g.][]{1985MNRAS.213..777A, 1991MNRAS.248..606R}. These molecules are excited by ultraviolet photons from young stars, but destroyed by the hard radiation of an AGN out to $\\sim$ 100 pc - 1 kpc scales \\citep[e.g.][]{1991ApJ...379..122V,2004A&A...414..123S}. Thus, detection of the PAH emission bands can only be attributed to star-formation and not to excitation by an AGN.  If obscured and unobscured quasars are identical types of objects viewed at a different orientation, then the properties of their host galaxies, such as the PAH luminosity, should be identical on average. In a detailed study of local AGN, \\citet {2000A&A...357..839C} found that indeed the PAH luminosities of Seyfert~1 and Seyfert~2 AGN are statistically indistinguishable, and that the observed differences between the Seyfert~1 and 2 spectra are due to the different contrast between the PAHs and the continuum, the latter being weaker for Seyfert~2s.    In this letter we present the results of mid-infrared spectroscopy of a sample of radio-intermediate obscured (type-2) quasars, which show either narrow lines or no emission lines in their optical (rest-frame ultraviolet) spectra. Throughout this paper we assume a $\\Lambda$CDM cosmology with $h = H_{0} / (100 ~ \\rm km ~ s^{-1} ~ Mpc^{-1}) = 0.7$; $\\Omega_{\\rm m} = 0.3$; $\\Omega_{\\Lambda} = 0.7$. ",
        "conclusions": "Two main conclusions can be extracted from this work.  Firstly, we have found that the mid-infrared spectra are continuum dominated, and 12 out of 18 spectra show a deep silicate absorption feature (with $\\tau_{9.7}$$\\sim$1-2).  We have confirmed spectroscopically the existence of sources with mid-infrared spectra characterstic of heavily obscured quasars, but that do not show rest-frame ultraviolet emission lines (observed in the optical). The number of these sources is similar to the number of obscured quasars that do show ultraviolet narrow emission lines. Using the new mid-infrared redshifts, we confirm that the population of radio-intermediate obscured (type-2) quasars outnumbers the unobscured population (type-1), and that 63$^{+14}_{-22}$\\% of $z\\sim2$ radio-intermediate quasars are obscured. This had already been suggested in \\citet{2005Natur.436..666M}, and in this paper we confirm the result spectroscopically: amongst radio-intermediate sources, the population of obscured quasars is responsible for the bulk of the accretion onto supermassive black holes.   Secondly, we have found that obscured quasars are sometimes hosted by galaxies undergoing vigorous star formation, with detectable PAHs comparable in strength to those found in SMGs. The accompanying SFR is expected to be very large ($\\grtsim$100 M$_{\\odot}$ yr$^{-1}$), so that it can form a significant fraction of an $L^{*}_{\\rm K}$ galaxy in $\\sim10^{7-8}$ yr.  The dust and gas fuelling this star formation are likely to be responsible for at least part of the obscuration in some of the sources, with the accreting black hole deeply embedded.  These properties are those expected from an obscured phase of black hole growth \\citep{1999MNRAS.308L..39F}, and could explain why in some obscured quasars the rest-frame ultraviolet lines are not detectable."
    },
    "0711/0711.4359_arXiv.txt": {
        "abstract": "The outflowing proper motions of fifteen knots in the dominant northwestern lobe of the high--excitation poly--polar planetary nebula NGC~6302 have been determined by comparing their positions relative to those of faint stars in an image taken at the San Pedro Martir Observatory in 2007 to those in a South African Astronomical Observatory archival plate obtained by Evans in 1956. The Hubble-type expansion of this lobe is now directly confirmed in a model--independent way from these measurements. Furthermore, an unambigous distance to NGC~6302 of 1.17 $\\pm$ 0.14 kpc is now determined. Also all the velocity vectors of the fifteen knots (and two others) point back to the central source. An eruptive event from within the central torus $\\approx$ 2200 yr previously must have created the high speed lobes of NGC~6302. ",
        "introduction": "The poly--polar planetary nebula NGC 6302 has always attracted interest (see Meaburn et al 2005 - hereafter Paper 1 - and papers therein)  because it has an extremely hot star in its central stellar system and because multiple bipolar lobes which emanate from this source reach outflow velocities of $\\geq$ 600 \\kms (Paper 1).  Furthermore, a surprisingly massive central CO emitting torus ($\\approx$ 2  \\msun\\ and  expanding at 8 \\kms) viewed nearly edge--on, obscures this central stellar system (Matsuura et al 2005; Peretto et al 2007). These authors estimate that this was ejected between 7500 yr and 2900 yr ago. The surface temperature of the hot central star has been estimated most recently from FIR ISO observations as 220,000 K by Wright et al. (2007). This is the lowest temperature yet estimated for the hydrogen deficient central star.   The high--speed outflow of the principal northwestern lobe is Hubble-type i.e. outflow velocity is proportional to distance from the source. This  has been conclusively illustrated by the kinematic--morphological modelling of the spatially resolved optical line profile observations in Paper 1, though it was also revealed in a more limited fashion in Meaburn \\& Walsh (1980). This behaviour alone favours a ballistic origin for the ejecta i.e. simply the gas more distant from the source has been ejected at the highest speed. This Hubble--type assertion would also predict that the velocity vectors of all of the outflowing material in a lobe point directly back to their source.  Estimations of the crucial distance to NGC~6302 (see Paper 1 for a summary) ranged from 0.15 to 2.4 kpc prior to its expansion proper motion (PM) measurement of 1.04 $\\pm$ 0.16 kpc reported in Paper 1. The latter value was based on the kinematic--morphological modelling of the northwestern lobe combined with a  measurement of the PM of a single compact  knot apparent between the 1956 image of Evans (1959) and a comparable image taken in 2001. The prediction that the dynamical age of the northwestern lobe is 1900 yr is also a natural consequence of this PM measurent. This would suggest that the lobe was ejected after the central torus had formed.  The original 1956 plate taken of NGC~6302 by Evans (1959) has now been obtained from the SAAO archive and scanned to permit comparison  with a new image of the northwestern lobe taken in 2007. The northwestern lobe is particularly suited to this analysis of its expansion PMs for its axis is close to the plane of the sky and its outflow velocities become high in its extremities and are now well--known (Paper 1). PMs of fifteen nebular knots distributed along the length of the northwestern lobe have now been measured. The confirmation of the Hubble--type nature of the outflow has now been established along this lobe by an independent method as a consequence, as has the ballistic nature of the ejection process. More refined values for the expansion--PM distance to NGC 6302 and the dynamical age of the northwestern lobe have also been derived. ",
        "conclusions": "A Hubble--type outflow has been shown to be occurring directly from PM measurements of 15 knots in the northwestern lobe of NGC~6302. This confirms the model--dependent prediction of the same behaviour from radial velocity measurements in Paper 1 (and see Corradi 2004 for similar behaviour in other PNe).  The velocity vectors of these 15 knots in NGC ~6302 point back to the central source.  The outflowing velocities of $\\geq$ 600 \\kms\\ at the furthest extremities of the northwestern lobe are confirmed.  The northwestern lobe of NGC~6302 must then have been created in an eruptive event $\\approx$ 2200 yr ago but probably after the creation of the central torus 2900 - 7500 yr ago as estimated by Matsuura et al (2005) \\& Peretto et al (2007).  The distance to NGC~6302 is now determined unambiguously as D = 1.17 $\\pm$ 0.14 kpc."
    },
    "0711/0711.1326_arXiv.txt": {
        "abstract": "We have developed a new method for post-Newtonian, high-precision integration of stellar systems containing a super-massive black hole (SMBH), splitting the forces on a particle between a dominant central force and perturbations. We used this method to perform fully collisional $N$-body simulations of inspiralling intermediate-mass black holes (IMBHs) in the centre of the Milky Way. We considered stellar cusps of different power-law indices and analysed the effects of IMBHs of different masses, all starting from circular orbits at an initial distance of 0.1\\,pc.  Our simulations show how IMBHs deplete the central cusp of stars, leaving behind a flatter cusp with slope consistent with what has recently been observed. If an additional IMBH spirals into such a flat cusp, it can take 50\\,Myr or longer to merge with the central SMBH, thus allowing for direct observation in the near future. The final merger of the two black holes involves gravitational wave radiation which may be observable with planned gravitational wave detectors.  Furthermore, our simulations reveal detailed properties of the hypervelocity stars (HVSs) created, and how generations of HVSs can be used to trace IMBHs in the Galactic Centre. We find that significant rotation of HVSs (which would be evidence for an IMBH) can only be expected among very fast stars ($v > 1000$\\,km\\,s$^{-1}$). Also, the probability of creating a hypervelocity binary star is found to be very small. ",
        "introduction": "In recent years, observations of the Galactic Centre revealed a number of interesting phenomena. Analysis of the motions of stars within 1.2 arcsec of the Galactic Centre, for example, showed that the Milky Way contains a super-massive black hole (SMBH) of mass $\\approx 3.5 \\times 10^6\\,\\msun$, surrounded by an extremely dense stellar cluster \\citep{sch03, ghe05}.  Since most massive galaxies are now believed to host such SMBHs \\citep{mag98, getal00}, due to its proximity ($\\approx$ 8\\,kpc), the centre of the Milky Way provides a unique environment to study the dynamics of SMBHs and the surrounding stellar systems.  While SMBHs have been directly observed as radio sources and compact massive objects, the existence of black holes of masses $10^2\\,\\msun < M_{\\rm BH} < 10^4\\,\\msun$, so-called intermediate-mass black holes (IMBHs), is still a matter of debate \\citep[e.g. ][]{hpa04, mc04, pbhmm04}.  IMBHs could form in galaxies through runaway collisions of stars in star clusters, giving rise to ultra-luminous X-ray sources \\citep{gfr04,pbhmm04,bhpm06}. They could later be brought into the centres of galaxies through dynamical friction of the star clusters \\citep{pbmmhe06}. Some of the dynamical processes at the Galactic Centre may be well explained by the presence of an IMBH, such as the existence and kinematics of young stars in the central 0.1\\,pc \\citep{hm03}. Mergers of IMBHs with the central SMBH would be an important source of gravitational waves, detectable with the next generation of gravitational wave detectors, like e.g. the \\emph{Laser Interferometer Space Antenna} \\citep[\\emph{LISA},][]{lisa}.   During the inspiral process of an IMBH, close encounters with stars from the central cluster eject some of them with velocities large enough to leave the gravitational potential of the Milky Way, so-called hypervelocity stars (HVSs). \\citet{h88} was the first to show that the ejection of HVSs is a natural consequence of galaxies hosting SMBHs. Recently, several HVSs have been discovered in the Galactic halo \\citep{b05, hi05, b06a, b06b, b07a, b07b}. Except for one star which might have been ejected from the Large Magellanic Cloud \\citep{edel05}, the travel times of all HVSs are short enough that the stars could have been ejected from the Galactic Centre within their lifetimes, confirming Hills' predictions. \\citet{bgp06} and \\citet{l06} have shown that HVSs may be ejected almost isotropically by short bursts of inspiralling IMBHs, explaining the near-isotropic distribution of the so far detected HVSs \\citep{b07a}. However, \\citet{shm07} found that the relatively small medium velocity of the (admittedly limited) sample of HVSs detected so far favours other ejection mechanisms like the tidal breakup of stellar binaries \\citep{h88, yt03, pha06} or the scattering of stars by stellar-mass black holes \\citep{mg00, oll06}.  In this work, we present new results on the inspiral of IMBHs, using a novel method for studying the dynamics of Galactic Centre-like systems where stars orbit a central SMBH on weakly perturbed Keplerian orbits. In Section \\ref{sec:bhint}, we describe the details of our integration method. In Section \\ref{sec:models}, we describe our set of N-body runs with IMBHs in the Galactic Centre. Our results are given in Section \\ref{sec:results} and discussed in Section \\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  We have developed a novel post-Newtonian, high-precision integrator \\textsc{bhint} for stellar systems containing a SMBH. The algorithm makes use of the fact that the Keplerian orbits in such a potential can be calculated directly, and are only weakly perturbed. We used this method to follow the inspiral of IMBHs into the Galactic Centre as a result of dynamical friction. We find that on their way to the centre, IMBHs deplete the central cusp of stars, leaving behind a shallower density profile compatible with the one recently observed in the centre of the Milky Way. Due to the decreased stellar density, IMBHs may get stuck at a distance of a milliparsec from the SMBH for a few 10\\,Myr. If the recently detected HVSs have been ejected by one or several IMBHs, it is likely that an IMBH still resides close to Sgr A*.  Such an IMBH may be observable by detecting the proper motion of the radio source Sgr A* with upcoming high-resolution telescopes. Our results suggest that an IMBH in the Galactic Centre may spend several Myr in the regime where it is detectable with an astrometric resolution of 0.1\\,mas, yielding good chance of direct observation in the near future.  Once an IMBH comes close enough to the centre, gravitational wave emission will take over and lead to the merger with the central SMBH. These mergers are among the most powerful sources of gravitational waves for the planned \\emph{LISA} mission \\citep{t87} and may well be detected if the IMBH inspiral rate is of the order of 1 per 25\\,Myr.   Our results show that there should be a population of HVSs with very high velocities ($v > 1000$\\,km\\,s$^{-1}$) which has not been observed so far, consistent with the findings of \\citet{shm07}. However, as the sample size of HVSs observed is relatively small, IMBHs cannot be excluded as ejection mechanism. Moreover, HVSs ejected by IMBHs are hard to distinguish from those ejected by other mechanisms such as binary disruptions as we do not expect significant rotation in either case, but possible observations of significant rotation of HVSs in the high velocity tail may favour the IMBH model."
    },
    "0711/0711.1110_arXiv.txt": {
        "abstract": "{B-type supergiants represent an important phase in the evolution of massive stars. Reliable estimates of their stellar and wind parameters, however, are scarce, especially at mid and late spectral subtypes. } {We apply the NLTE atmosphere code FASTWIND to perform a spectroscopic study of a small sample of Galactic B-supergiants from B0 to B9. By means of the resulting data and incorporating additional datasets from alternative studies, we investigate the properties of OB-supergiants and compare our findings with theoretical predictions.} {Stellar and wind parameters of our sample stars are determined by line profile fitting, based on synthetic profiles, a Fourier technique to investigate the individual contributions of stellar rotation and ``macro-turbulence'' and an adequate approach to determine the Si abundances in parallel with micro-turbulent velocities.} {Due to the combined effects of line- and wind-blanketing, the temperature scale of Galactic B-supergiants needs to be revised downward, by 10 to 20\\%, the latter value being appropriate for stronger winds.  Compared to theoretical predictions, the wind properties of OB-supergiants indicate a number of discrepancies. In fair accordance with recent results, our sample indicates a gradual decrease in \\vinf over the bi-stability region, where the limits of this region are located at lower \\Teff than those predicted. Introducing a distance-independent quantity $Q'$ related to wind-strength, we show that this quantity is a well defined, monotonically increasing function of \\Teff $outside$ this region. $Inside$ and from hot to cool, \\Mdot changes by a factor (in between 0.4 and 2.5) which is much smaller than the predicted factor of 5.} {The decrease in \\vinf over the bi-stability region is $not$ over-compensated by an increase of \\Mdot, as frequently argued,  provided that wind-clumping properties on both sides of this region do not differ substantially. } ",
        "introduction": "Hot massive stars are key objects for  studying and understanding many exciting phenomena in the Universe such as re-ionisation and $\\gamma$-ray bursters. Due to their powerful stellar winds hot massive stars are important contributors to the chemical and dynamical evolution of galaxies, and in the distant Universe they dominate the integrated UV radiation in young galaxies.  While the number of Galactic O and early B stars with reliably determined stellar and wind parameters has progressively increased during the last few years (e.g., \\citealt{Herrero02, repo, GB04, bouret05, martins05, crowther06}), mid and late B supergiants (SGs) are currently under-represented in the sample of stars investigated so far. Given the fact that B-SGs represent an important phase in the evolutionary sequence of massive stars, any study aiming to increase our knowledge of these stars would be highly valuable, since it would allow several important issues to be addressed (see below).  Compared to O-type stars the B-SG spectra are more complicated due to a larger variety of atomic species being visible, the most important among which is Silicon, the main temperature indicator in the optical domain. Thus, the reproduction of these spectra by methods of {\\it quantitative spectroscopy} is a real challenge for state-of-the art model atmosphere codes, since it requires a good knowledge of the physics of these objects, combined with accurate atomic data. In turn, any discrepancy that might appear between computed and observed spectral features would help to validate the physical assumptions underlying the model calculations as well as the accuracy of the adopted atomic models and data.  Numerical simulations of the non-linear evolution of the line-driven flow instability (for a review, see \\citealt{O94}), with various degrees of approximation concerning the stabilising diffuse, scattered radiation field (\\citealt{op96, OP99}) as well as more recent simulations concentrating on the outer wind regions \\citep{ro02, ro05}, predict that hot star winds are not smooth but structured, with clumping properties depending on the distance to the stellar surface. However, recent observational studies of clumping in O-SGs have revealed inconsistencies both between results originating from different wind diagnostics, such as UV resonance lines, \\Ha and the IR-/radio-excess \\citep{FMP06, puls06}, and between theoretical predictions and observed constraints on the radial stratification of the clumping factor \\citep{bouret05, puls06}. In addition, there are observational results which imply that clumping might depend on wind density.  Because of their dense winds, B-SGs might provide additional clues to clarify these points.  Due to their high luminosities, BA-SGs can be resolved and observed, both photometrically and spectroscopically, even in rather distant, extragalactic stellar systems (e.g., \\citealt{Kud99, Bresolin02, Urb03, Bianchi06}). This fact makes them potential standard candles, allowing us to determine distances by means of purely spectroscopic tools using the wind-momentum luminosity relationship (WLR, \\citealt{klp95}). Even though certain discrepancies between predicted and observed wind momenta of early B0 to B3 subtypes have been revealed \\citep{crowther06}, relevant information about later subtypes is still missing.  During the last years, the quantitative analyses of spectra in the far-UV/UV and optical domains (e.g., \\citealt{Herrero02,BG02,Crowther02,bouret03,repo,massey04,heap06}) have unambiguously shown that the inclusion of line-blocking and blanketing and wind effects (if present) significantly modifies the temperature scale of O-stars (for a recent calibration at {\\it solar} metallicity, see \\citealt{Markova04, martins05}). Regarding B-SGs, particularly of later subtype, this issue has not been addressed so far, mostly due to lacking \\Teff estimates.  The main goal of this study is to test and to apply the potential of our NLTE atmosphere code FASTWIND \\citep{Puls05} to provide reliable estimates of stellar and wind parameters of SGs with temperatures ranging from 30 to 11 kK. By means of these data and incorporating additional datasets from alternative studies, we will try to resolve the questions outlined above.  In Sects.~\\ref{obs} and \\ref{sample}, we describe the stellar sample and the underlying observational material used in this study. In Sect.~4 we outline our procedure to determine the basic parameters of our targets, highlighting some problems faced during this process. In Sect.~5 the effects of line blocking/blanketing on the temperature scale of B-SGs at solar and SMC metallicities will be addressed, and in Sect.~6 we investigate the wind properties for Galactic B-SGs (augmented by O-and A-SG data), by comparison with theoretical predictions. Particular emphasis will be given to the behaviour of the mass-loss rate over the so-called bi-stability jump. Sect.~7 gives our summary and implications for future work. \\begin{table} \\caption{Galactic B-SGs studied in this work, together with adopted photometric data. For multiple entries, see text.} \\label{log_phot} \\tabcolsep0.8mm \\begin{tabular}{llllllll} \\hline \\hline \\\\ \\multicolumn{1}{l}{Object} &\\multicolumn{1}{l}{spectral} &\\multicolumn{1}{c}{member-} &\\multicolumn{1}{c}{d} &\\multicolumn{1}{c}{$V$} &\\multicolumn{1}{l}{$B-V$} &\\multicolumn{1}{c}{$(B-V)_{\\rm 0}$} &\\multicolumn{1}{c}{$M_{\\rm V}$}\\\\ \\multicolumn{1}{l}{(HD\\#)} &\\multicolumn{1}{l}{type} &\\multicolumn{1}{c}{ship} \\\\ \\hline 185\\,859  &B0.5 Ia   &          &             &      &       &       &-7.0*\\\\ 190\\,603  &B1.5 Ia+  &        &1.57$^{c}$   &5.62  &0.760  &-0.16  &-8.21\\\\ &          &        &   &5.62         &0.54$\\pm$0.02& &-7.53\\\\  206\\,165  &B2 Ib     &Cep OB2   &0.83$^{a}$   &4.76  &0.246  &-0.19  &-6.19\\\\%SIMBAD0.219 198\\,478  &B2.5 Ia   &Cyg OB7   &0.83$^{a}$   &4.81  &0.571  &-0.12  &-6.93\\\\ &           &         &             &4.84  &0.40$\\pm$0.01 &&-6.37\\\\ 191\\,243  &B5 Ib     &Cyg OB3   &2.29$^{a}$   &6.12  &0.117  &-0.12  &-6.41\\\\%SIMBAD 0.103 &          &          &1.73$^{b}$   &      &       &       &-5.80\\\\ 199\\,478  &B8 Iae    &NGC 6991  &1.84$^{d}$   &5.68  &0.408  &-0.03  &-7.00\\\\%SIMBAD 0.367 212\\,593  &B9 Iab    &          &             &      &       &       &-6.5*\\\\ 202\\,850  &B9 Iab    &Cyg OB4   &1.00$^{a}$   &4.22  &0.098  &-0.03  &-6.18\\\\ \\hline \\end{tabular} \\newline $^{a}$ \\citet{Humphreys78}; $^{b}$ \\citet{GS}; $^{c}$ \\citet{BC}; $^{d}$ \\citet{DH}\\\\ $^{*}$ from calibrations \\citep{HM}\\\\ \\end{table} ",
        "conclusions": "\\label{summary}  In this study, we have presented a detailed investigation of the optical spectra of a small sample of Galactic B supergiants, from B0 to B9. Stellar and wind parameters have been obtained by employing the NLTE, unified model atmosphere code FASTWIND \\citep{Puls05}, assuming unclumped winds. The major findings of our analysis can be summarised as follows.  \\smallskip \\noindent 1. We confirm recent results \\citep{Ryans,dft06, simon} of the presence of a (symmetric) line-broadening mechanism in addition to stellar rotation, denoted as ``macro-turbulence''. The derived values of \\vmac are highly supersonic, decreasing from $\\approx$ 60~\\kms at B0 to $\\approx$ 30~\\kms at B9.  \\smallskip \\noindent 2.  We determined the Si abundances of our sample stars in parallel with their corresponding micro-turbulent velocities.\\\\ (i) For all but one star, the estimated Si abundances {\\bf were} consistent with the corresponding solar value (within $\\pm$0.1 dex), in agreement with similar studies \\citep{GL,Roll,Urb,Przb06}. For HD~202\\,850, on the other hand, an overabundance of about 0.4~dex has been derived, suggesting that this late-B supergiant might be a silicon star.\\\\ (ii) The micro-turbulent velocities tend to decrease towards later B subtypes, from 15 to 20~\\kms at B0 (similar to the situation in O-supergiants) to 7~\\kms at B9, which is also a typical value for A-SGs. \\\\ (iii) The effect of micro-turbulence on the derived effective temperature was negligible as long as Si lines from the two major ions are used to determine it.  \\smallskip \\noindent 3. Based on our \\Teff estimates and incorporating data from similar investigations \\citep{crowther06,Urb,Przb06,Lefever}, we confirm previous results (e.g., \\citealt{crowther06}) on a 10\\% downwards revision of the effective temperature scale of early B-SGs, required after incorporating the effects of line blocking/blanketing. Furthermore, we suggest a similar correction for mid and late subtypes. When strong winds are present, this reduction can become  a factor of two larger, similar to the situation encountered in O-SGs \\citep{Crowther02}.  \\smallskip \\noindent 4. To our surprise, a comparison with data from similar SMC objects \\citep{Trundle04, Trundle05} did not reveal any systematic difference between the two temperature scales. This result is interpreted as an indication that  the re-classification scheme as developed by \\citet{lennon97} to account for lower metal line strengths in SMC B-SGs also removes the effects of different degrees of line blanketing.  \\smallskip \\noindent 5. Investigating the wind properties of a statistically significant sample of supergiants with \\Teff between 10 and 45 kK, we identified a number of discrepancies between theoretical predictions \\citep{Vink00} and observations. In fair accordance with recent results \\citep{Evans2, crowther06}, our sample indicates a gradual decrease in \\vinf in the bi-stability (``transition'') region, which is located at lower temperatures than predicted: 18 to 23~kK (present study) against 22.5 to 27 kK.   By means of a $newly$ defined, distance independent quantity, $Q'= \\Mdote/\\Rstare^{1.5} \\, \\geff/\\vinfe)$ we have investigated the behaviour of \\Mdot as a function of \\Teff.  Whereas inside the transition zone a large scatter is present (coupled with a potential {\\it local} maximum in wind efficiency around 21~kK), $Q'$ remains a well defined function with low scatter in the hot and cool temperature region outside the transition zone. Combining the behaviour of $Q'$ and the modified wind-momentum rate, the change in \\Mdot over the bi-stability jump (from hot to cool) could be constrained to lie within the factors 0.4 to 2.5, to be conservative. Thus, \\Mdot either decreases in parallel with \\vinf/\\vesc (more probable), or, at most, the decrease in \\vinf is just balanced by a corresponding increase in \\Mdot (less probable). This finding contradicts the predictions by \\citet{Vink00} that the decrease in \\vinf should be {\\it over-compensated} by an increase in \\Mdot, i.e., that the wind-momenta should increase over the jump. Considering potential clumping effects, we have argued that such effects will not change our basic result, unless hotter objects turn out to be substantially more strongly clumped than cooler ones. In any case, at least in the low temperature region present theoretical predictions for \\Mdot are too large!  \\medskip \\noindent This finding is somewhat similar to the recent ``weak-wind problem'' for late O-dwarfs\\footnote{A detailed UV-analysis by \\citet{martins04} showed the mass-loss rates of young late-O dwarfs in N81 (SMC) to be significantly smaller (factors 10 to 100) than theory predicts (see also \\citealt{bouret03}). In the Galaxy, the same dilemma applies to the O9V 10~Lac(\\citealt{Herrero02}) and maybe also for $\\tau$ Sco (B0.2V), which show very low mass-loss rates.}, though probably to a lesser extent. Thus, it might be that our understanding of radiation driven winds is not as complete as thought only a few years ago. Thus, it is of extreme importance to continue the effort of constructing sophisticated wind models, including the aforementioned effects (wind-clumping, diffuse radiation field, micro-turbulence), both in terms of stationary and time-dependent simulations. With respect to the objects of the present study, a re-analysis of the ``peculiar'' mid-type B-supergiants from the KPL99 sample is urgently required as well. Finally, let us (once more) point to the unresolved problem of macro-turbulence, which implies the presence of rather deep-seated, statistically distributed and highly supersonic velocity fields. How can we explain such an effect within our present-day atmospheric models of hot, massive stars?"
    },
    "0711/0711.4503_arXiv.txt": {
        "abstract": "We present $B_J$, $R_{59F}$ and $I_N$ photometry, and optical and near-infrared spectroscopy, of a new ultracool white dwarf (UCWD) discovered in the SuperCOSMOS Sky Survey. The spectrum of SSSJ1556-0806 shows strong flux suppression due to the presence of collisionally induced absorption by molecular hydrogen (H2CIA), a feature characteristic of the cool, high density environments found in the atmospheres of ultracool white dwarfs. SSSJ1556-0806 therefore joins a list of $<$10 ultracool white dwarfs displaying extreme flux suppression. Synthetic model fitting suggests an effective temperature $<$3000K, which if true would make this one of the coolest white dwarfs currently known. We also exploit the similarity between the SEDs of SSSJ1556-0806 and the well-studied UCWD LHS 3250 to aid in the determination of the atmospheric parameters in a regime where models consistently fail to reproduce observations. SSSJ1556-0806 is relatively bright ($R \\sim 17.8$), making it particularly amenable to follow up observations to obtain trigonometric parallax and IR photometry. ",
        "introduction": "The Galactic disc white dwarf population is of importance in obtaining disc age estimates through the use of the luminosity function. Now devoid of the nuclear fuel that sustains these stars during their main sequence lifetime, white dwarfs shine at the expense of their residual thermal energy, slowly radiating away this fossil heat until they disappear from our telescopes as cold, compact objects. This is a very long process, overlapping the age of the galactic disc at the extreme. Age determination methods independent of WD cooling rates have been used to estimate the age of the galactic disc at around 8 Gyr (eg. \\citealt{jimenez1998}). \\cite{salaris2000} calculate a magnitude of M$_{\\textrm{bol}}\\sim15.5$ for 0.61M$_{\\odot}$ H-rich white dwarfs that have been cooling for 8 Gyr, based on an interpolation of results in their Table 1. Thus, there is predicted to be a shortfall of white dwarfs at fainter magnitudes, a consequence of the finite age of the galactic disc population - white dwarfs have simply not had the time to cool further. The local disk white dwarf luminosity function is thus expected to show a marked downturn at luminosities corresponding to cooling ages older than that of the disk.  Ultracool white dwarfs with effective temperatures below $\\sim$3700K fall on the cusp of this downturn in spatial density. They represent some of the oldest stars in the sky, fossils left over from periods of star formation in the very early history of the galaxy. Obtaining accurate age estimates using this method thus depends on finding sufficient numbers of these extremely faint stars to constrain the observational determination of the luminosity function at the faint end. Large white dwarf surveys have in the past been hindered by the technical difficulties in finding such a population, with the limiting magnitude of the survey determining the cut off in the luminosity function, rather than the intrinsic lack of white dwarfs at cooling ages older than that of the galaxy.  Modern survey class telescopes are now capable of observing the faintest white dwarfs in the disc, and \\cite{liebert1988} were the first to detect the downturn in the luminosity function, at around M$_{\\textrm{bol}}\\sim15.76$. \\cite{winget1987} used an early version of the observational results to place a constraint on the age of the galactic disc of $9.3 \\pm 2.0$ Gyr. Currently, the largest observational determination of the white dwarf luminosity function is that of \\citet{harris2006} using SDSS photometry combined with USNO-B data for improved proper motions. Out of a sample of over six thousand white dwarfs, only four objects fell on the downturn in the luminosity function, preventing any quantitative conclusions as to the age of the population. Evidently, the reliable identification of faint ultracool white dwarfs, even in small numbers, will enable tighter constraints to be placed on the spatial density of the oldest white dwarfs, and the statistical uncertainty in age estimates to be reduced. ",
        "conclusions": "We have presented observations of a new ultracool white dwarf found in the SuperCOSMOS Sky Survey, and made an initial effort to understand the unusual SED of this object using both atmosphere models for cool white dwarfs and comparison to the well-studied UCWD LHS 3250. Without better photometry, models are poorly constrained and can do no more than give a rough interpretation of the object. The current inconsistencies between the theoretical and observed spectral energy distributions of UCWDs rules out any precise determination of the atmospheric parameters. The most likely interpretation of this object is as a thick disk white dwarf, with a helium dominated atmosphere and an effective temperature $<$4000K.   Infrared photometry is essential for a good fit to theoretical models. Trigonometric parallax observations can also provide constraints by measuring both the absolute magnitude and, indirectly, mass. We currently have time allocated on the robotic Liverpool Telescope for a program to obtain parallax measurements of UCWDs. SSSJ1556-0806 has been added to our queue-scheduled list and has already had one epoch of observation carried out at high parallax factor."
    },
    "0711/0711.2505_arXiv.txt": {
        "abstract": "We describe a new radial velocity survey of T Tauri stars and present the first results.  Our search is motivated by an interest in detecting massive young planets, as well as investigating the origin of the brown dwarf desert.  As part of this survey, we discovered large-amplitude, periodic, radial velocity variations in the spectrum of the weak line T Tauri star LkCa 19.  Using line bisector analysis and a new simulation of the effect of star spots on the photometric and radial velocity variability of T Tauri stars, we show that our measured radial velocities for LkCa19 are fully consistent with variations caused by the presence of large star spots on this rapidly rotating young star.  These results illustrate the level of activity-induced radial velocity noise associated with at least some very young stars.  This activity-induced noise will set lower limits on the mass of a companion detectable around LkCa 19, and similarly active young stars. ",
        "introduction": "The first successful detection produced by radial velocity surveys of a planet orbiting a star occured in 1995 when a near Jupiter-mass planet was found around the main sequence star 51 Pegasus \\citep{1995Natur.378..355M}.  Subsequent observations verified the existence of the companion and refined its orbital parameters \\citep{1995AAS...187.7004M,1997ApJ...481..926M}.  By 2007, over 200 planetary systems have been discovered via the radial velocity method and the statistics of various properties of the planets have been catalogued \\citep{2006ApJ...646..505B}  Most radial velocity surveys have been conducted on main sequence stars typically older than 1 Gyr.  The use of the radial velocity method on younger stars is challenging because of increased activity, specifically in the form of signals caused by star spots.  Spots on the surface of the star rotate in and out of view of the observer, potentially creating spurious radial velocity signatures.  \\citet{2004AJ....127.3579P} looked for planetary companions to stars in the Hyades cluster ($\\sim$790 Myr).  They found activity induced jitter consistent with the empirical relationship between activity and radial velocity described in \\citet{1997ApJ...485..319S}.  \\citet{2004AJ....127.3579P} also found one binary ($K = 1152 $ m s$^{-1}$) and two candidate binaries with linear velocity trends, but no planetary companions.  More recently, \\citet{2006PASP..118..706P} examined 61 stars in the $\\beta$ Pic moving group ($\\sim$12 Myr) and the Ursa Majoris association ($\\sim$300 Myr) and found no Jupiter mass planets in short period orbits.  The authors found these results unsurprising, as planetary companions are rare around main sequence stars with the small orbital separations ($<$ 0.1 AU) that these surveys were able to probe.  They measured maximum radial velocity amplitudes of 30-100 m s$^{-1}$ for targets in Ursa Majoris and 250-600 m s$^{-1}$ for targets in $\\beta$ Pic.  These searches were motivated, as is ours, by an interest in understanding the formation of both stellar and planetary systems.  By probing the youngest stars for planetary companions, we can learn at what stage planets form.  Since the time scale for the core accretion model \\citep{1996Icar..124...62P} is longer than formation due to disk instability \\citep{1997Sci...276.1836B}, a detection of a young planet around a T Tauri star should also provide clues to the formation mechanism itself.  We can also explore the level of stellar activity in young stars and how this can effect our ability to measure radial velocity variations, and therefore detect planets.  Recently, \\citet{2007ApJ...660L.145S} have reported the detection of a planet around the 100 Myr star HD 70573.  An additional motivation to search for companions around young stars is to understand the origin of the so-called ``brown dwarf desert.\"  As more planetary companions were discovered in the last decade, and the completeness of the surveys increased, \\citet{2000PASP..112..137M} noted the dearth of substellar objects with masses greater than 20 M$_{\\rm Jup}$ in close orbits ($\\lesssim$ 3 AU).  This ``brown dwarf desert\" was further quantified by \\citet{2006ApJ...640.1051G} who analyzed an unbiased sample of nearby sun-like stars to measure the rate of stellar, brown dwarf, and planetary companions.  They found that $\\sim$16\\% of their sample stars had close (P $<$ 5 yr) companions with masses greater than that of Jupiter.  Of these, 11\\% $\\pm$ 3\\% were stellar, $<$ 1\\% were brown dwarfs, and 5\\%$\\pm$ 2\\% were giant planets, where a brown dwarf is defined as an object in the mass range between 13 M$_{\\rm Jup}$ to 80 M$_{\\rm Jup}$.  \\citet{2002MNRAS.330L..11A} suggested that the kinematic viscosity of a sufficiently massive circumstellar disk could cause brown dwarfs in close orbits, interacting dynamically with the disks, to catastrophically migrate towards the host star and be lost to mergers.  They predict that young stars should have an order of magnitude more brown dwarfs in close orbits than main sequence stars, as the time scale for catastrophic migration is $\\sim$ 1 Myr.  A search for substellar companions around T Tauri stars could confirm or refute this scenario, improving our understanding of the brown dwarf desert in older stars.  In this paper, we summarize our spectroscopic and photometric observational techniques for our sample of T Tauri stars and report a detection of periodic radial velocity variations in the weak line T Tauri star LkCa 19, for which we have a large number of observations.  LkCa 19 has a  $v \\sin i$ = 18.6 km s$^{-1}$ \\citep{1987AJ.....93..907H} and spectral type K0 \\citep{1988cels.book.....H}.  LkCa 19, with its short rotation period of 2.24 days \\citep{1993A&A...272..176B}, provides an excellent example of the effects of star spots on radial velocity measurements.  Using new simulations, we show that the radial velocity properties of LkCa 19 are well modeled by the presence of one large spot. ",
        "conclusions": "Analysis of early results from our radial velocity survey of young stars reveals large radial velocity variations in our target stars compared to radial velocity standards.  LkCa 19 shows a radial velocity signal with a significant periodic variation.  Initial analysis of the radial velocities of LkCa 19 led to twice-nightly observations in our February 2006 run to obtain better phase coverage.  Bisector analysis, however, indicates that the radial velocity variations in LkCa 19 are very likely induced by star spots, as does the similarity between LkCa19's 2.177$\\pm$0.0264 day radial velocity period and its known rotation period of 2.24 days \\citep{1993A&A...272..176B}.  While it is possible that a close planet could experience tidal locking \\citep{1997ApJ...481..926M} and thus have an orbital period identical to the rotation period of the star, it is more likely we are seeing radial velocity modulation arising in star spots.  A surprising feature of LkCa 19 is the phase coherence of the spot-induced radial velocity modulation.  The velocities remain in phase at the period of 2.177 days over a time span of approximately 13 months.  This result is similar to the false planet detection reported by \\citet{2001A&A...379..279Q}, also attributed to star spots, which showed reasonable phase-coherence for a time span of roughly 2 years.  We have also presented a new simulation of the effects of a rotating star with spots on line profiles and have measured the radial velocities and photometric variations from spectra produced by the simulation.  The models can reproduce the photometric and velocity amplitudes of LkCa 19.  All simulations show a strong and significant bisector correlation, summarized in Table 3. Further, each spot simulation shows a 0.25 lag in phase between photometry and radial velocity, regardless of spot configuration, which is expected for cool spots.  The spot geometry with the most regular or planet-like radial velocity curve is that of a single spot near the pole of a star with a low inclination angle.  A highly spotted star with the spots set randomly on the stellar surface, produces much more irregular radial velocity and photometric light curves.  If the magnitude of the photospheric activity induced radial velocity variations of LkCa 19 is typical of young stars, then our ability to find companions to these objects will be limited to the extremely massive and short-period sub-stellar companions.  LkCa19 has an exceptionally short rotation period and relatively high $v \\sin i$  value, which offers some hope that its spot-induced radial velocity variations are not necessarily typical for T Tauri stars."
    },
    "0711/0711.0500_arXiv.txt": {
        "abstract": "{In the late 1990s, the Optical Gravitational Lensing Experiment (OGLE) team conducted the second phase of their long-term monitoring programme, OGLE-II, which since has been superseded by OGLE-III. All the monitoring data of this second phase, which was primarily aimed at finding microlensing events, have recently been made public.} {Fields in the OGLE-II survey have typically been observed with a cadence of once per night, over a period of a few months per year. We investigated whether these radically differently sampled data can also be used to search for transiting extrasolar planets, in particular in the light of future projects such as PanSTARRS and SkyMapper, which will monitor large fields, but mostly not at a cadence typical for transit studies.} {We selected data for $\\sim$15700 stars with 13.0$<$I$<$16.0 in three OGLE-II fields towards the galactic disc in the constellation Carina, each with 500-600 epochs of I-band photometry. These light curves were first detrended using Sys-Rem, after which they were searched for low amplitude transits using the Box Least Squares algorithm.} {The detrending algorithm significantly decreased the scatter in the light curves, from an average of 0.5\\% down to 0.2$-$0.3\\% for stars with $I$$<$15. Several dozens of eclipsing binaries and low amplitude transits were found, of which 13 candidates exhibit transits with such depth and duration  that they are possibly caused by an object with a radius less than twice that of Jupiter. Eleven out of these thirteen candidates show significant ellipsoidal light variations and are unlikely to host a transiting extrasolar planet. However, OGLE2-TR-L9 (CAR\\_SC2\\_75679), is an excellent planet candidate comparable to the known OGLE-III transiting planets, and deserves further follow-up observations. } {} ",
        "introduction": "An extrasolar planet that transits its host star is of great scientific value. Analysis of its transit light curve can reveal the planet/star size ratio, and implicitly the radius of the planet after the star has been spectrally typed. Since the occurrence of a transit implies that the inclination of the planetary orbit is near 90$^\\circ$, radial velocity measurements will reveal the true mass of the planet, in contrast to only a lower limit in mass for non-transiting planets. This will yield not only the mean density of the planet, but can also give insight into its internal structure (eg. Sato et al. 2005). In addition, many interesting follow-up opportunities are within reach, such as atmospheric transmission spectroscopy, secondary eclipse and Rossiter-McLaughlin effect observations, and transit timing, providing information on the planet's atmosphere, orientation of its orbit, and on the presence of other planets in the system (eg. Charbonneau et al. 2002; Charbonneau et al. 2005; Deming et al. 2005; Winn et al. 2007; Snellen 2004; Steffen et al. 2007).  Several successful transit surveys are now in place. OGLE-III (the third phase of the Optical Gravitational Lensing Experiment; Udalski et al 2002a,b,c;2004) is the longest running survey and has yielded so far 5 transiting planets. While more than 150 planet candidates have been reported by this team, most of them turned out to be blended or grazing eclipsing binaries and/or eclipsing M-dwarfs (Bouchy et al. 2005; Pont et al. 2005). It shows that radial velocity follow-up is crucial for the identification of the genuine transiting extrasolar planets. Now, several other dedicated ground-based transit surveys, such as XO (McCullough et al. 2005), SuperWASP (Collier Cameron et al. 2007), TrES (O'Donovan et al. 2007), and HAT (Bakos et al. 2007), have found their first planets.  Therefore, the number of confirmed transiting planets has recently increased significantly (20 in total), with more than half of them discovered in the last 1.5 years. However, the number of planets found to transit the stars being monitored still lags behind the expectations from radial velocity surveys (Gaudi 2006). In addition, almost all of the planets found so far have radii larger than that of Jupiter. This indicates that a significant fraction of the transiting hot Jupiters are being missed. The most likely cause has been identified as correlated noise in the light curve data sets (e.g. Pont et al. 2006; Aigrain \\& Pont 2007). Variations in atmospheric (seeing, airmass) or instrumental/telescope properties result in photometric features on the time scale of a transit, causing the sensitivity of a survey to decrease significantly. The recently launched CoRoT satellite and future Kepler mission will likely resolve these issues.   In this paper we present a search for transiting planet candidates in the 1997$-$2000 OGLE-II survey (Udalski et al. 1997), a precursor of OGLE-III, which was not designed for transit searches. While it is in itself highly desirable to find more transiting exoplanets, it is also valuable in the light of the radically different  time sampling of OGLE-II compared to dedicated transit surveys (typically once per night compared to once per 15 minutes), which should significantly reduce correlated noise problems. This is particularly interesting in the light of future projects such as PanSTARRS (e.g. Afonso \\& Hennig 2007) and SkyMapper (eg. Bayliss \\& Sackett 2007); groundbased telescopes with huge field of view that will monitor millions of stars, but mostly at a cadence very different from transit surveys (although they will probably also spend some limited time performing transit surveys). In the next section the OGLE-II survey and our data analysis are presented. In section 3 the results are shown and discussed, and our conclusions are given in section 4. The folded light curves are presented in Figure \\ref{curves} in the appendix.  \\begin{figure} \\psfig{figure=mag_noiseLOW.ps,width=0.5\\textwidth} \\caption{\\label{dispersion} The dispersion in the stellar light curves as function of the median magnitude before (light grey points) and after (dark grey points) the use of Sys-Rem. The black points indicate the median dispersion per 0.1 magnitude bin before and after this detrending correction. It shows that the algorithm is most efficient for the brightest stars. For the bright stars (I$<$15) the photometric uncertainties are decreased by a factor of $\\sim$2. } \\end{figure} ",
        "conclusions": "In this paper we searched the online database of the second phase  of the OGLE project (OGLE-II; 1997-2000) for transiting extrasolar planets, concentrating on the galactic disc fields in the Carina constellation. The main goals were, next to finding possibly more transiting exoplanets, to investigate whether such a data set with a very different cadence compared to typical transit surveys, is actually sensitive to low amplitude transits.  We found that detrending the data using the Sys-Rem algorithm significantly reduced the scatter in the light curves by about a factor two, down to 0.2\\%. About 27000 stars were searched for transits with periods between 0.5 and 5 days using the BLS algorithm, and several dozens of eclipsing binaries were found. Thirteen stars were found to exhibit low amplitude transits, of which most show significant ellipsoidal light variability indicative of tidal interactions and stellar mass secondaries. One star, OGLE2-TR-L9 (P=2.4855 days), shows a transit shape consistent with a Jupiter size object orbiting a solar type star, also consistent with its optical to near-infrared colours. Comparing this candidate to those from the dedicated OGLE-III transit survey campaigns, we estimate a $\\sim$50\\% probability for this to be a genuine transiting planet, deserving follow-up observations."
    },
    "0711/0711.2219_arXiv.txt": {
        "abstract": "In two recent articles a detailed study has been presented of the out of equilibrium dynamics of an infinite system of self-gravitating points initially located on a randomly perturbed lattice. In this article we extend the treatment of the early time phase during which strong non-linear correlations first develop, prior to the onset of ``self-similar'' scaling in the two point correlation function. We establish more directly, using appropriate modifications of the numerical integration, that the development of these correlations can be well described by an approximation of the evolution in two phases: a first perturbative phase in which particles' displacements are small compared to the lattice spacing, and a subsequent phase in which particles interact only with their nearest neighbor. For the range of initial amplitudes considered we show that the first phase can be well approximated as a transformation of the perturbed lattice configuration into a Poisson distribution at the relevant scales. This appears to explain the ``universality'' of the spatial dependence of the asymptotic non-linear clustering observed from both shuffled lattice and Poisson initial conditions. ",
        "introduction": "Structure formation in the universe is currently addressed primarily using numerical simulations of purely self-gravitating particle systems, with initial configurations generated by displacing the particles slightly from a perfect lattice (see e.g. \\cite{efstathiou_88, couchman}). The physics of the strongly non-linear regime of the observed evolution is, in detail, very poorly understood.  Progress in understanding would be very useful in providing analytical guidance for numerical simulations, and in particular better control on their precision in representing the relevant continuum limit. In a series of recent papers \\cite{sl1,sl2} we have studied a reduced version of the full cosmological problem, considering a very simple class of randomly perturbed lattices as initial conditions, and evolution in a static universe\\footnote{See also \\cite{Baertschiger:2004tx, maurizio} for earlier studies of evolution from these initial conditions.}.  One of our primary results is that, despite the simplifications, the system we study has qualitative behavior very similar to that observed in the more complex cosmological simulations (correlated perturbations in initial conditions, expanding universe). Notably the evolution is clearly ``hierarchical'' (i.e. structures build up at successively larger scales driven by the linearized fluid theory growth of the initial perturbations), and asymptotically ``self-similar'' (i.e. the time dependence of the two point correlation function is given by a simple scaling of the spatial variable which can be inferred from the linearized fluid theory)\\footnote{To avoid any possible confusion we note that both these terms are used here with meanings different to those commonly ascribed to them in  condensed matter physics. In the latter context both are associated with invariance properties of the spatial correlations under spatial rescalings (see e.g. \\cite{book}). Such properties are not implied by their use in the present context.}. The {\\it functional form} of the spatial dependence of the non-linear correlation function is, on the other hand, just as in the cosmological simulations, a fundamental quantity characterizing the bresults which is determined numerically, but not currently understood (i.e. not predicted analytically or semi-analytically). We have noted in \\cite{sl1,sl2}, however, that this form emerges, to a good approximation, in our simulations {\\it prior} to the asymptotic bscaling regime, in the preceding transient phase in which strong non-linear correlations first develop. In this paper we extend and detail our analysis of this phase. We show in greater detail that the emergence of the observed non-linear two body correlations can be very well approximated by modeling the evolution as constituted of two subsequent phases, with an abrupt matching from one to the other. During the first phase the particles evolve as described by a perturbative analytical approximation we have introduced and studied in \\cite{joyce_05, marcos_06}; in the second phase the particles evolve under a force coming solely from their nearest neighbor (NN).  We relate our work here to previous work along similar lines concerning evolution from Poissonian initial conditions \\cite{bottaccio3, maurizio, Baertschiger:2004tx}. In this case it has been shown explicitly \\cite{Baertschiger:2004tx} that the emergence of the first strong non-linear correlations can be very well accounted for by an approximation, at the relevant early times, in which the full gravitational force on each particle is truncated to that due only to its initial NN\\footnote{The importance of NN interactions at early times starting from Poisson initial conditions has also been discussed previously by Saslaw. See \\cite{saslaw2000} and references therein.}. In the case of ``shuffled lattice'' (SL) initial conditions, which we consider in this work, such an approximation is not generically good: when the typical displacement of a particle is small compared to the lattice spacing, the high degree of symmetry gives that the force on a typical particle is the sum of comparable contributions from many particles. In this regime, however, we can describe the evolution very well by a simple perturbative approximation, which has been developed fully in \\cite{joyce_05, marcos_06}.  This latter approximation breaks down, roughly, when particles start to approach one another, which is precisely when one expects a NN approximation for the force may become appropriate.  We show here that this is indeed the case, and that an abrupt switch between the two phases gives a very good approximation to the evolution. Further this model allows us to explain the fact that the observed form of the non-linear correlations in our simulations is independent of the amplitude of the initial shuffling, and the same as that observed from Poisson initial conditions. This is the case because, for the range of amplitude of the initial perturbations we use, the evolution in the first phase brings the system to a distribution with correlation properties which, at the relevant scales, are essentially those of a Poisson distribution.  The main interest of our results here is that they give a semi-analytical understanding of the origin of the form of the observed non-linear two point correlations for this class of initial conditions, which are qualitatively similar to those used in cosmological type simulations. As remarked in \\cite{sl1,sl2} the form of this early time correlation function coincides, to a very good approximation, with that which is also observed in the asymptotic scaling regime attained by the system at longer times. This suggests strongly (but does not prove) that the physical mechanism leading to the former correlations, which we identify here, is also that which gives rise to the latter correlations. In the context of cosmological simulations such a conclusion, if appropriate also for that case, would be very important for the following reason. In this context the results derived from the numerical simulations which are physically relevant  are those which are representative of the Vlasov-Poisson (VP) limit of the simulated particle system. The mechanism we describe here for the generation of the non-linear correlations is, on the other hand, clearly {\\it not} representative of this limit: the effect of interactions with single NN particles are precisely of the kind which are discarded in the VP limit, which is a mean-field approximation. Therefore, if the form of the non-linear correlations in the long time evolution turns out actually to be determined in such a phase, this form would not be representative, as required, of the VP limit. We discuss this point a little further in our conclusions, and suggest numerical tests which could be performed to determine whether the long-time behavior is indeed linked to the early time mechanism we study here.  The paper is organized as follows. In Sec.~\\ref{sec2} we discuss some of the relevant properties of the initial conditions of our simulations.  In the next section we give the details of the simulations considered here and summarize briefly the main relevant results of \\cite{sl1,sl2}.  In Sec.~\\ref{sec3} we present in detail the two-phase model which captures the essential elements of the formation of the first non-linear correlations. We give results here also of numerical simulations.  Finally in Sec.~\\ref{sec4} we discuss the results and draw our main conclusions. ",
        "conclusions": "} \\end{table}  To characterize the correlation properties of the distributions we will use the same quantities as in \\cite{sl1,sl2}: the reduced two point correlation function $\\xi(\\ve r)$, the power spectrum $P(\\ve k)$ (which is related to $\\xi(\\ve r)$ by a Fourier transform\\footnote{The power spectrum is the Fourier transform of $\\tilde \\xi(\\ve r)$, which differs from what we refer to here as the ``correlation function'' by a delta function singularity at $r=0$. See \\cite{sl1,sl2} or \\cite{book}.}), the NN PDF $\\omega(r)$. We refer the reader to \\cite{sl1} for the precise definitions of these quantities. We will also consider the stochastic properties of the force, which we characterize using $P(F)$, the PDF for the modulus of the force $F$.  While $\\delta=\\infty$ corresponds exactly to the Poisson distribution, one expects any SL with $\\delta \\geq 1$ to approximate the correlation properties of a Poisson distribution up to a scale of order $\\Delta=\\delta \\ell$. Indeed, if $\\delta \\geq 1$, the effect of the short distance exclusion of the underlying lattice should disappear and the particles are, to a good approximation, randomly placed in a volume of order $\\Delta^3$. This can be seen explicitly for the power spectrum, of which the exact analytical expression may be written \\cite{book,sl1} in the form \\begin{equation} \\label{eq_pk} P(\\ve k)= \\frac{1}{n_0} + |\\tilde{p} (\\ve k) |^2 A(\\ve k) \\end{equation} where $n_0$ is the mean particle density and $\\tilde{p} (\\ve k)$ is the characteristic function of $p (\\ve u)$ (i.e its Fourier transform normalized so that $\\tilde{p} (0)=1$). The function $A(\\ve k)$ depends only on the initial unperturbed lattice distribution (and not on the shuffling). For any simple form of the PDF for the shuffling (such as the top-hat one considered here and in \\cite{sl1}, or, e.g., a Gaussian PDF as used in \\cite{sl2}) $\\tilde{p} (\\ve k)$ decreases toward zero for $k \\Delta > 1$, giving that the power spectrum tends to the Poissonian value (given by $1/n_0$). Thus for wave-numbers $k$ larger than of order $1/(\\delta \\ell)$ the power spectrum converges to this Poissonian behavior. We refer the reader notably to Fig. 2 of \\cite{sl1}, which show the power spectrum for the four initial conditions SL16, SL24, SL32 and SL64.  \\subsection{Nearest neighbor distribution} \\label{Nearest neighbor distribution}  We will consider below often the NN PDF, and it is useful to know its form in the initial conditions just described. For the Poissonian limit $\\delta=\\infty$ it is straightforward to show analytically \\cite{book} that it is \\begin{equation} \\omega_P (r)= 4\\pi n_0 r^2 \\exp \\left(-\\frac{4}{3}\\pi n_0 r^3 \\right) \\label{poisson-NNpdf} \\end{equation} which gives an average distance between NN $\\Lo\\approx 0.55 n_0^{-1/3}=0.55 \\ell$. Since the NN distribution characterizes the small scale properties we expect, following our discussion above, that this expression will be a good approximation for $\\delta \\geq 1 $. For $\\delta \\ll 1$, on the other hand, one may show\\footnote{The derivation of the expression given is straighforward, but tedious. For a given shuffling of a particle and its six NN, one must determine exhaustively the different combinations, and associated probabilities, which lead to a given NN separation. The approximation $\\delta \\ll 1$ is used in taking the inter-particle separations to linear order in $\\delta$.} that \\be \\ell \\omega(r) \\approx \\frac{1}{(2\\delta)^9} f \\left(\\frac{r}{\\ell}-1\\right) \\label{eq:omegasmall} \\ee where \\be f(x)= \\begin{cases} \\delta(2\\delta + x)\\left[2\\delta x^2 + 8\\delta^2 x \\right]^2 & \\text{if } x \\in [-2\\delta,-\\delta] \\\\ \\left[ \\delta^2 - x^2 + \\delta x \\right] \\\\ \\times \\left[ 2\\delta x^2 +\\frac{4}{3} \\delta^3-4\\delta^2x+\\frac{4}{3}x^3\\right]^2 & \\text{if } x \\in \\ ]-\\delta,0] \\\\ \\frac{16}{9}(\\delta - x)^8 & \\text{if } x \\in \\ ]0,\\delta] \\\\ 0 & \\text{otherwise}, \\end{cases} \\nonumber \\ee corresponding to average distance between NN $\\Lo = \\ell - (86827/80640) \\Delta$.  In Fig.~\\ref{omegat0} we show the behavior of $\\omega(r)$ for most of the SL studied here. For the SL with very small shuffling ---~SL16 or SL24~--- this PDF is strongly peaked around the average distance between NN (which is approximately equal to $\\ell$), and in very close agreement with the analytical approximation given by Eq.~(\\ref{eq:omegasmall}). For SL32 (with $\\delta=0.177$) a small discrepancy with this approximation is visible, while for SL64 (with $\\delta=1$) the NN PDF is, as expected, in very good agreement with that in the Poisson case given by Eq.~(\\ref{poisson-NNpdf}). \\begin{figure} \\includegraphics[width=0.5\\textwidth]{FIG1.eps} \\caption{NN PDF of the SL considered in this paper (the name of each SL is indicated above the corresponding curve). For SL32, SL24 and SL16, the function~\\eqref{eq:omegasmall} is shown for comparison. In the insert panel we show enlarged the SL64 (SL) case together with the behavior of $\\omega_p(r)$ for a Poisson distribution (P) with same number density [i.e. Eq.\\eqref{poisson-NNpdf}]. \\label{omegat0}} \\end{figure}  \\subsection{Force distribution} \\label{Force distribution}  The PDF of the modulus of the force $W(F)$ is a useful quantity in our analysis. Notably if the forces on particles are dominated by that coming from their NN the simple relation \\begin{equation} W(F)dF=\\omega (r)dr \\label{NNdomination} \\end{equation} must hold. For the case of a Poisson distribution the analytical expression for $W(F)$ was first given by Chandrasekhar \\cite{chandra43}. It is proportional to the so-called Holtzmark distribution (see \\cite{book} for the explicit result and a simple derivation). In Fig.~\\ref{forceHNN} \\begin{figure} \\includegraphics[width=0.5\\textwidth]{FIG2.eps} \\caption{Holtzmark distribution and the PDF inferred if only NNs contribute, i.e., $W_{NN}(F)dF=\\omega (r) dr$. The agreement is very good in the strong field limit where $W(F) \\sim F^{-5/2}$. At weak field the PDF due to the NN has a sharp cut-off while the Holtzmark distribution shows a more gentle decrease (see discussion in \\cite{book,gabrielli_06}). \\label{forceHNN}} \\end{figure} we show a plot of the full (Holtzmark) distribution, and the PDF inferred if only NNs contribute $W_{NN}(F)dF=\\omega (r) dr$. The domination of the NN in the force is clearly seen at stronger values of the force. The relation is not valid at weaker values of the force as these correspond to the (rare) particles for which the force picks up comparable (and possibly canceling) contributions from more than one particle. Note that the tail of the PDF at large $F$ decays in proportion to $F^{-5/2}$, which means that the variance (i.e. second moment) of the PDF is infinite.  In \\cite{gabrielli_06} we have studied in detail the statistical properties of the force in an SL. As one might expect, one can show that the force PDF is very well approximated by that of a Poisson distribution when the typical displacement is larger than the inter-particle spacing, i.e., $\\delta > 1$. At small values of the displacements, on the other hand, the force PDF is very different to that in a Poisson distribution, decaying much more rapidly at large values of the force: the strong forces due to NN are completely absent as the typical particle feels a comparable effect from its {\\it six} NN when the configuration is close to a perfect simple cubic lattice. More precisely, for a top-hat PDF of the displacements (as used here), the functional behavior of the PDF at large $F$ changes qualitatively, from exponential decay to a $F^{-5/2}$ power law decay, at $\\delta=0.5$. For $\\delta \\gtapprox 0.5$ the amplitude of the latter tail is lower than that in the Poisson, with this difference becoming negligible as $\\delta$ increases to of order unity.   \\label{sec4}  In this article we have studied in detail the early time evolution of infinite self-gravitating shuffled lattices, as well as the limiting case given by the Poisson distribution. We have shown that a very good description of the evolution of two point correlations in this phase is given by a simple approximation in which the force on particles abruptly switches from that given by the PLT approximation developed in \\cite{joyce_05, marcos_06} to the force due only to NN particles. Further in the first phase the system evolves at small scales to always resemble closely the Poisson distribution, explaining the universality of the form of the non-linear correlation function which emerges. We have noted, however, that this universality will not extend to SL initial conditions with arbitrarily small initial shuffling. In this limit effects come into play in the very long first (PLT) phase leading to a strongly correlated evolution at all scales, which is different from (and unrelated to) that in the Poisson distribution.  We have thus given, for this specific class of initial conditions, an explanation of the non-linear correlations which emerge at early times. As underlined in \\cite{sl1,sl2} this non-linear correlation function coincides with that which is observed at later times, when the system manifests a simple spatio-temporal scaling (or ``self-similarity''). Thus the model appears to explain this asymptotic form of the non-linear correlations. As described in \\cite{sl2} this can be understood also in the following way: these non-linear correlations in the system evolving at any later time can be well approximated by those in an evolved coarse-grained ``daughter'' distribution. In the latter the system may be in the early time phase studied here, while the original distribution is not. The non-linear correlation functions of the two systems nevertheless coincide.  Our results are of relevance to simulations of structure formation in the universe in cosmology. In this case the goal of numerical simulation is to recover the non-linear correlations in the Vlasov-Poisson (VP) limit of the evolution of a self-gravitating $N$-body system. These initial conditions are different to those used here --- with initially {\\it correlated} displacements and an expanding space --- but, as shown in \\cite{sl1,sl2}, the evolution is qualitatively similar to our simpler case. {\\it If} the same kind of model can be used to explain non-linear correlations in an early time regime in this case then these correlations clearly are not described by the VP limit: the forces due to NN particles are neglected in this (mean field) limit, while in this model they are the dominant ones. In this context this would mean that results in this regime would need to be discarded as unphysical (i.e. not representing the required physical limit, but just a numerical effect arising from the method of discretization). Further, if the form of the asymptotic non-linear correlation function is really determined by this early time evolution, these effects of discreteness are then important at all times and the system never represents the VP limit. This resemblance of the early time and asymptotic correlation function may, however, be a simple coincidence. The evolution at longer times may then indeed be representative of the VP limit while discreteness effects are important only at early times.   We will study these issues further in future work. Firstly the application of our simple two phase description to cosmological simulations should be investigated. We expect the expansion of space to change nothing qualitatively in our model: the description of the PLT phase is qualitatively unchanged \\cite{marcos_06}, and NN domination during the formation of the first non-linear structures in such simulations has been explicitly shown in \\cite{Baertchiger:2002tk}. The fact that the typical initial conditions of such simulations have much more long wavelength power than those considered here may, however, be important [typically one has $P(k) \\sim k^{m}$ with $m\\leq -1$, compared to $n=2$ (SL)and $n=0$ (Poisson) considered here]. Secondly, as we have discussed at some length in the conclusions of \\cite{sl2}, it would be instructive to study numerically the relation of the asymptotic regime to the early time regime by completely modifying the former using a large smoothing in the force, i.e., a smoothing scale $\\varepsilon \\gg \\ell$. In this case the dynamics we have described, between $NN$ particles, will not occur. However at longer times the system should still evolve the same correlations if the VP limit of the system is indeed that represented by the simulations with $\\varepsilon \\ll \\ell$."
    },
    "0711/0711.2958_arXiv.txt": {
        "abstract": "We present the initial results from a study of the SNR population in a sample of six nearby galaxies (NGC 2403, NGC 4214, NGC 4449, NGC 5204, NGC 3077, NGC 4395) based on Chandra archival data. We discuss the analysis of the Chandra data and we present candidate SNR sources selected on the basis of their X--ray colours. We also present deep [S {\\sc ii}] 6716 \\& 6731 ${\\rm \\AA}$~and H$\\alpha$ line images for most of the galaxies in our sample, which provide optically selected samples of SNRs. Comparison of the X--ray results with the complementary optical observations provides a more complete picture of the SNR population and allows us to address their X--ray emission. Our preliminary analysis of the [S {\\sc ii}]/H$\\alpha$ images show that 48 X--ray sources are typically associated with H$\\alpha$ sources, 7 of which are SNR candidates based on their [S {\\sc ii}]/H$\\alpha$ ratio and one is an already known radio SNR. ",
        "introduction": "Supernova Remnants (SNRs) are an important component of the X--ray source populations in galaxies (Blair \\& Long 1997), especially in luminosities below $10^{37}$ erg sec$^ {-1}$. The expanding SNR shocks heat the surrounding interstellar medium (ISM) to temperatures of $10^6$ -- $10^7$ K and produce X--ray emission. This way SNRs provide a significant fraction of the mechanical energy that heats, shapes and chemically enriches the ISM. Despite extensive studies of individual galactic SNRs, a complete understanding of these sources remains elusive. It is well known that different wavebands provide a picture of different evolutionary stages of the SNR populations. On the other hand, although it is almost certain that environment plays a major role in the evolution and the multiwavelength properties of SNRs (e.g. Pannuti et al. 2007), the details of this connection are poorly understood. In order to address this connection, we initiated an X--ray and optical study of six nearby galaxies. Here we present the preliminary results from this work focusing on the identification and classification of the SNR candidates.  \\begin{figure} \\centerline{\\psfig{file=leonidakifig2.ps,width=7.7cm,clip=} } \\caption{X-ray color--luminosity plot. \\label{image}} \\end{figure} ",
        "conclusions": ""
    },
    "0711/0711.0994_arXiv.txt": {
        "abstract": "We extend our earlier model of the small scale structure of cosmic string networks through an improved treatment of the separation of long and short scales.  We find that the production of small loops (at the gravitational radiation scale) is a robust feature of string networks, in addition to a population of loops near the horizon scale.  We obtain quantitative agreement with the scaling of loop production functions as found in simulations by two groups. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.4265_arXiv.txt": {
        "abstract": "{We use the high magnification event seen in the 1999 OGLE campaign light curve of image C of the quadruply imaged gravitational lens Q2237+0305 to study the structure of the quasar engine. We have obtained g'- and r'-band photometry at the Apache Point Observatory 3.5m telescope where we find that the event has a smaller amplitude in the r'-band than in the g'- and OGLE V-bands. By comparing the light curves with microlensing simulations  we obtain constraints on the sizes of the quasar regions contributing to the g'- and r'-band flux. Assuming that most of the surface mass density in the central kiloparsec of the lensing galaxy is due to stars and by modeling the source with a Gaussian profile, we obtain for the Gaussian width $1.20\\times 10^{15} \\sqrt{{M}/{0.1M_\\odot}}~cm \\lesssim \\sigma_{g'} \\lesssim 7.96\\times 10^{15} \\sqrt{{M}/{0.1M_\\odot}}~cm$, where M is the mean microlensing mass, and a ratio $\\sigma_{r'}/\\sigma_{g'}=1.25^{+0.45}_{-0.15}$. With the limits on the velocity of the lensing galaxy from Gil-Merino et al. (2005) as our only prior, we obtain $0.60\\times 10^{15} \\sqrt{{M}/{0.1M_\\odot}} ~ cm \\lesssim \\sigma_{g'} \\lesssim 1.57 \\times 10^{15} \\sqrt{{M}/{0.1M_\\odot}} ~ cm$ and a ratio $\\sigma_{r'}/\\sigma_{g'}=1.45^{+0.90}_{-0.25}$ (all values at 68 percent confidence). Additionally, from our microlensing simulations we find that, during the chromatic microlensing event observed, the continuum emitting region of the quasar crossed a caustic at $\\geq72$ percent confidence.} ",
        "introduction": "One of the best studied quasar lensing systems yet is \\2237. This quasar was discovered by the Center for Astrophysics (CfA) redshift survey \\cite[]{huchra85}. The system is made up of a barred spiral galaxy (z=0.0394) and four images of a quasar at a redshift of z=1.695 that appear in a nearly symmetric configuration around the core of the spiral galaxy at a distance of approximately 1'' from the center.  \\2237 is an ideal system to study quasar microlensing. At the position of the quasar images the optical depth due to microlensing by the stars is high \\cite[]{kayser86,kayser89,wambsganss90}. Furthermore the expected time-delay between the images of this quasar is of the order of a day or less \\cite[]{schneider88, rix92,wambsganss94}. As microlensing events for this system are of the order of months or years, it is easy to distinguish intrinsic variability from microlensing. Another big advantage is the fact that the quasar is $\\sim$10 times farther from us than the lensing galaxy. This leads to a large projected velocity of the stars in the source plane and thus a short timescale of caustic events. Not long after its discovery in 1985, \\cite{irwin89} observed the first microlensing signature in this quasar.  The quasar has been monitored for almost two decades by different surveys \\cite[e.g.,][]{corrigan91,pen93,ostensen96,alcalde02,schmidt02}. The one that has the longest and best sampled data is that of the Optical Gravitational Lensing Experiment (OGLE) team \\cite[]{wozniak00, udalski06}. They have monitored \\2237 since 1997 delivering the most complete light curves for this system in which many microlensing events can be seen. These data have been used for various studies on the system: \\cite{wyithe00f,wyithe00g,wyithe00b,wyithe00d,wyithe00e}; \\cite{yonehara01}; \\cite{shalyapin02} and \\cite{kochanek04} used this data together with microlensing simulations not only to obtain limits on the properties of the quasar such as size and transversal velocity, but also on the mass of the microlensing objects.    The OGLE team has monitored this quasar with a single V band filter. However, \\cite{wambsganss91} showed that color variations should be seen in quasar microlensing events and how they would provide additional information on the background source. We know from thermal accretion disk models that emission in the central regions of a quasar should be bluer than in the outer regions \\cite[e.g.,][]{SS73}. As an approximation we can assume the central engine of the quasar to be of the order of 1000 Schwarzschild radii: \\begin{equation} R_{\\rm ad}=1000\\times r_s=2.8\\times10^{16} \\left(\\frac{M_{\\rm BH}}{10^8M_{\\odot}}\\right){\\rm cm} \\end{equation} where $r_s$ is the Schwarzschild radius, $R_{\\rm ad}$ is the radius of the accretion disk and $M_{\\rm BH}$ is the mass of the black hole. We can compare this value with the Einstein Radius ($R_E$) - the length scale in quasar microlensing - which can be given in the source plane by \\cite[]{SEF92}:  \\begin{equation} R_E= \\sqrt{ \\frac{4GM}{c^2} \\frac{D_s D_{ls}}{D_l}} = 5.68 \\times 10^{16} \\sqrt{\\frac{M}{0.1 M_{\\odot}}} {\\rm cm} \\label{eradius} \\end{equation} where we have replaced the values of the distances $D_l$, $D_s$ and $D_{ls}$ to the ones corresponding to \\2237. The accretion disk scales are thus comparable to those of the Einstein rings of stars in the lensing galaxy. This shows that differential magnification of the different emission regions of the quasar can be observed.   Some multi-band observations of this system have been carried out \\cite[]{corrigan91,vakulik97,vakulik04,koptelova05}, however most of them have either large gaps between the observations, or no obvious microlensing event.  A very interesting event, that has been one of the bases for predictions \\cite[]{wyithe00g} and study \\cite[]{yonehara01,shalyapin02, vakulik04}, is the high magnification event of image C in the year 1999 as observed by OGLE.  In this paper we present data of this particular event obtained in two filters at Apache Point Observatory (APO). We analyze the event using both the OGLE light curve and the APO data and study the implications for the size of the quasar emission region using microlensing simulations. We use a flat cosmology with $\\Omega_m=0.3$ and $H_0=70~km~s^{-1} Mpc^{-1}$.  \\begin{figure}[] \\centering \\includegraphics[width=8cm,bb=0 0 732 598]{8221fig1.eps} \\caption{APO g' image of the quasar system \\2237.} \\label{full2237} \\end{figure} ",
        "conclusions": "We present two band (g' and r') APO data covering the high magnification event seen in image C of the quadruple quasar \\2237 in the year 1999. We find that the amplitude of the brightness peak is more pronounced in the g'-band than in the r'-band. This is also consistent with the observations by \\cite{vakulik04} (see also the recent analysis by \\citealt{koptelova07}).  By using this data together with the well known OGLE data \\cite[]{wozniak00,udalski06} and combining it with microlensing simulations, we have been able to obtain limits on the size of regions of the quasar's central engine emitting in these bands: Gaussian width $\\sigma_{g'}=4.60^{+3.36}_{-3.40}\\times 10^{15} \\sqrt{{M}/{0.1M_\\odot}}$ cm and $\\sigma_{r'}/\\sigma_{g'}=1.25^{+0.45}_{-0.15}$.  Because of the degeneracy between source size and transverse velocity we use a prior on the velocity obtained from the work of \\cite{gilmerino05} in order to improve the results: Gaussian width $\\sigma_{g'}=1.33^{+0.24}_{-0.73}\\times 10^{15} \\sqrt{{M}/{0.1M_\\odot}}$ cm and $\\sigma_{r'}/\\sigma_{g'}=1.45^{+0.90}_{-0.25}$.  Both values for the ratio between the source sizes are close to the ratio obtained for a face-on \\cite{SS73} accretion disk (f$\\sim$1.5). Recent studies \\cite[e.g.,][]{poindexter07,morgan07,pooley07} suggest microlensing yields a slightly bigger value for the ratio than that obtained analytically with the thin disk model (see Section 5.2), also in agreement with our results.  We also show that this event was probably produced by the source directly interacting with a caustic as we obtain probabilities of 97 percent and 94 percent that the r'-band emitting region touches a caustic without and with the G-M05 prior, respectively, and of 75 percent and 72 percent that the source center (regardless of size) crossed a caustic without and with the G-M05 prior, respectively."
    },
    "0711/0711.1440_arXiv.txt": {
        "abstract": "The QSO luminosity function at $z>5$ provides strong constraints on models of joint evolution of QSO and their hosts. However, these observations are challenging because the low space densities of these objects necessitate surveying of large areas, in order to obtain statistically meaningful samples, while at the same time cosmological redshifting and dimming means that rather deep Near Infrared (NIR) imaging must be carried out. Several upcoming and proposed facilities with wide-field NIR imaging capabilities will open up this new region of parameter space. In this paper we present predictions for the expected number counts of $z>5$ QSOs, based on simple empirical and semi-empirical models of QSO evolution, as a function of redshift, depth and surveyed area.  We compute the evolution of observed-frame QSO magnitudes and colors in a representative photometric system covering the wavelength range $550 \\, \\rm{nm} < \\lambda < 1800 \\, \\rm{nm}$, and combine this information with different estimates for the evolution of the QSO luminosity function. We conclude that planned ground-based surveys such as Pan-STARRS and VISTA should be able to detect a large number of luminous QSOs up to $z \\la 7.5$, but that space-based missions such as EUCLID (formerly SPACE/DUNE) or SNAP are probably required in order to obtain substantial samples at higher redshift. We also use our models to predict the expected number counts for future X-ray space missions (such as XEUS and Constellation-X), and show that because of their small field-of-view, these telescopes are unlikely to discover significant numbers of AGN at very high redshift. However, X-ray follow-up of objects detected at longer wavelength will be an important means of confirming their identity as AGN and constraining obscuration. ",
        "introduction": "A number of observations point towards a tight relationship between the properties of the AGN/QSO\\footnote{We refer to bright AGNs as QSOs, with no reference to their radio properties.}  population and their host galaxies. This evidence includes the local relation between the mass of the central supermassive black hole (SMBH) and the mass of the spheroidal component of the host (e.g. H\\\"aring \\& Rix, 2004), and the apparent common ``downsizing'' behavior of star formation and black hole accretion (Hasinger, Miyaji \\& Schmidt, 2005). As well, it is now becoming fairly widely accepted that the energy released by accretion onto SMBH is an important mechanism in regulating galaxy growth and star formation, although the details of how this ``AGN feedback'' process works remain unclear.  In recent years, several theoretical studies have attempted to understand the complex interplay between the physical mechanisms that lead to the observed properties of galaxies and SMBH and their evolution with redshift (for recent implementations see e.g., Monaco, Fontanot \\& Taffoni, 2007; Hopkins et al., 2007; Somerville et al., 2007; Croton et al., 2006; Bower et al., 2006).  It has become evident that the redshift evolution of the QSO population provides some of the strongest constraints on this class of models (Fontanot et al., 2006; Bromley, Somerville \\& Fabian 2004). It is therefore of fundamental importance to determine the statistical properties of high-redshift QSOs, and in particular their luminosity function (LF), up to the highest possible redshifts (Richards et al., 2006; Jiang et al., 2007; Fontanot et al., 2007, hereafter F07). In addition, samples of high redshift QSOs are important in order to place constraints on when and how SMBH formed, and are of interest for studying the reionization epoch (Gallerani et al. 2007).  A considerable difficulty in determining the high-z QSO LF is due to the very low space-density of these objects, which are detected in considerable numbers only in very large area surveys such as the Two Degree Field QSO Redshift Survey (2QZ, Croom et al.  2004) and the third edition of the Sloan Digital Sky Survey (SDSS) Quasar Catalog (SDSSqso3, Schneider et al.  2005). As well, cosmological redshifting and dimming, and the severe IGM absorption at such redshifts require fairly deep Near-Infrared (NIR) observations. This combination of area and depth in the NIR has not been achievable up until now.  However, a number of upcoming and proposed projects will begin to change this situation. For example, several deep-wide NIR surveys from the ground are in progress or planned, such as the UKIDSS\\footnote{http://www.ukidss.org/} Large Area Survey (LAS), Deep Extragalactic Survey (DXS), and Ultra Deep Survey (UDS), the Panoramic Survey Telescope and Rapid Response System (Pan-STARRS)\\footnote{http://www.ps1sc.org/index.htm} and VISTA\\footnote{http://www.eso.org/sci/observing/policies/PublicSurveys/ \\, sciencePublicSurveys.html\\#VISTA} Ultra-VISTA, VIKING, VHS, and VIDEO surveys.  Hopefully, Wide Field Camera 3 (WFC3) will soon be installed on the Hubble Space Telescope, and with its 4.8 arcmin$^2$ Field of View (FOV), it will greatly improve our current abilities to survey relatively large areas in the NIR from space. However, major progress in characterizing the very high redshift QSO population will probably have to wait for two kinds of future space telescopes. The James Webb Space Telescope (JWST) will have a large aperture (6.5 m) and high sensitivity ($\\sim 3.5$ nJy), but its NIR imager NIRCAM will have a relatively small FOV (2 $\\times$ 4.7 arcmin$^2$). On the other hand, several new space telescopes have been proposed with smaller apertures ($\\sim 1-2$ m) but with very large ($\\sim 0.5$--1 sq. deg.) optical and NIR cameras, such as SNAP\\footnote{http://snap.lbl.gov/} and EUCLID (formerly DUNE\\footnote{http://www.dune-mission.net/} and SPACE\\footnote{http://www.spacesat.info/}). While the main motivation of these latter kinds of missions is to constrain dark energy through Supernovae, weak lensing, or baryon oscillations, as we will show in this paper, they are also very well suited for the important goal of studying the QSO population at very high redshift.  Complementary information on the evolution of the QSO population will be provided by the next generation of X-ray space observatories.  Two proposed missions are of particular interest in this regard: XEUS\\footnote{http://sci.esa.int/science-e/www/area/index.cfm?fareaid=103} and Constellation-X\\footnote{http://constellation.gsfc.nasa.gov/}. The former is a 4.2 m telescope (minimum effective area 3 m$^2$ from 2 to 10KeV), with an imaging resolution better than 5 arcsec; the latter consists of four coaligned 1.3 m telescopes on a single spacecraft (minimum effective area 4 $\\times$ 0.1 m$^2$), reaching 30 arcsec resolution in the hard band.  Both have a relatively small FOV ($\\sim 7$--10 and 5 arcmin$^2$, respectively).   The goal of this paper is to present predictions for the observed-frame near-infrared colors and expected number density of $z>5$ QSOs, which may help to guide the planning for future surveys with these kinds of facilities. For comparison we also present predictions in terms of the flux limit in the hard X-ray band. Of particular interest is the trade-off between area and depth. Although it is possible to make predictions for the relevant quantities using physically motivated semi-analytic models set within the hierarchical structure formation paradigm (e.g. Volonteri \\& Rees 2006; Salvaterra et al. 2007; Rhook et al. 2008), there are extremely large uncertainties in these predictions at high redshift, associated with our lack of knowledge about such factors as the nature of seed black holes, the efficiency with which early black holes can grow, and whether these holes can be ejected from their host galaxies by effects such as the gravitational rocket (see e.g. Volonteri \\& Rees 2006). We therefore adopt a more empirical approach, in which we make use of observed AGN luminosity functions at the highest available redshifts ($z\\sim 5.5$--6) and template AGN spectral energy distributions (SEDs). We combine these with an array of simple assumptions about how the QSO population evolved with time to predict the observable properties back to $z\\sim 10$. In our simplest and most optimistic model, we assume that the intrinsic properties of the underlying population do not evolve at all. In the absence of strong evolution in the SEDs of QSOs, this should be an upper limit to the number of objects that can be detected at high redshift. In an intermediate model, we extrapolate the observed evolution in the redshift interval $3.5 < z < 5.2$ to higher redshift using a simple parameterization of luminosity/density evolution developed by Fontanot et al. (2007). Finally, in a semi-empirical model, we compute the backwards evolution of the luminosity function assuming that QSOs at $z\\sim 6$ have grown by steadily accreting at their Eddington luminosity. Although simple, these three models are empirically motivated and probably provide a reasonable bracketing of the expected results. As well as providing a good guess for the parameters of the populations to be targetted by the future NIR and X-ray space missions that we have mentioned, they will provide a useful foil for predictions from more physically motivated formation models.  The structure of the paper is as follows. In section~\\ref{clone} we describe our technique for computing observed-frame QSO colors as a function of redshift. In section~\\ref{counts} we then combine this information with different estimates for the evolution of the LF, in order to bracket the expected QSO counts as a function of redshift, area and depth. In section~\\ref{x} we use a similar approach to predict the corresponding quantities as a function of X-ray flux in the $2-10 \\, {\\rm keV}$ band. Finally in Section~\\ref{fin} we present our conclusions.  Throughout this paper we assume a cosmology with $h$, $\\Omega_{\\rm tot}$, $\\Omega_m$, $\\Omega_\\Lambda$ = 0.7, 1.0, 0.3, 0.7; magnitudes and colors are in the AB system. ",
        "conclusions": "\\label{fin}  In this paper we present empirical predictions of the observed colors and number densities of very high redshift ($z>5$) QSOs. We combine a representative photometric system with a QSO spectral template library to predict the evolution of QSO colors at $z>5$. Using a set of four optical-NIR filters, we show that it is possible to define color-color criteria that select high-z QSO candidates on the basis of their photometric properties. We then combine the estimated k-corrections with different models for the LF evolution in order to estimate the expected number of QSOs as a function of surveyed area, magnitude limit, and redshift.  We confront our findings with the parameters of existing, planned and proposed surveys.  First, the SDSS has already covered the bright end of the LF to $M_{145}\\approx-26.5$ at $z<6.4$, and is pursuing efforts to push another 2 magnitudes fainter over the 300 square degrees of its deep southern stripe (Jiang et al.\\ 2007).  The Pan-STARRS survey, which is about to begin data-taking, will reach $z$ magnitudes that are up to 1\\,mag fainter still over 20,000 square degrees of extragalactic sky in its $3\\pi$ survey, putting of order 3000 low-luminosity quasars at $6< z<6.4$ within its reach (from an extrapolation of the Shankar \\& Mathur LF).  Since the $3\\pi$ survey includes deep $z$ and $y$ band imaging, it is also ideally suited for the discovery of quasars at redshifts $6.6 \\le z \\le 7.3$. Again integrating the Shankar \\& Mathur LF to the $y$-band flux limit, we expect that Pan-STARRS will contain of order 150 quasars at $6.6 \\le z \\le 7.3$.  In this paper we address the prospects for finding fainter quasars than are accessible with SDSS and Pan-STARRS at $5<z<7.5$, or \\emph{any} quasars at $z>7.5$, with currently planned or proposed surveys. It is evident from fig.~\\ref{fig:nc2} that up to $z\\la 6.5$, the ground-based near-infrared UKIDSS and VISTA surveys should be able to produce samples of roughly tens to hundreds of low-luminosity QSOs, depending on the actual evolution of the LF. Finding quasars in the Pan-STARRS Medium Deep Survey (MDS) requires complementary near-infrared imaging (ideally in $J$) at the depth of VIDEO (3 of whose fields are included in 3 of the MDS fields); if such imaging is available, the larger area would make about 5-6 times more quasars accessible than VIDEO alone.  Even in the absence of a decline in the space density of QSOs, however, our results suggest that in order to obtain samples of substantially more than a hundred low-luminosity QSOs at $z\\ga 6.5$, we will require the larger survey volumes that are accessible only from space (because of the much lower NIR background, a 2m telescope in space is more efficient at surveying large areas than an 8--10m on the ground). It is also possible to use our predictions to estimate the magnitude of the brightest QSO we expect, as a function of redshift, at a fixed surveyed area (fig.~\\ref{fig:newplot}). We define this quantity as the magnitude corresponding to an integrated space density of 1 QSO over the surveyed volume (in order to compute the volumes we consider the same redshift intervals as in fig.~\\ref{fig:nc2}, left panel).  We consider the evolving and non-evolving prescriptions (solid and dashed line respectively) applied to the F07a and SM07 LFs (thick and thin line respectively). In the same plot we also show the faintest magnitude reached at a given depth in the $H$-band (dotted lines). Combining the two informations, we are then able to define an accessible magnitude interval. \\begin{figure} \\centerline{ \\includegraphics[width=9cm]{brightQSO.ps} } \\caption{The magnitude of the brightest QSO in a surveyed area of 10 deg$^2$ (left), 100 deg$^2$ (middle), 500 deg$^2$ (right). Solid and dashed lines correspond to the evolving LF and non-evolving LF prescription respectively. Thick and thin lines correspond to the F07a and SM07 LF respectively. Dotted lines show the faintest absolute limiting magnitude as a function of redshift for a magnitude-limited sample.} \\label{fig:newplot} \\end{figure} Our results highlight that, while JWST will be able to detect extremely faint QSOs, because of its small FOV it is unlikely that it will be possible to survey large enough areas to detect significant numbers of the most luminous QSOs at very high redshift. For this important goal, we need the wider-area surveys that could be carried out by missions like EUCLID or SNAP.  Similar arguments hold if we consider planned X-ray surveys with missions such as XEUS and Constellation-X. The very small FOV of these surveys, combined with the relatively long exposure times needed to reach high enough sensitivity, probably preclude the possibility of detecting large samples of $z>6$ QSOs using the X-ray information alone. Targetted X-ray observations of objects detected in the NIR will, however, be crucial for confirming their identity as AGN and constraining obscuration.  The predictions presented here are complementary to those based on semi-analytic models set within the hierarchical structure formation paradigm. For example, Salvaterra, Haardt \\& Volonteri (2007) have presented predictions for the properties of the AGN population that could be observed in deep X-ray surveys with future surveys, based on the Volonteri \\& Rees (2006) models with different assumptions.  They consider a merger tree based model based on the Press-Schechter formalism and Eddington accretion onto the SMBH. An alternative model has been presented by Rhook \\& Haehnelt (2008). Their hybrid models also assume that QSO activity is triggered by major merger events, and they combine this information with different sets of assumptions for the decline of the accretion rate onto the SMBH. As in our empirical models, their predicted counts span several orders of magnitude. This underlines the importance of the observational determination of the statistics of the QSO population in order to put constraints on models of SMBH growth. We plan to compare the empirical results presented here with additional semi-analytic, merger-tree based models for SMBH evolution (Fontanot et al. 2007; Somerville et al. 2008) in a forthcoming paper.  We expect that only for the brightest objects will spectroscopic follow-up be feasible. The development of adequate photometric redshift estimators, as well as a larger multi-wavelength coverage, is therefore fundamental for the study of the faint population. Moreover a simple drop-out technique is not able to disentangle a priori faint QSOs from Lyman Break Galaxies at comparable redshift.  A key goal in the study of faint high-z QSOs is therefore the definition of reliable criteria to separate QSOs from galaxies. Here, we simply assume that we can perfectly disentangle the faint QSOs from the contaminants such as galaxies and stars. Under this hypothesis we consider a simple test to assess the impact of the uncertainties in the photometric redshift determination on the overall estimate of the QSO-LF and its space density evolution.  We consider the F07 Monte Carlo algorithm for the estimate of the LF, which combines the GOODS and SDSS samples. We then introduce a random relative error in the redshifts of the faint (GOODS) objects, perform 1000 bootstrap realizations of LF fitting, and consider the new estimate of the LF parameters. We conclude that the F07 technique is still able to recover the statistical properties of the LF in the presence of these photometric redshift errors: assuming a $5\\%$ ($10\\%$) error for the photometric redshifts, we add only $0.3\\%$ ($0.5\\%$) to the error on the redshift evolution and $1.0\\%$ ($1.5\\%$) to the error on the faint-end slope. Moreover, if we assume a $5\\%$ ($10\\%$) error on the photometric redshift of all objects (both faint and bright) we add $5\\%$ ($10\\%$) to the error on the redshift evolution; $10\\%$ ($20\\%$) to the error on the faint-end slope and $5\\%$ ($10\\%$) to the error on the bright-end slope.  It seems very likely that at least one ``dark energy'' mission will fly in the next decade. We have shown here that several of the proposed designs for such missions are also very well-suited to addressing a completely different but important science goal: constraining the formation of massive black holes and luminous QSOs at the earliest epochs in the history of our Universe."
    },
    "0711/0711.1506_arXiv.txt": {
        "abstract": "Once a parity is introduced in unparticle physics, under which unparticle provided in a hidden conformal sector is odd while all Standard Model particles are even, unparticle can be a suitable candidate for the cold dark matter (CDM) in the present universe through its coupling to the Standard Model Higgs doublet. We find that for Higgs boson mass in the range, 114.4 GeV $\\lesssim  m_h \\lesssim$ 250 GeV, the relic abundance of unparticle with mass 50 GeV $\\lesssim  m_{\\cal U}  \\lesssim$ 80 GeV can be consistent with the currently observed CDM density. In this scenario, Higgs boson with mass $m_h \\lesssim 160$ GeV dominantly decays into a pair of unparticles and such an invisible Higgs boson may be discovered in future collider experiments. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.4719_arXiv.txt": {
        "abstract": "{We discuss the young population of stars and clusters in the Magellanic Clouds. We present the discovery of pre-main sequence candidates in the nebula N~11 in the Large Magellanic Clouds. The comparison of the Colour-Magnitude diagram with pre-main sequence tracks and the presence of Spitzer objects YSO I and II suggest that the star formation has been active for a long period in the region, from  a few $10^5$ yrs  to several Myr ago} ",
        "introduction": "The process of star formation in different environments is far from being understood. In particular, it is difficult to reconcile the prominent influence of the local environment (turbulence, compression, initial trigger) on small scales with the universality of the Schmidt and Kennicut law on Galactic scales which suggests that Galactic-scale gravity is involved in the first stages of star formation. Detailed understanding of how an entire molecular cloud converts into stars is still lacking. To analyze  star forming regions can cast light on the formation scenarios, since  spatial distributions of different components (molecular clouds, ionized gas), and of stars in different evolutionary stages can help to distinguish between the proposed models. Local Group galaxies, and in particular the Magellanic Clouds (MCs), are very promising regions to study the process of star formation. Their young population of stars and clusters allow to derive hints about the formation process in low metal content environments. This is particularly important since the formation of stars depends on the balance between the gas heating and the cooling, on which  the presence of metals has a significant effect. In this paper, in Section~\\ref{forma} we first discuss the formation mechanism of field stars and clusters, in section~\\ref{trigger} the  effects on the star formation of the gravitational interaction between the MCs and the Milky Way are discussed, in section~\\ref{n11a} and \\ref{n11b} the discovery of pre-main sequence candidates in N~11 in the LMC is presented. Final remarks are  given in section \\ref{conclu}. ",
        "conclusions": "\\label{conclu} In this paper we discuss the young population of stars and clusters in the Magellanic Clouds. The last gravitational interaction between LMC and SMC has triggered a star formation episode in both  the cluster and field stars. We find that the SMC was relatively quiescent at ages older than 6-8 Gyr. This result suggests that at older ages, the tidal interaction between the Magellanic Clouds and the Milky Way was not able to trigger significant star formation events. Field star and cluster formation are not completely correlated, as expected on the basis of theoretical models and different star formation modes are expected to act in the field and in cluster populations. Young stars are well known to  trigger star formation on local scale by injecting energy in the surrounding interstellar medium.   N~11 region in the LMC is one of the most promising candidate of triggered star formation. In this paper we discuss the star formation process in this region. We report on the discovery of a PMS candidates associated with N~11 from HST ACS/WFC photometry. The PMS are consistent with an age going from 1 to 10 Myr, although a more precise determination is not possible due to the uncertainties on the interstellar extinction. The comparison with the isochrones  shows that pre-main sequence stars having masses from 1.3 ~M$_\\odot$ to 2.0  ~M$_\\odot$ are present. Spitzer IRAC and MIPS observations of the region, made in the framework of the SAGE project reveals the presence of  young stellar object (YSO) candidates. YSO type I and II having ages $< 1 Myr$ are found at the same location than the candidate PMS stars. The data seems to suggest that the star formation in the region is a long lasting process where stars from 0.1 to 7 Myr are widely distributed, as expected in a turbulent scenario of star formation."
    },
    "0711/0711.3908_arXiv.txt": {
        "abstract": "We carry out a comprehensive joint analysis of high quality HST/ACS and Chandra measurements of A1689, from which we derive mass, temperature, X-ray emission and abundance profiles. The X-ray emission is smooth and symmetric, and the lensing mass is centrally concentrated indicating a relaxed cluster. Assuming hydrostatic equilibrium we deduce a 3D mass profile that agrees simultaneously with both the lensing and X-ray measurements. However, the projected temperature profile predicted with this 3D mass profile exceeds the observed temperature by $\\sim 30\\%$ at all radii, a level of discrepancy comparable to the level found for other relaxed clusters. This result may support recent suggestions from hydrodynamical simulations that denser, more X-ray luminous small-scale structure can bias observed temperature measurements downward at about the same ($\\sim 30\\%$) level. We determine the gas entropy at $0.1r_{\\rm vir}$ (where $r_{\\rm vir}$ is the virial radius) to be $\\sim 800$ keV cm$^2$, as expected for a high temperature cluster, but its profile at $>0.1r_{\\rm vir}$ has a power-law form with index $\\sim 0.8$, considerably shallower than the $\\sim 1.1$ index advocated by theoretical studies and simulations. Moreover, if a constant entropy ``floor'' exists at all, then it is within a small region in the inner core, $r<0.02r_{\\rm vir}$, in accord with previous theoretical studies of massive clusters. ",
        "introduction": "\\label{introduction}  As the largest gravitationally bound systems displaying a range of distinct observational phenomena, clusters of galaxies provide information of central importance in cosmology. Total gas and dark matter masses and their profiles provide significant insight into the formation and evolution of clusters and the relationship between baryons and dark matter. It has become increasingly clear that much more can be learned from careful comparisons of cluster observables including galaxy motions, gas properties from X-ray and Sunyaev-Zel'dovich (SZ) measurements, and mass profiles from lensing distortions and multiple images, providing new insight and tighter constraints on the dynamical state of a cluster and the nature of dark matter.  Chandra and XMM observatories have provided very detailed information on the physical state of the ubiquitous X-ray emitting plasma found in clusters. It has become clear that, broadly speaking, there are two classes of clusters, those showing some evidence of interaction, evidenced by complex structures in the gas, and those for which the gas emission is symmetric with a smooth radial variation in temperature and abundance, indicating the gas is probably relaxed and in a hydrostatic equilibrium with the overall gravitational potential. For lensing work the HST/ACS allows the inner caustic structure and central mass distribution to be examined in detail (Gavazzi et al.\\ 2002; Broadhurst et al.\\ 2005a; Sharon et al. 2005). The wide field imagers such as on the Subaru and CFHT telescopes permit a statistically significant detection of weak lensing distortion and magnification effects on the background galaxies to be traced out to the outskirts of the cluster (Gavazzi et al.\\ 2004; Kneib et al.\\ 2005; Broadhurst et al.\\ 2005b).  In the case of interacting clusters it has proved very interesting to compare their lensing-based mass distribution with their disturbed gas distribution. Clear evidence has emerged in the most favorable case of 1E0657-56 (the \"bullet cluster\") that two massive clusters have recently collided in the plane of the sky with a high relative velocity, leaving the gas lying in-between a bimodal distribution of galaxies and dark matter, traced by weak lensing (Markevitch et al.\\ 2002; Clowe, Gonzalez \\& Markevitch 2004; Bradac et al.\\ 2006). This particular case demonstrates that the bulk of the mass is dark and relatively collisionless, as anticipated in CDM dominated cosmogonies (Markevich et al.\\ 2004, Clowe et al.\\ 2006, Milosavljevic et al.\\ 2007, Randall et al.\\ 2007), although the estimated relative velocity between the two massive components may be exceptional in the context of $\\Lambda$CDM simulations (Hayashi \\& White 2006). Other such examples are coming to light, with collisions closer to the line of sight (Czoske et al.\\ 2002, Jee et al.\\ 2007, Dupke et al.\\ 2007).  A detailed lensing and X-ray study of a larger sample of interacting clusters by Okabe \\& Umetsu (2007) spans the full range of dynamical interaction, from premergers where the gas is clearly unaffected by the mutual gravitational attraction, to cases where both mass components are still readily distinguishable but the gas heavily disrupted and shock heated, and finally the postmerger phase where the gas shows only local signs of interaction, with a relatively small degree of substructure visible in the dark matter as probed by lensing.  Strong lensing based masses of the central regions of massive clusters have often been significantly higher than central masses deduced from X-ray analysis, by factors of $\\sim 2-4$ (Miralda-Escude \\& Babul 1995, Wu \\& Fang 1997; Allen 1998, Wu et al.\\ 1998, Voigt \\& Fabian 2006). Obvious reasons for the different masses - in addition to modeling and intrinsic observational uncertainties - include possible deviations from hydrostatic equilibrium, sphericity and gas isothermality (see Allen 1998). Lensing estimates are naturally biased upward in the case where the gravitational potential is significantly elliptical with the major axis preferentially aligned towards the observer. Evaluation of this effect in the context of $\\Lambda$CDM simulations shows that the inferred concentration of the cluster profile, measured in terms of the Navarro, Frenk \\& White (1996, hereafter NFW) model, can be enhanced by up to 20\\% by this bias in the worst cases (Hannawi et al.\\ 2005, Oguri et al.\\ 2005), falling well short of the reported discrepancies.  Motivated by the need to improve the precision of measurements of gas and DM profiles of clusters, we have begun a program of joint X-ray, strong and weak lensing analyses of several clusters for which we can combine high quality resolved observations. The advantages of a simultaneous analysis of X-ray and lensing data are clear, given that strong lensing measurements yield the total mass profile in the inner cluster core while X-ray and weak lensing measurements cover a much larger region of the cluster. The increasing quality and degree of detail of such data allows a more model-independent determination of the relevant profiles along with redundancy so that self-consistency can be checked. Under the assumption of spherical symmetry and hydrostatic equilibrium the projected temperature and gas density profiles as well as the total surface mass density profile may now be derived directly without resorting to assumed models or simple parameterizations of the profiles.  Here we apply a model independent approach to derive the density profile of a relaxed cluster from a simultaneous fit to both the X-ray and lensing data.  We apply our technique to A1689 ($r_{vir} \\sim 2$ h$^{-1}$ Mpc), a rich and moderately-distant ($z=0.183$) cluster that has been extensively observed in the optical, near IR and X-ray regions, and is the first cluster in our sample.  The cluster has a cD galaxy whose center is within $\\sim 1.5\"$ of the X-ray centroid; this fact, and the low degree of X-ray ellipticity ($\\epsilon\\simeq 0.08$, Xue \\& Wu 2002, hereafter XW02) indicate that the cluster is likely to be well relaxed and nearly spherical.  Previous estimates of its mass were obtained from the analysis of observed arcs and arclets produced by strong lensing (Broadhurst et al.\\ 2005a), from the distortion of the background galaxy luminosity function and number density (Taylor et al.\\ 1998, Dye et al.\\ 2001), and from weak lensing observations (Broadhurst et al.\\ 2005a,b, Medezinski et al.\\ 2007).  Here we derive the X-ray surface brightness and temperature profiles from an analysis of the full set of Chandra observations, which can be compared with a smaller subset of Chandra observations analyzed by XW02, and an independent study based on XMM (Andersson \\& Madejski 2004; hereafter AM04).  The X-ray and lensing observations and data reduction are described in Section~2, followed by a detailed account of the spectral and spatial data analysis in Section~3. In Section~4 we describe the methodology of deriving the gas and mass profiles, and in Section~5 we present the results of our deduced gas, total mass, and entropy profiles. Our results are discussed and assessed in Section~6. ",
        "conclusions": "The increasing quality of X-ray and optical imaging data motivates renewed and more thorough examinations of the physical nature of galaxy clusters as revealed by the very different processes of bremsstrahlung radiation and the gravitational deflection of light. In this paper we have examined the apparently relaxed cluster A1689, where only minimal substructure is evident from the dark matter, galaxy, and X-ray distributions and where the X-ray emission is smooth and symmetric. The longstanding claims of discrepancies in the total cluster mass estimated from these different kinds of analysis can now be investigated with greater precision, fewer assumptions, and in a more model-independent way.  There is a strong incentive to describe in analytical form the general nature of the total mass, gas, and temperature profiles. These quantities are inextricably bound up with the process of structure formation, including the nature of dark matter and the cooling history of the gas including interaction and merging of substructure. Thus, having an analytical form of the profiles is very useful.  The most commonly examined profile for the dark matter, based on N-body simulations of collisionless dark matter, is the NFW profile. In contrast the most commonly fitted gas mass density or surface brightness profile, the $\\beta$ model, is essentially empirically based. The temperature profile is often derived from the analytical expression of the gas density using the polytropic equation of state, or the gas is taken to be isothermal since the temperature gradient was undetectable with the old X-ray satellites. This way of analyzing the cluster can lead to substantial errors in the DM density, gas density, and temperature values and the derived quantities such as the overall mass.  As we discussed in the Introduction, the total mass profile can be independently determined from lensing measurements, or from X-ray measurements of the gas density and temperature profiles based on the assumption of hydrostatic equilibrium. (Clearly, the latter can be done also with spatially resolved S-Z measurements, when available.) The second basic assumption adopted in our analysis is spherical symmetry. Obviously, elongation along the line of sight is possible, but - as we have mentioned in the Introduction - this typically can introduce a 20\\% bias in the mass estimate. For a detailed discussion of the impact of triaxial cluster morphology, see Gavazzi (2005).  We have suggested here a model-independent method which uses free parameters and does not assume a specific profile for any mass component. The only hidden assumption is a linear (in log-log) interpolation between the free parameters; this just assumes a reasonable degree of smoothness in the profiles. The profiles are also extrapolated beyond the data, but we showed that the results are insensitive to the detailed extrapolation. We specifically used a simple power law extrapolation for both the total mass density and for the gas mass density, with the same power law index assumed in order to approach a constant gas fraction at large radii. Our model-independent method is the best way to obtain the values of the important parameters of clusters at radii where there are good quality data. Our model-independent method can and will be applied to joint analyses of measurements of other clusters.   Within the data range, the highest value of the 3D temperature (fig.~\\ref{T3D profile}) is 2--3 times higher than the lowest value. Specifically, the temperature profile peaks at $r\\sim 0.1r_{500}$ and declines at both smaller and larger radii. The denser environment at the center of a cluster should naturally cause a decrease in the temperature there. We indeed find in this cluster that the 2D temperature decreases towards the center of the cluster, at $r< 50-100$ $h^{-1}$ kpc (figure~\\ref{T2D comparison}). This can be due to several mechanisms. First, the high gas number density in the center of the cluster causes a rapid loss of energy and a decrease in the temperature. As pressure support decreases, the gas should gravitate towards the bottom of the cluster potential well, i.e., the center of the cluster, in a so-called \"cooling flow\" (Fabian 1994).  A cooling flow should manifest in a detectable surface brightness enhancement in the X-ray emission from the central region of a cluster, $\\sim 100-200$ kpc. Many studies have failed to find these cooling flows (MME99), at least not as expected in a simple or a standard cooling flow model (Fabian et al.\\ 2001; Peterson et al.\\ 2001). An explanation for this failure can be a more complicated mechanism which maybe includes reheating of the gas (e.g., by energetic particles; Rephaeli \\& Silk 1995), mixing, differential absorption, efficient conversion of cooling flow gas into low-mass stars, an inhomogeneous metallicity distribution (which is not the situation in A1689, at least not in terms of a radial gradient: fig.~\\ref{Abundance gradient}), or disruption of cooling flows by a recent subcluster merger (see discussion in Fabian et al.\\ 2001; Peterson et al.\\ 2001).   There are uncertainties about the criteria for finding cooling flows, and even with fixed criteria the decision whether cooling flows exist is still data dependent.  At times the \"solution\" has been to exclude the inner, problematic region.  Finding the value of the gas mass density in the model-independent approach bypasses these two problems, the unknown physical processes and the data dependence, as it allows for an analysis with fewer prior assumptions. Another way to reduce the core temperature is through the effect of the cD galaxy that is often anchored at the cluster center.  The cD galaxy has a lower temperature and is denser than the cluster so it can in principle be the main cause for the low temperature at the center.  However, based on our our initial, limited study it seems unlikely that the cD has a major effect in A1689, since the emissivity of the cD dominates that of the cluster only below $r=4$ h$^{-1}$ kpc (fig.~\\ref{emissivity cD Vs cluster}).  We have found a good fit to the observed 2D profiles of the lensing surface density and the X-ray surface brightness. Due to the smoothness of the profiles, we were able to fit the data with a relatively small number of free parameters. We have found a good agreement with previous results for all the parameters we have checked for A1689, including gas mass density, gas fraction profiles, measured temperature, adiabatic index, and abundance (except compared to AM04). The total mass density profile we obtained was essentially determined directly by the lensing data alone, with $\\lesssim 1\\%$ differences introduced by using the X-ray data as well. This is the case since we did not use the X-ray temperature as part of the fit.  We have shown that it is possible to obtain a model independent 3D mass profile for which very good agreement is found between the lensing mass profile and the X-ray emission profile. However, there is still a discrepancy between the temperature derived directly from (X-ray) measurements, and that deduced from a solution to the HEE, with the latter $\\sim30\\%$ higher at all radii, as has already been determined by Mazzotta et al.(2004) and Rasia et al. (2005). Other phenomena, such as bulk motions, turbulence, and nonthermal degree of freedom may contribute to the pressure, and thus reduce the temperature with respect to the one obtained from assuming that thermal gas pressure is the only contributor. Specifically, Faltenbacher et al. (2005) claim that based on their high-resolution cosmological simulations - in which they identified and analyzed eight clusters at $z=0$ - about 10\\% of the total pressure support may be contributed by random gas bulk motions, which may affect temperature by up to 20\\%. It has also been suggested that X-ray luminous clumps of relatively low temperature may bias projected temperature measurements downward (Kawahara et al.\\ 2006). In any case, as we have already mentioned this temperature discrepancy is much smaller than the mass discrepancy by a factor of $2-4$ encountered in previous (separate X-ray and lensing) analyses.  The derived entropy profile for A1689 has a power law form with no obvious central flattening, at least at $r>0.02r_{\\rm vir}$, as expected for massive clusters from theoretical work (Tozzi \\& Norman 2001). The existence and the level of the entropy floor is still not completely clear in simulations, since different simulation methods give different results (Voit, Kay, \\& Bryan 2005). This is probably due to the limited resolution of simulations which is especially critical in the core of the cluster. Upcoming simulations with spatial resolution of $\\sim 5$ kpc should give a better understanding of the entropy floor. Our derived radial slope is around $0.8$ rather than $1.1$, flatter than the prediction. Since using simple parametrized models for the DM and the gas gave a power law index of $1.1$, we believe the value of $0.8$ might be a more accurate result, which might be a reflection of a more complex gas dynamics and preheating history. Indeed, it has been suggested from simulations that preheating decreases the power law index in the region of our data (Borgani et al.\\ 2005)."
    },
    "0711/0711.4469_arXiv.txt": {
        "abstract": "In this paper, we study the possibility of model building of cyclic universe with Quintom matter in the framework of Loop Quantum Cosmology. After a general demonstration, we provide two examples, one with double-fluid and another double-scalar field, to show how such a scenario is obtained. Analytical and numerical calculations are both presented in the paper. ",
        "introduction": "Quintom model of dark energy \\cite{Feng:2004ad} has a salient feature that its equation-of-state (EoS) crosses smoothly  over the cosmological constant barrier $w = -1$. Phenomenologically, this kind of model is mildly favored by the current observational data fitting \\cite{Zhao:2006qg}, but theoretically, the model building of Quintom dark energy is a challenge due to the No-Go theorem \\cite{Xia:2007km} (also see Refs. \\cite{Feng:2004ad, Vikman:2004dc, Hu:2004kh, Caldwell:2005ai, Zhao:2005vj, Kunz:2006wc}). This No-Go theorem forbids a traditional scalar field model with Lagrangian of the general form ${\\cal L} = {\\cal L}(\\phi,\\partial_{\\mu}\\phi\\partial^{\\mu}\\phi)$ from having its EoS across the cosmological constant boundary. Consequently, in order to realize a viable Quintom model in the framework of Einstein's gravity theory, we need to add extra degrees of freedom to the conventional single field theory. The simplest Quintom model involves two scalars with one being Quintessence-like and another Phantom-like \\cite{Feng:2004ad, Guo:2004fq}. This model has been studied in detail in \\cite{Quintom_tf, Cai:2006dm}. In recent years there have been a lot of theoretical studies of Quintom-like models. For example, motivated from string theory, the authors of Ref. \\cite{Cai:2007gs} realized a Quintom scenario by considering the non-perturbative effects of a DBI action. Moreover, there are models which involve higher derivative terms for a single scalar field \\cite{Li:2005fm}, models with vector field \\cite{ArmendarizPicon:2004pm}, making use of an extended theory of gravity \\cite{Cai:2005ie}, non-local string field theory \\cite{Aref'eva:2005fu}, and others (see e.g. \\cite{Quintom_1, Onemli:2004mb, Quintom_others}). Due to this property, Quintom scenario has some interesting applications in cosmology. For example, a recent study has shown that a universe dominated by Quintom matter can provide a bouncing cosmology which allows us to avoid the problem of the initial singularity. Quintom Bounce, proposed in \\cite{Cai:2007qw} with its perturbation developed in \\cite{Cai:2007zv}, takes place in the early time of the universe when the corresponding energy scale becomes very high. Therefore it may be accompanied by quantum effects of gravity.    One possible approach to investigating quantum gravity effects in the early universe is Loop Quantum Cosmology (LQC) \\cite{LQC} which is based on the theory of Loop Quantum Gravity (LQG) \\cite{LQG}. Beyond the usual knowledge of spacetime, LQC predicts that the underlying geometry is discrete and the big bang singularity is replaced by a big bounce when the energy density of the universe approaches the Planck scale \\cite{Ashtekar-old, Ashtekar-improved}. Therefore, the quantum geometry plays a significant role on determining the evolution of the universe in the early time. Interestingly, one notes that effects of loop quantum gravity also affect the evolution of Quintom universe in late time \\cite{Zhang:2007an}, since the energy density of phantom component usually increases during the expansion of the universe. In this paper we extend the idea of Quintom Bounce by considering a universe filled with Quintom matter in the frame of LQC. We interestingly find that the universe possesses an exactly cyclic evolution such that its scale factor undergoes contracting and expanding periodically.   In the context of LQC there exist two distinct quantum geometry modifications, namely the inverse volume modification \\cite{Quantization-ambiguity, inversevolume} and the quadratic density correction \\cite{Ashtekar-improved}. Based on the inverse volume modification authors of Ref. \\cite{oscillating-universe} obtained an cosmological solution with oscillating behavior for the closed universe(see Ref. \\cite{loop-branecyc} for cyclic universe in flat model). The inverse volume modification appears by choosing a large SU(2) representation for the holonomy in the matter part of Hamiltonian constraint. As pointed out in \\cite{inversevolume}, it is natural to take the same spin representation for the gravitational sector as for matter sector in the whole Hamiltonian constraint. In this paper, we work with the fundamental spin representation both for the matter sector and gravity sector. So, the quantum geometry modification only comes from the quadratic density correction which arises from the ``minimal area gap\" in LQG. In this paper we consider quadratic density correction of LQC for the spatially flat universe.   This paper is organized as follows. In Sec. II, we study the possibility of building the models of the cyclic universe. Firstly, we will review briefly on LQC, then based on the modified Friedmann equation from LQC we will present the general picture of the cyclic universe, followed by two examples. The last section is the discussion and conclusion. ",
        "conclusions": "In this paper, we have studied an application of Quintom matter in the framework of LQC. Our results show that due to the specific property of Quintom and with the help of modified Friedmann equation, a scenario of cyclic universe can be realized naturally.  Before conclusion we should point out that:  1) in the scenario of cyclic universe with Quintom matter in LQC, there could be two possibilities causing the turnaround or the bounce, one being $\\rho=\\rho_c$ as studied in this paper, the other being $\\rho=0$ from Eq.(\\ref{b2}). If the energy density reaches zero during the evolution, the turnaround or bounce can be realized as well which is pointed out in Ref. \\cite{Cai:2007qw} (also see Ref. \\cite{Brown:2004cs}). However, this possibility will not happen in the examples studied in this paper. For the double-fluid model, we can see from the formula that the energy density cannot approach zero; while for the double-field model, the ``ghost\" field $\\psi$ has a minus mass squared term, which prohibits the abnormal kinetic term to be so large to drive the total energy density to reach zero.  A cyclic universe is a non-standard scenario of cosmology which, by having the scale factor oscillating and the universe expanding and contracting alternately, is expected to solve the Big-Bang singularity and coincidence problem. In the literature there have been many discussions on such a topic and a number of models have been proposed, among which there are cyclic models in braneworld scenario \\cite{Steinhardt:2001vw,Piao:2004hr,Piao:2004me}, cyclic cosmologies with spinor matter \\cite{ArmendarizPicon:2003qk}, closed oscillating universe \\cite{oscillating-universe, Clifton:2007tn}, cyclic braneworld in LQC \\cite{loop-branecyc, singh-bounce}, and see Refs. \\cite{Zhang:2007an, Saridakis:2007cf} for recent developments, and so on. The difference of our model from these models is that our model is in the framework of 4-dimensional spatially flat FRW universe with LQC.  2)  our model provides a qualitative picture of cyclic evolution of the universe. In this paper we have not included the radiation and matter, however, the general behavior of our solution will not be changed since the epochs for radiation and matter only last for a short time in one period during the whole evolution. There are some open issues in the scenario of cyclic universe, such as the entropy problem \\cite{Turok:2004yx, Baum:2006nz}, blackhole problem \\cite{Steinhardt:2004gk, Brown:2004cs, Zhang:2007yu}, the problem of causality \\cite{Turok:2004yx} and so on, which is still in discussion in the literature."
    },
    "0711/0711.3555_arXiv.txt": {
        "abstract": "Presuming weak collisional interactions to exchange the kinetic energy between dark matter and baryonic matter in a galaxy cluster, we re-examine the effectiveness of this process in several `cooling flow' galaxy clusters using available X-ray observations and infer an upper limit on the heavy dark matter particle (DMP)$-$proton cross section $\\sigma_{\\rm xp}$. With a relative collisional velocity $V-$dependent power-law form of $\\sigma_{\\rm xp}=\\sigma_0(V/10^3\\ {\\rm km\\ s^{-1}})^a$ where $a\\leq 0$, our inferred upper limit is $\\sigma_0/m_{\\rm x}\\lsim 2\\times10^{-25}\\ {\\rm cm}^2\\ {\\rm GeV}^{-1}$ with $m_{\\rm x}$ being the DMP mass. Based on a simple stability analysis of the thermal energy balance equation, we argue that the mechanism of DMP$-$baryon collisional interactions is unlikely to be a stable nongravitational heating source of intracluster medium (ICM) in inner core regions of `cooling flow' galaxy clusters. ",
        "introduction": "Astrophysical and cosmological measurements together with numerical simulation experiments indicate that the cold dark matter (CDM) constitutes most of the matter in the Universe, even though the fundamental physical nature of such dark matter particles (DMPs) remains unknown. In the simplest scenario, these cold DMPs are presumed to be collisionless and they interact with each other or with other baryons only through the mutual gravity at the present epoch. While the collisionless CDM model is successful in explaining the formation of large-scale structures in the Universe, observational contradictions with numerical simulations appear inevitable regarding structures on sub-cluster scales, e.g., the prediction of a higher number of dwarf galaxies than that actually observed. Spergel \\& Steinhardt (2000) revitalized the concept of strongly interacting massive particles (SIMPs) to confront this issue and suggested a self-interacting cross section per unit DMP mass $\\sigma_{\\rm xx}/m_{\\rm x}\\sim 10^{-24}\\ {\\rm cm}^2\\ {\\rm GeV}^{-1}$ where $\\sigma_{\\rm xx}$ is the collisional cross section among DMPs and $m_{\\rm x}$ is the DMP mass. Along this line of reasoning, if DMPs are strongly self-interacting, then similar strong interactions would also be equally expected to exist between DMPs and baryons. Elastic scatterings between DMPs with nuclei could generate a recoil energy of the order of $\\sim 10$ keV, which would then be detectable in underground and underwater particle experiments. Such `direct detection' experiments currently limit the cross section to the order of $10^{-42}-10^{-40}$ cm$^2$ for a cold DMP of mass in the range of $\\sim 10-10^3$ GeV (e.g., Akerib et al. 2004). Limits on such collisional interactions in a wider mass range were investigated by various physical experiments and observations of several astrophysical processes, such as $\\beta\\beta$ decays, cosmic-ray detections, the galactic-halo stability, the cooling of molecular clouds, proton decay experiments, the existence of old neutron stars and the Earth (e.g., Starkman et al. 1990), satellite experiments in space (e.g., Wandelt et al. 2001), primordial nucleosynthesis and cosmic rays (e.g., Cyburt et al. 2002), and cosmic microwave background anisotropy and large-scale structure power spectrum (e.g., Chen et al. 2002). These experiments and observations have provided a limit on the mass-dependent cross section per unit DMP mass between DMPs and protons as $\\sigma_{\\rm xp}/m_{\\rm x}\\lesssim10^{-26}-10^{-24}$ cm$^2$ GeV$^{-1}$ for a DMP mass $m_{\\rm x}\\gtrsim 1$ GeV.  Recent high-resolution observations of the hot X-ray emitting gaseous intracluster medium (ICM) by {\\it Chandra} (e.g., Peterson et al. 2001) and {\\it XMM-Newton} (e.g., Kaastra et al. 2001; Tamura et al. 2001) satellites have revealed deficits of cool gases (with gas temperatures much less than the virial temperature $T_{\\rm vir}$) in the core of so-called `cooling flow' galaxy clusters, inconsistent with predictions of the conventional radiative cooling models (e.g., Cowie \\& Binney 1977; Fabian \\& Nulsen 1977; Mathews \\& Bregman 1978; Stewart et al. 1984; Fabian 1994 and extensive references therein). Several heating mechanisms have been proposed to resolve this `cooling flow dilemma', such as the inward thermal conduction from hot outer regions (e.g., Narayan \\& Medvedev 2001 and references therein), energy injections associated with central activities by an active galactic nucleus (e.g., Churazov et al. 2002), and outward or inward acoustic wave heating (e.g., Pringle 1989; Fabian et al. 2003a, b; Fujita et al. 2004; Feng, Zhang, Lou \\& Li 2004). The resonant excitation of internal gravity modes ($g-$modes) in the ICM by orbiting galaxies was explored by Balbus \\& Soker (1990); they examined the processes of excitation, propagation, amplification, damping of such galaxy cluster $g-$modes in the context of providing thermal energy in `cooling flow' galaxy clusters.  Qin \\& Wu (2001) proposed collisional interactions between baryons and heavy DMPs ($m_{\\rm x}\\gg m_{\\rm p}$ with $m_{\\rm p}$ being the proton mass) as a major nongravitational heating mechanism for ICM in the core of a galaxy cluster. Assuming DMPs and baryons have comparable velocity dispersions, the kinetic energy of a single DMP would then be much larger than that of a baryon because of the presumed mass difference $m_{\\rm x}\\gg m_{\\rm p}$. Therefore in numerous elastic collisions, kinetic energies of DMPs can be systematically transferred to the baryonic ICM to balance the radiative cooling in X-ray bands by hot electrons. By equating heating and cooling rates, they estimated a specific cross section of $\\sigma_{\\rm xp}/m_{\\rm x}\\sim 10^{-25}$ cm$^2$ GeV$^{-1}$ for $m_{\\rm x}/m_{\\rm p}>10^5$. This is actually a requirement of compensating the radiation cooling and should not be regarded as a limit of any sort. It should also be noted that they took strong lensing cluster CL0024+1654 to infer the DMP-baryon interaction cross section for balancing the cooling, which may not be proper for the following reason. The significant discrepancy by a factor of $3\\sim 4$ between the mass profiles derived from X-ray observations and gravitational lensing effects shows that the ICM may not be in a static equilibrium and might be still collapsing (e.g., Kneib et al. 2003; Zhang et al. 2005), and the bimodal velocity distribution of cluster galaxies indicates a merger of two systems with a mass ratio of $1$ to $2$ (e.g., Czoske et al. 2002). The main heating mechanism of CL0024+1654 could well be the gravitational collapse.  Chuzhoy \\& Nusser (2006) re-considered the ICM heating scenario of Qin \\& Wu (2001), corrected their calculations and derived a similar cross section for the heavy DMP-proton elastic collisional interaction. They found that, if $\\sigma_{\\rm xp}$ is independent of the relative velocity $V$ of colliding particles, a thermal equilibrium state between heating by DMPs and radiative cooling by hot electrons of ICM would be always unstable. However, in galaxy clusters with $T>2$ keV, a stable energetic balance may be achieved for a relative velocity $V-$dependent cross section $\\sigma_{\\rm xp}\\propto V^a$ with $a\\lesssim -3$.  There are two major simplifications in both analyses of Qin \\& Wu (2001) and Chuzhoy \\& Nusser (2006). First, they took the temperature and density of the ICM and the density of DMPs for typical values, rather than the relevant distributions determined by X-ray observations with high angular and spectral resolutions. Secondly, they estimated the velocity dispersion of DMPs either simply similar to that of baryons or by the results of numerical simulations, rather than a dynamically self-consistent `true' value obtained by solving the Jeans equation (Binney \\& Tremaine 1987; Subramanian 2000; see also Ikebe et al. 2004 for the case of galaxy cluster A1795). Since the specific cross section limit they derived (at the centre or at the virial radius) is highly sensitive to the chosen parameters [see equation (8) of Qin \\& Wu (2001)], it is crucial to investigate this ICM heating mechanism more carefully using an actual sample of `cooling flow' galaxy clusters, well observed in X-ray bands by the {\\it XMM-Newton} and {\\it Chandra} satellites in space.  This paper is structured as follows. In \\S 2, we set the upper limits for DMP-proton elastic collisional cross section per unit DMP mass. In \\S 3, we demonstrate that the DMP-proton collisional interaction alone is unlikely to be a stable ICM heating mechanism to compensate the radiative cooling of ICM. Discussion and conclusions are contained in \\S 4. Details on the X-ray cooling function can be found in Appendix A.  In our theoretical model consideration, we have adopted the currently favoured standard $\\Lambda$CDM cosmology with the cosmological parameters $h=0.7$, $\\Omega_{\\rm m}=0.3$ and $\\Omega_\\Lambda=0.7$ in conventional notations. ",
        "conclusions": "Based on two different elastic collision models and X-ray observations of galaxy clusters, we have re-examined effects of collisional interactions between heavy DMPs and protons in `cooling flow' clusters of galaxies and estimated a more reliable upper limit for the specific cross section between DMPs and protons by solving the equation of energy conservation for five ``cooling flow\" galaxy clusters using available X-ray data. In the regime of $m_{\\rm x}\\gg m_{\\rm p}$, the upper limit for $\\sigma_{\\rm xp}=\\sigma_0(V/10^3\\ {\\rm km\\ s^{-1}})^a$ with $a\\leq0$ can be expressed as $\\sigma_0/m_{\\rm x}\\lsim 2\\times 10^{-25}$ cm$^2$ GeV$^{-1}$, which is fully consistent with the earlier results. Similar to other astrophysical constraints, this upper limit is independent of the underlying model for particle physics governing collisional interactions.  In our model calculations, we have assumed that the total gravitational potential $\\phi$ remains invariable during the process of galaxy cluster evolution. The gravitational potential $\\phi$ of a galaxy cluster is dominantly determined by the amount of DMPs. Besides the weak energy-momentum transfer between the DMPs and baryons, the density distribution of dark matter is only affected by the mild accumulation of cool gas in the central region of clusters. The exact solution of such complex co-evolution of DMPs and baryons should be treated with Fokker-Planck equation (e.g., Binney \\& Tremaine 1987). We will pursue the relevant problems in contexts of galaxy clusters in separate papers.  We have also demonstrated an intrinsic instability in explaining the problem of `cooling flow' galaxy clusters via the heat transfer by heavy DMP-proton collisional scattering. According to expression (\\ref{transfer1}), the DMP-proton collisional heating of ICM works only when $m_{\\rm p}v_{\\rm p}^2<m_{\\rm x}v_{\\rm x}^2$. In `cooling flow' galaxy clusters, $v_{\\rm p}^2/v_{\\rm x}^2\\sim\\mu\\simeq 0.6$ (Hu \\& Lou 2007) and thus $m_{\\rm x}>0.6$ GeV is required in this model. However, as the collisional cross section in the DMP mass range 0.5 GeV$<m_{\\rm x}<10^5$ GeV is less than $5\\times10^{-28}$ cm$^2$ (e.g., Wandelt et al. 2001), the light DMP collisional heating mechanism can then be completely ignored. Therefore, we rule out the possibility of DMP collisional scattering as a major stable non-gravitational heating sources of the ICM in `cooling flow' galaxy clusters."
    },
    "0711/0711.1550_arXiv.txt": {
        "abstract": "Type IIn supernovae (SNe IIn) dominate the brightest supernova events in observed FUV flux ($\\sim$$1200-2000$\\AA).  We show that multi-band, multi-epoch optical surveys complete to m$_{r}=27$ can detect the FUV emission of $\\sim$25 $z>2$ SNe IIn deg$^{-2}$ yr$^{-1}$ rest-frame ($\\sim$10 SNe IIn deg$^{-2}$ yr$^{-1}$ observed-frame) to $4\\sigma$ using a technique monitoring color-selected galaxies.  Moreover, the strength and evolution of the bright emission lines observed in low-redshift SNe IIn imply that the \\lya emission features in $\\sim$70\\% of $z>2$ SNe IIn are above 8m-class telescope spectroscopic thresholds for $\\sim$2 yr rest-frame.  As a result, existing facilities have the capability to both photometrically detect and spectroscopically confirm $z>2$ SNe IIn and pave the way for efficient searches by future 8m-class survey and 30m-class telescopes.  The method presented here uses the sensitivities and wide-field capabilities of current optical instruments and exploits (1) the efficiency of $z>2$ galaxy color-selection techniques, (2) the intrinsic brightness distribution ($\\langle M_B\\rangle=-19.0\\pm0.9$) and blue profile of SNe IIn continua, (3) the presence of extremely bright, long-lived emission features, and (4) the potential to detect blueshifted SNe Ly-$\\alpha$ emission shortward of host galaxy Ly-$\\alpha$ features. ",
        "introduction": "%  Type II supernovae (SNe) are explosive, highly luminous events that result from the core collapse of massive stars after exhausting their nuclear fuel.  Observations of the sites of Type II SNe have identified the progenitors of Type II-P, IIn, and SN 1987A-type events [see \\citet{gal07}, references therein, and \\citet{li07}].  Because these SN types are linked to massive stars, they are sensitive probes of star formation and fundamental tools for the study of galaxy formation and evolution.  In this paper, we focus on the Type IIn SNe [SNe IIn; \\citet{sch90,f97}] identified by the very bright, narrow emission lines observed in their spectra.  The emission lines are believed to be caused in part by the strong interaction of the SN ejecta with dense circumstellar material\\footnote{Spectroscopic observations reveal that the emission lines are often comprised of multiple components; a narrow feature (FWHM $\\lesssim200$ km s$^{-1}$) superposed on an intermediate (FWHM $\\sim$1000-2000 km s$^{-1}$) or sometimes broad base [FWHM $\\sim$5000-10000 km s$^{-1}$].  A discussion of the mechanism behind SNe IIn emission-line profiles is beyond the scope of this paper.  The aim here is the detection and confirmation of high redshift SNe IIn via their continua and emission-line flux.}.  SNe IIn exhibit very blue continua (\\S~\\ref{detection}), and although the continua fade over a period of weeks, the emission lines are observed to remain bright for years.  The luminosities of most types of SNe near maximum brightness provide rare opportunities to study galactic and individual stellar processes out to very high redshift.  However, the properties and extinction of most SNe, and capabilities of current facilities, have prevented detections beyond $z>2$.  For example, although Type Ia SNe are intrinsically bright, they have high UV extinction that makes $z>2$ detections in deep optical imaging surveys difficult.  In addition, the deep wide-field IR surveys necessary to find the brighter rest-frame optical flux of such events are difficult to perform.  Type II SNe are intrinsically fainter than Type Ia SNe as a whole, but a subset of the Type IIL and IIP SNe show less UV extinction and are bright enough to be detected in deep wide-field optical surveys using current facilities.  However once detected, the short outburst duration and relative faintness of their continua make a reliable spectroscopic confirmation prohibitive.  This is true for Type Ia detections as well.  Therefore the identification and classification of these SNe at high redshift must use methods such as photometric data and model template fitting.  We show that SNe IIn are an exception.  The intrinsic brightness and blue continua of SNe IIn not only enable photometric detection at $z>2$ in deep optical surveys but their bright long-lived emission lines provide a means for spectroscopic confirmation.  Because their properties enable detection out to $z>2$, they have great utility as a diagnostic tool to probe cosmic scales.  The progenitors of SNe IIn are believed to be luminous blue variable stars (LBVs) with main sequence masses of $\\gtrsim80M_{\\odot}$ [\\citet{gal07,vd06}].  Detonations of these short-lived, high mass stars are direct indicators of recent star formation and can be used to trace the universal star formation rate (SFR) out to high redshift. However, to do this accurately, the dust extinction of SNe IIn as compared to the entire SN population needs to be quantified. Detections at $z\\ge2$ would extend the measurements of the universal supernova rate (SNR) [e.g., \\citet{dahlen04,calura06}], help confirm or refute observed trends with redshift, and improve subsequent measurements that rely on these values.  As the progenitor mass range becomes more refined, SNe IIn detections can provide a constraint on the upper-end of the initial mass function for targeted $z>2$ galaxy populations with an assumption of little mass evolution with redshift. In addition, SNe IIn have the potential to provide an independent statistical discriminator between the SFRs and star formation histories of $z>2$ galaxies modeled from photometric data and stellar synthesis codes [e.g., \\citet{aes03,pap01}].  Because these are core-collapse SNe, their density at high redshift will place constraints on the relic neutrino count [e.g.,\\citet{f04,ando04,louis04}], important for future facilities such as Hyper-K \\citep{nak03} and UNO \\citep{jung00}.  Furthermore, a refinement of the kinetic energy output and fraction of SNe IIn to the entire SNe population will enable $z>2$ SNe IIn to act as an anchor to extrapolate the SN contribution to galactic-scale winds observed in high redshift galaxies [e.g., \\citet{a03,c05}] and help estimate mass loss.  In this paper, we present a photometric technique monitoring color-selected galaxies that can detect $z>2$ SNe IIn in reasonable numbers.  Moreover, we analyze the strength and evolution of bright optical and FUV emission lines observed in low-redshift SNe IIn and show that the emission lines of a large fraction of $z>2$ SNe IIn are above the thresholds of 8m-class telescopes for $\\sim$2 years rest-frame.  Consequently, SNe IIn (and for similar reasons Type IIa, see \\S~\\ref{contaminants}) are the only SNe that can be spectroscopically confirmed at $z>2$ using existing facilities.  The results of deep optical surveys for $z>2$ SNe IIn will carve out a new area of high redshift research and provide the framework for routine surveys using next generation facilities.  Deep wide-field imaging capabilities of future 8m-class survey telescopes and the imaging and spectroscopic sensitivities of 30m-class telescopes will not only greatly increase the density of SNe IIn detections but extend confirmed detections to higher redshifts (\\S~\\ref{predictions}).  A description of our photometric technique and tests using data from the Canada France Hawaii Telescope Legacy Survey (CFHTLS) are presented in \\S~\\ref{detection}.  An analysis of the observed SNe IIn spectroscopic features, \\lya behavior, and contaminants to SNe IIn surveys are discussed in \\S~\\ref{spec}.  We estimate in \\S~\\ref{expectations} the \\zrange SNe IIn SNR, the number of detectable SNe IIn for surveys complete to $m_{r}=27$, and the observational expectations for the CFHTLS.  We conclude with a brief summary in \\S~\\ref{summary}.  We adopt $\\Omega_{M}$=0.3, $\\Omega_{\\Lambda}$=0.7, $H_0$=70 km s$^{-1}$ Mpc$^{-1}$ cosmology.  All magnitudes, unless otherwise noted, are in the AB system \\citep{f96}.  Finally, we work in the rest-frame of the SNe but note that effect of time-dilation in the observed frame is important in detection expectations. ",
        "conclusions": "\\label{summary}%  We have presented a photometric and spectroscopic technique to detect $z>2$ SNe IIn in reasonable numbers using current technology. Although this technique can be applied to any redshift, we focus on SNe at $1.9<z<3.3$ because of (1) the sensitivities and wide-field capabilities of current optical instruments, (2) the success and efficiency of $z>2$ photometric selection techniques, and (3) the inherent properties of SNe IIn and $z>2$ galaxies.  SNe IIn have $\\langle M_B\\rangle=19.0\\pm0.9$ and very blue continua with an estimate FUV (1700\\AA) flux $\\sim1$ mag fainter at maximum brightness.  In addition, SNe IIn have bright emission features that make a non-negligible contribution to the total flux that can be used for detection and contaminant rejection.  We estimate that there exist $\\sim145^{+375}_{-20}$ SNe IIn deg$^{-2}$ yr$^{-1}$ rest-frame ($45^{+160}_{-6}$ observed-frame) in color-selected $1.9<z<3.3$ LBGs to m$_r=27$.  We show from our simulations using real data obtained from the CFHTLS Deep Synoptic Survey that deep, multi-band, multi-epoch wide-field imaging surveys complete to m$_{r}\\sim27$ can detect $\\sim$25 SNe IIn deg$^{-2}$ yr$^{-1}$ rest-frame ($\\sim$8 observed-frame) at $1.9<z<3.3$ to $\\ge4\\sigma$ flux uncertainty. These surveys can be performed independently or in conjunction with a targeted narrow- to intermediate-band deep imaging program. Contaminants, such as AGN, can be largely eliminated by the design of the photometric program, selection criteria, and by follow-up deep spectroscopy.  The general evolution of H-$\\alpha$ and Ly-$\\alpha$ emission lines at low redshift predicts that the Ly-$\\alpha$ features of $\\gtrsim85$\\% of $z>2$ SNe IIn are above the spectroscopic thresholds of 8m-class facilities for $\\sim2$ years rest-frame.  Therefore, sensitive low-resolution optical spectrographs on 8m-class telescopes have the capability to spectroscopically identify $1.9<z<3.3$ SNe IIn photometric candidates.  Positive confirmation of spectroscopic candidates can be obtained at subsequent epochs as the line flux evolves.  The blueshift ($\\sim1500-2500$ km s$^{-1}$) of the \\lya emission lines observed in low redshift SNe IIn would result in a clean separation from an LBG host galaxy Ly-$\\alpha$ emission and absorption if observed at $z>2$.  Under the assumption that this property is inherent to SNe IIn, the findings of \\citet{aes03} and the analysis of the LBG spectroscopic sample of \\citet{c05} result in an estimate that the Ly-$\\alpha$ emission lines from SNe IIn can escape $\\gtrsim$90\\% of all $z>2$ LBGs.  Combined with the strength of SNe IIn \\lya emission, this results in the detectability of $\\sim$70\\% of all $z>2$ SNe IIn.  To pursue the technique presented here, we have initiated a program to better characterize SNe IIn features at low to intermediate redshift using the {\\it Hubble Space Telescope} and a complete program to search for $1.9<z<3.3$ SNe IIn using the CFHTLS Deep Synoptic Survey data.  Analyzing seasonal stacked images of the CFHTLS Deep Synoptic survey, we estimate the photometric detection of $\\sim4^{+6}_{-1}$ SNe deg$^{-2}$ season$^{-1}$ in the observed frame, or $\\sim60-200$ and spectroscopic confirmation of $\\sim40-140$ \\zrange SNe IIn in the four square-degree pointings over five years of the survey.  These, and future programs, will pave the way for routine surveys using the next generation of 8m-class survey and 30m-class telescopes that will detect a greater density of SNe IIn and potentially extend SNe IIn detection and Ly-$\\alpha$ confirmation to $z\\sim6$.  In addition, such surveys could confirm $0<z<2$ SNe IIn using MgII, H-$\\alpha$, and other bright lines.  The detection of $z>2$ SNe IIn is important to many areas of astrophysics, such as stellar physics and galaxy formation and evolution, and will become more so as we learn more about the mechanism behind SNe IIn and the details of their properties and environments."
    },
    "0711/0711.2553_arXiv.txt": {
        "abstract": "A detailed study of the lowest states $1s_0, 2p_{-1} ,2p_0$ of the hydrogen atom placed in a magnetic field $B\\in(0-4.414\\times 10^{13}\\,{\\rm G})$  and their  electromagnetic transitions ($1s_{0}   \\leftrightarrow 2p_{-1}$ and  $ 1s_{0}  \\leftrightarrow 2p_{0}$) is carried out in the Born Oppenheimer approximation. The variational method is used with a physically motivated  recipe to design simple trial functions applicable to the whole domain of magnetic fields. We show that the proposed functions yield very accurate results for the ionization (binding) energies. Dipole and oscillator strengths are  in good agreement with results by Ruder {\\em et al.}~\\cite{Ruderbook} although we observe deviations up to $\\sim 30\\%$   for the oscillator strength of the (linearly polarized)  electromagnetic transition $1s_{0}   \\leftrightarrow 2p_{0}$  at strong magnetic fields $B\\gtrsim 1000$\\,a.u. ",
        "introduction": "\\label{sec:intro}  Contemporary  X-ray space observatories, like Chandra, XMM-Newton and their predecessors, have  collected a considerable amount of observational data of the thermal emission  coming from surface layers of neutron stars which are characterized by enormous  magnetic fields  $B\\sim 10^{12}-10^{13}\\,$G  (see \\cite{Pavlov2001},\\cite{Zavlin:2007}). In particular, the observation of  absorption features in the X-ray spectrum of some isolated neutron stars (see {\\em e.g.} ~\\cite{Sanwal:2002,Kerkwijk:2004}) has suggested  possible models of atmospheres which allow the presence of Coulomb systems ~\\cite{Turbiner:2004m, Turbiner:2005m,Mori:2007,Heyl:2007}. The hydrogen atom is the simplest and most studied   Coulombic system in weak and strong magnetic fields (see for example the early review~\\cite{Garstang} and references therein, and \\cite{Ruderbook,K&L:ET,Karnakov:2003} for more recent studies).  In the present study our goal is to apply a physics recipe (described in full generality in~\\cite{Tur}) for choosing variational trial functions to the study the hydrogen atom in a magnetic field and its electromagnetic transitions between the lowest bound states $1s_{0}$, $2p_{-1}$ and $2p_{0}$. The study is intended as a test of the methodology developed in~\\cite{Tur}. Electromagnetic transitions in the hydrogen atom in the absence of a magnetic field is a widely described subject (see {\\em e.g.} \\cite{Sobelman}). In a strong magnetic field such electromagnetic transitions have been studied by a number of authors, specially focusing on the effects of the transverse motion across the magnetic field direction (see {\\em e.g.} \\cite{CuvelliezBaye,PP95,Potekhin94,Potekhin98}).  Our consideration is non-relativistic, based on a variational solution of the Schroedinger equation. Thus, the magnetic field strength is restricted by the Schwinger limit $B=4.414\\times 10^{13}$\\,G. Our study is also based on the Born-Oppenheimer approximation of zero order: the proton  is assumed to be infinitely massive. Thus we neglect the effects of the CM motion {\\it i.e.}  the effects of the transverse motion of the atom with respect to the magnetic field orientation.   The study is realized in two steps: (i) a variational calculation of the states $1s_{0}$, $2p_{-1}$ and  $2p_{0}$  is done with suitable trial functions (selected according to the physics recipe), and (ii) with the variationally obtained approximate wavefunctions  we calculate the allowed radiative transitions among these states in the electric dipole approximation (see below).  Atomic units are used throughout ($\\hbar$=$m_e$=$e$=1), albeit energies are expressed in Rydberg (Ry). The magnetic field $B$ is given in a.u. with $B_0= 2.35 \\times 10^9$\\,G, although frequently we will also use  magnetic fields strengths given in Gauss for convenience.     \\subsection{Hamiltonian} \\label{sec:Hamiltonian} The Hamiltonian which describes the Coulomb system formed by an infinitely massive proton and one electron $(pe)$  placed in a homogeneous constant magnetic field directed along the $z$-axis, $\\nobreak{{\\bf B}=(0,0,B)}$ is given by\\ \\begin{equation} \\label{Ham} \\hat{\\mathcal {H}}  = ({\\hat {\\mathbf p}}+{\\cal A})^2  -\\frac{2}{r} \\ = \\ \\hat{\\mathbf p}^2  -\\frac{2}{r} + (\\hat{\\mathbf p} \\cdot {\\cal A}+{\\cal A}\\cdot\\hat{\\mathbf p}  ) + {\\cal A}^2 \\ , \\end{equation} where $\\hat{\\mathbf p}  =-i \\nabla$ is the electron momentum,  $r$ is the distance between the electron and the proton fixed at the origin, and ${\\cal A}$ is a vector potential  corresponding to the magnetic field $\\bf B$. A contribution to the energy coming from the coupling between the electron intrinsic magnetic moment and the magnetic field $\\sim {\\hat {\\mathbf S}}\\cdot {\\mathbf B}$, being  constant, has been  dropped from (\\ref{Ham}). Now, if we choose the vector potential in the  symmetric gauge \\[ {\\cal A}= \\frac{B}{2} (-y,\\ x,\\ 0)\\ , \\] the Hamiltonian acquires the form \\begin{equation} \\label{ham-H} \\hat{\\mathcal {H}}= -\\mathbf{\\Delta} - \\frac{2}{r} + {\\hat{l}}_{z}B + \\frac{B^2\\rho^2}{4}\\ , \\end{equation} where $\\mathbf{\\Delta}$ is the Laplacian operator, and ${\\hat{l}}_{z}$ is the conserved $z$-component of the electron angular momentum. The Hamiltonian (\\ref{ham-H}) is also invariant with respect to  the $z$-parity $\\pi_z$   ({\\em i.e.} reflections $z\\to -z$). Thus the eigenstates can be classified by $(m,\\pi_z)$: the magnetic quantum number $m$, corresponding to the conservation of ${\\hat{l}}_{z}$, and the $z$-parity quantum number $\\pi_z=\\pm 1$.   \\subsection{Choice of trial functions} \\label{sec:trialfchoice}  The procedure which we use to explore the problem, is the variational  method  with a well defined recipe for choosing trial functions. This recipe is based on physical arguments, described in full generality in \\cite{Tur}. The basic ingredients are (i) for a given trial function $\\psi_{trial}$, the potential $V_{trial}=\\mathbf{\\Delta}\\psi_{trial}/\\psi_{trial}$, for which such function is an exact eigenfunction, should reproduce as many as possible the basic properties of the original potential, {\\em e.g.} in the present case it should reproduce the Coulomb singularities and the harmonic oscillator behavior at small and large $r$ distances  respectively, and (ii)  the trial function $\\psi_{trial}$ should include the symmetries of the problem. For example, if the ground state is studied, the trial function $\\psi_{trial}$ has to be a nodeless function.  Adhering to this recipe, in~\\cite{PotekhinTurbiner} it was proposed  the following function for the hydrogen ground state $1s_0$ \\begin{eqnarray} \\label{trialgs} \\lefteqn{ \\Psi_{1s_{0}}=} \\\\[5pt] && e^{-\\sqrt {\\gamma_{1} r^{2} + (\\gamma_{2}r^{3} + \\gamma_{3}\\rho^{2}r+ \\gamma_{4}\\rho^{3}+\\gamma_{5}\\rho r^{2})B^2  + \\frac{\\gamma_{6}B^{2}\\rho^{4}}{16} + \\frac{\\gamma_{7}B^{2}\\rho^{2}r^{2}}{16} }} \\nonumber\\,, \\end{eqnarray} where $\\gamma_1,\\ldots \\gamma_7$ are variational parameters. The potential associated with the function (\\ref{trialgs}) reproduces the Coulomb singularity at $r\\to 0$ as well as the harmonic oscillator for $r\\to \\infty$ of the original Hamiltonian ({\\ref{ham-H}}). The function (\\ref{trialgs}) has no nodes. In the following we will use the function (\\ref{trialgs}) for the variational study of the ground state $1s_{0}$.  For the lowest excited states $2p_{-1}$, $2p_{0}$ the presence of the magnetic field does not modify the nodal structure of the field-free exact eigenfunctions\\footnote{{\\em i.e.} the nodal surfaces defined by the condition $\\psi_{trial}(r,\\theta,\\phi)=0$.}. A hint for this can be obtained in first order degenerate perturbation theory in $B^2$, where one can see that the states $2p_{-1}$, $2p_{0}$ are not mixed. Thus   the following variational functions for the  $2p_{-1}$, $2p_{0}$ states are proposed  \\begin{subequations} \\begin{eqnarray} \\label{trial2pm1} \\Psi_{2p_{-1}}&=& \\rho e^{-i\\,\\phi}\\,\\psi_0 \\,, \\\\ \\label{trial2p0} \\Psi_{2p_{0}}&=& z\\,\\psi_0 \\,, \\end{eqnarray} \\end{subequations} where the functions $\\psi_0$ have the same functional form as (\\ref{trialgs}) but with their own $\\gamma$-parameters. Thus, in the  functions (\\ref{trial2pm1}), (\\ref{trial2p0})  we keep the same polynomial prefactor  as in the corresponding field-free wavefunctions and multiply them by  a nodeless function. It is easy to check that the functions  (\\ref{trialgs}), (\\ref{trial2pm1}) and (\\ref{trial2p0})  are orthogonal. These functions describe the lowest energy states, $1s_0, 2p_{-1}, 2p_{0}$, among the states with quantum numbers $(m,\\pi_z)=(0,+),(-1,+),(0,-)$ respectively, in all the range of magnetic fields studied. ",
        "conclusions": "Along the present study we have used a variational approach with a physics recipe for choosing simple trial functions, as a test for an alternative method to study  electromagnetic transitions in the hydrogen atom placed in a constant magnetic field. We assume that the proton is infinitely massive (Born-Oppenheimer approximation of zero order). It was found that the method yields very accurate results for the binding energies, in particular, for the lowest states studied $1s_0, 2p_{-1},2p_0$, in all the range of  magnetic fields $B=0-4.414\\times 10^{13}$\\,G. The results for binding energies  show that the accuracy given by  the simple 7-parametric trial functions (\\ref{trialgs}),(\\ref{trial2pm1}) and (\\ref{trial2p0}) is excellent for {\\em small} magnetic fields $B\\lesssim 0.1$\\,a.u., however it  decreases monotonously as the magnetic field grows: the relative differences between the results given by the proposed trial functions (\\ref{trialgs}),(\\ref{trial2pm1}) and (\\ref{trial2p0}) and the today's most accurate results \\cite{Ruderbook} and \\cite{K&L:ET} reach about $\\sim 1\\%$ for $B=1000\\,$a.u.  A possible explanation for  this reduction in  accuracy is the  fact that transverse and longitudinal degrees of freedom ($\\rho$ and $z$ respectively)  in  the trial functions (\\ref{trialgs}), (\\ref{trial2pm1}) and (\\ref{trial2p0}) appear {\\lq isotropically\\rq} in the electron-proton distance $r$, preventing their adiabatic separation at $B\\to \\infty$.    Dipole $d^{(q)}(B)$, and oscillator strengths $f^{(q)}(B)$, of the   electromagnetic transitions ($1s_{0}   \\leftrightarrow 2p_{-1}$ and  $ 1s_{0}  \\leftrightarrow 2p_{0}$)  were computed  with the approximate  wave functions (\\ref{trialgs}),(\\ref{trial2pm1}), (\\ref{trial2p0}) as functions of the magnetic field strength $B$. Our results for the $1s_{0}   \\leftrightarrow 2p_{-1}$ transition are in very good agreement  with the results of ~\\cite{Ruderbook}, with small deviations varying rather monotonously in $0.001-0.1\\%$ for the interval of magnetic fields $B=0.1-1000$\\,a.u.    Major deviations between the present results and the results in~\\cite{Ruderbook}  for the dipole and oscillator strengths were observed in the case of the $1s_{0}   \\leftrightarrow 2p_{0}$  transition, where we have differences up to $\\sim 30\\%$ in the oscillator strength $f^{(0)}_{1s_{0}   \\leftrightarrow 2p_{0}} (B)$ at $B=1000$\\,a.u. A similar difference $\\sim 30\\%$ is obtained for the corresponding  transition probability  $w^{(0)}_{1s_{0}   \\leftrightarrow 2p_{0}} (B)$. It is worth to emphasize that for a strong magnetic field $B\\simeq 1000$\\,a.u. the   transition probability $w^{(0)}_{1s_{0}   \\leftrightarrow 2p_{0}} (B)$ is 3 orders of magnitude larger than the transition probability  $w^{(+1)}_{1s_{0}   \\leftrightarrow 2p_{-1}} (B)$   corresponding to the $1s_{0}   \\leftrightarrow 2p_{-1}$ transition, so  the difference in results for the transition probability  $w^{(+1)}_{1s_{0}   \\leftrightarrow 2p_{-1}} (B)$ might be of relevance for the analysis of the X-ray spectra of neutron stars.  It is important to mention that if for the variational binding energies there exists a criterion to decide which results given by different approximate wavefunctions are better (since variational binding energies approach from below to the true binding energies),  there exists no similar criterion for other expectation values or matrix elements. Thus, it is not clear so far which  results for dipole strengths  are better. Therefore, more investigations on the electromagnetic transitions in the hydrogen atom in a magnetic field, specially in the domain of strong magnetic fields,  would be desirable in order to answer this question."
    },
    "0711/0711.2848_arXiv.txt": {
        "abstract": "Velocity structures of jets in a coronal hole have been derived for the first time. {\\it Hinode} observations revealed the existence of many bright points in coronal holes. They are loop-shaped and sometimes associated with coronal jets. Spectra obtained with the Extreme ultraviolet Imaging Spectrometer (EIS) on board {\\it Hinode} are analyzed to infer Doppler velocity of bright loops and jets in a coronal hole of the north polar region. Elongated jets above bright loops are found to be blue-shifted by 30 km s$^{-1}$ at maximum, while foot points of bright loops are red-shifted. Blue-shifts detected in coronal jets are interpreted as upflows produced by magnetic reconnection between emerging flux and the ambient field in the coronal hole. ",
        "introduction": "Coronal holes are dark regions of the corona seen in soft X-ray or coronal spectral lines. In a coronal hole, there is open magnetic field where solar wind is flowing out, while bright corona of the surrounding area is composed of closed field. It is obvious that magnetic field plays an important role in forming varying structure of corona, but the mechanism responsible for energy and mass supply to the corona remains a mystery. Study of the dynamics in coronal holes is, therefore, crucial for understanding coronal heating and the acceleration of the solar wind.  \\cite{shimojo1996} and \\cite{shimojo2000} derived physical parameters of X-ray jets using the soft X-ray telescope \\citep{tsuneta1991} on board {\\it Yohkoh} spacecraft \\citep{ogawara1991}. They claimed that the observed properties of X-ray jets were consistent with the theory, in which jets were produced by evaporation flows due to reconnection heating.  \\cite{moses1997} studied coronal jets in a sequence of 195\\AA~ images obtained with EIT on {\\it SOHO}. They found apparent rising velocity of the jets is 100 -- 400 km s$^{-1}$.  But the measurement of Doppler velocity of coronal jets is needed to firmly establish the model, because their analysis is based on the apparent motion of the jets which may differ from the actual plasma motion.  \\cite{wang1998} identified white-light jets in 2 -- 6 R$_{\\odot}$ range of LASCO images, which corresponds to EUV jets observed with EIT. They were rooted in bright points The leading edge of white-light jets propagated with speeds of 400 -- 1000 km s$^{-1}$.  Coronal plume is a stationary structure extending from coronal holes observed in the EUV and coronagraphs in visible light. \\cite{deforest1997} traced plume structures from the surface to 15 R$_{\\odot}$ by combining EIT and coronagraph data. While coronal jets are transient phenomena, plume structures are stable, at least for 24 hours. The study of temporal variation is important to distinguish jets and plumes.  The Extreme ultraviolet Imaging Spectrometer (EIS; \\citealt{culhane2007}) on board {\\it Hinode} \\citep{kosugi2007} is capable of obtaining EUV spectra with high resolution and high efficiency. It is designed to derive properties of the solar corona, such as temperature, line-of-sight velocity, and density, through spectroscopic observation. Although features in coronal hole are dark in coronal spectral lines, the high throughput of the EIS allows us to determine Doppler shifts in coronal holes. In the present work, we studied the velocity structure of bright points in a coronal hole in the north polar region. The EIS observations are described in \\S 2. The data reduction procedure and the method of data analysis are shown in \\S 3. Results are presented in \\S 4 and their interpretation is discussed in \\S 5. ",
        "conclusions": "We studied the properties of jets in a coronal hole by using a raster scan of EIS on board {\\it Hinode}. The high resolution and sensitivity of EIS allowed us to study the Doppler velocity of coronal jets. Our observation showed that bright points in coronal holes are loop-shaped at coronal temperatures, and have bright foot points at transition region temperatures.  Determination of Doppler shifts in Fe {\\textsc XII} line revealed that jets extending from bright points are blue-shifted by 30 km s$^{-1}$ at maximum. Taking the projection effect into account, the observed velocity is consistent with that of jets expected from chromospheric evaporation flow caused by magnetic reconnection. Bright loops at coronal temperature are explained as loops formed after reconnection between emerging flux and ambient fields in the coronal hole. But observed red-shifts near the foot points of loops are not fully understood in the model.  Collaboration of all three telescopes on board {\\it Hinode} is important for studying the connection between the photosphere and the corona. XRT is capable of obtaining a sequence of X-ray images with high cadence, which is suitable for studying the temporal variation of the coronal bright points. The spectro-polarimeter of solar optical Telescope (SOT) obtains very high resolution transverse and longitudinal magnetic fields in the photosphere which are needed for a better understanding of the formation of bright points and jets in the corona.  \\vspace{12pt} {\\it Hinode} is a Japanese mission developed and launched by ISAS/JAXA, collaborating with NAOJ as domestic partner, NASA and STFC (UK) as international partners. Scientific operation of the {\\it Hinode} mission is conducted by the Hinode science team organized at ISAS/JAXA. This team mainly consists of scientists from institutes in the partner countries. Support for the post-launch operation is provided by JAXA and NAOJ (Japan), STFC (U.K.), NASA (U.S.A.), ESA, and NSC (Norway).  This work is carried out at the NAOJ {\\it Hinode} Science Center, which is supported by the Grant-in-Aid for Creative Scientific Research \"The Basic Study of Space Weather Prediction\" (17GS0208, Head Investigator: K. Shibata) from the Ministry of Education, Science, Sports, Technology, and Culture (MEXT) of Japan, donations from Sun Microsystems, and NAOJ internal funding."
    },
    "0711/0711.2881_arXiv.txt": {
        "abstract": "We compare both the Milky Way and M31 to local external disk galaxies within the same mass range, using their relative locations in the planes formed by V$_{flat}$ vs. M$_K$ (the Tully-Fisher relation), j$_{disk}$ (specific angular momentum) and the average Fe abundance of stars in the galaxy outskirts. We find, for all relationships, that the MW is systematically offset by 1 $\\sigma$ or more, showing a significant deficiency in stellar mass, angular momentum, disk radius and [Fe/H] in the stars in its outskirts at a given V$_{flat}$. Our Galaxy appears to have escaped any significant merger over the last 10-11 Gyr which may explain its peculiar properties. As with M31, most local spirals show evidence for a history shaped mainly by relatively recent merging. ",
        "introduction": "In Fig. \\ref{tfjfe}, we show the local Tully-Fisher Relation (TFR) in K-band, as well as the specific angular momentum of the disk j$_{disk}$ vs. the rotational velocity V$_{flat}$. In both relations, V$_{flat}$ is taken as a proxy for the total mass. Noteworthy, the local TFR was carefully derived comparing three different local samples (see top-left inset and \\citealt{hammer07}). Strikingly, the MW systematically falls at $\\sim$ 1 $\\sigma$ of these relations, conversely to M31: the MW has a smaller stellar mass and specific angular momentum, compared to local galaxies of similar total mass. On the basis of their location in the (M$_K$, V$_{flat}$, R$_d$) volume, the fraction of spirals like the MW is $\\sim$7\\%, while M31 appears to be a much more typical spiral (see \\citealt{hammer07}). The representativity of the MW as a typical spiral galaxy is actually even smaller, if one takes into account the iron abundance of the stellar halo (see right panel).   \\begin{figure}[!h] \\centering \\includegraphics[width=12cm]{puech1.eps} \\caption{\\emph{Left and middle panels:} K-band TFR and j$_{disk}$ vs. V$_{flat}$ relation for three local galaxy samples (blue triangles from \\citealt{ver01}, red dots from \\citealt{pizagno07}, and black squares from \\citealt{courteau97}). Solid lines show the fits for the combined three samples, and the dashed lines the residual 1 $\\sigma$ dispersion. The locations of the MW and M31 are marked by two large green dots. In the top-left panel insert, we show the individual fits of the TF relations for each of the three samples (see more details in \\citealt{hammer07}). In the middle panel, open stars show results of simulations by \\cite{steinmetz99}, illustrating the difficulty in reproducing the angular momentum of spiral galaxies when using the standard model for disk formation. \\emph{Right panel:} Iron abundances of stellar haloes of eight spiral galaxies (from \\citealt{mouhcine06}) vs. log(V$_{flat}$) (small red points). Large green points mark the position of the Milky Way and M31. The full line assumes M$_{star}\\sim$V$^4$ \\citep{McGaugh05} and M$_{star}\\sim$Z$^2$ following the prescription of \\cite{dekel87}. Red dashed lines identify a box including the color measurements of \\cite{zibetti04}, after stacking 1047 edge-on SDSS galaxies, and assuming that their colors are dominated by red giant stars.} \\label{tfjfe} \\end{figure} ",
        "conclusions": ""
    },
    "0711/0711.0760_arXiv.txt": {
        "abstract": "We investigate DBI inflation using $N$ multiple branes and show how the configuration is equivalent to a single wrapped $D5$-brane with flux. We then show how $1/N$ corrections can be implemented, and we examine the sound speed and levels of non-Gaussianities in two distinct cases. For models with constant warping we find that the non-Gaussian amplitude is bounded from above (as a function of $\\gamma$). For $AdS$ backgrounds we find that the signature is generally large and positive, although is no longer globally defined over the full phase space. We then investigate an inflationary mechanism using a representation cascade, whereby the transition from a reducible representation to the irrep drives inflation. ",
        "introduction": "In the absence of any direct test of string theory, cosmology remains the best laboratory with which to test string theoretic models \\cite{lectures}. Over the past few years we have witnessed cosmology become a precision science, with COBE, WMAP and SDSS \\cite{data} providing crucial support for the flatness of the universe, the existence of dark energy and for a period of cosmic inflation. Whilst the dark energy puzzle remains an outstanding problem for theoretical physics in general \\cite{darkenergy}, inflation has been a carefully developed paradigm with many explicit models. Unfortunately as far as inflationary model building is concerned, there are still many problems to be resolved. Particularly since many of the models suffer from super-Planckian VEV's for the inflaton field \\cite{lectures}, and therefore find themselves in a region where quantum gravity effects are non-negligible. Conversely the lack of a background independent formulation of string theory has prevented the explicit construction of top-down models, and much of string cosmology has been done explicitly at the field theory level. Whilst there is nothing wrong with this in principle, many of the models are somehow missing much of the underlying string theoretic structure which is where we would expect the more interesting physics to emerge.  There are indications that this picture is about to change \\cite{lectures}. Our understanding of both geometric \\cite{geometricflux} and non-geometric flux compactifications \\cite{nongeomflux} of type II string theory has increased immeasurably in recent years, allowing for the construction of more realistic inflationary models \\cite{realisticmodels}. Additionally models emerging from heterotic M-theory \\cite{assistedinflationmtheory} can also now be placed on a more secure footing, and may yet unify both the standard model and inflaton sectors.  Of course there remains much work to be done, but the general prognosis is that inflationary model building will only improve.  One of the simpler models of string inflation relies on the motion of branes, where the inflaton is now reflected in terms of geometry. Either as the distance between a pair of $D3-\\bar{D}3$-branes \\cite{dbarinflation}, or as the distance between a single brane and some reference point in a warped throat \\cite{braneinflation}. These models are especially appealing, not just because of their simplicity, but because the inflaton is an open string mode which will vanish at the end of inflation and therefore one doesn't need to worry about how it interacts with the standard model sector \\cite{reheating}. In the light of recent developments in type IIB flux compactifications, and the existence of a potential multitude of warped throats with which to resolve the hierarchy problem, these models have become even more appealing. Given the vast number of free parameters that we often find in string model building, it is relatively easy to construct a model that satisfies the WMAP data. Therefore we should be interested in predictions that can be ruled out. These should not be regarded as being deficiencies of string theory, on the contrary in fact, as we are narrowing the parameter space with each one eliminated. Much of the community is now involved in determining which signatures of a particular model can be tested. Indeed many of them rely on bounds placed on cosmic (super)string formation during or after inflation \\cite{cosmicstrings}.  One particular model based on the non-linear structure of the DBI action itself, named DBI inflation \\cite{dbiinflation}, has an interesting signature in that it predicts large levels of non-Gaussian perturbations during inflation \\cite{nongaussianities}. This is important since the result is apparently background independent \\cite{lectures}. There has been much work on this model and its implications \\cite{irinflation, related, spinflation, slowroll}, but the general consensus now is that the simplest scenario is no longer viable. This means it is essential for us to develop more realistic variations of this model\\cite{assistedinflationdbi, multibrane, ward}. There have been several proposals for extending this work, ranging from multi-brane configurations to branes wrapping non-trivial cycles. Although these extensions are able to satisfy the experimental bounds, there is still some concern about the range of validity of such models.  In this paper we will initiate an investigation into the multi-brane proposal of \\cite{ward}, and begin to institute higher order corrections to the action \\cite{finiten}. In this case the corrections we are interested in are the $1/N$ corrections in the large $N$ limit. These corrections are important as the primary constraint on the model is that $N<<M$, where $N$ is the number of branes and $M$ is the total flux in a throat. In the compact case we know that $M$ is bounded from above, by considering compactification of F-theory onto Calabi-Yau (CY) four-folds \\cite{geometricflux}, and therefore this restricts the number of dynamical branes accordingly. The large $N$ limit ensures that the action simplifies, moreover the relevant physical scales are either suppressed or enhanced by this large number allowing us to evade many of the tight constraints. However assuming large $N$ also means that the backreaction could be dangerously out of control. What we would like is therefore to keep $N$ relatively large but also understand how some of the $1/N$ corrections alter things. Given the highly non-linear structure of the DBI even at leading order, we expect that these corrections will be analytically complicated. Therefore rather than search for concrete models of inflation we will restrict ourselves to analysis of the interesting observables associated with DBI inflation. In short we will be interested in i) how the corrections alter the sound speed and ii) how the corrections alter the prediction of large non-Gaussianity. We leave a more detailed investigation of the model dynamics to another publication.  In section II we will introduce the action for the multi-brane configuration we are describing and its features. We will also demonstrate how this is equivalent to a model based on a $D5$-brane wrapping a two-cycle and carrying flux along its internal directions - and therefore overlapping with the model proposed in \\cite{wrappedbranes}. In section III we will show how these $1/N$ corrections can be implemented, and how they alter the leading order behaviour of the Lagrangian. We will then investigate how the corrections alter the predictivity of the model. In the final section we will consider an alternative model using a similar set-up, but we will use the group representation space as the inflationary phase space. ",
        "conclusions": "Cosmology has entered into a new era of precision data \\cite{data}, and it is therefore imperative that top-down models make some falsifiable predictions in order to distinguish them from simple field theory phenomenology. The simplest models in string theory belong to the class of DBI inflation, and their defining characteristic is that the sound speed of fluctuations is greatly suppressed leading to potentially observable non-Gaussian signatures in the CMB \\cite{dbiinflation}. Whilst this has been an important result, recent work has determined that we need to develop more intricate models in order to completely satisfy the current observational data. Extending these scenarios to include wrapped branes, or multiple branes allows us to evade these constraints at the cost of losing control over the low energy theory. The previous work \\cite{ward} demonstrated new physical effects when one uses a finite number of coincident branes, this is the best of both worlds in some sense - since we can still control the backreaction of the branes upon the warped geometry, but also capture more interesting world-volume effects. In this paper we have investigated the effect of the $1/N$ corrections to the large $N$ solution, since this is essentially a combinatorics issue \\cite{finiten} and also overlaps with much of the recent work on wrapped configurations \\cite{wrappedbranes}. Specifically we have seen how these corrections, suppressed in the large $N$ limit, affect the speed of sound and the non-Gaussianity (at least for the equilateral triangle modes) in backgrounds with constant warping, and backgrounds of the form $AdS_\\times X_5$. In both cases we have seen that the spectrum of non-Gaussianities has new features present in the non-relativistic (constant warping) and intermediate ($AdS_5$) limits. This indicates that the $1/N$ correction plays an interesting role in the inflationary dynamics. We have also started to develop an alternative inflationary scenario using a cascade through representation space to drive inflation. This appears to be sensitive to the charge of the $D3$-brane and also imposes tight restrictions on the inflaton velocity. For the usual $D3$-brane solution, we find that the field prefers to tunnel from the false vacuum to the true vacuum and that the decay rate for such a vacuum is relatively short and goes like $e^{-N}$ at large $N$. This is very much reminiscent of old inflation, although because the field will still roll towards the minimum it may evade the graceful exit problem. For $\\bar{D}3$-branes on the other hand, slow roll inflation appears to be preferred - although we estimate that in order for the model to satisfy the COBE bounds we require multiple transitions which may not be feasible.  There remains much work to do on building viable models of DBI inflation. The results shown here and in \\cite{assistedinflationdbi} have shown that the non-Gaussianities can be significantly different from the leading order term once you start to include sub-leading effects. This suggests that other corrections could also become important, even in the slow roll regime \\cite{slowroll} of DBI inflation. These corrections could be particularly interesting for the finite $N$ solutions in \\cite{ward}, since the backreaction of the branes on the warped geometry is still under control - however the sound speed has dramatically different behaviour to the large $N$ and single brane models. We must also develop better mechanisms for reheating \\cite{reheating} in these models. Since the non-linear form of the action captures all the terms in the $\\alpha'$ expansion, one would hope that there could be some stringy signature present in the standard model which lies just beyond the current collider physics scale. This is important not only for aesthetics, but also because the string signature can again be tested. We hope to return to these issues in the future.  \\begin{center} \\textbf{Acknowledgements} \\end{center} We wish to thank Steve Thomas, Larus Thorlacius, Adam Ritz, Shinji Tsujikawa, Shinji Mukohyama and Maxim Pospelov for useful remarks and comments."
    },
    "0711/0711.4908_arXiv.txt": {
        "abstract": "We show that high-resolution observations of resolved Zeeman split lines can be used to obtain new constraints on the stellar magnetic field geometry. In particular, the contrast of the field strength distribution over the stellar surface can be deduced from the differential measurement of the second moment of the $\\pi$ and $\\sigma$ Zeeman components. Our analysis of the triplet lines in slowly rotating cool magnetic CP stars uncovers a surprisingly homogeneous field structure, inconsistent with any low-order multipolar geometry. ",
        "introduction": "A number of slowly rotating magnetic A and B stars have surface magnetic fields strong enough to produce observable Zeeman splitting of line profiles in high-resolution spectroscopic observations. This effect is routinely used to diagnose the surface magnetic field strength (e.g. Mathys et al. \\cite{mathys}). But, at the same time, information provided by the shapes of the resolved $\\pi$ and $\\sigma$ components is largely neglected. Here we propose a novel technique to measure the {\\it relative width} of the resolved Zeeman components and to infer an important global field characteristic from the resulting new magnetic observable. ",
        "conclusions": ""
    },
    "0711/0711.1370_arXiv.txt": {
        "abstract": "The rare phenomenon of ribbon-like hard X-ray (HXR) sources up to 100~keV found in the 2005 May 13 M8.0 flare observed with the \\textit{Reuven Ramaty High Energy Solar Spectroscopic Imager} provides detailed information on the spatial distribution of flare HXR emission. In this Letter, we further investigate the characteristics of HXR emission in this event using imaging spectroscopy, from which we obtain spatially resolved HXR spectral maps during the flare impulsive phase. As a result we found, along a flare ribbon, an anticorrelation relationship between the local HXR flux and the local HXR spectral index. We suggest that this can be regarded as a spatial analog of the well-known temporal soft-hard-soft spectral evolution pattern of the integrated HXR flux. We also found an anticorrelation between HXR spectral index and local electric field along the ribbon, which suggests the electron acceleration by the electric field during flares. ",
        "introduction": "It has long been recognized that hard X-ray (HXR) emission is a powerful diagnostic of accelerated energetic electrons produced by flares. The HXR spectrum emitted by nonthermal electrons often appears as a power-law distribution in photon energy, which implies a characteristic in the energy distribution of the electron flux bombarding the target under the bremsstrahlung emission mechanism. Temporal evolution of the electron energy distribution is highly important for identifying the dominant acceleration process.  As known from the early results of the traditional scintillation-counter spectrometers, the HXR spectral index generally follows a soft-hard-soft (SHS) spectral pattern in the rise-maximum-decay phase of flares \\citep{parks69,benz77,brown85,dennis85}. It was further corroborated by the results of \\textit{Reuven Ramaty High Energy Solar Spectroscopic Imager} \\citep[\\textit{RHESSI};][]{lin02} that can measure HXR spectra with much higher resolution. \\citet{hudson02} illustrated the consistency of the SHS pattern derived using higher energy resolution \\textit{RHESSI} data (of order 1~keV) with that obtained with the HXRS scintillation-counter spectra. A systematic study by \\citet{grigis04} using \\textit{RHESSI} revealed that the SHS pattern appears even in sub-peaks of HXR with durations of one minute to shorter than 8~s. With \\textit{RHESSI}'s capability of imaging spectroscopy, the study of SHS has been extended to spatially resolved HXR sources. \\citet{emslie03} and \\citet{battaglia06} examined spectral evolution of both coronal and footpoint HXR sources in the course of flares and find that they commonly show the SHS behavior. Based on the results they suggested that SHS can be an intrinsic feature of the electron acceleration in the flare impulsive phase.  The SHS can also be regarded as a relation between HXR spectral index and flux because the hardest spectrum usually appears in the period of maximum flux \\citep{grigis06}. As a possible extension of this relation to spatial characteristics, we may consider a spatial analog of the SHS pattern in which a stronger HXR region has a harder spectrum. Such idea has been proposed in a couple of studies. For instance, \\citet{masuda01} determined spectral index maps of an extended HXR event using \\textit{Yohkoh} data to find spectral hardening toward the ribbon edge where the most intense energy release is expected. \\citet{hudson04} proposed that the regions of weaker HXR emission would correspond to softer HXR spectra, in an attempt to explain the footpoint-like HXR morphology in contrast with extensive ribbon structure seen at H$\\alpha$ and UV wavelengths. However, it will need a more thorough examination to see whether the relation between HXR flux and spectral index in time can simply be transformed to such a spatial analog. This is one of the goals of this paper that we intend to achieve using \\textit{RHESSI} imaging spectroscopy.  In addition to spectral evolution, the HXR flare ribbon motion can also manifest the progression of flare energy release due to magnetic reconnection. This was enabled after \\citet{forbes84} derived a relationship between the ribbon expansion velocity $u$ and the electric field $E$ in the reconnecting current sheet (RCS) in the form of $E = uB$, where $B$ is the local vertical magnetic field strength in the footpoint. It is important to have an indirect measurement of the electric field, because direct acceleration of electrons by the electric field is a candidate mechanism for creating high energy electrons in flares \\citep{litvinenko96}. Some studies showed that there is a temporal correlation between HXR flux and the electric field derived this way \\citep[e.g.,][]{qiu02}. It is then of next interest whether there is not only a temporal but also a spatial correlation between HXR spectral index and electric field in the RCS. Recently, \\citet{wood05} presented a particle simulation with prescribed electric and magnetic fields, which shows a hardening of the electron energy spectrum with increasing electric field. Such a model prediction can be compared with the above mentioned measurements of the HXR spectrum and electric field.  In this Letter, we investigate spatial distribution of HXR spectrum and electric field during the 2005 May 13 M8.0 flare. This is a particular event \\citep{liu07c}, where the HXR sources are so extended as to be suitable for the imaging spectroscopy, and is thus adequate for studying the possible spatial relationship between the spectral index and electric field. ",
        "conclusions": "Exploiting \\textit{RHESSI}'s capability of imaging spectroscopy, we have determined spectral index of the local photon spectrum in the unusually extended HXR source observed in the 2005 May 13 M8.0 flare. We then find a spatial anticorrelation relationship between both the local HXR flux and electric field corresponding to the local HXR spectral index at several time intervals. We discuss the present results in comparison with similar works on solar HXR imaging spectroscopy.  The present approach is closest to that of \\citet{masuda01} who used \\textit{Yohkoh} images at two energy bands to find a spectral index change across flare ribbons. In their result, the HXR spectrum at the outer edge of the ribbon is found to be harder than in other regions of the ribbon and the hardness of the spectrum gradually changes across the ribbon width. This is consistent with the physical picture that the ribbon edge is connected to the most recently reconnected field lines and thus shows the most energetic electrons. The present study shows a similar result with more spatial details. It clearly shows the hardest HXR sources lying along the edge of flare ribbon in UV \\cite[c.f. Fig~\\ref{f2} and][Fig.~2]{liu07a} and a smooth transition of spectral index across the HXR sources. Particularly, we find that the HXR spectral index exhibits a strong spatial anticorrelation with the HXR flux in all time intervals during the flare impulsive phase. We call this \\textit{spatial SHS} behavior in analogy with the well-known temporal SHS pattern of integrated HXR flux. The spatial SHS may also be an essential feature of solar flare electron acceleration, and may help explaining the confined nature of HXR sources compared with extended H$\\alpha$ and UV ribbons \\citep[c.f.][]{hudson04}.  We must note that the spatial SHS implies a more strict relationship between HXR flux and spectral index than that found in the previous \\textit{RHESSI} studies on the temporal SHS behavior in individual sources \\citep{emslie03,battaglia06}. In the latter results, each isolated source exhibits the temporal SHS pattern independent of each other and thus the normalization of the flux--spectral index relation ($F$--$\\delta$ relation) may differ from one source to another. On the other hand, our spatial SHS implies that the same normalization of the $F$--$\\delta$ relation should hold in all regions. A major difference between those works and ours lies in that our spatial SHS refers to correlations among local sub-regions within one footpoint side of the magnetic arcade, whereas \\citet{emslie03} and \\citet{battaglia06} compared spectra integrated over individual footpoint sources. We suspect that almost all local regions in this event were subject to a common acceleration and transport process to share the same $F$--$\\delta$ relation. For this reason we do not believe that our result is in conflict with the previous results.  An entirely new result in this study is the spatial anticorrelation between HXR spectral index and electric field in the RCS. This property was found because we could measure both the spectral index and the electric field as functions of position within an extended HXR ribbon. As a comparison, we note that the numerical simulation for the direct electric field acceleration of electrons predicted electron energy distribution with power law index $\\delta=4.0, 2.6, 1.5$ for electric field strength $E =$~0.1, 1.0, 10~V~cm$^{-1}$, respectively \\citep{wood05}. We are not sure how to convert this electron power law index $\\delta$ into the observed photon spectral index $\\gamma$, because further assumptions need to be made on the nature of the numerical solutions and the radiation. Within the scope of this paper, we note that there is a qualitative agreement between our observational result (Fig.~\\ref{f4}) and the model result \\citep{wood05}, in that both indicate a hardening of electron energy distribution with increasing electric field strength. It appears that the present result suggests the dominance of direct electric field acceleration of flare electrons. It however does not exclude other possibilities. As discussed by \\citet{hudson02}, an explicit theoretical demonstration of the SHS behavior under the framework of the thick-target model can be made with the stochastic acceleration mechanism discussed by \\citet{benz77}. Alternatively it is still possible that the efficient electron acceleration is confined in the regions of strong HXR flux and the weaker HXR flux regions nearby result from some propagation effect. A more systematic survey of the relationship between these two physical quantities, both temporally and spatially, will be needed in order to ascertain such association under the context of electron acceleration mechanism."
    },
    "0711/0711.3696_arXiv.txt": {
        "abstract": "The deep homogeneous survey of the large Local-Group spiral galaxy M~31 is a milestone project for X-ray astronomy, as it allows a detailed X-ray inventory of an archetypal low-star-formation-rate galaxy like our own. We present first results of the survey, which covers the entire $\\mathrm{D_{25}}$ ellipse. Information from different X-ray energy bands are combined in an X-ray colour image of M~31. In the first 15 observations we found about 1000 sources, the full survey will yield about 2000 X-ray sources. \\\\ Sources will be classified using hardness ratios, extent, high quality spectra and time variability. In addition the sources will be correlated with catalogues in optical, infra-red and radio wavelengths.\\\\ Our goal is to study M~31 X-ray binaries and globular cluster sources, supersoft sources, supernova remnants and the hot interstellar medium and separate them from foreground stars and background objects. ",
        "introduction": "The goal of this survey is to improve our understanding of the different source classes in \\m31\\ and possible connections to the general structure of the galaxy. Additionally the interstellar medium (ISM) will be studied. The results achieved by this survey will help us to better understand the source classes of the Milky Way and of more distant galaxies. Due to its relatively short distance \\citep[780 kpc, ][]{1998AJ....115.1916H,1998ApJ...503L.131S} and its moderate foreground absorption \\citep[\\nh = 7\\hcm{20}, ][]{1992ApJS...79...77S}, the bright, massive spiral (SA(s)b) galaxy \\m31\\ is an ideal target for such a study. \\\\%The external face on look on \\m31\\ allows us to observe a great part of this galaxy with only one observation.\\\\ Early observations in the X-ray wavelength regime were performed by the \\ein\\ observatory \\citep[][]{1979ApJ...234L..45V,1990ApJ...356..119C,1991ApJ...382...82T} and ROSAT \\citep[][]{1993ApJ...410..615P,1997A&A...317..328S,2001A&A...373...63S}. In the last few years many observations with \\chandra, \\xmm, Swift and other X-ray satellites covered mainly the central region of \\m31. Individual references will be given in Section~\\ref{Ssclass}, where the different source classes are discussed. \\chandra\\ results are discussed in a different paper of this meeting \\citep{2007Chan..M31....1W}. ",
        "conclusions": "The deep \\xmm\\ survey of \\m31\\ consists of fifteen pointing directions, which are placed in such a way that together with the observations in the \\xmm\\ archive the whole $\\mathrm{D_{25}}$ ellipse of the galaxy is covered.  The observations were taken in AO5 between June 2006 and January 2007. Some of the observations affected by high background were reapproved in AO6. The full survey will yield about 2000 sources in the field of \\m31. This survey will separate sources belonging to \\m31\\ (supersoft sources, quasisoft sources, supernova remnants and X-ray binaries) from background objects, especially AGNs, and foreground stars.\\\\ We produced an RGB colour image in which SSSs appear as red, fg stars and SNRs as red to yellow sources. All \"hard\" sources are blue. A more qualitative form of the selection criterion, which forms the basis of this colour selection, are hardness ratios. In addition X-ray time variability can be helpful to distinguish between different source classes. Further information is achieved by cross correlations with catalogues in other wavelengths. For bright sources X-ray light curves, X-ray spectra and short term time variability can be examined.  \\\\ The next steps will be an improvement of the source positions and the preparation of a source catalogue."
    },
    "0711/0711.0140_arXiv.txt": {
        "abstract": "%  We present first results from evolutionary simulations aimed at exploring the Hot-Flasher scenario for the formation of H-deficient subdwarf stars. The two types of late hot flashers that lead to He-enriched surfaces, ``deep'' and ``shallow'' mixing cases, are investigated for different metallicities. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.2145_arXiv.txt": {
        "abstract": "{ The observation of high-energy neutrinos from astrophysical sources would substantially improve our knowledge and understanding of the non-thermal processes in these sources, and would in particular pinpoint the accelerators of cosmic rays. The sensitivity of different design options for a future cubic-kilometre scale neutrino telescope in the Mediterranean Sea is investigated for generic point sources and in particular for some of the galactic objects from which TeV gamma emmission has recently been observed by the H.E.S.S.\\ atmospheric Cherenkov telescope. The effect of atmospheric background on the source detection probabilities has been taken into account through full simulation. The estimated event rates are compared to previous results and limits from present neutrino telescopes. }  \\begin{document} ",
        "introduction": "The KM3NeT consortium {\\cite{KM3NeT-1}} is currently working on a conceptual design for a future Mediterranean neutrino telescope, which will have an instrumented volume of a scale of one km$^3$.  One of the goals of the KM3NeT Design Study is the optimisation of the detector configuration with respect to the observation of astrophysical neutrino sources. The observation of high-energy neutrino point sources will substantially improve our knowledge and understanding of the non-thermal processes in these sources, and would in particular pinpoint the accelerators of cosmic rays. Good candidates for high-energy neutrino emission are sources of high energy gamma-rays, which have e.g.\\ been observed with the H.E.S.S.\\ atmospheric Cherenkov telescope {\\cite{HESS}}. In spite of the recent significant progress in the field of high-energy neutrino astronomy, the observation of cosmic neutrinos from such sources is likely to require cubic-kilometre scale neutrino telescopes, such as IceCube {\\cite{IceCube}} and KM3NeT.  In this article two possible configurations of the KM3NeT neutrino telescope are studied to estimate the sensitivity to generic point sources (assuming a neutrino flux proportional to $E^{-2}$, where $E$ is the neutrino energy), and to selected H.E.S.S.\\ sources (estimating the neutrino fluxes from the observed gamma spectra). The full chain from simulation to reconstruction has been taken into account in this study. ",
        "conclusions": "Sensitivity to point sources of the future Mediterranean KM3NeT neutrino telescope have been considered for two different configurations. Neutrino effective areas for these configurations have been obtained using a full simulation and reconstruction chain. In the considered energy interval (about 0.1--100\\,TeV) the effective areas of both configurations are larger than the IceCube effective area. The study indicates that the KM3NeT telescope can significantly improve the current limits on generic point sources. The expected number of events from given sources rises with the instrumented volume and the total photocathode area assumed. With the larger of the two KM3NeT configurations studied, the brightest H.E.S.S.\\ sources would be detectable within less than a decade. However, one has to note that for these studies the atmospheric muon background was neglected, and the energy reconstruction has been treated as perfect."
    },
    "0711/0711.0376_arXiv.txt": {
        "abstract": "We consider gauge vortices in symmetry breaking models with a non-canonical kinetic term. This work extends our previous study on global topological $k$-defects (hep-th/0608071), including a gauge field. The model consists of a scalar field with a non-canonical  kinetic term, while for the gauge field the standard form of its kinetic term is preserved. Topological defects arising in such models, $k$-vortices, may have quite different properties as compared to ``standard'' vortices. This happens because an additional dimensional parameter enters the Lagrangian for the considered model --- a ``kinetic'' mass. We briefly discuss possible consequences for cosmology, in particular, the formation of cosmic strings during phase transitions in the early universe and their properties. ",
        "introduction": "% Vortices are a class of topological defects which may form as a result of symmetry-breaking phase transitions. In condensed matter physics linear defects arise rather commonly. Well-known examples are flux tubes in superconductors \\cite{Abr} and vortices in superfluid helium-4. In cosmology topological defects attract much interest because they might appear in a rather natural way during phase transitions in the early universe. The breaking of discrete symmetries leads to the appearance of domain walls, while the breaking of a global or a local $U(1)$-symmetry is associated with global \\cite{glstr} and local \\cite{lostr} cosmic strings, respectively. Localized defects or monopoles may form in  gauge models possessing a $SO(3)$ symmetry which is spontaneously broken to $U(1)$ \\cite{Pol,tHo}.  Many properties of topological defects arising in symmetry-breaking models with a canonical kinetic term are well-known, see, e.g. \\cite{Vilenkin book,Rubakov}. Adding non-linear terms to the kinetic part of the Lagrangian has interesting consequences for topological defects. For example, defects can exist without a symmetry-breaking potential term \\cite{Sky}. Non-standard kinetic terms in the form of some non-linear function of the canonical term may arise in string theory, due to the presence of higher-order corrections to the effective action for the scalar field. Non-canonical kinetic structures appear also commonly in effective field theories.  During the last years Lagrangians with non-canonical fields were intensively studied in the cosmological context. So-called $k$-fields were first introduced in the context of inflation \\cite{k-inflation} and then $k$-essence models were suggested as solution to the cosmic coincidence problem \\cite{k-essence,kessence others}. Tachyon matter \\cite{tachyon} and ghost condensates \\cite{ghost} are other examples of non-canonical fields in cosmology. An interesting application of $k$-fields is the explanation of dark matter as a self-gravitating coherent state of $k$-field matter \\cite{Halo}. The production of strong gravitational waves in models of inflation with nontrivial kinetic term was considered in \\cite{MukVik}. The effects of scalar fields with non-canonical kinetic terms in the neighborhood of a black hole were investigated in \\cite{BH}.  Recently, symmetry-breaking models with $k$-essence-like terms have been discussed in literature. General properties of global topological defects appearing in such models were studied in \\cite{kdefect}. It was shown that the properties of such defects (dubbed $k$-defects) are quite different from ``standard'' global domain walls, vortices and monopoles. In particular, depending on the concrete form of the kinetic term, the typical size of such a defect can be either much larger or much smaller than the size of a standard defect with the same potential term. A detailed study of global defect solutions for one space dimension was carried out in \\cite{Bazeia}. A self-gravitating $k$-monopole was considered in \\cite{Jin}. In \\cite{Adam} the authors argued that a special type of $k$-defects may be viewed as ``compactons'', i.e. solutions with a compact support. Global strings with a Dirac-Born-Infeld (DBI) term were considered in \\cite{DBIstring}.  In this paper we study properties of gauge vortices arising in a model with a $k$-essence-like kinetic term and a symmetry breaking potential. We dubb such defects ``gauge $k$-vortices'', in analogy to global $k$-defects. We extend our previous investigation on $k$-defects \\cite{kdefect} including a gauge field into the model. As for the global $k$-defects, the scalar field has a non-canonical structure of the kinetic term, while for the gauge field we keep the canonical form of the kinetic term. The existence of non-trivial configurations is ensured by the symmetry-breaking potential term. The generic feature of the model with a non-canonical kinetic term is the appearance of a new scale --- the kinetic ``mass''. The presence of a new mass scale in the model radically changes basic properties of vortices.  We show that generally the size of the scalar core of the gauge vortex solution is almost independent on the presence of the gauge field. Its value can be approximated by the core' size of the global $k$-vortex with the same kinetic structure \\cite{kdefect}. With an additional, natural assumption we find that the vector core has roughly the same size as a standard vortex. A particularly interesting result is that the mass of a vortex radically vary depending on the choice of kinetic term. As the concrete examples we study numerically the vortex solutions for the models with DBI and with a power-law kinetic terms.  The paper is organized as follows. In Sec.~\\ref{SModel}, we describe our model and derive its equations of motion and the energy functional for a vortex solution. General properties of $k$-vortices are studied in Sec.~\\ref{SGeneral}. In Sec.~\\ref{SConstraints} we find constraints on the parameters of the model. Numerical solutions for particular choices of the non-canonical kinetic term  are presented in Sec.~\\ref{Sec Numerics}. We summarize and discuss results and cosmological applications in the concluding Sec.~\\ref{SDiscussion}. ",
        "conclusions": "\\label{SDiscussion}% We have studied topological linear gauge defects (gauge vortices), in the model with a non-canonical kinetic term. The action for the model (\\ref{act}) contains  kinetic terms for the scalar and gauge vector fields and a symmetry-breaking potential. The principal difference of the studied model from the standard one is the presence of a non-standard kinetic part for the scalar field. The term $K(X)$ in the action is in general some non-linear function of canonical kinetic term $X$.  A remarkable feature of the model is that the non-linearity of the kinetic term inevitably leads to the appearance of a new scale in the action, a ``kinetic'' mass, in addition to the ``usual'' scalar and vector masses. The presence of another mass scale in the model changes radically the basic properties of a vortex: the size of the scalar core and the energy of a vortex per unit length vary considerably as compared to the standard case.  We have investigated general properties of $k$-vortices and found restrictions on the parameters of the model, having in mind a rather general form of a kinetic term with the asymptotic behavior $K(X)\\sim X^n$. Also, for the sake of simplicity and clarity of results we assumed that the scalar and vector masses are of the same order, $e\\eta \\sim \\sqrt\\lambda\\eta$. We can summarize our general estimations as follows. The size of the scalar core, $l_H$, depends on the coupling $\\lambda$, and mass scales $\\eta$ and $M$, Eq.~(\\ref{lH1}). A remarkable point is that $l_H$ roughly coincides with the characteristic size of the core in the case of a global $k$-vortex \\cite{kdefect}. The size of the vector core $l_V$ does not depend on the kinetic mass $M$ and is roughly the same as in the standard case, Eq.~(\\ref{lV1}).  Having the values for the core' sizes one can estimate the energy of $k$-vortex per unit length, see Eq.~(\\ref{Eest}). An important result is that the mass of a vortex radically vary depending on the choice of kinetic term. In the case $n>1$ and the limit $\\lambda v^4\\gg 1$, we have found a simple exact expression for the energy functional of $k$-vortex, Eqs.~(\\ref{Efunc}), (\\ref{F}).  As particular examples, we studied numerically two concrete models having non-canonical kinetic terms: A DBI-like term, Eq.~(\\ref{BI}), and a power-law term, Eq.~(\\ref{X3}). The field profiles of domain walls for different choices of $K(X)$ are shown in Fig.~\\ref{Fig num}: The numerical solutions are in agreement with our general estimations.  As we already discussed in our previous work \\cite{kdefect}, interesting properties of $k$-defects may have important consequences for cosmological applications. Standard cosmic strings which might have been formed during phase transitions in the early universe have a mass scale directly connected to the temperature of a phase transition $T_c$, $\\mu\\sim \\eta^2 \\sim T_c^2$. By contrast, the mass scale of a resulting $k$-string depends both on $T_c$ and the kinetic mass $M$. This means that the tension of $k$-strings may not be close to $T_c^2$, thus helping to avoid constrains on cosmic strings $G\\mu\\lesssim 10^{-7}$ \\cite{csc1,csc3} (or even stronger, $G\\mu<3\\times 10^{-8}$, see \\cite{csc2}). Meanwhile theoretical predictions give $G\\mu\\sim 10^{-6}-10^{-7}$ for GUT strings. If, however, physics at the GUT scale involves non-standard kinetic terms, then the GUT phase transition may have lead to the formation of cosmic strings with smaller tension, $G\\mu \\ll 10^{-6}$, thereby evading conflicts with the present observations."
    },
    "0711/0711.2690_arXiv.txt": {
        "abstract": "We present an updated timing solution and an analysis of the profile evolution - including precession and beam shape - of the young, relativistic binary pulsar J1906+0746. The 144-ms pulsar, in a 3.98-hour orbit with eccentricity 0.085 \\cite{lorimer}, was initially discovered during the early stages of the ALFA (Arecibo L-band Feed Array) pulsar survey \\cite{cordes} using the 305-metre Arecibo telescope and was subsequently found in archival Parkes Multibeam Survey data. We have since been regularly monitoring the system using the Arecibo and Green Bank telescopes, and include data from the Jodrell Bank, Parkes, Nancay and Westerbork telescopes. The nature of the binary companion will also be discussed based on improved estimates of the total and companion masses obtained from the updated timing solution. ",
        "introduction": "In order to better understand the profile evolution and in turn remove some timing noise from the system, we have created standard profiles for each epoch of data, and fitted sets of 1,2, or 3 (as necessary) gaussians to each of the pulse and interpulse (seperately). All fitting was done using the gaussian fitting package bfit \\cite{kramer}. After experimenting with several different methods of aligning the gaussian templates, we decided that the smoothest alignment was achieved by keeping the phase of the tallest gaussian constant. This also seems to produce a fairly monotonic behaviour in the phase of the interpulse. The full collection of modelled profiles is shown in Figure 1. These were used as standard profiles for timing.    \\begin{figure} \\includegraphics[height=.3\\textheight]{ltempalign} \\caption{The fits of the pulse and interpulse of J1906+0746, based on ASP (Arecibo; coherent dedispersion), GASP (GBT; coherent dedispersion) and WAPP (Arecibo; spectrometer) observations, each modelled with several gaussian components. The pulses are spaced vertically based on observation date, and the interpulse is magnified by a factor of ten relative to the main pulse. Different machine dispersion methods account for some but not all of the apparent interpulse differences in nearby observations; the re maining differences are under investigation. The vertical line for the main pulse (left) illustrates the chosen alignment, whereas the vertical line for the interpulse demonstrates the marked phase shift of the interpulse relative to the main pulse (see proceedings by Desvignes et al.)}   \\end{figure} ",
        "conclusions": ""
    },
    "0711/0711.0006_arXiv.txt": {
        "abstract": "\\noindent The functional Schrodinger equation is used to study the quantum collapse of a gravitating, spherical domain wall and a massless scalar field coupled to the metric. The approach includes backreaction of pre-Hawking radiation on the gravitational collapse. Truncating the degrees of freedom to a minisuperspace leads to an integro-differential Schrodinger equation. We define a ``black hole'' operator and find its eigenstates. The black hole operator does not commute with the Hamiltonian, leading to an energy-black holeness uncertainty relation. We discuss energy eigenstates and also obtain a partial differential equation for the time-dependent gravitational collapse problem. ",
        "introduction": "\\label{introduction}  Gravitational collapse is expected to result in the formation of a black hole, which then evaporates by Hawking radiation \\cite{Hawking:1974sw}. During this process, any information contained in the initial state is lost, since black hole evaporation leads to thermal radiation which is uncorrelated with the initial state. If the initial state is chosen to be a pure quantum state, the final state would be described by a density matrix and would not be a pure state. The evolution of a pure quantum state into a mixed state violates the tenets of quantum mechanics ({\\it e.g.} \\cite{Susskindbook}). These problems of black hole formation, evaporation and information loss have been central to discussions of combining general relativity and quantum mechanics over the last several decades.  A number of approaches have been made to resolve the issues that arise in combining black holes and quantum field theory, including more sophisticated calculations of the emitted radiation, modifications of quantum mechanics, modifications of quantum field theory, lower dimensional calculations, loop gravity, and string theory ({\\it e.g.} for some reviews see \\cite{Susskindbook,frolov,Strominger:1994tn,Wald:1998de, Ashtekar:2004eh}). Depending on the approach and the new ingredient in the calculation, the conclusions have varied, and no concensus has emerged so far. The issues have become ever more pressing with discussion of black hole formation in highly energetic collisions in LHC \\cite{Banks:1999gd,Giddings:2001bu,Dimopoulos:2001hw} and by ultra-high energy cosmic rays colliding with the atmosphere \\cite{Anchordoqui:2001cg}.  In contrast to earlier studies in which a black hole spacetime is adopted as an arena for quantum field theory, we wish to study the temporal development of an initial state that does not contain a black hole. From the laws of quantum mechanics, the initial state will evolve unitarily in time. If unitary evolution leads to black hole formation then, as has been described by Hawking and others \\cite{Hawking:1974sw,Susskindbook}, unitarity will be lost.  The 3+1 dimensional analysis in Ref.~\\cite{Vachaspati:2006ki} used the functional Schrodinger formalism to study the temporal evolution of a scalar field in the {\\it classical} background of a gravitating, collapsing, spherical domain wall. It was found that the collapse was accompanied by quantum radiation of scalar particles, called ``pre-Hawking'' radiation, with a non-thermal spectrum. The analyses in Refs.~\\cite{Vachaspati:2006ki,Vachaspati:2007hr} were clearly limited by their use of the semiclassical approximation, in which the backreaction of the scalar radiation on the gravitational collapse spacetime is ignored. The purpose of the present work is to attempt to remedy this shortcoming within the functional Schrodinger formalism. The aim of the present work is to provide a framework to study quantum gravitational collapse beyond the semiclassical approximation. The framework includes backreaction of the scalar radiation on the collapse dynamics (see Fig.~\\ref{systemfig}).  The specifics of the domain wall plus scalar field system are not expected to be very important for the questions we are interested in. For example, the Hamiltonian for a collapsing wall that is close to forming a black hole reduces to the form for an ultra-relativistic particle \\cite{Vachaspati:2006ki}, which is expected to be true for any form of matter. So the choice of a collapsing domain wall does not affect the near-horizon dynamics. The scalar field can also be thought of as representing one degree of freedom of a photon or a graviton. One assumed property that is special though, is the assumed masslessness of the scalar field. If the scalar field has a mass and/or carries a global or gauge charge, we do expect some differences that we discuss briefly in Sec.~\\ref{discussion}.  \\begin{figure} \\scalebox{0.6}{\\includegraphics{system.eps}} \\caption{A picture of the system under consideration. A spherical domain wall with radius $R(t)$ collapses towards its Schwarzschild radius, $R_S$, and emits scalar radiation in the process. Due to the emission, $R_S$ also decreases. } \\label{systemfig} \\end{figure}   In general, the Schrodinger formulation yields a functional differential equation for the wavefunctional, $\\Psi [g_{\\mu\\nu}, X^\\mu , \\Phi, t]$, where $g_{\\mu\\nu}$ is the metric, $X^\\mu$ the domain wall position, and $\\Phi$ the scalar field. We will truncate this problem to minisuperspace, considering only spherical domain walls described by a radial function, $R(t)$, the metric fixed to the Schwarzschild form though with variable Schwarzschild radius, $R_S$, and $\\Phi$ decomposed into modes with mode coefficients $\\{ a_k \\}$. Then the wavefunctional gets replaced by an ordinary wavefunction, $\\Psi (R,\\{ a_k \\},t)$. We are interested in finding the time evolution of this minisuperspace wavefunction, starting from black hole free initial conditions, and in determining if there is a breakdown of the evolution at some time.  In any treatment of the gravitational collapse problem, it is necessary to choose a time coordinate since the collapse, by its very definition, is an evolution problem. In the current analysis, we have chosen to work with the Schwarzschild time coordinate. However, there is the possibility that we may miss some portion of spacetime due to this choice of slicing. If this were the case, we would expect some sickness ({\\it e.g.} geodesic incompleteness) to show up in the temporal evolution of the wavefunction. If the temporal evolution remains well-behaved at all times, the evolution is no different from that of an ordinary star for which Schwarzschild coordinates are valid and commonly used. In other words, the Schrodinger formulation simply evolves the system forward in time, without assuming the presence or absence of an event horizon in the future. This is different from other approaches where quantum field theory is used on a spacetime background, where the entire spacetime must be assumed at the onset and then sliced using some coordinate system. However, it would be worthwhile to re-work gravitational collapse in the Schrodinger formalism with a different choice of time coordinate.  The backreaction problem, even for classical point charges, is notorious for its non-local nature, since the trajectory depends on the radiative losses over the entire past. However, this problem only occurs in a perturbative treatment of a point charge, since there is then a zeroth order trajectory due to which there is radiation, and then backreaction to first order in some coupling constant, then a first order correction to the trajectory, then radiation and backreaction to second order, {\\it ad infinitum}. Instead, if the classical point charge is replaced by a regular solution to some field theory, as in a 't Hooft-Polyakov magnetic monopole, a classical solution of the field theory will include the full dynamics and radiation of the point (magnetic) charge. Similarly, the wavefunction for the spherical wall and radiation in the Schrodinger formalism, includes the full dynamics of the wall and the radiation and non-locality is absent for this reason. However, there are two other reasons that make the analysis difficult and lead to non-locality. The first is that the pre-Hawking scalar radiation not only affects the dynamics of the collapse but also contributes to the precise form of the metric. In the minisuperspace approach we have adopted, a form of the metric needs to be specified. This is chosen to be of the Schwarzschild form and thus amounts to the assumption that the energy-momentum tensor in the pre-Hawking radiation only causes negligible departures from the minisuperspace. We expect this approximation to be justified for large collapsing mass where we know that the energy-momentum density of the radiation is comparitively small. This has also been explicitly verified in the semiclassical calculation \\cite{Davies:1976ei}. Also, the present approach is similar to that taken in calculating radiation backreaction on an accelerating charge in electrodynamics where the backreaction is taken to affect the dynamics of the charge but the effects on the Coulomb electric field of the charge is not considered. The second factor that complicates the analysis is the non-linear nature of gravity. In particular, the minisuperspace Hamiltonian depends on the mass of the collapsing object, and this is itself related to the Hamiltonian. In order to isolate the Hamiltonian, we need to invert the momentum operator and this leads to an integro-differential form for the Hamiltonian. Solutions of the corresponding integro-differential Schrodinger equation are hard to find in general but we are able to transform the problem to a purely differential equation in the ``incipient limit'' where the collapse approaches black hole formation.  We discuss energy eigenstates in Sec.~\\ref{stationary} and the time-dependent gravitational collapse problem in Sec.~\\ref{gravcollapse}.  In order to solve the gravitational collapse problem, we need to solve the Schrodinger equation to obtain a wavefunction that describes the collapsing wall and radiation. The first task, however, is to specify an initial value for the wavefunction, such that the initial state itself does not contain a black hole. This means that we need to specify a criterion for deciding if a given wavefunction is black hole free. To answer this question, we propose a ``black hole operator'' in Sec.~\\ref{bhfstate}. Eigenvalues of the black hole operator signify the ``black holeness'' of the state. Interestingly we find that the black hole operator does not commute with the Hamiltonian, implying an uncertainty relation between energy and black holeness.  We start by discussing the classical Hamiltonian (Sec.~\\ref{classicalhamiltonian}), which is then promoted to a quantum Hamiltonian in Sec.~\\ref{quantumhamiltonian}. We then discuss explicit representations of an important operator in Sec.~\\ref{hatB} and use it to define a black hole free state in Sec.~\\ref{bhfstate}. We then discuss energy eigenstates in Sec.~\\ref{stationary} and the gravitational collapse problem in Sec.~\\ref{gravcollapse}. We discuss our results and conclude in Sec.~\\ref{discussion}. A discussion of Hermiticity of certain operators can be found in Appendix~\\ref{hermiticity}. ",
        "conclusions": "\\label{discussion}  We have set up a minisuperspace version of the Schrodinger formalism for studying quantum gravitational collapse of a spherical domain wall in the presence of a massless scalar field coupled to the metric. The description automatically includes backreaction of the quantum radiation on the quantum dynamics of the domain wall. Although the passage to minisuperspace involves a drastic truncation of the system degrees of freedom, it still retains those that are relevant to describe black hole formation and evaporation. A clear advantage of the present approach is that the action for the system greatly simplifies in the interesting limit of an incipient black hole and raises the hope that a solution, even to the notorious back-reaction problem, may be within reach.  In the incipient limit and in the approximation that the collapsing mass is large (see the discussion in Sec.~\\ref{introduction}), the Hamiltonian is given by Eq.~(\\ref{quantumH}). A striking feature of the Hamiltonian is that it involves the operator ${\\hat B}$ which is an integral operator, for which we are also able to find an explicit representation (Eq.~(\\ref{alphaBoperator})).  Our analysis in the incipient limit should be sufficient to study the problem of quantum collapse and pre-Hawking radiation. However, unlike in the semiclassical approximation where the wall radius takes on a definite value, the wavefunction is defined over the entire range of possible wall radii, including very large radii. In a fuller treatment of the problem, it may become necessary to extend the Hamiltonian that we have found in the incipient limit to large values of the radius. The precise extension is not expected to be important, as long as the incipient limit of the extended Hamiltonian matches the Hamiltonian found here.  Once we have the minisuperspace Schrodinger equation, we need the relevant solutions to it that represent a solution to the gravitational collapse problem. Such a solution would contain the fate of matter that is collapsing towards forming a black hole. However, the initial conditions have to be chosen to be ``black hole free'' and some measure of ``black holeness'' has to be defined. We have discussed some possible black hole operators, an example of which is the operator ${\\cal B}$ that coincides with ${\\hat B}$ in the incipient limit. We find eigenstates of ${\\hat B}$ and the eigenvalue $\\beta$ may be used as a measure of black holeness. States with $\\beta > 0$ can be said to be black hole free. However, the black hole operator does not commute with the Hamiltonian, thus implying an energy-black holeness uncertainty relation.  It is interesting to note that the eigenfunction $\\xi_\\beta$ of the operator ${\\hat B}$ oscillates infinitely fast as $\\beta \\to 0$ which corresponds to the domain wall tending to the event horizon (see Eq.~(\\ref{Beigenstate})). Naively, if we start with a wavefunction that overlaps only with states with some finite range with $\\beta > 0$, to get an overlap with states with $\\beta < 0$, it would seem that we would need to go through the infinitely oscillating state at $\\beta =0 $. Quantum mechanically, though, this is not clear. An analogy might help clarify this situation. In the case of a non-realativistic Schrodinger particle, our initial state may be a wavepacket having overlap with only positive momentum eigenstates $e^{ikx}$ with $k > 0$. Suppose at some later time, we find that there is non-zero overlap with a negative momentum eigenstate, $e^{ipx}$ with $p < 0$. Classically this would mean that the particle's momentum has reversed and so, at some point in time, the particle had to pass through zero momentum. In quantum mechanics, however, this is not essential since the particle need not even have a well-defined momentum at intermediate times. So it seems that the singular state at $\\beta =0$ cannot be used to argue, at least straightforwardly, that there is an obstruction to the formation of a black hole.  While we have limited ourselves in this paper to a massless scalar field, it is interesting to consider how the analysis might change if the scalar field has a non-vanishing mass. In that case, the spectrum of eigenmodes of the scalar field would contain bound states in addition to scattering states. Gravitational collapse of the domain wall would then populate both the scattering and bound states, transferring energy from the wall to pre-Hawking radiation and also to an atmosphere of self-gravitating scalar particles. Presumably, the bound states will result in a boson star, though the details are not clear since boson stars as solutions of free or interacting scalar field theory are themselves known to have instabilities if their mass is large \\cite{Ruffini:1969qy,Colpi:1986ye,Jetzer:1991jr}.  In Sec.~\\ref{stationary} we have discussed eigenstates of the Hamiltonian and in Sec.~\\ref{gravcollapse} we have set up a differential equation that describes the gravitational collapse problem. The differential equation contains mixed $t$ and $R$ derivatives, and it may also become singular along curves in the $(R,b)$ plane. These characteristics make the equation hard enough that we have postponed attempts at its solution for future work. In its solution, though, might lie answers to some of the questions we have raised in the introduction, and we may be able to see if the unitary evolution contained in the Schrodinger equation is self-limiting. On the other hand, it may be that the Schrodinger equation always yields unitary evolution, and since it is expected that black holes violate unitarity, it may be impossible to get to the black hole state as suggested by the semiclassical calculation."
    },
    "0711/0711.0230_arXiv.txt": {
        "abstract": "% Modelling global disc galaxies is a difficult task which has previously resulted in the small scale physics of the interstellar medium being greatly simplified. In this talk, I compare simulations of galaxies with different ISM properties to determine the importance of the ISM structure in the star formation properties of the disc. ",
        "introduction": "One of the main problems in understanding galaxy formation is that star formation is extremely complex. The gas out of which stars form is a turbulent mix of forces in which gravitational collapse, thermal pressure, magnetic fields, cosmic ray pressure and energy from supernovae all fight for dominance \\citep{MacLow2004}.  This leaves galaxy simulators with a choice; either to model the interstellar gas in detail but restrict their study to a small patch of the galaxy \\citep[e.g.][]{Slyz2005} or to simulate the entire galaxy but use a greatly simplified model for the ISM \\citep[e.g.][]{Li2005}. The former approach allows the inclusion of many more of the important physical processes but is unable to tell us anything about the global evolution of the disc. The latter, meanwhile, reveals properties of the whole galaxy including star formation histories and global structures, but it is impossible to judge the effect the simple ISM model is having on these results.  However, recent numerical simulations are now able to include a more complex multiphase ISM in these global models \\citep{Tasker2007, Tasker2006, Wada2007}. This allows us to test the impact of modelling the ISM in a more realistic way in galaxy formation. This has particular baring on cosmological simulations, where the resolution of the ISM of individual galaxies is still beyond our reach. ",
        "conclusions": "One of the most surprising results in this research is that despite significant structural differences in the ISM, the star formation properties remain largely unaffected. The reason for this can be seen in Figure~\\ref{fig:pdfs}, where we see that the biggest differences between our three ISM types occur in the low to medium density gas. The high density, star-forming gas, meanwhile, forms a consistent lognormal profile in all cases. This result is largely positive; it suggests that a decent subgrid model for star formation can be used where a detailed ISM model is not possible, such as large-scale cosmological simulations. However, it also implies that a through understanding of star-formation cannot be found from the Kennicutt-Schmidt law, which is largely insensitive to the input physics. This work is presented in greater detail in \\citet{Tasker2007}."
    },
    "0711/0711.2017_arXiv.txt": {
        "abstract": "{} {Our aim was to measure and characterize the short-wavelength radio emission from young stellar objects (YSOs) in the Orion Nebula Cluster and the BN/KL star-forming region.} {We used the NRAO Very Large Array at a wavelength of $\\lambda=1.3$~cm and we studied archival X-ray, infrared, and radio data.} {During our observation, a strong outburst (flux increasing $> 10$ fold) occurred in one of the 16 sources detected at $\\lambda=1.3$~cm, while the others remained (nearly) constant. This source does not have an infrared counterpart, but has subsequently been observed to flare in X-rays. Curiously, a very weak variable \\textsl{double} radio source was found at other epochs near this position, one of whose components is coincident with it. A very high extinction derived from modeling the X-ray emission and the absence of an infrared counterpart both suggest that this source is very deeply embedded.} {} ",
        "introduction": "A large number of compact  (sub-arcsecond size) radio sources have been found associated with stars in the Orion Nebula Cluster (ONC) and/or infrared sources in the Becklin-Neugebauer/Kleinmann-Low (BN/KL) region, an active star-forming site located behind the optically visible nebula \\citep[ hereafter GMR]{gmr87} and \\citep{chu87}.  Non-variable radio emission from many of the compact ONC sources comes from ionized circumstellar disks (``proplyds'') photoevaporated by the intense ultraviolet field of the ``Trapezium'' star $\\theta^1$C, the exciting star of the nebula. Other sources, such as the low-mass companion to the Trapezium O-type star $\\theta^1$A \\citep{fel89,fel91,men07}, show pronounced variability, indicating non-thermal emission. \\citet{fel93} present a comprehensive study of the Orion sources' variability on time scales from two weeks to months.  Protostellar variability has recently been extensively studied at X-ray wavelengths (most notably in the Chandra Orion Ultra-deep Project, or COUP\\footnote{see the ApJS COUP Special Issue, October 2005}: \\citealp{get05,fav05,fla05,wol05}), yet little is known about short-term radio variability and even less about multi-wavelength correlations. \\citet{bow03} presented the first observation of a millimeter wavelength flare of a YSO in Orion (GMR~A). This observation was simultaneously covered in X-rays. The X-ray flux increased by a factor of about ten, two days before the radio flare observation of the source. During the radio flare, the X-ray flux was still at around half of the maximum level reached previously. In observations starting a few days after this outburst, the source was found to show millimeter wavelength activity for at least 13 days, while the infrared $K$-band magnitude remained unchanged \\citep{fur03}.  While longer-term radio variability of YSOs has been known for some time (e.g. \\citealp{fem99}), there are few radio observations of a YSO that cover a period of strong variability, i.e., that document an outburst. Here, we present a new event of this kind occurring in the BN/KL region. In an 8 hour VLA observation at a wavelength of 1.3~cm, we serendipitously found extremely strong variability in a source that had shown only very weak, if any, emission at other occasions and/or at other wavelengths. While this source is detected in X-rays it was not detected at infrared wavelengths. This implies that is is deeply embedded, contrary to GMR~A, which is a bright near-infrared source \\citep{bow03}.  The special observing mode employed and the data reduction are described in Sect.~\\ref{obs}. The outburst is described in Sect.~\\ref{results}, which also presents observations of this source at other wavelengths and discusses the physical scenario. We give a summary in Sect.~\\ref{summary}. ",
        "conclusions": "\\label{summary}  We report the observation of a 1.3~cm wavelength radio flare of a deeply embedded source in the % Orion BN-KL region. The flare source was observed to rise within four hours to become the strongest source in the field. This was followed by a decrease in flux density.  The source does not have a near-infrared counterpart in VLT and HST data but does have a hard X-ray counterpart (a flaring COUP source), implying an absorbing column density corresponding to $A_V \\approx 160$~mag. Consistent with this $A_V$, the source remains undetected at mid-IR wavelengths, although deep observations of the region are hampered by the bright neighboring BN object.  There are 15 additional sources detected at $\\lambda=1.3$~cm, including GMR~A, towards which a 86~GHz flare was reported by \\citet{bow03}. Apart from BN and the RBS, variability was studied for the bright sources GMR~A, GMR~12, GMR~6, and GMR~F. These sources were found to be marginally variable, but instrumental effects appear to play a role. All 16 sources detected at $\\lambda=1.3$~cm have counterparts at $\\lambda=3.5$~cm, including the flare source whose 3.5~cm wavelength counterpart matches one component of a double structure with a separation of  $0\\farcs45$. The very weak components of the double are variable and have been detected on two different observations, but undetected at other times.  The RBS is the second YSO in Orion known to show burst behavior at radio as well as at X-ray wavelengths. It is much deeper embedded than GMR~A \\citep{bow03} and thus may be in an earlier evolutionary stage. In our data, GMR~A is detected at a flux density level of about 8~mJy, roughly consistent with the post-flare flux density observed by \\citet{bow03}. While for the RBS, two separate flares have been observed in X-rays and at radio wavelengths, both observations ended before a quiescent level was reached. During the COUP observations, GMR~A showed a more complex lightcurve, but we may see only the beginning of an activity phase in the COUP data of the RBS."
    },
    "0711/0711.2367_arXiv.txt": {
        "abstract": "In recent years, X-ray observations with \\textit{Chandra} and \\textit{XMM-Newton} have significantly increased our understanding of rotation-powered (radio) millisecond pulsars (MSPs).  Deep \\textit{Chandra} studies of several globular clusters have detected X-ray counterparts to a host of MSPs, including 19 in 47 Tuc alone. These surveys have revealed that most MSPs exhibit thermal emission from their heated magnetic polar caps.  Realistic models of this thermal X-ray emission have provided important insight into the basic physics of pulsars and neutron stars. In addition, intrabinary shock X-ray radiation observed in ``black-widow'' and peculiar globular cluster ``exchanged'' binary MSPs give interesting insight into MSP winds and relativistic shock.  Thus, the X-ray band contains valuable information regarding the basic properties of MSPs that are not accesible by radio timing observations. ",
        "introduction": "Rotation-powered (``recycled'') millisecond pulsars (MSPs) were discovered at X-ray energies during the \\textit{ROSAT} all-sky survey \\citep{Beck93}. In its lifetime, \\textit{ROSAT} detected X-ray emission from 10 MSPs in the field of the Galaxy and one in the globular cluster M28 \\citep{Beck99}. These observations established that MSPs are very faint X-ray sources with typical X-ray luminosities in the range $L_{X}\\sim10^{30-31}$ ergs s$^{-1}$. In recent years, the superb sensitivity of the \\textit{Chandra}\\footnote{\\url{http://cxc.harvard.edu/}} and \\textit{XMM-Newton}\\footnote{\\url{http://xmm.vilspa.esa.es/}} X-ray observatories have permitted much more detailed studies of these objects. \\textit{Chandra}, in particular, with its unprecedented sub-arcsecond angular resolution, has proven to be ideally suited for studies of globular clusters, where MSPs are found in great abundance. In addition, both facilities have allowed detailed spectroscopic and timing studies of the nearby MSPs in the field of the Galaxy \\citep{Zavlin06,Zavlin07}.  Herein, we summarize the current state of affairs and examine future prospects in the X-ray study of MSPs. ",
        "conclusions": "To date, nearly 40 MSPs have been detected at X-ray energies, 19 of which are found in the globular cluster 47 Tuc \\citep{Bog06a}. It has become apparent that X-ray spectroscopic and timing observations of MSPs represent a very promising approach towards elucidating important pulsar properties that are not measurable by other observational means such as radio timing observations. Further study of these sources will provide more important information regarding the X-ray properties of the Galactic population of MSPs, as well as valuable insight into the fundamental properties of NSs.     \\begin{theacknowledgments}    \\end{theacknowledgments}  We would like to thank G. Rybicki, C. Heinke, M. van den Berg, F. Camilo, P. Freire, A. Harding, and W. Becker for their valuable contribution to the work described herein. We also extend our thanks to the conference organizers for their hospitality."
    },
    "0711/0711.0648_arXiv.txt": {
        "abstract": "We present spectral measurements made in the soft (20--100\\,keV) gamma-ray band of the region containing the composite supernova remnant G11.2-0.3 and its associated pulsar \\hbox{PSR~J1811$-$1925}. Analysis of INTEGRAL/IBIS data allows characterisation of the system above 10\\,keV. The IBIS spectrum is best fitted by a power law having photon index $\\Gamma =$ 1.8 $_{-0.3}^{+0.4}$ and a 20--100\\,keV flux of 1.5 $\\times $10$^{-11}$ erg cm$^{-2}$ s$^{-1}$. Analysis of archival \\textit{Chandra} data  over different energy bands rules out the supernova shell as the site of the soft gamma-ray emission while broad band (1--200\\,keV) spectral analysis strongly indicates that the INTEGRAL/IBIS photons originate in the central zone of the system which contains both the pulsar and its nebula. The composite X-ray and soft gamma-ray spectrum indicates that the pulsar provides around half of the emission seen in the soft gamma-ray domain; its spectrum is hard with no sign of a cut off up to at least 80\\,keV. The other half of the emission above 10\\,keV comes from the PWN; with a $\\Gamma$=1.7 its spectrum is softer than that of the pulsar.  From the IBIS/ISGRI mosaics we are able to derive 2$\\sigma $ upper limits for the 20--100\\,keV flux from the location of the nearby TeV source \\hbox{HESS~J1809$-$193} to be 4.8 $\\times $10$^{-12} $ erg cm$^{-2}$ s$^{-1}$. We have also examined the likelihood of an association between \\hbox{PSR~J1811$-$1925} and \\hbox{HESS~J1809$-$193}. Although PSR J1811$-$1925 is the most energetic pulsar in the region, the only one detected above 10 keV and thus a possible source of energy to fuel the TeV fluxes, there is no morphological evidence to support this pairing, making it an unlikely counterpart. ",
        "introduction": "Recently, \\textit{Chandra~} and \\textit{XMM-Newton~} X-ray observations have allowed detailed morphological studies of Pulsar Wind Nebulae (PWN), resolving complex structures within the general envelope of the plerion, including toroidal features, axial jets and wisps on arc second angular scales; for recent reviews see Kaspi et al. (2006) and Gaensler \\& Slane (2006). The PWN are generally associated with young energetic pulsars with ages of less than a few tens of thousands of years such as the cases of the Crab and Vela systems (Weisskopf et al. 2004, Kirsch et al. 2006; Helfand et al 2001; Pavlov et al. 2003; Mangano et al. 2005 ). It is generally thought that the radiation mechanism from radio to X-ray wavelengths is due to synchrotron emission as a result of electrons injected by the pulsar, with the radio emission due to the relic electrons and X-rays resulting from freshly injected electrons. The diminishing angular sizes of the spectral images with increasing photon energy provide strong evidence for synchrotron cooling in a number of cases. However the \\textit{Chandra~} spectral images show that the separate morphological regions within the overall nebular configuration, such as a jet, counter jet, torus, extended PWN, moving hot spots and so on are characterised by distinct spectral indices in the X-ray domain. This indicates that the transport of the energetic electrons is not simply ``isotropically outwards'' from a central source, but can be the result of collimated beams, sometimes containing moving blobs with considerably harder spectra than the surrounding emission structures (e.g. Mori et al. 2004 for the case of the Crab nebula).  The study of these objects in the soft gamma-ray domain, adjacent to the X-ray band, can provide key information to disentangle the mechanisms active in different emitting regions. In view of this, the IBIS gamma-ray imager on board INTEGRAL is a powerful tool for the study of PWN systems, allowing source detection above 20\\thinspace keV with sensitivity of the order of mCrab. Moreover, the recent survey performed with the HESS (High Energy Stereoscopic System) telescope array revealed the existence of emission at TeV energies from the vicinity of a number of pulsar wind nebulae (Aharonian et al. 2006; 2007). The combined analysis of X-/gamma-ray and Very High Energy (VHE) emission is of key importance to estimate the relative contribution of synchrotron and inverse Compton mechanisms and thus to derive information on the magnetic field of the nebula and the spatial energy distribution of the accelerated particles. ",
        "conclusions": "The INTEGRAL/IBIS telescope has detected soft $\\gamma $-ray emission from the site of \\hbox{PSR~J1811$-$1925}. Comparison of the IBIS/ISGRI and \\textit{Chandra} spectra suggests that the soft $\\gamma $-ray emission is not coming from the SNR but originates within the complex structure associated with the pulsar and its PWN system. The composite X-ray and soft $\\gamma $-ray spectrum indicates the pulsar itself provides a greater fraction of the emission seen by IBIS; its spectrum is a hard power law showing no sign of a cut off up to at least 80\\,keV; the PWN contribution to the flux above 10\\,keV is smaller and well described by a steeper power law than the pulsar.  Hard ($>2$\\,keV) pulsed X-ray emission is believed to come from the magnetosphere of the rotating neutron star and is dominated by non-thermal radiation. Two kinds of models can explain this non-thermal pulsed X/$\\gamma $-ray emission: the polar cap model (Zhang \\& Harding 2000) and the outer gap model (Cheng, Ho \\& Ruderman 1986a,b; Zhang \\& Cheng 1997; Wang et al. 1998; Cheng et al. 1998).  In the outer gap model, the high energy emission would originate from the tail of the synchrotron component, while in the polar cap model it would stem from curvature radiation. In the latter case a spectral cut-off is expected at GeV energies while in the former a break is predicted at lower MeV frequencies.  Either way, the spectral turnover is expected above the IBIS/ISGRI energy band, and broader spectral sampling is required to place constraints on any spectral break.  It is interesting to note, that a fraction of the 20--100\\,keV emission, roughly 50\\%, is coming from the PWN associated to \\hbox{PSR~J1811$-$1925}; i.e. this is one of the few cases (notably Crab and Vela) where emission above 10\\,keV is detected from a PWN.  The nature of this hard X-ray emission is not yet clear, as both the synchrotron and inverse Compton processes are likely emission mechanisms. Perhaps, the Spectral Energy Distribution (SED) of INTEGRAL detected PWN could be used to tackle this issue, provided that in all cases we can separate the pulsar from the PWN as we have done here for \\hbox{PSR~ J1811$-$1925}. Comparison with better studied objects like Crab and Vela can help in discriminating between competing models and provide more insight into the physics of PWN.  To reach this objective the close relationship between VHE gamma-ray  emission and energetic young pulsars, and in a number of cases INTEGRAL sources, can be of valuable help. It is a pity that in the particular case of \\hbox{PSR~J1811$-$1925} the association with \\hbox{HESS~J1809$-$193} is highly improbable. \\hbox{PSR~J1811$-$1925} is the most energetic young pulsar within the complex region surrounding the TeV source, but not the only one capable of providing the requisite energy.  \\hbox{PSR~J1809$-$1917}, on balance, provides a more plausible counterpart, although it is still quite possible that neither fuels the VHE emitter, which could instead be powered by one of the supernova remnants which exist in the region."
    },
    "0711/0711.3682_arXiv.txt": {
        "abstract": "{Dwarf Spheroidal galaxies are amongst the best targets to search for a Dark Matter (DM) annihilation signal. The annihilation of WIMPs in the center of Sagittarius dwarf spheroidal (Sgr dSph) galaxy would produce high energy $\\gamma$-rays in the final state. Observations carried out with the H.E.S.S. array of Imaging Atmospheric Cherenkov telescopes are presented. A careful modelling of the Dark Matter halo profile of Sgr dwarf was performed using latest measurements on its structural parameters. Constraints on the velocity-weighted cross section of Dark Matter particles are derived in the framework of Supersymmetric and Kaluza-Klein models. }     \\begin{document} ",
        "introduction": "Astrophysical and cosmological observations provide a substantial body of evidences for the existence of Cold Dark Matter (CDM) although its nature remains still unknown. It is commonly assumed that CDM is composed of yet undiscovered non-baryonic particles for which plausible candidates are Weakly Interacting Massive Particles (WIMPs). In most theories, candidates for CDM are predicted in theories beyond the Standard Model of particle physics~\\cite{bertone}. The annihilation of WIMPs into $\\gamma$-rays may lead to detectable very high energy (VHE, E $>$ 100 GeV) $\\gamma$-ray fluxes above background via continuum emission from the hadronization and decay of the cascading annihilation products, predominantly from $\\pi^0$'s generated in the quark jets. Among the best-motivated CDM candidates are the lightest neutralino $\\tilde{\\chi}$ provided by R-parity conserving supersymmetric extensions of the Standard Model~\\cite{jungman}, and the lightest Kaluza-Klein particle (LKP)~\\cite{servant} % in universal extra dimension theories which is most often the first KK mode of the hypercharge gauge boson, $\\tilde{B}^{(1)}$. The H.E.S.S. array of Imaging Atmospheric Cherenkov Telescopes (IACTs), designed for high sensitivity measurements in the 100 GeV - 10 TeV energy regime, is a suitable instrument to detect VHE $\\gamma$-rays and investigate their possible origin.\\\\ Dwarf Spheroidal galaxies such as Sagittarius or Canis Major, discovered recently in the Local Group, are among the most extreme DM-dominated environments. Indeed, measurements of roughly constant radial velocity dispersion of stars imply large mass to luminosity ratios~\\cite{wilkinson}. The core of the Sgr dSph at l=5.6$^{\\circ}$ and b=-14$^{\\circ}$ in galactic coordinates at a distance of about 24 kpc from the Sun~\\cite{majewski}. Latest velocity dispersion measurements on M giant stars with 2MASS yields a mass to light ratio of about 25~\\cite{majewski2}. The luminous density profile of Sgr dSph has two components~\\cite{monaco}. The compact component, namely the core, is characterized by a size of about 3 pc FWHM, which corresponds to a point-like region for H.E.S.S. This is the DM annihilation region from which $\\gamma$-ray signal may be expected. A diffuse component is well fitted by a King model with a characteristic size of 1.6 kpc.\\\\ We present in this paper the observations of the Sgr dSph galaxy by the H.E.S.S. array of Imaging Atmospheric Cherenkov Telescopes. A careful modeling of the Dark Matter halo using the latest measurements on the structural parameters of Sagittarius is presented to derive constraints on the WIMP velocity-weighted annihilation rate. ",
        "conclusions": "The observations of Sgr dSph with H.E.S.S. reveal no significant $\\gamma$-ray excess at the nominal target position. The Sagittarius dwarf DM halo profile has been modeled using latest measurements of its structure parameters. Constraints have been derived on the velocity-weighted cross section of the DM particle in the framework of supersymmetric and Kaluza-Klein models."
    },
    "0711/0711.1364_arXiv.txt": {
        "abstract": "A general class of cosmological models driven by a nonlocal scalar field inspired by the string field theory is studied. Using the fact that the considering linear nonlocal model is equivalent to an infinite number of local models we have found an exact special solution of the nonlocal Friedmann equations. This solution describes a monotonically increasing Universe with the phantom dark energy. ",
        "introduction": "Recently string theory and brane cosmology have been intensively discussed as promising candidates for the theoretical explanation of the obtained experimental data  (see for example \\cite{string-cosmo}--\\cite{Biswas}).  The purpose of this paper is to present new results concerning studies of nonlocal linear models in the Friedmann--Robertson--Walker Universe. These models are inspired by the string field theory (SFT) (for review of the SFT see~\\cite{review-sft}). A distinguished feature of nonlocal linear and nonlinear models~\\cite{IA1}--\\cite{Jukovskaya0710} is the presence of infinite number of higher derivative terms (note also nonlocal models in the Minkowski space-time~\\cite{Woodard}--\\cite{LJ}). For special values of the parameters these models describe linear approximations to the cubic bosonic or nonBPS fermionic SFT nonlocal tachyon models, p-adic string models or the models with the invariance of the action under the shift of the dilaton field to a constant. The NonBPS fermionic string field tachyon nonlocal model has been considered as a candidate for the dark energy~\\cite{IA1}.  Present cosmological observations \\cite{cosmo-obser} do not exclude an evolving dark energy (DE) state parameter $w$, whose current value can be less than  $-1$, that means the violation of the null energy condition (NEC) (see \\cite{DE1,DE2} for a review of the DE problems and \\cite{0312430} for a search  for a super-acceleration phase of the Universe).  Field theories, which violate the NEC~\\cite{Hawking-Ellis,Venziano}, are of interest not only for the construction of cosmological dark energy models with the state parameter $w<-1$, but also for the solution of the cosmological singularity problem. A possible way to avoid cosmological singularities consists of dealing with nonsingular bouncing cosmological solutions. In this scenario the Universe contracts before the bounce~\\cite{Turok}. Such models have strong coupling and higher-order string corrections are inevitable. It is important to construct nonsingular bouncing cosmological solutions in order to make a concrete prediction of bouncing cosmology.  A simple possibility to violate the NEC is just to deal with a phantom field. In the present paper we consider nonlocal models which are linear and admit solutions, which are linear combinations of local fields. Some of these local fields are phantoms. Namely due to the presence of these ghost excitations such nonlocal models present an interest for cosmology.  At the same time there are well known problems with instability of quantum models with phantoms, namely a lost of unitarity and so on. We believe that nonlocal SFT models in true vacua are stable with respect to quantum fluctuations. This question has to be consider in the full SFT framework and demands further investigations. We also believe that due to these string theory origin the corresponding nonlocal cosmological models, which are nonlinear in matter fields, have no problem with instability in the quantum case. In this paper we consider only the classical case and models, which are linear in a nonlocal scalar field.  In our previous paper~\\cite{AJV0701184} as well as in paper~\\cite{AVzeta} nonlocal linear models already have been studied. In~\\cite{AJV0701184} nonlocal linear model has been studied in the flat space-time and we have proposed special deformations of the potential, which allows us to get the same scalar field solutions in flat and nonflat (the FRW metric) cases. As result we have obtained nonlinear models in the FRW metric. In~\\cite{AVzeta} few exact solutions to linear model in the FRW metric have been found. In this paper we present a systematic method that permits us to transform the initial nonlocal system into infinity set of local systems. The choice of a local system is equivalent to the choice of a special solution of the nonlocal system. This approach allows us to use the standard method of analysis of the differential equations and in particular to find exact solutions.  The paper is organized as follows. In Section 2 we describe string inspired models with quadratic nonlocal potentials. In Section 3 we assume that the metric is given and consider the equation of motion as an equation for the nonlocal scalar field. We construct solutions, using eigenfunctions of the $\\square_g$-operator with eigenvalues, belonging to the set of roots of the characteristic equation. In Section 4 we find values of the energy density and pressure for these solutions. In Section 5 we consider the Friedmann--Robertson--Walker Universe and find local models, which correspond to particular solutions of the initial nonlocal model. In the case of dilaton massless scalar field we construct the general solutions for the corresponding local model, which are the special exact solution for the initial nonlocal model as well. We analyze cosmological properties of the obtained solutions. ",
        "conclusions": "We have studied the SFT inspired linear nonlocal model. This model has an infinite number of higher derivative terms and are characterized by two positive parameters: $\\xi^2$ and $c$. For particular cases of the parameters $\\xi^2$ and $c$ the corresponding actions describe linear approximations to either the bosonic or nonBPS fermionic cubic SFT as well as  to the nonpolynomial SFT.  Roots of the characteristic equation do not depend on the form of the metric and this property allows us to study properties of energy density and pressure. We have found that in an arbitrary metric the energy-momentum tensor for an arbitrary N-mode solution is a sum of the energy-momentum tensors for the corresponding one-mode solutions. In the Friedmann--Robertson--Walker spatially flat metric the pressure for a one-mode solution corresponding to a real root can be positive or negative, depending on parameters of our nonlocal model. Namely, for $c\\leq 1$ the one mode pressure is positive and for $c>1$ it could be negative or positive.  The investigation performed in this paper shows that the general field equations in  linear nonlocal  models admit an equivalent description in terms of local theory and as a consequences we have representations (37) and (38) for the energy and pressure. This calculation also supports the use of the Ostrogradski representation for our system in the case of arbitrary metric.  In distinguish to our previous paper~\\cite{AJV0701184} we do not use any approximation scheme and do not add any terms in the potential. We have shown that our linear model with one nonlocal scalar field generates an infinite number of local models. These models can be studied numerically and we plan to present this analysis in future papers. Some of these models have been solved explicitly and, hence, special exact solutions for nonlocal model in the Friedmann--Robertson--Walker metric have been obtained. In particular we have constructed an exact kink-like solution, which correspond to monotonically increasing Universe with phantom dark energy. Note that the obtained behaviour of the Hubble parameter is close to behavior of the Hubble parameter in the nonlinear nonlocal model~\\cite{IA1}, which recently has been obtained numerically~\\cite{Jukovskaya0707}."
    },
    "0711/0711.1708_arXiv.txt": {
        "abstract": "I discuss the essential features of the QCD axion: the strong CP solution and hence its theoretical necessity. I also review the effects of the QCD axion on astrophysics and cosmology, in particular with emphasis on its role in the dark matter component together with its supersymmetric partner axino. It is pointed out that string theory may or may not give a detectable QCD axion. ",
        "introduction": "Cosmology with cold dark matter (CDM) was the leading candidate of the early universe in 1980s \\cite{Blum}. But this view has given a way to the new cosmology with the discovery of dark energy (DE) in 1998 \\cite{DE98}. The current view of the dominant components of the universe is $\\Omega_{\\rm CDM}\\simeq 0.23$ and $\\Omega_\\Lambda\\simeq 0.73$ \\cite{WMAP3}. The plausible DM candidates at present are the lightest SUSY particle (LSP), axion, axino, and gravitino. Here, we will review mostly on axion and its CDM-related possibility. The need for DM was first suggested by Zwicky in 1933 \\cite{Zwicky}.  There are numerous cases supporting the nonluminous dark matter in the universe: flat rotation curves, Chandra satellite photo, and gravitational lensing effects. If the galactic bulge is the dominant mass in the galaxy, the velocity $v$ of a star located at $r$ from the center behaves as $v\\sim r^{-1/2}$. But the observed flat rotation curve \\cite{flatrot} violates this expectation and implies the extended mass into the halo as $\\rho(r)\\sim 1/r^2$. Also, the Chandra observation of X-ray and gravitational lensing images implies the matter profile around the bullet galaxy \\cite{Chandra}. The circular gravitational lensing photo \\cite{gravlens} also supports the idea that DM exist. The DM density around us here is usually taken as $\\rho_{\\rm DM}\\simeq 0.3$ GeV/cm$^3$.  The proposed CDM candidates belong to either incoherent particles or coherent oscillations of spin-0 fields. The incoherent CDM particles are usually called weakly interacting massive particles(WIMP) or decay products of WIMPs. A coherent oscillation arises from the potential describing a bosonic particle. WIMPs are the massive particles with weak interaction cross sections, which was first discussed by B. W. Lee and S. Weinberg  in 1977 in terms of a heavy neutrino as shown in Fig. \\ref{LeeWein} \\cite{LeeWein}. \\begin{figure} \\begin{center} \\begin{picture}(400,260)(0,-30) \\Line(30,-20)(30,219.1) \\Line(230,-20)(230,219.1) \\Line(30,-20)(230,-20) \\Line(30,219.1)(230,219.1) \\Text(130,-40)[c]{Log($m_\\nu$[GeV])} \\rText(6,120)[t][l]{Log($\\Omega_{\\nu}h^2$)}  \\Text(27,0)[r]{$-1$} \\Text(27,40)[r]{$0$} \\Text(27,80)[r]{$1$} \\Text(27,120)[r]{$2$} \\Text(27,160)[r]{$3$}\\Text(27,200)[r]{$4$} \\Line(30,212)(33,212) \\Line(230,212)(227,212) \\Line(30,200)(36,200)  \\Line(230,200)(224,200) \\Line(30,188)(35,188) \\Line(230,188)(225,188) \\Line(30,184.1)(33,184.1) \\Line(230,184.1)(227,184.1) \\Line(30,179.1)(33,179.1) \\Line(230,179.1)(227,179.1) \\Line(30,172)(33,172) \\Line(230,172)(227,172) \\Line(30,160)(36,160)  \\Line(230,160)(224,160) \\Line(30,148)(35,148) \\Line(230,148)(225,148) \\Line(30,144.1)(33,144.1) \\Line(230,144.1)(227,144.1) \\Line(30,139.1)(33,139.1) \\Line(230,139.1)(227,139.1) \\Line(30,132)(33,132) \\Line(230,132)(227,132) \\Line(30,120)(36,120)  \\Line(230,120)(224,120) \\Line(30,108)(35,108) \\Line(230,108)(225,108) \\Line(30,104.1)(33,104.1) \\Line(230,104.1)(227,104.1) \\Line(30,99.1)(33,99.1) \\Line(230,99.1)(227,99.1) \\Line(30,92)(33,92) \\Line(230,92)(227,92) \\Line(30,80)(36,80)  \\Line(230,80)(224,80) \\Line(30,68)(35,68) \\Line(230,68)(225,68) \\Line(30,64.1)(33,64.1) \\Line(230,64.1)(227,64.1) \\Line(30,59.1)(33,59.1) \\Line(230,59.1)(227,59.1) \\Line(30,52)(33,52) \\Line(230,52)(227,52) \\Line(30,40)(36,40)  \\Line(230,40)(224,40) \\Line(30,28)(35,28) \\Line(230,28)(225,28) \\Line(30,24.1)(33,24.1) \\Line(230,24.1)(227,24.1) \\Line(30,19.1)(33,19.1) \\Line(230,19.1)(227,19.1) \\Line(30,12)(33,12) \\Line(230,12)(227,12) \\Line(30,0)(36,0)  \\Line(230,0)(224,0) \\Line(30,-12)(35,-12) \\Line(230,-12)(225,-12) \\Line(30,-15.9)(33,-15.9) \\Line(230,-15.9)(227,-15.9)  \\Text(50,-27)[c]{${-7}$}\\Text(70,-27)[c]{${-6}$} \\Text(90,-27)[c]{${-5}$}\\Text(110,-27)[c]{${-4}$} \\Text(130,-27)[c]{${-3}$}\\Text(150,-27)[c]{${-2}$} \\Text(170,-27)[c]{${-1}$}\\Text(190,-27)[c]{$0$} \\Text(210,-27)[c]{${1}$}  \\Line(44,-20)(44,-17) \\Line(44,219.1)(44,216.1) {\\SetWidth{1}\\SetColor{Blue} \\Line(50,-20)(50,-14) \\Line(50,219.1)(50,213.1) \\Text(50,-9)[c]{$10^2$eV}} \\Line(64,-20)(64,-17) \\Line(64,219.1)(64,216.1) \\Line(70,-20)(70,-14)  \\Line(70,219.1)(70,213.1) \\Line(84,-20)(84,-17) \\Line(84,219.1)(84,216.1) \\Line(90,-20)(90,-14)  \\Line(90,219.1)(90,213.1) \\Line(104,-20)(104,-17) \\Line(104,219.1)(104,216.1) \\Line(110,-20)(110,-14)  \\Line(110,219.1)(110,213.1) \\Line(124,-20)(124,-17) \\Line(124,219.1)(124,216.1) \\Line(130,-20)(130,-14)  \\Line(130,219.1)(130,213.1) \\Line(144,-20)(144,-17) \\Line(144,219.1)(144,216.1) \\Line(150,-20)(150,-14)  \\Line(150,219.1)(150,213.1) \\Line(164,-20)(164,-17) \\Line(164,219.1)(164,216.1) \\Line(170,-20)(170,-14)  \\Line(170,219.1)(170,213.1) \\Line(184,-20)(184,-17) \\Line(184,219.1)(184,216.1) {\\SetWidth{1}\\SetColor{Blue} \\Line(190,-20)(190,-14) \\Line(190,219.1)(190,213.1) \\Text(190,-9)[c]{GeV}} \\Line(204,-20)(204,-17) \\Line(204,219.1)(204,216.1) \\Line(210,-20)(210,-14)  \\Line(210,219.1)(210,213.1) \\Line(224,-20)(224,-17) \\Line(224,219.1)(224,216.1)  {\\SetWidth{1} \\Line(30,2)(125.6,191.1) \\Curve{(125.6,191.1)(136,196.8)(145.6,191.1)} {\\SetColor{Blue} \\Curve{(145.6,191.1)(157.6,164)(172.3,120) (206,0)(211,-20)}}% {\\SetColor{Red} \\Curve{(157.2,191.1)(178,124)(202,42)(213,0) (217.5,-20)}% \\Curve{(125.6,191.1)(139.8,203.9)(157.1,191.1)} } {\\SetColor{Brown} \\DashLine(30,2)(230,2){4} } } \\Text(185,70)[r]{Dirac} \\Text(176,140)[l]{Majorana}  \\end{picture} \\end{center} \\caption{A schematic view of the Lee-Weinberg calculation of $\\Omega_\\nu h^2$. For the larger mass region, the freezeout temperature is $T_f\\approx m_\\nu/20$ [8]. The dash line is for $\\Omega_\\nu h^2=0.113$ [9].}\\label{LeeWein} \\end{figure} This implies that the usage $\\lq\\lq$weak\" is involved in WIMP. SUSY with the R-parity conservation allows the LSP as such a particle. The LSP interaction is $\\lq\\lq$weak\" since the interaction mediators (SUSY particles) are supposed to be in the 100 GeV range. For a WIMP to be a successful CDM candidate, usually the WIMP interaction cross section at the time of decoupling is needed around \\cite{WMAPsigma} \\begin{align} &\\langle\\sigma_{\\rm int} v\\rangle|_{\\rm at\\ decoupling}\\approx 0.2\\times 10^{-26} ~{\\rm cm^3s}^{-1},\\nonumber\\\\ &\\quad\\quad {\\rm with}\\ \\Omega_m h^2\\simeq 0.113\\pm 0.009.\\label{WIMPcond} \\end{align} This is roughly explained by the LSP with the low energy SUSY, which is the reason that we are so much concerned about the WIMP LSP. The proposed particles are shown in the $\\sigma_{\\rm int}$ versus mass plane in Fig. \\ref{sigmavsmass} with minor modification from that of \\cite{Roszfig}. Currently, there are experimental efforts to discover the LSP as predicted by SUSY models. Also, direct cosmological searches are going on \\cite{Dama}. At the LHC, the probable LSP mass ranges will be looked for. \\begin{figure} \\begin{center} \\begin{picture}(400,240)(10,-10) \\SetScale{0.8} \\Line(50,10)(300,10) \\Line(50,10)(50,285) \\rText(10,150)[t][l]{Log~$\\sigma_{\\rm int}$ [cm$^2$]} \\Text(140,-15)[c]{$m_{\\rm DM}$ [GeV]} \\Line(300,10)(300,285)\\Line(50,285)(300,285)  \\Line(50,280)(55,280)\\Line(295,280)(300,280) \\Line(50,275)(55,275)\\Line(295,275)(300,275) \\Line(50,270)(55,270)\\Line(295,270)(300,270) \\Line(50,265)(55,265)\\Line(295,265)(300,265) \\Line(50,260)(57,260)\\Text(38,207.5)[r]{${-35}$} \\Line(293,260)(300,260) \\CBox(110,248.5)(119,244){Red}{Red} \\Line(50,255)(55,255)\\Line(295,255)(300,255) \\Line(50,250)(55,250)\\Line(295,250)(300,250) \\Line(50,245)(55,245)\\Line(295,245)(300,245) {\\SetWidth{1.5}\\SetColor{Blue} \\Line(50,240)(55,240)\\Line(295,240)(300,240)} \\Text(48,194)[l]{fb} \\Line(50,235)(57,235)\\Text(38,187.5)[r]{${-40}$} \\Line(293,235)(300,235) \\Line(50,230)(55,230)\\Line(295,230)(300,230) 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\\Line(268.75,10)(268.75,15)\\Line(268.75,285)(268.75,280) \\Line(275,10)(275,17)\\Line(275,285)(275,278) \\Line(281.25,10)(281.25,15)\\Line(281.25,285)(281.25,280) \\Line(287.5,10)(287.5,15)\\Line(287.5,285)(287.5,280) \\Line(293.75,10)(293.75,15)\\Line(293.75,285)(293.75,280) \\Line(300,10)(300,17)\\Line(300,285)(300,253) \\Text(240,0)[c]{$10^{20}$}   \\SetWidth{0.5} \\Text(75,190)[l]{neutrino $\\nu$} \\Text(50,137)[l]{axion $a$} \\Text(120,109)[l]{axino ${\\tilde a}$} \\Text(111,59)[l]{gravitino ${\\tilde g_{3/2}}$} \\rText(167,183)[t][l]{WIMP} \\Text(180,210)[r]{neutralino ${\\chi}$} \\rText(218,150)[t][l]{wimpzilla} {\\SetWidth{1.2}\\SetColor{Brown}\\DashLine(187,252)(187,235){3} \\SetColor{Yellow}\\DashLine(187,235)(187,219){3} \\LongArrow(187,222)(187,218)}  \\end{picture} \\end{center} \\caption{Some proposed particles in the interaction cross section versus mass plane.}\\label{sigmavsmass} \\end{figure}    It is known that the density perturbation grew much earlier than the time of recombination. If it grew after the recombination time, the density perturbation grown afterward was not enough to make galaxies. For galaxy formation, therefore, DM is needed since proton density perturbation could not grow before the recombination time, but DM could. With DM, the equality point of radiation and matter energy densities can occur much earlier than the recombination time since DM is not prohibited in collapsing by Silk damping \\cite{Silkdamp}. If the WIMP mass and interaction cross section falls in the region allowed by Eq. (\\ref{WIMPcond}), it can be CDM. In this talk, we consider CDM as the LSP and in addition axion also. But if the LSP is the only the CDM component, then the LSP mass would give one number for the DM density, which may not be the case. Thus, even if the LSP is contributing to the CDM density, we may need axion to account for the right amount of CDM around us. ",
        "conclusions": "The popular CDM candidates are WIMPs and very light axions. We know that they must fill the universe based on the observational grounds and by the fact that we exist here in a planet. Direct searches for WIMPs in the universe use the WIMP cross section in our environment. The LHC machine will tell whether the LSP mass falls in the CDM needed range or not. The other candidate a very light axion, whether or not it is the dominant CDM component, is believed to exist from the need for a solution of the strong CP problem. So, I reviewed the axion and the related issues: solutions of the strong CP problem, the axion CDM possibility, the axino CDM possibility with  order GeV axino mass, and an observable QCD axion possibility from superstring.  Maybe solar axions are easier to be detected than cosmic axions. Then, axion is not the dominant component of CDM.  Most exciting however would be that axion is discovered and its discovery confirms instanton physics of QCD (by experiments).\\\\  \\noindent {\\bf Acknowledgments\\ :} This work is supported in part by the KRF Grants, No. R14-2003-012-01001-0 and No. KRF-2005-084-C00001."
    },
    "0711/0711.3227_arXiv.txt": {
        "abstract": "White dwarfs represent the endpoint of stellar evolution for stars with initial masses between approximately 0.07 \\msun~and 8-10 \\msun, where \\msun~is the mass of the Sun (more massive stars end their life as either black holes or neutron stars). The theory of stellar evolution predicts that the majority of white dwarfs have a core made of carbon and oxygen, which itself is surrounded by a helium layer and, for $\\sim$80 per cent of known white dwarfs, by an additional hydrogen layer$^{1-3}$. All white dwarfs therefore have been traditionally found to belong to one of two categories: those with a hydrogen-rich atmosphere (the DA spectral type) and those with a helium-rich atmosphere (the non-DAs). Here we report the discovery of several white dwarfs with atmospheres primarily composed of carbon, with little or no trace of hydrogen or helium. Our analysis shows that the atmospheric parameters found for these stars do not fit satisfactorily in any of the currently known theories of post-asymptotic giant branch evolution, although these objects might be the cooler counterpart of the unique and extensively studied PG1159 star H1504+65 (refs 4-7). These stars, together with H1504+65, might accordingly form a new evolutionary sequence that follow the asymptotic giant branch. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.0297_arXiv.txt": {
        "abstract": "\\noindent We present phase resolved optical photometry and spectroscopy of the accreting millisecond pulsar HETE~J1900.1$-$2455. Our $R$-band light curves exhibit a sinusoidal modulation, at close to the orbital period, which we initially attributed to X-ray heating of the irradiated face of the secondary star. However, further analysis reveals that the source of the modulation is more likely due to superhumps caused by a precessing accretion disc. Doppler tomography of a broad H$\\alpha$ emission line reveals an emission ring, consistent with that expected from an accretion disc. Using the velocity of the emission ring as an estimate for the projected outer disc velocity, we constrain the maximum projected velocity of the secondary to be 200 km~s$^{-1}$, placing a lower limit of 0.05 M$_{\\sun}$ on the secondary mass. For a 1.4 M$_{\\sun}$ primary, this implies that the orbital inclination is low, $\\lesssim$ 20\\degr. Utilizing the observed relationship between the secondary mass and orbital period in short period cataclysmic variables, we estimate the secondary mass to be $\\sim$0.085 M$_{\\sun}$, which implies an upper limit of $\\sim$2.4 M$_{\\sun}$ for the primary mass. ",
        "introduction": "Low mass X-ray binaries (LMXBs) are close binary systems consisting of a degenerate primary -- black hole or neutron star (NS) -- accreting matter from a low mass ($<$ 1 M$_{\\sun}$) secondary, via Roche lobe overflow. X-ray transients (XRT) are LMXBs which undergo periodic outbursts, explained by the disc instability model \\citep[e.g.][]{dubus2001}, while persistently bright systems are permanently in the outbursting state. One subclass of LMXBs are the accretion-powered millisecond pulsars (AMSPs), the first of which, SAX J1808.4$-$3658, was discovered in 1998 \\citep{wijnands1998}. Seven more of these systems have been discovered since then \\citep{poutanen2006, wijnands2006, markwardt2007}.  \\vspace{1mm}  In general, the outburst duration for AMSPs ranges from a few weeks to months. AMSPs generally contain a weakly magnetized ($\\sim$10$^8$--10$^9$ G) NS with spin frequencies between 180 and 600 Hz. The orbital periods range between 40 min and 5 h. The secondary star in these systems is either a white or brown dwarf \\citep[][and references therein]{falanga2007}. \\cite{falanga2005} have measured the predicted decrease in the NS spin period for the AMSP system IGR J00291$+$5934, supporting the idea that these AMSPs are in fact old NS, which have over time been spun up to millisecond periods by acquiring angular momentum from the accretion of material from the secondary. As such, they provide the missing link between LMXBs and old, isolated millisecond radio pulsars.  \\vspace{1mm}  In comparison to brighter LMXBs, relatively little is known about the optical properties of AMSPs. Even in outburst, the eight so far discovered are comparatively faint, the brightest being SAX J1808.4$-$3658 at $R$ $\\simeq$ 16.2 mag \\citep{wang2001}. In quiescence, these systems are generally very faint optically, the brightest measured again being SAX J1808.4$-$3658 at $R$ $\\simeq$ 20.9 mag.  \\vspace{1mm}  HETE~J1900.1$-$2455 is the seventh AMSP, discovered by \\cite{vanderspek2005} on 2005 June 14 by the {\\textit{High Energy Transient Explorer II}} ({\\textit{HETE II}}). Several type-I bursts have since been observed, and assuming that the peak flux observed during the brightest of these is Eddington-limited, HETE~J1900.1$-$2455 is at a distance of $\\sim$4.3 kpc \\citep{suzuki2007}.  \\vspace{1mm}  The pulsation frequency was determined with the {\\textit{Rossi X-ray Transient Explorer}} to be 377.3 Hz \\citep{morgan2005}. Pulse timing analysis by \\cite{kaaret2006} revealed a circular orbit with a period ($P_{\\mathrm{orb}}$) of $\\sim$83.3 minutes (4995.258 $\\pm$ 0.005 s), and a projected primary semimajor axis, $a_{1}\\mathrm{sin}i$, of 18.41 $\\pm$ 0.01 $\\times$ 10$^{-3}$ light second. These parameters combine to provide a mass function $f(M) = 2.004 \\pm 0.003 \\times10^{-6}\\ \\mathrm{M_{\\sun}}$. \\cite{galloway2006} report that pulses were detected intermittently in the first two months after discovery. No pulses have been reported since then.  \\vspace{1mm}  Unusually for an XRT, this system has remained active for more than 2 years, at a level of $\\sim$2 $\\times$ 10$^{36}$ erg s$^{-1}$. Several times during 2007 March/February, the X-ray flux fell by an order of magnitude, but each time returned to outburst levels within $\\sim$1 week \\citep{galloway2007}. In 2007 May, the source was observed to decline to its lowest level since discovery (by a factor of $>$ 10$^{3}$), lasting for several weeks, prompting speculation that HETE~J1900.1$-$2455 was returning to quiescence \\citep{degenaar2007a,galloway2007,torres2007}. However, the system returned to its outburst state within $\\sim$2--3 weeks \\citep{degenaar2007b,garnavich2007}.  \\vspace{1mm}  The optical counterpart was found by \\cite{fox2005} with an $R$-band magnitude of $\\sim$18.4 mag. Subsequent observations by \\cite{steeghs2005} found the $R$-band magnitude to be 18.02 $\\pm$ 0.03 mag, and the $V-R$ colour (dereddened) to be $-0.16$ mag, with spectroscopy revealing a broad \\ion{He}{ii} $\\lambda$4686 \\AA\\ emission line.  \\vspace{1mm}  In this paper we present optical photometry and spectroscopy of HETE~J1900.1$-$2455 in its outburst state. These observations were motivated by several goals, key amongst them the aim of measuring the radial velocity of the secondary, and hence constraining the mass of the compact object. For the longer period transient systems, such measurements are possible via radial velocity studies of absorption lines from the secondary once the system has returned to quiescence. For the shorter period AMSPs however, the secondary is likely to be so faint as to make such observations extremely difficult, if not impossible. However, \\cite{steeghs2002} and \\cite{casares2004} have shown how radial velocity measurements of the secondary are still possible for X-ray bright systems, by studying the Bowen blend emission from the irradiated face of the secondary: this technique may offer the only way of measuring the velocity of the secondary in AMSPs (i.e. whilst they are still X-ray bright), and usefully constraining the mass of the neutron star. ",
        "conclusions": "We have presented a detailed optical study of the AMSP HETE~J1900.1$-$2455. The $R$-band light curves reveal a modulation whose phase is inconsistent with simple heating of the secondary star/accretion disc: we interpret this as due to superhumps associated with a precessing accretion disc. Fitting to the wings of the averaged H$\\alpha$ emission line gives the systemic velocity of 132~$\\pm$~8~km~s$^{-1}$, and a similar value is found from the \\ion{He}{ii} line.  Doppler tomography of the H$\\alpha$ emission line suggests a upper limit of $\\sim$200~km~s$^{-1}$ on the projected secondary velocity, implying that HETE~J1900.1$-$2455 has a low inclination, $\\lesssim$ 20\\degr, and a minimum secondary mass of 0.05~M$_{\\sun}$.  Assuming that the secondary in HETE~J1900.1$-$2455 is similar to those seen in short period CVs, we estimate a secondary mass of 0.085~M$_{\\sun}$, which implies an upper limit to the NS mass of 2.4 M$_{\\sun}$ (again for $K_2$ $\\lesssim$ 200~km~s$^{-1}$). If the secondary is a brown dwarf, then this constraint is relaxed."
    },
    "0711/0711.4361_arXiv.txt": {
        "abstract": "We present three-dimensional, self-consistent radiative transfer solutions with a new Monte Carlo radiative equilibrium code.  The code, RADISHE ($\\bf{RAD}$iative transfer $\\bf{I}$n $\\bf{S}$moothed particle $\\bf{H}$ydrodynamics and $\\bf{E}$ulerian codes), can be applied to calculate the emergent spectral energy distributions (SEDs) and broadband images from optical to millimeter wavelengths of arbitrary density geometries with distributed sources of radiation.  One of the primary uses of this code has been to interface with hydrodynamical codes to calculate emergent SEDs along a simulation time sequence.  We focus the applications of this code in this paper on infrared bright dusty galaxies, but RADISHE is also ideal for calculating emission from star clusters or protostellar environments.  The primary methodological focus of this paper is on the radiative equilibrium temperature calculation.  We find that an iterative calculation of the temperature, which takes as the Monte Carlo estimator for the mean free intensity the sum of photon flight paths, is significantly faster than relaxation temperature calculation methods, particularly when large numbers of grid cells are required, i.e., in modeling three-dimensional geometries such as the dust envelopes of turbulent massive protostellar cores or infrared bright galaxies.  An accurate long wavelength SED and corresponding temperature calculation will be essential for analyzing upcoming $\\it{Herschel~Space~Observatory}$ observations.  We present simulated color-color plots for infrared bright galaxies at a range of redshifts, and unfold these plots as color vs the fractional AGN luminosity, to demonstrate that $\\it{Herschel}$ will be able to effectively discriminate between submillimeter galaxies where the energy source is dominated by AGN and those where star formation dominates.  We demarcate in particular the ``Class II'' or energetically active AGN evolutionary phase in $\\it{Herschel}$ color-color plots. ",
        "introduction": "Ultraluminous Infrared Galaxies (ULIRGs) radiate more than $1\\times 10^{12}L_{\\odot}$ between $8-1000~\\micron$ (Rieke \\& Low 1972; Soifer et al. 1984; Aaronson et al. 1984; Sanders \\& Mirabel 1996).   Submillimeter galaxies (SMGs), a high redshift population of dusty, infrared-bright galaxies, ($F_{850~\\micron} \\ga 1~\\rm mJy$ (Smail et al. 1997; Blain et al. 2002), of which Chapman et al. (2005) were able to obtain spectroscopic redshifts for, guided by optical and radio associations, have been studied with SCUBA (Smail et al. 1997; 2002; Ivison et al. 2000; Ivison et al. 2007), more recently with $\\it{Spitzer}$ (Pope et al. 2006), with $\\it{Chandra}$ (Alexander et al. 2005), and SHARC-2 (Kovacs et al. 2006).  These observations reveal that SMGs are a cosmologically significant (Blain et al. 2002), diverse (Ivison et al. 2000) population of galaxies, with some SMGs harboring obscured AGN (Alexander et al. 2005; Pope et al. 2006).  The Kovacs et al. (2006) sample of $350~\\micron$ observations is currently the most direct probe of the rest-frame far-IR of the SMG population at $z \\sim 2$.  Aside from this set of observations, there are few rest-frame far-IR observations of $z \\sim 2$ SMGs.  As such, the bolometric luminosities of this population remain uncertain for the majority of this population, as the far-IR luminosity is comparable to the bolometric luminosity.  The $\\it{Herschel~Space~Observatory}$ is expected to secure rest-frame far-IR observations of high redshift infrared bright galaxies, which combined with existing data sets can provide for the first time a panchromatic view of the dust emission from this population.  A new development in the last several years is the use of Smoothed Particle Hydrodynamics (SPH) and Eulerian codes to routinely calculate the time evolution of the gaseous and collionless components of galaxies undergoing mergers or evolving quiescently (Kravtsov 2003; Chakrabarti et al. 2003; Springel et al. 2005; Governato et al. 2007; Cox et al. 2006; Kazantzidis et al. 2007).  Calculating the emergent SEDs and images along a time sequence of simulations of merging galaxies using a self-consistent radiative transfer code can allow us to build an interface between simulations of galaxy mergers and observations of dusty galaxies like ULIRGs.  To aid in the interpretation of current and upcoming infrared observations, we have developed a self-consistent, three-dimensional Monte Carlo radiative transfer code.  It is nicknamed $\\bf{RADISHE}$ (RADiative transfer In Smooth particle Hydrodynamics and Eulerian codes), as it has been primarily employed to interface with the time outputs of hydrodynamical simulations.  Chakrabarti et al. (2007a;b) employed RADISHE to calculate the emergent SEDs of galaxy mergers.  In particular, Chakrabarti et al. (2007a) used an early version of RADISHE to demonstrate that the cold-warm IRAS classfication of galaxies (de Grijp et al. 1985) can be reproduced by galaxy merger simulations and can be explained due to AGN feedback.  Chakrabarti et al. (2007b) demonstrated that trends in observed color-color plots can be explained by decomposing simulated color-color plots as a function of time or some intrinsic galaxy parameter, such as the ratio of the black hole to stellar bolometric luminosity.  Chakrabarti et al. (2007b) also proposed an evolutionary model for SMGs, where they evolve from high ratios of $L_{\\rm IR}/L_{\\rm x}$ (Class I) through a bright x-ray bright phase (Class II), to low values of $L_{\\rm IR}/L_{\\rm x}$ typical of merger remnants (Class III).  The Monte Carlo method is ideal for treating inhomogeneous density distributions where there is no clear symmetry.  Boundary conditions, which must be dealt with explicitly in finite-difference codes (Steinacker et al. 2003), and which often become difficult to treat in multi-dimensional finite-difference codes, are naturally accomodated in the Monte Carlo method as long as the photon statistics are robust.  Analytic methods are essential to reduce the solution of the radiative transfer problem to a few key parameters under simplifying assumptions (Chakrabarti \\& McKee 2005; Chakrabarti \\& McKee 2007), while numerical methods are needed to model the effects of inhomogeneous geometry and dust grain distributions so as to calculate emergent SEDs in a panchromatic manner.  The turbulent envelopes of massive protostars and disordered morphologies of merging galaxies cannot be modeled in detail by treating them as purely axisymmetric systems, although this is a useful first approximation (Efstathiou \\& Rowan-Robinson 1995; Silva et al. 1998).  The Monte Carlo method relies on sampling probability distributions, for which it is necessary to use large numbers of photons.  Although this has been previously computationally restrictive, advances in computing speed and memory, along with the development of new algorithms and codes, have allowed this method to mature into a sophisticated and rigorous means of solving the radiative transfer problem.   Self-consistent Monte Carlo codes have been used to model the dust emission from protostellar regions (Whitney et al. 2003a;200b; Indebetouw et al.2006; Robitaille et al. 2006), to interface with SPH density fields (Stamatellos \\& Whitworth 2005), to study protoplanetary disk evolution (Wood et al. 2002), to compute SEDs and images from arbitrary density and stellar distributions (Misselt et al. 2001), and to compute the SEDs of clumpy toroidal models of AGN (Dullemond \\& van Bemmel 2005).  The Monte Carlo method has also been used recently to study the effects of dust attenuation and scattering on the emergent UV-near-IR SED from galaxy merger simulations (Jonsson 2006), and more recently to self-consistently compute the multi-wavelength emergent SEDs, images, and photometric time history of merger simulations to explain trends (cold-warm IRAS classfication; IRAC color-color plots) in observed data (Chakrabarti et al. 2007a; 2007b), and to calculate the demographics of infrared bright galaxies locally and as a function of redshift (Chakrabarti, Huang et al., in preparation).  In this paper, we focus our attention on algorithms designed to enforce radiative equilibrium in Monte Carlo codes.  The radiative equilibrium condition is the self-consistency condition that ensures that the flux of radiation is divergence free.  This is the basic condition that needs to be satisified to calculate the global radiation field and equilibrium dust temperature within the envelope.  Solving for the dust temperature can be accomplished locally only when the envelope is optically thin to its own reprocessed radiation.  There are two general classes of algorithms that have been used to enforce radiative equilibrium in Monte Carlo codes, namely what we refer to as relaxation methods (Bjorkman \\& Wood 2001) and iterative methods (Lucy 1999).  We review the details of these methods in \\S 2.  The organization of the paper is as follows.  In \\S 2, we review two basic types of algorithms designed to enforce radiative equilibrium in Monte Carlo codes. In \\S 3 we contrast the efficiency of these algorithms for calculating the SEDs and temperatures of homogeneous dust envelopes in one and three-dimensional grids.   We discuss two separate conditions that users may enforce, namely, the number of photons required for the temperature to converge ($N_{\\rm temp}$) and for noise requirements on the SED ($N_{\\rm SED}$).  In \\S 3.3, we present radiative transfer solutions for three-dimensional inhomogeneous envelopes and also compare the emergent SEDs to those of equivalent homogeneous envelopes.  In \\S 4, we present the basic setup for RADISHE, convergence studies, and comparisons of our calculated SEDs to observed SEDs of $z \\sim 2$ SMGs.  We show in \\S 5 that Herschel color-color plots will be able to discern the energy sources of SMGs.  We conclude in \\S 6. ",
        "conclusions": "$\\bullet$  We find that estimating the mean free intensity from the sum of photon flight paths and iteratively imposing radiative equilibrium (Lucy 1999) is a versatile and efficient temperature calculation algorithm, and particularly so for modeling three-dimensional dust envelopes.  We find that it is on average a factor of 10 times faster than relaxation algorithms for three-dimensional envelopes.  We have done extensive convergence studies to assess the viability of radiative equilibrium temperature algorithms, comparing to standard homogeneous test cases in one and three-dimensional grids.  $\\bullet$ We discuss how many photons are needed to obtain a converged temperature profile and a high fidelity SED, which we have termed $N_{\\rm temp}$ and $N_{\\rm SED}$ respectively.  We find that the total number of photons required to get a converged temperature profile througout the entire dust envelope that is accurate to 2 \\% can be accomplished with $N_{\\rm temp,per~iter} \\sim N_{\\rm grid}$, where $N_{\\rm temp,per~iter}$ is the number of photons per iteration and $N_{\\rm grid}$ the number of grid cells.  Typically a total of four iterations are sufficient.  $N_{\\rm SED} \\sim 5 N_{\\rm temp,per~iter}$ is sufficient to reach a noise-free long wave inclination averaged SED.  $N_{\\rm SED} \\sim 24 N_{\\rm temp,per~iter}$ is generally sufficient to produce a low level of noise over viewing angles out to submillimeter wavelengths.  $\\bullet$  We apply the code RADISHE to calculate the emergent SEDs of dusty galaxies along time sequences of hydrodynamical simulations.  Since RADISHE treats attenuation, scattering, and reemission self-consistently, it can generate a suite of accurate SEDs and corresponding dust temperatures calculated throughout the dust envelope.  We have compared our calculated SEDs during the Class I and Class II evolutionary phases to observed multi-wavelength data of SMGs to find good agreement.  $\\bullet$  Performing radiative transfer calculations along a time sequence of SPH simulations allows us to construct simulated color-color plots, which we can unfold as a function of some intrinsic galaxy parameter, such as time or AGN luminosity.  We have previously used this decomposition to explain trends in IRAC color-color plots (Chakrabarti et al. 2007b).  Here, we predict that $\\it{Herschel}$ color-color plots will be able to effectively discriminate between galactic energy sources in the $z \\sim 2-3$ SMG population.  $\\bullet$ We have demarcated the energetically active AGN region - the Class II phase, in $\\it{Herschel's}$ bands.  Specifically, we predict that (observed-frame) $F_{110~\\micron}/F_{170~\\micron}$-$F_{170~\\micron}/F_{360~\\micron}$ color-color plots will be able to most effectively select energetically active AGN in the $z \\sim 2$ SMG population.  We emphasize the viewing angle dependence of these colors, and demarcate the Class II phase in a conservative manner.  Although the viewing angle dependence weakens the correlation in either of these colors with $L_{\\rm BH}/L_{\\rm total}$, the coincidence of $\\it{both~colors}$ having high values is firmly correlated with $L_{\\rm BH}/L_{\\rm total} \\ga 0.5$.  Since SMGs undergo a brief excursion into the Class II region, wide-area surveys will be needed to detect a statistically significant sample of $z \\sim 2$ SMGs with energetically active AGN.  \\bigskip \\bigskip"
    },
    "0711/0711.4682_arXiv.txt": {
        "abstract": "\\baselineskip=0.6 cm  Quintessence field is a widely-studied candidate of dark energy. There is ``tracker solution'' in quintessence models, in which evolution of the field $\\phi$ at present times is not sensitive to its initial conditions. When the energy density of dark energy is neglectable~($\\Omega_\\phi\\ll1$), evolution of the tracker solution can be well analysed from ``tracker equation''. In this paper, we try to study evolution of the quintessence field from ``full tracker equation'', which is valid for all spans of $\\Omega_\\phi$. We get stable fixed points of $w_\\phi$ and $\\Omega_\\phi$ (noted as $\\widehat w_\\phi$ and $\\widehat\\Omega_\\phi$) from the ``full tracker equation'', i.e., $w_\\phi$ and $\\Omega_\\phi$ will always approach $\\widehat w_\\phi$ and $\\widehat\\Omega_\\phi$ respectively. Since $\\widehat w_\\phi$ and $\\widehat\\Omega_\\phi$ are analytic functions of $\\phi$, analytic relation of $\\widehat w_\\phi\\sim\\widehat\\Omega_\\phi$ can be obtained, which is a good approximation for the $w_\\phi\\sim\\Omega_\\phi$ relation and can be obtained for the most type of quintessence potentials. By using this approximation, we find that inequalities $\\widehat w_\\phi<w_\\phi$ and $\\widehat\\Omega_\\phi<\\Omega_\\phi$ are statisfied if the $w_\\phi$ (or $\\widehat w_\\phi$) is decreasing with time. In this way, the potential $U(\\phi)$ can be constrained directly from observations, by no need of solving the equations of motion numerically. ",
        "introduction": "Present astronomical observations require the existence of dark energy, a significant component of the universe with a negative pressure \\cite{Riess:1998cb,Riess:2004nr,Riess:2006fw,Spergel:2003cb,Komatsu:2008hk}. Though it has been more than ten years since its discovery, one is yet to tell what the dark energy is. We are still analyzing properties of dark energy from observational data and seeking suitable candidates. Most properties of dark energy depend on two parameters: the equation of state $w_{de}$ and the fractional energy density $\\Omega_{de}$. Once the $w_{de}\\sim\\Omega_{de}$ relation is obtained, we know almost all we need. At present, it is still not possible to constrain the evolution of dark energy from observations \\cite{Daly:2004gf,Mignone:2007tj,Sullivan:2007pd}. There are only definite constraints of present values of $w_{de}$ and $\\Omega_{de}$ from observations: $w_{de}^{(0)}$ is rather close to $-1$ and $\\Omega_{de}^{(0)}$ is dominating (about 70\\%) \\cite{Lazkoz:2007zk,Copeland:2006wr,Mantz:2007qh}. More constraints on $w_{de}$ and $\\Omega_{de}$ will be forthcoming from future observations, to get the evolution of the $w_{de}\\sim\\Omega_{de}$ relation from the observations, more theoretical efforts should be made.  At present, the most economical candidate of dark energy is still the cosmological constant $\\Lambda$, whose equation of state $w_\\Lambda=-1$. There is only a free parameter $\\Omega_\\Lambda$ in the flat $\\Lambda$CDM model. But it suffers from several problems, such as the coincidence problem and the fine tuning problem. Another well studied candidate is the quintessence $\\phi$, a slowly rolling scalar field, analogous to the inflaton. Its equation of state is $w_\\phi=(\\dot{\\phi}^2/2-U)/(\\dot{\\phi}^2/2+U)$ so one has $-1\\leq w_\\phi\\leq1$. In quintessence models, the coincidence problem and the fine tuning problem can be alleviated \\cite{Copeland:2006wr}. For example, there are tracker solutions for certain type of quintessence models, in which the evolution of $\\phi$ today is not sensitive to its initial conditions at early times \\cite{Steinhardt:1999nw}. The coincidence problem thus becomes less severe. But it is difficult to find quintessence models with analytic solutions of equation-of-motion, due to the existence of background matters (dark matter, baryon and radiations). To study evolutions of quintessence models and to be compared with observations, one usually has to solve the equations numerically. There are efforts to find analytic approximations for solutions of equations of motions, such as \\cite{Watson:2003kk} which gives a first order approximation solution for inverse power law potentials.  In this paper, we will try to approximate the $w_\\phi\\sim\\Omega_\\phi$ relation at the recent $\\Omega_\\phi$ dominating period in a semi-analytic way. To make sure that the evolution of $\\phi$ at present only depends on $U(\\phi)$, we assume there was tracking solution at early times. In \\cite{Steinhardt:1999nw}, conditions for the existence of tracker solution was given by the ``tracker equation'', which is a differential equation for $w_\\phi$. But this ``tracker equation'' are only valid as $\\Omega_\\phi\\ll1$. For our purpose, we need a full tracker equation that is valid for all $\\Omega_\\phi$ without conditions attached. Such an equation has been obtained \\cite{Scherrer:2005je,Chiba:2005tj,Lee:2006gx} and will be used here to study evolutions of quintessence models.  The paper is organized as follows. In section II, we introduce two new functions $\\widehat w_\\phi$ and $\\widehat \\Omega_\\phi$ which are fixed points of the full tracker equation. Assuming that $\\Gamma\\equiv U''U/U'^2$ and $\\epsilon\\equiv (U'/U)^2/2$ are nearly constant, we find that the fixed points are stable for $w_\\phi$ and $\\Omega_\\phi$ if $\\Gamma\\geq1$. If $\\Gamma$ and $\\epsilon$ do not evolve extremely fast, the relation of $w_\\phi\\sim\\Omega_\\phi$ will always approach to that of $\\widehat w_\\phi\\sim\\widehat \\Omega_\\phi$. In section III we show comparisons between \\{$\\widehat w_\\phi(\\phi)$, $\\widehat \\Omega_\\phi(\\phi)$\\} and \\{$w_\\phi(\\phi)$, $\\Omega_\\phi(\\phi)$\\} numerically for several typical quintessence models. The relation of $\\widehat w_\\phi\\sim\\widehat \\Omega_\\phi$ is shown to be a good approximation for the $w_\\phi\\sim\\Omega_\\phi$ relation. In section IV we show how to constrain $U(\\phi)$ directly from observational conditions on $w_\\phi$ and $\\Omega_\\phi$ through $\\widehat w_\\phi$ and $\\widehat \\Omega_\\phi$. Observational conditions are converted to simple inequalities for $U(\\phi)$. We conclude in section V with discussions. ",
        "conclusions": "We have gotten stable fixed points $\\widehat w_\\phi$ and $\\widehat\\Omega_\\phi$ from the full tracker equation, and shown that they are good approximations for $w_\\phi$ and $\\Omega_\\phi$ even in the $\\Omega_\\phi$ dominating period. $\\widehat w_\\phi$ and $\\widehat\\Omega_\\phi$ are analytic functions of $U(\\phi)$. The relation of $\\widehat w_\\phi\\sim\\widehat\\Omega_\\phi$ thus is gotten from the parametric functions $\\widehat w_\\phi(\\phi)$ and $\\widehat\\Omega_\\phi(\\phi)$, which is also a good approximation to the relation of $w_\\phi\\sim\\Omega_\\phi$.  Formally, functions of $\\widehat{w}_\\phi$ and $\\widehat{\\Omega}_\\phi$ with respect to expansion factor $a$ can also be obtained. Substituting Eq.(\\ref{e6}),(\\ref{e7}) into the equation \\be \\frac{d\\Omega_\\phi}{d\\ln a}=-3w_\\phi\\Omega_\\phi(1-\\Omega_\\phi) \\ee one gets the function of the field $\\phi$ with respect to $a$ upon integration. For example, for $U=U_0/\\phi^2~~(\\phi>0)$ one has: \\be \\phi(a)=\\frac{\\sqrt{14}}{3\\sqrt{5}}(120a^3+49a^6)^{1/4} \\ee where we have set present $\\widehat \\Omega_\\phi^{(0)}=0.7$ and $a_0=1$. Substituting $\\phi(a)$ into Eq.(\\ref{e6}) and Eq.(\\ref{e7}) one gets: \\ba \\widehat\\Omega_\\phi&=&\\frac{7}{60}(\\sqrt{120a^3+49a^6}-7a^3)\\nonumber\\\\ \\widehat w_\\phi&=&-\\frac{1}{2}-\\frac{7}{2\\sqrt{120a^{-3}+49}} \\ea For most potentials, it is not easy to get explicit functions of $\\phi(a)$, $\\widehat w_\\phi(a)$ and $\\widehat\\Omega_\\phi(a)$.  The critical points $\\widehat w_\\phi$ and $\\widehat\\Omega_\\phi$ can also be used to constrain the potential of quintessence directly from observational conditions on ($w_{de},\\Omega_{de}$). We have adopted two conditions on present ($w_{de}^{(0)},\\Omega_{de}^{(0)}$) for illustration. Further astronomical observations will yield more properties of dark energy. It may give conditions on ($w_{de},\\Omega_{de}$) at other redshifts, or even the exact shape of the $w_{de}\\sim\\Omega_{de}$ relation. In that case, our method can be still usable to constrain the potential and study the properties of the quintessence models that are fit with observations directly.    \\begin{figure}[t] \\includegraphics[width=8cm,height=5cm]{cross.eps} \\caption{The curve of the quintessence model with $U(\\phi)=U_0(e^{-\\phi/2}+e^{-20\\phi})$ in the $w'-w$ phase space. This curve crosses the the boundary for thawing and freezing fields \\cite{Caldwell:2005tm} and the lower bound $w'=-(1-w)(1+w)$ in \\cite{Scherrer:2005je,Chiba:2005tj}.} \\label{f2} \\end{figure}  In this paper, we have only studied the case that $w_\\phi$~(and $\\widehat w_\\phi)$ keeps on decreasing monotonously, from which the inequality (\\ref{e12}) is obtained. In fact, there are quintessence models in which $w_\\phi$ is increasing at present. One example is the case with $U(\\phi)=U_0(e^{-\\phi/2}+e^{-20\\phi})$ \\cite{Barreiro:1999zs}. In this type of models, $w_\\phi$ will decrease to a minimum close to $-1$ and then begin to increase. So the boundary for thawing and freezing fields in \\cite{Caldwell:2005tm} will be crossed, as shown in Fig. \\ref{f2}. It can be shown that when $$\\frac{d\\ln{(\\Gamma-1)}}{d\\ln{U}}>\\frac{3}{2\\epsilon}$$ $\\widehat w_\\phi$ will be increasing, so will be $w_\\phi$. It requires a rapid decrease of $\\Gamma$. As $\\Gamma$ at early times must be close to 1 to get enough tracking, usually there is a rapid increase of $\\Gamma$ at recent times. In this case the lower bound $w'>-(1-w)(1+w)$ for quintessence models \\cite{Scherrer:2005je,Chiba:2005tj} may be crossed too. It is because $w=(w_b-2\\Gamma+2)/(2\\Gamma-1)$ will no longer be larger than $w_\\phi$ if the increase of $\\Gamma$ is too fast. It can be seen in Fig. \\ref{f2} that the line of $w'\\sim w$ with the double exponential potential is very close to the strict lower bound $w'>3w(1+w)$ given in \\cite{Caldwell:2005tm}. For this type of potential, as $w_\\phi$ and $\\widehat w_\\phi$ are increasing, there is an inequality similar to (\\ref{e12}): \\be \\widehat w_\\phi>w_\\phi,~~~~\\widehat \\Omega_\\phi>\\Omega_\\phi \\ee This inequality can be used to constrain $U(\\phi)$ from conditions on ($w_{de},\\Omega_{de}$). The methods used in this paper can also be extended to Phantom and K-essence models."
    },
    "0711/0711.3407_arXiv.txt": {
        "abstract": "{} {We describe various expansion schemes that can be used to study gravitational clustering. Obtained from the equations of motion or their path-integral formulation, they provide several perturbative expansions that are organized in a different fashion or involve different partial resummations. We focus on the two-point and three-point correlation functions, but these methods also apply to all higher-order correlation and response functions.} {We present the general formalism, which holds for the gravitational dynamics and for any similar models, such as the Zeldovich dynamics, that obey similar hydrodynamical equations of motion with a quadratic nonlinearity. We give our explicit analytical results up to one-loop order for the simpler Zeldovich dynamics. For the gravitational dynamics, we compare our one-loop numerical results with numerical simulations.} {We check that the standard perturbation theory is recovered from the path integral by expanding over Feynman's diagrams. However, the latter expansion is organized in a different fashion, and it contains some UV divergences that cancel out as we sum all diagrams of a given order. Resummation schemes modify the scaling of tree and one-loop diagrams, which exhibit the same scaling over the linear power spectrum (contrary to the standard expansion). However, they do not significantly improve over standard perturbation theory for the bispectrum, unless one uses accurate two-point functions (e.g. a fit to the nonlinear power spectrum from simulations). Extending the range of validity to smaller scales, to reach the range described by phenomenological models, seems to require at least two-loop diagrams.} {} ",
        "introduction": "\\label{Introduction}   According to standard cosmological scenarios, the large-scale structures of the present Universe arise from the amplification by gravitational instability of small primordial fluctuations (Peebles 1980). Then, from observations of the recent Universe at redshifts $z<5$, through galaxy surveys (Tegmark et al. 2006; Cole et al. 2005), weak-lensing studies (Massey et al. 2007; Munshi et al. 2007), or studies of baryonic oscillations (Eisenstein et al. 1998, 2005), we can derive constraints on the cosmological parameters and on the properties of those primordial fluctuations. This requires accurate theoretical predictions for the evolution of gravitational clustering to compare cosmological models with observations. On large scales, where density fluctuations are small, linear theory is sufficient. On small scales, deep into the nonlinear regime associated with virialized structures such as galaxy clusters, one can use numerical simulations or phenomenological models, such as the halo model (Cooray \\& Sheth 2002), which allow for an accuracy near $10\\%$. However, many observational probes (such as weak lensing or baryonic oscillations) are mostly sensitive to intermediate scales (wavenumbers $k \\sim 0.1 h$ Mpc$^{-1}$ at $z \\la 1$) that might be studied through analytical methods, such as perturbative expansions.  Thus, there has recently been a renewed interest in building analytical methods that could improve over the standard perturbation theory and provide reliable predictions in this weakly nonlinear regime. Crocce \\& Scoccimarro (2006a,b) present a partial resummation of the diagrammatic series associated with the response function, which recovers the expected decay into the nonlinear regime, in agreement with numerical simulations. This allows one to obtain a more accurate prediction for the matter power spectrum than with the usual perturbation theory (Crocce \\& Scoccimarro 2007). On the other hand, Valageas (2004, 2007a) described the path-integral formalism associated with the cosmological gravitational dynamics. This allows one to apply the usual tools of field theory, such as large-$N$ and semi-classical expansions. This provides several methods to perform partial resummations of the usual perturbation theory. Next, Matarrese \\& Pietroni (2007a,b) proposed to use the dependence of the system on a high-$k$ cutoff to derive some resummation schemes through the path-integral formalism. An alternative approach developed by McDonald (2007), based on a renormalization group technique, also permits going beyond the standard perturbative expansion.  These works have focused on two-point functions, such as the matter power spectrum or two-point correlation, since the latter are the main quantities obtained from observations. However, the additional information included in higher-order statistics, such as the three-point correlation, can prove very useful in breaking degeneracies (Sefusatti et al. 2006; Munshi et al. 2007). Therefore, we describe in this paper how such methods can be used to compute higher-order functions. We focus on the bispectrum, which is the quantity of most practical interest, since higher-order statistics obtained from observational surveys are increasingly noisy. We consider both the exact gravitational dynamics and the simpler Zeldovich dynamics. Indeed, the equations of motion associated with both systems are similar and these expansion schemes apply in the same manner to both (Valageas 2007b). Then, the simpler Zeldovich dynamics can be used as a convenient benchmark to track possible errors in the derivation of the equations associated with various schemes.  This article is organized as follows. First, in Sect.~\\ref{Equations-of-motion} we recall the equations of motion associated with cosmological gravitational clustering. We focus on weakly nonlinear scales hence we restrict ourselves to the hydrodynamical description. We write the path-integral formulation of this system and we briefly present two large-$N$ methods which can be derived within this formalism. We also recall the main properties of the simpler Zeldovich dynamics. Next, we present in a general manner various perturbative expansions that can be applied to both cosmological dynamics (or any equation of motion of the same form, i.e. with a quadratic nonlinearity). Thus, we first recall the standard perturbative expansion over powers of the linear density field in Sect.~\\ref{Standard-perturbation-theory}, we describe the equivalent expansion that can be derived from the path-integral formalism by expanding over the interaction term of the action in Sect.~\\ref{Expansion-over-cubic}, we present the resummations associated with the large-$N$ methods in Sect.~\\ref{LargeN}, and we briefly discuss in Sect.~\\ref{large-k-limit} the resummation associated with the random advection of small-scale structures by the larger-scale coherent flow (involving a high-$k$ limit). Then, we apply these methods to the Zeldovich dynamics in Sect.~\\ref{Zeldovich}, up to one-loop order. Indeed, the expressions obtained in this case are simpler than for the gravitational dynamics. This allows us to show explicitly the connection between the different methods and to point out the possible divergences which may appear in intermediate steps as diagrams are gathered in different manners. Finally, we apply these methods to the actual gravitational dynamics in Sect.~\\ref{Gravitational-dynamics}, again up to one-loop order, and we compare their predictions with numerical simulations. We conclude in Sect.~\\ref{Conclusion}.  The reader who is not directly interested in the theoretical aspects of gravitational clustering (e.g. the discussion of the various formalisms and perturbative methods that may be built, and the problems, such as UV divergences, that one may encounter) might first read Sect.~\\ref{nonlinear-power-spectrum} and Figs.~\\ref{figBksdNL_z0}-\\ref{figBtsdNL}. There, we compare the standard perturbation theory with the most efficient resummation scheme considered in this paper (at one-loop order), as well as a recent phenomenological model and numerical simulations. ",
        "conclusions": "\\label{Conclusion}   In this article, we have described several methods to obtain systematic expansions for two-point and three-point correlation functions for the cosmological large-scale density field. These methods also extend to higher-order correlations and response functions (e.g., the response of the two-point and three-point correlations to external perturbations or to the initial conditions). They also apply to related systems, defined by similar hydrodynamical equations, such as the Zeldovich dynamics. We showed how the simplest perturbative expansion derived from the the path-integral formulation of the system (i.e. expanding over the cubic part of the action) corresponds to the equivalent standard expansion directly obtained from the hydrodynamical equations of motion. Next, we described two resummations, associated with ``large-$N$'' limits, that can be conveniently derived from the path integral (Valageas 2007a,b), and a high-$k$ resummation that is most easily derived from the equation of motion (Crocce \\& Scoccimarro 2006a,b; Valageas 2007b). Thus, we showed how to derive expansions for high-order correlations (or responses) and we compared the resummations involved in these various schemes.  We first applied these methods up to one-loop order to the simpler case of the Zeldovich dynamics, where explicit expressions are simple enough to be given here. We explicitly checked that the perturbative expansion of the path integral recovers the standard perturbative expansion over powers of the linear density field. However, since these two expansions are organized in different manners, we found that two diagrams, obtained within the former method for the three-point correlation, exhibit an additional UV divergence (for $n(\\infty)>-3$). This divergence cancels out as we sum over all diagrams but it shows explicitly that some care must be taken in using these expansions. For numerical purposes, it is best to first regularize each diagram by subtracting the divergent parts that cancel out. Next, we applied the large-$N$ steepest-descent method to this Zeldovich dynamics. We showed how the time-dependence can be factorized, through functions $T(\\omega_j D)$, that measure the decay of different-time correlation and response functions in the nonlinear regime (within simple approximations). We found that, contrary to the standard expansions, tree diagrams and one-loop diagrams obey the same scaling laws over the linear amplitude $D \\sigma_8$. This is due to the cutoff $T(\\omega_j D)$ associated with the partial resummations. We also explained how it is again possible to regularize the expansion, as for the previous expansion.  Second, we applied these expansion schemes to the gravitational dynamics. All properties obtained for the Zeldovich dynamics remain valid and we have compared our results with numerical simulations. We checked that on quasi-linear scales all expansion schemes agree with one another and with the numerical simulations, up to $0.2 h$ Mpc$^{-1}$ (where $\\Delta^2_L(k)\\simeq 1$ for a $\\Lambda$CDM power spectrum with $\\sigma_8=0.9$). At smaller scales they deviate from one another and from the numerical simulations. We recovered the known failure of the standard perturbative expansion to reproduce the observed flattening in the nonlinear regime of the reduced bispectrum $Q(k_1,k_2,k_3)$ as a function of the angle $\\theta$ between wavenumbers $\\bk_1$ and $\\bk_2$. We found that the steepest-descent resummation does not fare much better. Even substituting the Gaussian decay obtained in a high-$k$ limit for the different-time two-point functions does not help (it actually increases somewhat further the variation with $\\theta$). Nevertheless, the agreement with numerical simulations significantly improves if we use the ``exact'' nonlinear power spectrum (obtained from a fit to simulations) for the two-point correlation within this formalism (together with the Gaussian decay). This does not extend the range of validity for equilateral configurations towards higher $k$, but it manages to partly recover the observed flattening of the reduced bispectrum as a function of $\\theta$.  We have also compared the predictions obtained from these expansion schemes with the simple phenomenological model of Pan et al. (2007). It appears that the validity of the latter extends to smaller scales, up to $0.8 h$ Mpc$^{-1}$ where $\\Delta^2_L(k) \\simeq 17$ in the case studied here. Since this model is also much simpler to use, it seems that the scale transformations devised by Hamilton et al. (1991) may indeed provide remarkably efficient tools to investigate various statistics associated with gravitational clustering. At quasi-linear scales $k \\sim 0.1 h$ Mpc$^{-1}$, the systematic expansion schemes studied in this article may be somewhat more accurate (since they arise from rigorous systematic expansions) whereas the phenomenological model of Pan et al. (2007) may lead to a slight overestimate of the bispectrum (which would not be surprising since its expansion is not expected to agree with perturbation theory).  The results obtained in this article suggest that, in order to make further progress along these lines, one needs to include higher-order diagrams (two-loop diagrams or higher) to obtain reliable predictions for the bispectrum that accurately reproduce the weakening of the dependence on the angle $\\theta$ in the nonlinear regime. Moreover, using accurate two-point functions provides a significant improvement. This requires both the use of a Gaussian decay for different-time quantities (or a similar cutoff that agrees with numerical simulations) and of a reasonable equal-time correlation (e.g. using the nonlinear power spectrum measured in simulations). On the other hand, the path-integral formulation might serve as a basis for other approximation methods than the expansion schemes discussed here.  Alternative methods, based on phenomenological models, such as the halo model (Cooray \\& Sheth 2002) or the scaling transformation of Hamilton et al. (1991), may provide other means to describe the statistical properties of gravitational clustering. Although simpler and more efficient in the nonlinear regime, they do not have the rigorous and systematic character of expansion schemes. Moreover, they often involve several free parameters that must be fitted to simulations and they are not derived from first principles. It would be interesting to develop a connection between such phenomenological approaches and more systematic methods as those developed in this article. However, it is not clear whether such a connection can be achieved at a detailed level. This issue is left for future works."
    },
    "0711/0711.4820_arXiv.txt": {
        "abstract": "We propose an extension of the standard model with a B-L global symmetry that is broken softly at the TeV scale. The neutrinos acquire masses through a type-II seesaw while the lepton (L) asymmetry arises in the {\\it singlet sector} but without B-L violation. The model has the virtue that the scale of L-number violation ($\\Lambda$) giving rise to neutrino masses is independent of the scale of leptogenesis ($\\Lambda'$). As a result the model can explain {\\it neutrino masses, singlet scalar dark matter and leptogenesis at the TeV scale}. The stability of the dark matter is ensured by a surviving $Z_2$ symmetry, which could be lifted at the Planck scale and thereby allowing Planck scale-suppressed decay of singlet scalar dark matter particles of mass $\\approx 3$ MeV to $e^+ e^-$ pairs in the Galactic halo. The model also predicts a few hundred GeV doubly charged scalar and a long lived charged fermion, whose decay can be studied at Large Hadron Collider (LHC) and International Linear Collider (ILC). ",
        "introduction": "Within the standard model (SM), neutrinos are massless. On the other hand, the current low energy neutrino oscillation data~\\cite{solar_data, atmos_data,kamland_data} indicate that at least two of the physical left-handed (LH) neutrinos have tiny masses and therefore mix among themselves. However, as yet we do not know if the neutrinos are Dirac or Majorana. If the neutrinos are assumed to be Majorana then the sub-eV neutrino masses can be generated through the dimension five operator~\\cite{weinberg.79} \\be \\mathcal{O}_\\nu=\\frac{\\phi \\phi \\ell \\ell}{\\Lambda}\\,, \\label{dim-5-operator} \\ee where $\\Lambda$ is the scale of lepton (L) number violation ($\\Delta L=2$). The dimension five operator (\\ref{dim-5-operator}) can originate through the celebrated see-saw mechanism.  In type-I seesaw models, three right-handed (RH) neutrinos ($N$'s) are added without extending the gauge group of the SM. The canonical seesaw (or type-I seesaw)~\\cite{canonical_seesaw} then gives the light neutrino mass matrix: \\be m_\\nu=m_\\nu^I=-m_D M_N^{-1} m_D^T\\,, \\label{typeI-seesaw} \\ee where $m_D$ is the Dirac mass matrix of the neutrinos connecting the LH neutrinos ($\\nu_L$) with the RH neutrinos and $M_N$ is the Majorana mass matrix of the RH heavy neutrinos, which also sets the scale of L-number violation ($\\Lambda$). The Dirac mass terms determine the L-numbers of the RH neutrinos to be $+1$ and hence the Majorana mass of the RH neutrinos violates L-number by two units. The decays of the RH neutrinos would then violate L-number and their CP violating out-of-equilibrium decay to SM fields can be a natural source of L-asymmetry~\\cite{fukugita.86} in the early Universe. The CP-violation, which comes from the Yukawa couplings that determine the Dirac mass mass of neutrinos via the one-loop radiative vertex correction, requires at least two RH neutrinos. The masses of the RH neutrinos producing the final L-asymmetry then satisfy~\\cite{type_I_group} ~\\footnote{This could be next-to-lightest RH neutrino if flavor leptogenesis is considered~\\cite{N2_leptogenesis}.}. \\be M_N\\geq O(10^{9}) GeV\\,. \\ee If the corresponding theory of matter is supersymmetric (SUSY) then this bound, being dangerously close to the maximum reheat temperature, poses a problem. A modest solution was proposed in ref.~\\cite{ma&sahu.06} by introducing an extra singlet heavy fermion. However, the model only achieves a reduction of above bound~\\cite{type_I_group} by an order of magnitude.  In the type-II seesaw models, on the other hand, $SU(2)_L$ triplet Higgses ($\\Delta$'s) are added to the SM gauge group. Explicit breaking of the L-number by trilinear couplings of the triplet Higgs scalar then induces a tiny vacuum expectation value (VEV) of the heavy triplet Higgs scalars~\\cite{triplet_seesaw}, generating a light neutrino mass matrix: \\be m_\\nu=m_\\nu^{II}=f\\mu \\frac{v^2}{M_{\\Delta}^2}\\,, \\label{typeII-seesaw} \\ee where $M_{\\Delta}$ is the mass of the triplet Higgs scalar $\\Delta$, $\\mu$ is the coupling constant with mass dimension 1 for the trilinear term with the triplet Higgs and two standard model Higgs doublets, and $f$ is the Yukawa coupling of the triplet Higgs to the light leptons. $M_{\\Delta}$ and $\\mu$ are of the same order of magnitude and set the scale ($\\Lambda$) of L-number violation. $v$ is the VEV of the SM Higgs doublet. In these models, the L-asymmetry is generated through the L-number violating decays of the $\\Delta$ to SM leptons and Higgs~\\cite{ma&sarkar_prl}. The CP-violation, originating from the one-loop self-energy correction, requires at least two triplets. The scale of L-number violation is determined by $M_{\\Delta}$ and $\\mu$ and is required to be very high. In most cases the scale of L-violation in type-II seesaw models is larger than the type-I seesaw models~\\cite{type_II_group}.  From the above discussion we see that the L-number violating scales $M_N\\sim \\Lambda$ in type-I see-saw models or $M_\\Delta\\sim \\Lambda$ in type-II seesaw models comes out to be large because the same L-number violation gives rise to both neutrino masses and mixings and to the L-asymmetry. While sub-eV neutrino masses require large values of $\\Lambda$, the large values of $M_N$ and $M_\\Delta$ can easily satisfy the out-of-equilibrium decay condition, a necessary condition for the generation an L-asymmetry, without any fine-tuning of the Yukawa couplings. From this perspective, these models are attractive. However, they cannot be verified in future colliders since the scale of L-number violation is very high. Alternatives to these models are provided by mechanisms which work at the TeV scale, either in SUSY extensions of the SM~\\cite{susy_tev} or by introducing an additional source of CP violation into the model~\\cite{extra_tev}.  On the other hand, one could ask if the origin of neutrino masses and leptogenesis is different? It is well-known that there is no one-to-one correspondence between the parameters in the neutrino mass matrix and those involved in leptogenesis~\\cite{ leptogenesis-mnu-3N}. In particular, in type-I see-saw model with 3 RH neutrinos, while 15 parameters enter into leptogenesis, there are only 9 parameters: 3 masses, 3 mixing angles and 3 phases (one L-number conserving phase called the Dirac phase and two L-number violating phases called Majorana phases) in the low energy neutrino mass matrix. Obviously there is no connection. Conservatively, if one considers type-I seesaw model with 2 RH neutrinos for leptogenesis as well as for low energy sub-eV neutrino masses, then the number of parameters in both cases are 9. However, one can show that there is no one-to-one correspondence between the Majorana phases responsible for L-number violation in leptogenesis and Majorana phases in neutrino mass matrix~\\cite{leptogenesis-mnu-2N}. The connection between leptogenesis and low energy neutrino mass matrix is even worse in case of type-II see-saw models.   Motivated by the fact that there is no one-to-one correspondence between leptogenesis and the effective low energy neutrino mass matrix, we propose a new mechanism of leptogenesis which occurs completely in the {\\it singlet sector} at the TeV scale~\\cite{singlet_lep}. As we will show, the origin of neutrino masses is different from the origin of L-asymmetry. The L-asymmetry then arises without any B-L violation. We will show that the B-L violation required for neutrino masses does not conflict with the leptogenesis. This model is then extended to incorporate a singlet scalar dark matter particle. While the nature of dark matter is still a mystery, we will show that the singlet dark scalar can either be collisionless cold dark matter (CCDM) or self-interacting dark matter (SIDM). The stability of the dark matter is ensured by a surviving $Z_2$ symmetry which could be lifted at the Planck scale and thereby allow Planck scale-suppressed decay of the singlet dark scalars to $e^+e^-$ pairs in the Galactic halo. The most important feature of the model is that it predicts a few hundred GeV doubly charged scalar and a long lived singly charged fermion whose decay can be studied at the LHC/ILC.  The paper is organized as follows. In section II we introduce a model which simultaneously explains neutrino masses, singlet scalar dark matter and singlet leptogenesis. In section III, the type-II seesaw model of neutrino masses and the viability of testing it at colliders are discussed. In section IV we give a brief description of singlet leptogenesis, which arises from a conserved B-L symmetry. Section V is devoted to singlet scalar dark matter and its Planck scale suppressed decay to $e^+e^-$ pair in the Galactic halo. In section VI we give a brief description of collider signatures and section VII concludes. ",
        "conclusions": "We have introduced a new leptogenesis mechanism in a $U(1)_{B-L}$ extension of the SM which allowed us to explain simultaneously neutrino masses, dark matter and leptogenesis. The important message is that the L-asymmetry arises without any B-L violation. Since the L-number violation required for leptogenesis and neutrino masses are different, the leptogenesis scale can be lowered to as low as a few TeV.  Neutrino masses arise through a modified type-II seesaw which predicts a few hundred GeV triplet scalar. Note that in conventional type-II seesaw models the mass scale of the triplets is required to be at least $O(10^{10})$ GeV to produce sub-eV neutrino masses. Since one of the triplets in our model, namely $\\xi$, has a mass of a few hundred GeV, the proposed model can be tested in the near future at colliders through same sign dilepton decay of $\\xi$.  The model also predicts a singly charged fermion $\\eta^-$ of mass ranging from 200 GeV to 1 TeV. $\\eta^-$ can be produced in high energy neutrino collision with matter. Since it is long-lived, it can travel large distances and be observed in neutrino telescopes.  We proposed a singlet scalar, $\\chi$, with mass $\\approx 3$ MeV as a candidate for dark matter, whose stability is ensured by a $Z_2$ symmetry. This $Z_2$ discrete symmetry can be broken at the Planck scale. We then showed that a possible origin of 511 KeV Galactic line detected by INTEGRAL could be the Planck scale-suppressed decay of $\\chi$ to $e^- e^+$ pairs in the Galactic halo."
    },
    "0711/0711.0033_arXiv.txt": {
        "abstract": "We produce models of early WN, WC, and WO stars as a function of metallicity $Z$ using an analytic CAK-type approach. At log~$Z / Z_\\odot \\geq -2$ both WN and WC stars have the approximate dependence $\\dot{M} \\propto Z^{0.5}$.  For a WN wind the mass-loss rate drops rapidly below log~$Z / Z_\\odot = -2$, and no wind solution can be found for log~$Z / Z_\\odot < -3$. For WC and WO winds the mass-loss rate plummets in the range $-3 \\leq \\mathrm{log}~Z / Z_\\odot \\leq -2$ and tends to flatten due to the self-enrichment of C and O to around $10^{-8} \\: M_\\odot$ yr$^{-1}$ for log~$Z / Z_\\odot \\leq -4$. No significant difference in $\\dot{M}$ was found for WC versus WO stars at low metallicity. ",
        "introduction": "Near solar metallicity the mass-loss rates of Wolf-Rayet (WR) stars can be strongly dependent on $Z$, but the self-enrichment of CNO elements creates an effective $Z$ floor that can allow large $\\dot{M}$ even in a low-$Z$ environment. Using the modified CAK approach described in \\citet{onifergayley06}, we calculate $\\dot{M}$ as a function of $Z$ for early-type WR winds from WN, WC, and WO stars.  To ensure a complete line list at the high critical-point temperatures of these winds, we use data from the Kurucz list \\citep{kurucz79}, the Opacity Project \\citep{seaton05}, and the opacity group at Los Alamos \\citep{mazevetabdallah06}.  The results are compared to the $\\dot{M}$-$Z$ dependences for late-type stars by \\citet[hereafter VdK05]{vinkdekoter05}. ",
        "conclusions": "Figure \\ref{fig:wnwc} shows the $\\dot{M} \\propto Z^m$ relationship for WN (left) and WC/WO (right) stars.  For WN stars $m = 0.54$ at log~$Z / Z_\\odot > -2$ and steepens to $m = 1.1$ at lower metallicity.  No models could be calculated for log~$Z / Z_\\odot < -3$.  For WC stars $m = 0.57$ for log~$Z / Z_\\odot \\geq -2$. $\\dot{M}$ then falls rapidly, finally flattening in the region log~$Z / Z_\\odot \\leq -4$ to around $1.2\\times10^{-8} \\: M_\\odot$ yr$^{-1}$ for $L_{th} = 0$ or $6.5\\times10^{-9} \\: M_\\odot$ yr$^{-1}$ for finite $L_{th}$.  The WC and WO mass-loss rates are very similar.  \\begin{figure}[!ht] \\plottwo{aonifer01.eps}{aonifer02.eps} \\caption{Left: WN $\\dot{M}$ vs $Z$ for finite $L_{th}$ and $L_{th}$ = 0 Right: WC and WO $\\dot{M}$ with finite $L_{th}$ and $L_{th}$ = 0.  The VdK05 results (solid lines) are also shown. \\label{fig:wnwc}} \\end{figure}"
    },
    "0711/0711.1454.txt": {
        "abstract": "We present results of 120~ks observation of a compact group of galaxies HCG~62 ($z=0.0145$) with Suzaku XIS and HXD-PIN\\@. The XIS spectra for four annular regions were fitted with two temperature {\\it vapec} model with variable abundance, combined with the foreground Galactic component. The Galactic component was constrained to have a common surface brightness among the four annuli, and two temperature {\\it apec} model was preferred to single temperature model.  We confirmed the multi-temperature nature of the intra-group medium reported with Chandra and XMM-Newton, with a doughnut-like high temperature ring at radii 3.3--6.5$'$ in a hardness image. We found Mg, Si, S, and Fe abundances to be fairly robust. We examined the possible ``high-abundance arc'' at $\\sim 2'$ southwest from the center, however Suzaku data did not confirm it. We suspect that it is a misidentification of an excess hot component in this region as the Fe line. Careful background study showed no positive detection of the extended hard X-rays previously reported with ASCA, in 5--12~keV with XIS and 12--40~keV with HXD-PIN, although our upper limit did not exclude the ASCA result. There is an indication that the X-ray intensity in $r<3.3'$ region is $70\\pm 19$\\% higher than the nominal CXB level (5--12~keV), and Chandra and Suzaku data suggest that most of this excess could be due to concentration of hard X-ray sources with an average photon index of $\\Gamma=1.38\\pm 0.06$.  Cumulative mass of O, Fe and Mg in the group gas and the metal mass-to-light ratio were derived and compared with those in other groups. Possible role of AGN or galaxy mergers in this group is also discussed. ",
        "introduction": "Groups of galaxies play a key role in the formation of the universe. Most importantly, they act as building blocks in the framework of hierarchical formation of structures (e.g.\\ \\cite{Navarro1995}). In this sense, properties of groups should be critically compared with those in clusters of galaxies, to test the general hypothesis that groups of galaxies indeed represent the condition in clusters when they have not been evolved. For example, baryon to dark matter ratio and metal abundance in groups of galaxies can be a measure to be tested whether rich clusters are indeed simple superposition of many groups or some other mechanism may be involved.  The groups are characterized by a short dynamical time scale, and we expect galaxy encounters to take place frequently in such a high density environment. The X-ray search for an evidence of dynamical processes in the form of hard non-thermal emission would be an important subject of observational studies. The relatively low gas temperature compared with rich clusters greatly helps us find such emission even with standard CCD instruments. Also it is well known that group gas tends to contain significantly high entropy if one extrapolates the scaling relation holding for clusters \\citep{Ponman1999}. A simulation study indicated that simple galactic winds were unable to significantly raise the entropy in cluster size hot gas \\citep{Borgani2005}. This indicates that a certain heat input or preheating occurred over a widespread region and groups of galaxies offer us an opportunity to closely look into such an effect.  Another aspect concerning groups of galaxies is their role in the chemical enrichment of cluster hot gas. It has been shown that rich clusters maintain all the metals produced by the cluster galaxies \\citep{Fukazawa1998}, however early-type galaxies are thought to have released most of the metals formed through past supernova explosions \\citep{Makishima2001}. We should note that these galaxies almost always reside in groups and clusters of galaxies and that their X-ray haloes may not be clearly separable from the surrounding gas. Groups of galaxies lie in the region where the metal confinement can be marginally possible. This means that by looking into the distribution of various metals in groups, we can measure the efficiency of metal confinement and further estimate how each metal has been injected into the intergalactic space. The low temperature ($< 2$ keV) of the group gas also greatly helps us study emission lines from various metals ranging from oxygen to iron.  HCG~62 is one of the brightest and well-studied groups of galaxies among the Hickson compact groups of galaxies \\citep{Hickson1989} located at $z = 0.0145$ \\citep{Mulchaey2003}. \\citet{Zabludoff2000} extensively surveyed a $1.5\\times 1.5$ deg$^2$ region around the group and measured redshifts of 154 galaxies. They identified 63 of them as members of the group within a radius of $50'$ (900~kpc), and the measured velocity dispersion, $\\sigma_r = 390^{+37}_{-34}$ km~s$^{-1}$, is a typical value for galaxy groups. The central region within $r<1.1'$ is dominated by three galaxies, HCG~62a (NGC~4778), HCG~62b (NGC~4776), and HCG~62c (NGC~4761), and they are all classified as S0 galaxies. Kinematics of these galaxies suggests possible interaction among them \\citep{Spavone2006,Rampazzo1998}. Such a high galaxy density and the low velocity dispersion are considered to result in a galaxy merger, and \\citet{Ponman1993} suggested that they should culminate in a final merger within a few billion years, consisting of a large elliptical galaxy embedded in an extended X-ray halo.  HCG~62 is also known as the first compact galaxy group that was detected in X-ray to have an extended hot gas (intra-cluster or group medium; ICM), and the ICM properties have been extensively studied using ROSAT and ASCA (e.g., \\cite{Ponman1993,Pildis1995,Finoguenov1999,Mulchaey2003}). Furthermore, spatially extended hard X-ray excess over the ICM emission was detected with ASCA \\citep{Fukazawa2001,Nakazawa2007}, and two symmetrical cavities in the central region at a projected radius of $\\sim 25''$ (7.4~kpc) were discovered by Chandra \\citep{Vrtilek2002}. Detailed hot-gas properties including the cavities and the gravitational mass have been studied based on the combined data from Chandra and XMM-Newton \\citep{Morita2006}. Recently, \\citet{Gu2007} reported that a region at $2'$ offset from the central galaxy showed a twice higher metal abundance than the surrounding gas (the ``high-abundance arc'').  In this paper, we report the results from Suzaku observation of HCG~62\\@. Owing to the low and stable background as well as good sensitivity to emission lines below $\\sim 1$~keV of the X-ray Imaging Spectrometer (XIS; \\cite{Koyama2007}), accurate determination of O and Mg abundances to outer regions have become feasible. XIS is also effective to constrain the extended hard X-ray emission in combination with the Hard X-ray Detector (HXD; \\cite{Takahashi2007}). This paper is organized as follows: In sections~\\ref{sec:obs}~and~\\ref{sec:data}, we describe the Suzaku observation and the data reduction, and show images of HCG~62 obtained with Suzaku in section~\\ref{sec:image}. In sections~\\ref{sec:spec}--\\ref{sec:abun}, we describe spectral analysis and the derived temperature and abundance profiles, and examine the existence of the ``high-abundance arc'' in section~\\ref{sec:arc}. In section~\\ref{sec:hard-emission}, significance of the extended hard X-ray emission with Suzaku is carefully investigated. We discuss temperature and metallicity distributions, metal mass-to-light ratio, and galaxy mergers in section~\\ref{sec:discussion}, and a summary is given in section~\\ref{sec:summary}.  We use $H_{0} = 0.7\\; h_{100} = 70$ km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_{\\rm \\Lambda}=1-\\Omega_{\\rm M} = 0.73$ throughout this paper. At the redshift of $z = 0.0145$, 1$'$ corresponds to 17.8~kpc. The virial radius, $r_{180}\\equiv 1.95\\; h_{100}^{-1} \\sqrt{k\\langle T\\rangle /10~\\rm keV} = 1.08$~Mpc \\citep{Markevitch1998,Evrard1996} for the average temperature of $k\\langle T\\rangle = 1.5$~keV\\@. We adopt a Galactic hydrogen column density of $N_{\\rm H} = 3.03 \\times 10^{20}$~cm$^{-2}$ in the direction of HCG~62 \\citep{Dickey1990}. As for the definition of the solar abundance ratio, we followed \\citet{Anders1989}. Errors are 90\\% confidence range for a single interesting parameter. ",
        "conclusions": "\\label{sec:discussion}  \\subsection{High-Temperature Ring}  We found the doughnut-like high-temperature ring at 3.3--6.5$'$ surrounding the group center, as shown in the hardness image of figure~\\ref{fig:image}(b). Spectral fits, shown in figure~\\ref{fig:kT}(b), indicated an excess in the intensity ratio of hot to cool ICM components in this annulus. We note that the co-existence of the hot and cool components was already observed with Chandra and XMM-Newton \\citep{Morita2006} based on the deprojection analysis. However, as shown in figure~11(a) of \\citet{Morita2006}, the hot component should occupy more than $\\gtrsim 90$\\% of the volume if two components are under a pressure balance. In other words, the total thermal energy given by a product of pressure and volume, $PV$, is dominated by the hot component in this annulus.  We therefore assume that the 3.3--6.5$'$ (60--120~kpc) annulus is filled with the hot ICM with temperature $kT_{\\rm Hot}\\simeq 1.6$~keV and the electron density $n_{\\rm e}\\simeq 7\\times 10^{-4}$ cm$^{-3}$ for simplicity. Then the gas pressure and the total energy are $P = 1.92\\; n_{\\rm e} kT_{\\rm Hot} = 2.2$ eV~cm$^{-3}$ and $E = PV = 5.8\\times 10^{59}$ erg~s$^{-1}$, respectively. The bolometric luminosity is $L_{\\rm bol} = 1.1\\times 10^{42}$ erg~s$^{-2}$ using the {\\it vapec} model with the elemental abundances given in table~\\ref{tab:best-fit} (Gal 2T) for this annulus. This gives the radiation cooling time as \\begin{equation} \\tau_{\\rm cool} = E/L_{\\rm bol} = 17~\\makebox{Gyr}, \\end{equation} which is longer than the Hubble time. The thermal conduction time is calculated as \\begin{eqnarray} \\tau_{\\rm cond} &=& r^2/\\kappa_{\\rm S} = 0.33~\\rm Gyr \\\\\\nonumber &\\times& \\Big(\\frac{r}{100~\\rm kpc}\\Big)^2 \\Big(\\frac{kT_{\\rm Hot}}{1.6~\\rm keV}\\Big)^{-{\\frac{5}{2}}} \\Big(\\frac{n_{\\rm e}}{7\\makebox{$\\cdot$} 10^{-4}~\\rm cm^{-3}}\\Big), \\end{eqnarray} assuming the Spitzer thermal conductivity \\citep{Spitzer1962}. Though this time scale is much shorter than the Hubble time, the thermal conduction may be suppressed by a large factor if there are turbulent magnetic fields, and the time scale may become comparable to or longer than the group age of a few Gyr.  \\citet{Morita2006} suggest that the hydrostatic equilibrium is broken at $r\\sim 5'$, and that the outflow of the hot ICM may occur, based on the result of steep temperature drop at $r\\sim 5$--10$'$. In this case, the hot ICM may be expanding with nearly the sound velocity, $v_{\\rm s} = \\sqrt{\\gamma kT_{\\rm Hot}/(\\mu m_{\\rm p})} = 640$~km~s$^{-1}$, where $\\gamma=5/3$ and $\\mu = 0.62$. The expanding time scale is \\begin{equation} \\tau_{\\rm s} = r/v_{\\rm s} = 0.15\\; \\Big(\\frac{r}{100~\\rm kpc}\\Big) \\Big(\\frac{kT_{\\rm Hot}}{1.6~\\rm keV}\\Big)^{-0.5} ~\\rm Gyr. \\end{equation} % In this case, some non-thermal pressure % (e.g., magnetic field) or a heat source would be necessary to % maintain the hot ICM\\@. % Otherwise, the hot ICM becomes quickly cooler % due to the adiabatic expansion. The required heating power is Thus, the total power input required to form the high-temperature ring can be roughly estimated as, \\begin{equation} W = E/\\tau_{\\rm s} = 1.2\\times 10^{44}~\\makebox{erg~s$^{-1}$}. \\end{equation} This power is by about two orders of magnitude higher than that required to generate the two cavities in the central region of HCG~62\\@. The large implied power and the absence of strong AGN activity in this group suggest that the hot ring may be caused by some large-scale dynamical process in the central region of the group.  \\subsection{Metallicity Distribution in ICM}  The present Suzaku observation of HCG~62 showed abundance distribution of O, Ne, Mg, Si, S, and Fe out to a radius of $10'\\simeq 180$~kpc. Ne abundance has large ambiguity due to a strong coupling with Fe-L lines, as mentioned in section~\\ref{sec:abun}. Distributions of Mg, Si, S, and Fe are quite similar to each other, while O profile in the outer region has a large uncertainty. We plotted abundance ratios of O, Mg, Si, and S against Fe as a function of the projected radius in figure~\\ref{fig:fe_ratio}. Here, the values in the outermost region ($r>9.8'$) were excluded because of large uncertainties. The ratios Mg/Fe, Si/Fe and S/Fe are consistent to be a constant value around 1.5--2, while O/Fe ratio for the innermost region ($r<3.3'$) is significantly lower around 0.5. In the Gal 2T fit, the O/Fe ratio seems to increase with radius. These features have been seen with Chandra and XMM-Newton by \\citet{Morita2006}, and the present result gave a good confirmation in the outer region.  %\u00c2\u00be\u00a4\u00ce\u00a4\u00b9\u00a4\u00b6\u00a4\u00af\u00a4\u00c7\u00a4\u00ce\u00b7\u00eb\u00b2\u00cc\u00a1\u00a3 Recent Suzaku observations have presented abundance profiles in several other systems: an elliptical galaxy NGC~720 \\citep{Tawara2007}, the Fornax cluster and NGC~1404 \\citep{Matsushita2007b}, and a cluster of galaxies Abell~1060 \\citep{Sato2007}. While Si/Fe ratio is almost the same among all the systems, Mg/Fe ratio is slightly higher in HCG~62 and Abell ~1060 than in NGC~720, Fornax cluster, and NGC~1404\\@. We compare the abundances in HCG~62 with those in the Fornax cluster as shown in figure~\\ref{fig:fe_ratio}. Regarding the Fornax cluster, we use $r_{180}= 1.00 \\sqrt{k\\langle T\\rangle / 1.3~\\rm keV}$~Mpc, and $z=0.00429$ corrected to the reference frame defined by the 3~K microwave background radiation (NASA/IPAC Extragalactic Database; NED)\\@. The solar abundance by \\citet{Feldman1992} with $\\rm [Fe/H] = 3.24\\times 10^{-5}$ employed by \\citet{Matsushita2007b} was scaled by a factor of 0.7 to match the \\citet{Anders1989} value of $\\rm [Fe/H] = 4.68\\times 10^{-5}$. HCG~62 shows lower Fe abundance than the Fornax cluster in the central region at $r\\lesssim 0.05\\; r_{180}$, but abundances become similar at $r\\sim 0.1\\; r_{180}$. Abundance ratios of O/Fe, Mg/Fe, Si/Fe, and S/Fe are quite similar between HCG~62 and the Fornax cluster at $r\\sim 0.1\\; r_{180}$. %However, As pointed out in \\citet{Sato2007a}, when the difference is %considered due to the influence of the OBF contaminant, the ratio %seem to drop, the difference becomes insignificant. %\u00a4\u00c8\u00a4\u00cf\u00b8\u00c0\u00a4\u00a6\u00a4\u00e2\u00a4\u00ce\u00a4\u00ce\u00a1\u00a2Fornax\u00a4\u00c8\u00a4\u00ce\u00c8\u00e6\u00b3\u00d3\u00a4\u00f2\u00b5\u00c4\u00cf\u00c0\u00a1\u00a3 %\u00bc\u00e7\u00bb\u00dd\u00a4\u00c8\u00a4\u00b7\u00a4\u00c6\u00a4\u00cf\u00a1\u00a2\u00a5\u00a2\u00a5\u00d0\u00a5\u00f3\u00a5\u00c0\u00a5\u00f3\u00a5\u00b9\u00a5\u00d1\u00a5\u00bf\u00a1\u00bc\u00a5\u00f3\u00a4\u00cf\u00bb\u00f7\u00a4\u00c6\u00a4\u00a4\u00a4\u00eb\u00a4\u00ce\u00a4\u00cb\u00a1\u00a2\u00bc\u00a1\u00a4\u00c7\u00a4\u00a4\u00a4\u00a6 %IMLR\u00a4\u00ac\u00b7\u00eb\u00b9\u00bd\u00b0\u00e3\u00a4\u00a6\u00a4\u00b3\u00a4\u00c8\u00a4\u00f2\u00a4\u00a4\u00a4\u00a6\u00a4\u00bf\u00a4\u00e1\u00a4\u00ce\u00c9\u00db\u00c0\u00d0 %The gradient of the radial Fe profile of HCG~62 is similar %to that of Fornax cluster.  %\u00c5\u00c4\u00c2\u00bc\u00a4\u00b5\u00a4\u00f3\u00cf\u00c0\u00ca\u00b8(2004)\u00a1\u00a2\u00be\u00be\u00b2\u00bc\u00a4\u00b5\u00a4\u00f3\u00cf\u00c0\u00ca\u00b8(2003,2007)\u00a4\u00ca\u00a4\u00c9\u00a4\u00f2\u00b0\u00fa\u00cd\u00d1 \\citet{Tamura2004} reported abundance ratios for 19 clusters studied with XMM-Newton, and the mean Si/Fe ratio in cool and medium temperature clusters with $kT < 6$~keV was $\\sim 1.4$, consistent with our HCG~62 result. Their O/Fe ratio, $\\sim0.6$, in the cluster core also agrees with our result. \\citet{Matsushita2003,Matsushita2007a} also reported abundance ratio for M87 and the Centaurus cluster, respectively. M87 showed the Mg/O ratio to be $\\sim 1.3$ in the central region, and the Centaurus cluster indicated the O/Fe and Si/Fe ratios within $8'$ consistent with our results. The Mg/O ratio for HCG~62 with XMM-Newton \\citep{Morita2006} is $\\sim 3.3$ within $1'$, which is almost the same as our result of $\\sim 3.6$ within 3.3$'$. Mg/O ratios in other groups are $\\sim2.5$ for NGC~5044 \\citep{Tamura2003}, and $\\sim1.3$ for NGC~4636 \\citep{Xu2002} both measured with XMM-Newton RGS\\@.  \\citet{Sato2007c} studied contributions of type Ia and II supernovae to the metal enrichment, based on Suzaku results of HCG~62, Abell~1060, AWM~7, NGC~570 \\citep{Sato2007,Sato2007b}.  \\begin{figure} \\centerline{ \\FigureFile(0.45\\textwidth,1cm){fig/fig11.eps} } \\caption{ Comparison of the Gal 2T results for HCG~62 (black) with results for the Fornax cluster (gray; \\cite{Matsushita2007b}). Fe abundance and the O/Fe, Mg/Fe, Si/Fe, and S/Fe ratios in solar unit \\citep{Anders1989} are plotted against the projected radius scaled by the virial radius, $r_{180}$, from upper to lower panels. Results in $r>9.8'$ for HCG~62 are not plotted for the Fe ratios due to the large uncertainty. The innermost point of the O/Fe ratio is not plotted because of the unsatisfactory fit. }\\label{fig:fe_ratio} \\end{figure}  \\begin{figure*} \\centerline{ \\FigureFile(0.45\\textwidth,1cm){fig/fig12a.eps}% \\hspace*{0.05\\textwidth} \\FigureFile(0.45\\textwidth,1cm){fig/fig12b.eps}% } \\caption{ (a) Cumulative mass, $M(<R)$, within the 3-dimensional radius, $R$, for O, Fe, and Mg are plotted for HCG~62, combining the abundance determination with Suzaku and the gas mass profile with Chandra and XMM-Newton. The cross and diamond markers show the mass profiles with the Gal 2T model, and the O mass profile and its error region with the Gal 1T model are shown in solid and dotted lines. (b) Ratios of the oxygen, iron, and magnesium mass in unit of $M_\\odot$ to the $B$ band optical luminosity in unit of $L_\\odot$ (OMLR, IMLR, and MMLR, respectively) are plotted against the 3-dimensional radius. Black markers represent our results for HCG~62, and gray markers of OMLR and IMLR show the results for the Fornax cluster by \\citet{Matsushita2007b}. %The OMLR points for HCG~62 (black filled circles) denote the Gal 2T results, %and their error bars cover the Gal 1T error regions. }\\label{fig:mass_mlr} \\end{figure*}  \\subsection{Metal Mass-to-Light Ratio}  Metal mass-to-light ratios for oxygen, iron, and magnesium (OMLR, IMLR, and MMLR, respectively), were examined. First, we show the metal mass profiles in figure~\\ref{fig:mass_mlr}(a), based on the 3-dimensional gas mass profile by \\citet{Morita2006} and the abundance profile measured with Suzaku. The derived iron mass within the 3-dimensional radius of $R< 200$~kpc is (1.5--$3)\\times 10^8$ $M_\\odot$, which is quite consistent with the previous measurement in figure~11(e) of \\citet{Morita2006}. We obtained the magnesium mass to be (0.7--$1.1)\\times 10^{8}$ $M_\\odot$. Though the oxygen mass has large errors as seen in figure~\\ref{fig:mass_mlr}(a), it is unlikely that oxygen has smaller mass than iron. Therefore, the Gal 2T model (crosses) is preferred to the Gal 1T one (solid and dotted lines). The estimated oxygen mass is $\\sim 10^9$ $M_\\odot$, within $R< 200$~kpc.  Secondly, we adopted \\citet{Zabludoff2000} results as the member galaxy catalog of HCG~62 (12 galaxies in $r < 13'\\sim 230$~kpc), and their redshifts were used to estimated the 3-dimensional distribution. Since only the $R$-band optical magnitudes are provided in this catalog, we converted them into the $B$-band magnitudes using the color of $B-R = 2.0$~mag for the HCG~62a galaxy at the center \\citep{Hickson1989}, and $A_B = 0.224$ and $A_R = 0.139$ after NASA/IPAC Extragalactic Database (NED) in the direction of HCG~62\\@.  We thus calculated the integrated values of OMLR, IMLR, and MMLR within $r\\lesssim 230$~kpc as shown in figure~\\ref{fig:mass_mlr}(b), and they turned out to be $\\sim 4\\times 10^{-2}$, $\\sim 4.6\\times 10^{-3}$, and $\\sim 1.5\\times 10^{-3}$ $M_{\\odot}/L_{\\odot}$, respectively. %as shown in figure \\ref{fig:mass_mlr}(b). %The differencial OMLR, MMLR, and IMLR for each region is also shown %in figure \\ref{fig:mass_mlr}(c). The errors are only based on the statistical errors of metal abundances in the spectral fit, and the uncertainties of the gas mass profiles and the luminosities of the member galaxies are not considered. The IMLR values are consistent with the collective results with ASCA by \\citet{Makishima2001}. The MMLR and IMLR show similar steep increase with radius up to $r \\sim 100$~kpc and seem to reach almost a plateau at 100--200~kpc. This feature is not apparent for the OMLR, due partly to the large uncertainty. The behavior of MLR curves would be related to different enrichment processes, as discussed in \\citet{Morita2006}.  We also point out that the derived OMLR and IMLR for HCG~62 are much larger than those of the Fornax cluster \\citep{Matsushita2007b} as shown in figure~\\ref{fig:mass_mlr}(b). Particularly the difference in IMLR is significant, although that in OMLR is marginal due to the uncertainty in the Galactic emission. These two systems have a similar potential depth with $k\\langle T\\rangle \\sim 1.3$~keV\\@. This feature may indicate that the Fornax cluster is a younger system with less number of SN~Ia for the Fe production since the cluster formation, and/or that the metal distribution in the Fornax may be much more extended. Possible difference of the initial mass function (IMF) between the two may account for some fraction of the discrepancy in the MLRs. \\citet{Ikebe1996} discovered that there are two distinct length scales of dark matter concentration in the Fornax cluster, and the cD galaxy NGC~1399 is off-centered by $\\sim 50$~kpc with respect to the cluster hot gas. There is also an X-ray luminous elliptical galaxy NGC~1404 at 10$'$ south-west of NGC~1399\\@. These features suggest that the Fornax system may be dynamically young and that galaxy interactions in ICM may have caused extended metal distribution.  \\ifnum1=0 \\begin{table} \\caption{O, Mg, and Fe Masses for the observed clusters} \\label{tab:mass} \\begin{center} \\begin{tabular}{lccl} \\hline & $M_{\\rm O}\\ (M_\\odot)^\\ast $ & $M_{\\rm Mg}\\ (M_\\odot)$ & $M_{\\rm Fe}\\ (M_\\odot)$\\\\ & $\\times10^{9}$ & $\\times10^{7}$ &  $\\times10^{8}$ \\\\ \\hline HCG~62 & 2.2/0.5 & 8.9 & 2.7\\\\ \\hline \\multicolumn{4}{l}{\\parbox{0.45\\textwidth}{\\footnotesize \\footnotemark[$\\ast$] These values correspond to the case with 2 or 1 {\\it apec} models as the Galactic emission}} \\\\ \\end{tabular} \\end{center} \\end{table} \\fi  \\subsection{AGN versus Mergers}  HCG~62 exhibits several interesting activities in both the central and the outer regions: namely, two cavities \\citep{Vrtilek2002, Morita2006}, extended hard X-ray emission \\citep{Fukazawa2001,Nakazawa2007}, multi-phase ICM (\\cite{Morita2006}; this work), possible ``high-abundance arc'' \\citep{Gu2007}, and the doughnut-like high-temperature ring (this work). One might relate the central features with an AGN activity, however, the central galaxy (HCG~62a = NGC~4761) currently shows little evidence of AGN activities in both optical \\citep{Coziol1998,Shimada2000,Coziol2004} and radio bands, as summarized in section~8.1 of \\citet{Morita2006} and section~3.1 of \\citet{Gu2007}. \\citet{Morita2006} also have placed an upper limit on the X-ray luminosity of the AGN to be $L_{\\rm X}\\lesssim 10^{39}$~erg~s$^{-1}$ (0.5--4~keV)\\@. Since the Suzaku HXD-PIN shows no excess over the CXB flux (figure~\\ref{fig:xis-pin-spec}), heavily absorbed AGN with an intrinsic luminosity larger than $\\sim 3\\times 10^{42}$ erg~s$^{-1}$ (15--40~keV) is also ruled out. Another possibility related with the HCG~62 activity can be recent mergers, as suggested in the above mentioned previous works. Energy inputs from magneto-hydrodynamic interactions of the member galaxies with the ICM \\citep{Makishima2001} can be another possible mechanism.  We found  a concentration of hard sources in $r<3.3'$ with average photon indices $\\Gamma\\sim 1.4$. As shown in section~8, this concentration itself is within the range of CXB fluctuation, however, the spatial coincidence with the group center is remarkable as seen in figure~\\ref{fig:image}(b). Since the source luminosities are (3.3--$39)\\times 10^{39}$ erg~s$^{-1}$ (2--10 keV) at HCG~62, we suggest that some of the sources may be remnants of minor mergers which were previously central black holes of the merged galaxies. The overdensity of X-ray sources with $L_{\\rm X} = (4.0$--$250)\\times 10^{39}$ erg~s$^{-1}$ (2--10~keV) in the fields of A~194 and A~1060 was also reported by \\citet{Hudaverdi2006}. Optical identifications of these objects will be of much interest.    \\begin{itemize}\\setlength{\\itemsep}{1ex}  \\item Suzaku confirmed the multi-phase nature of the ICM out to a radius $\\sim 10'$. We found a doughnut-like high-temperature ring at 3.3--6.5$'$ in the hardness image, which is caused by higher hot component intensity over the cool component as shown by the spectral fit. Possible ICM heating by the mass accretion is suggested.  \\item We could not confirm the ``high-abundance arc'' in SW arc in the 1.1--3.3$'$ annulus reported by \\citet{Gu2007}, and possible misidentification of an excess hot component as the line was suggested.  \\item Mg to Fe ratio showed enhancement at the center confirming the previous Chandra and XMM-Newton result. Temperature, surface brightness, and O abundance at $r > 3.3'$ were subject to modeling of the Galactic component, while Mg, Si, S, and Fe abundances were fairly robust. The Gal 2T model was preferred in terms of the surface brightness profile and the integrated O mass.  \\item O abundance was $\\sim 0.4$ solar at the center and less than 0.5 at $r>3.3'$. Abundance ratios, O/Fe, Mg/Fe, Si/Fe, and S/Fe, showed similar values with those in the Fornax cluster in $r\\sim 0.1\\; r_{180}$. Comparison of 19 clusters with HCG~62 showed consistent levels of $\\rm O/Fe\\sim 0.6$ and $\\rm Si/Fe\\sim 1.4$. On the other hand, HCG~62 showed $\\rm Mg/O\\sim 3.6$ at $r<3.3'$, significantly higher than in other groups and giant ellipticals.  \\item The OMLR and IMLR values in HCG~62 are by about order of magnitude higher than the Fornax cluster results at $r\\sim 130$~kpc, while our IMLR agrees with collection of the ASCA results.  \\item Thermal fit and the 6--10~keV image indicated an excess above $\\sim 5$~keV in most of the radii, however it was not significant considering uncertainties in the NXB and CXB fluctuations. ASCA detection of the hard excess is still consistent with our results.  \\item We found an excess X-ray emission of $70\\pm 19$\\% times the nominal CXB intensity (5--12~keV) within $r<3.3'$, and most of it could be explained by a concentration of hard X-ray sources detected with Chandra. We suggested some of the sources could be remnants of minor mergers.  \\end{itemize}  \\bigskip Part of this work was financially supported by the Ministry of Education, Culture, Sports, Science and Technology of Japan, Grant-in-Aid for Scientific Research No.\\ 14079103, 15340088, 15001002, 16340077, 18740011."
    },
    "0711/0711.1729_arXiv.txt": {
        "abstract": "In the presence of self-gravity, we investigate the self-similar dynamics of a relativistically hot gas with or without shocks in astrophysical processes of stellar core collapse, formation of compact objects, and supernova remnants with central voids. The model system is taken to be spherically symmetric and the conservation of specific entropy along streamlines is adopted for a relativistic hot gas whose energy-momentum relation is expressed approximately by $\\varepsilon=c\\mathbbm{p}$ with $\\varepsilon$ and $\\mathbbm{p}$ being the energy and momentum of a particle and $c$ being the speed of light. In terms of equation of state, this leads to a polytropic index $\\gamma=4/3$. The conventional polytropic gas of $P=\\kappa\\rho^\\gamma$, where $P$ is the thermal pressure, $\\rho$ is the mass density, $\\gamma$ is the polytropic index, and $\\kappa$ is a global constant, is included in our theoretical model framework. Two qualitatively different solution classes arise according to the values of a simple power-law scaling index $a$, each of which is analyzed separately and systematically. With explicit conditions, all sonic critical lines appear straight. We obtain new asymptotic solutions that exist only for $\\gamma=4/3$. Global and asymptotic solutions in various limits as well as eigensolutions across sonic critical lines are derived analytically and numerically with or without shocks. By specific entropy conservation along streamlines, we extend the analysis of Goldreich \\& Weber for a distribution of {\\it variable} specific entropy with time $t$ and radius $r$ and discuss consequences in the context of a homologous core collapse prior to supernovae. As an alternative rebound shock model, we construct an Einstein-de Sitter explosion with shock connections with various outer flows including a static outer part of a singular polytropic sphere (SPS). Under the joint action of thermal pressure and self-gravity, we can also construct self-similar solutions with central spherical voids with sharp density variations along their edges. ",
        "introduction": "Radiation pressure (e.g., Chandrasekhar 1939, 1960; Rybicki \\& Lightman 1979), trapped neutrino pressure deep in the stellar interior of extremely high nuclear density, relativistically degenerate materials (e.g., Chandrasekhar 1939), extremely hot interior materials of stars, and processes likely involved in gamma-ray bursts (GRBs) etc. may be approximated by an equation of state with a polytropic index of $\\gamma=4/3$. Statistical physics indicates that the state for a stationary radiation field with particles of no rest mass such as photons is described by a polytropic relation with an index $\\gamma=4/3$. It is proven that $\\gamma=4/3$ is also a very good approximation for relativistically hot particles with negligible rest mass. Moreover for the stellar structure, Chandrasekhar (1939) noted that for an infinitesimal uniform expansion or contraction of a gas sphere, it involves precisely a polytropic process of an index $\\gamma=4/3$. For a static equilibrium configuration and a presumed $P=\\kappa\\rho^{\\gamma}$ with a globally constant $\\kappa$, the virial theorem indicates that $\\gamma<4/3$ situations are unstable and $\\gamma=4/3$ corresponds to a transition from unstable to stable configurations as $\\gamma$ increases. When $\\gamma=4/3$ for a static equilibrium configuration, the equilibrium condition is referred to as the Lane-Emden equation with the total enclosed mass $M$ being independent of the system radius but dependent upon the value of $\\kappa$.  On the other hand, based on the conventional polytropic equation of state $P=\\kappa\\rho^{\\gamma}$, where $\\kappa$ is a global constant and $\\gamma$ varies from $1$ for an isothermal case to $5/3$ for an adiabatic process of monatomic gas, astrophysicists explored properties of hydrodynamic behaviours in diverse contexts, such as star formation, core formation in molecular clouds and supernova explosions etc. For catching the basic physics and theoretical simplicity, most works on gravitational stellar core collapses or stellar explosions were usually carried out under the spherical symmetry. Hunter (1962) considered the stability of an equilibrium system and the collapse process based on a polytropic hydrodynamics. He demonstrated how a dynamical instability during the collapse leads to a breakup of the spherically symmetric radial inflow of gas. Shu (1977) constructed the isothermal expansion-wave collapse solution (EWCS) with a weak discontinuity; and this self-similar hydrodynamic model has been further developed in the past three decades, from the isothermal case (e.g., Shu 1977; Lou \\& Shen 2004) to the polytropic case (e.g., Suto \\& Silk 1988; Yahil 1983; Lou \\& Wang 2006), as well as to the logotropic case (e.g., Mclaughlin \\& Pudritz 1997). Observationally, spectral lines of CS, H$_2$CO and other molecules in star forming regions show that the single peak of each molecular line splits into double peaks with the blue peak brighter than the red peak, which has been regarded as a supporting evidence to the Shu model (e.g. Zhou 1992; Walker, Narayanan \\& Boss 1994; Myers et al. 1996). It is generally expected that shock waves are also involved in self-similar collapse or expansion profiles (e.g., Tsai \\& Hsu 1995; Shu et al. 2002; Bian \\& Lou 2005).  Note that all above studies were carried out on either the isothermal case or the $\\gamma\\neq4/3$ polytropic case. In one case, the polytropic case of $\\gamma=4/3$ is treated as a limiting case (Yahil 1983). In contrast, Goldreich \\& Weber (1980) directly considered homologous core collapse for a conventional polytropic gas with $\\gamma=4/3$ by making use of the time reversal invariance. They concluded that when the pressure decreases by a fraction of no more than about $3\\%$ from a static polytrope in equilibrium, a homologous core collapse would occur in the stellar interior. On the other hand, numerical simulation of Bethe et al. (1979) indicated a fractional reduction of pressure by $26\\%$ in order to initiate a core collapse for a supernova explosion. Goldreich \\& Weber (1980) tried to reduce this difference by introducing an inner core of a progenitor star; Yahil (1983) performed his polytropic analysis and treated their result as a limit of $\\gamma\\rightarrow 4/3^+$.  Meanwhile, specific entropy conservation along streamlines does not necessarily mean a constant specific entropy everywhere at all times. A more general distribution would be a variable specific entropy in time and radius (e.g. Cheng 1977, 1978). Fatuzzo et al. (2004) introduced a self-similar transformation to formulate a more general problem, which involves a scaling index $a$ and another index $q$. The more general equation of state appears to be $P\\propto M^q\\rho^\\gamma$. The $q=0$ case corresponds to the conventional polytropic gas. In fact, according to this more general equation of state, the conservation of mass implies the conservation of specific entropy along streamlines. Nevertheless, Fatuzzo et al. (2004) mainly focused on the isothermal cases with nonzero inward flow speeds far away in molecular clouds (Shen \\& Lou 2004).  Our consideration is on a more general polytropic gas with $\\gamma=4/3$ with the specific entropy conserved along streamlines. By a self-similar transformation, we can approach the resulting nonlinear ordinary differential equations (ODEs) systematically. Solution properties depend on the scaling index $a$. Given a distribution of variable specific entropy with time and radius, the result of Goldreich \\& Weber (1980) can be substantially extended. Meanwhile, many counterparts of previously known solutions in the isothermal and conventional polytropic cases can also be derived. In particular, several new asymptotic solutions unique to $\\gamma=4/3$ are also obtained. An important and interesting result of our analysis is that under the joint action of thermal pressure force and self-gravity, a central spherical void can form and evolve in a self-similar manner; this is to be compared with the central spherical void solution of Fillmore \\& Goldreich (1984b) which considered a collection of collisionless particles under self-gravity in the Einstein-de Sitter expanding universe.  This paper is structured as follows. Nonlinear adiabatic hydrodynamic equations in spherical symmetry and self-similar transformation are described in Section 2 for a polytropic gas with a polytropic index $\\gamma=4/3$. The extensions of the classical analysis of Goldreich \\& Weber (1980) are presented in Section 3 and further discussed for a homologous stellar core collapse in Section 6.1. We mainly focus on cases of $q=2/3$ for various solution properties such as the requirement on the scaling index $a$, the property of scaling invariance, global analytic solutions, the sonic singular surface, the straight sonic critical lines, eigensolutions across the sonic critical line, various asymptotic solutions, and shock jump conditions in Sections 4. We analyze various semi-complete numerical solutions with or without shocks and corresponding results in Section 5. Finally, we conclude and discuss our main results in Section 6. Three Appendices A, B and C are included at the end for technical details of derivations and analyses. ",
        "conclusions": "We have explored and examined the similarity flow solution structures in a general polytropic gas with a polytropic index $\\gamma=4/3$. Previously, Goldreich \\& Weber (1980) considered a special case of $\\gamma=4/3$ with a constant specific entropy. By their assumptions and analysis, only homologous collapse solutions exist by invoking the time reversal invariance, i.e., the radial flow velocity $u$ takes on the form of $2r/(3t)$ until the mass density vanishes at a certain point. Yahil (1983) mentioned the case of Goldreich \\& Weber (1980) as a special limit. In reference to earlier work and based on a self-similar transformation, we systematically examined the case of $\\gamma=4/3$ with specific entropy conservation along streamlines. We have substantially generalized the earlier analyses, discovered new asymptotic solutions, and constructed various self-similar solutions without or with shocks.  In reference to earlier analyses of Goldreich \\& Weber (1980) and Yahil (1983), our model framework mainly focuses on $\\gamma=4/3$ with the conservation of specific entropy along streamlines, which is more general and perhaps, closer to reality than the conventional polytropic gas of a constant specific entropy everywhere at all times. Of course, the case of a constant specific entropy is also possible and can be properly accommodated and treated within our polytropic model framework of $\\gamma=4/3$. Under our more general formalism, we extend the work of Goldreich \\& Weber (1980) and obtain many interesting results. The solutions are divided into two broad classes: solutions with $a=-2/3$ precisely belong to Class I and solutions with $a<-2/3$ belong to Class II. For the situation of $-2/3<a<0$ as mentioned at the beginning of deducing asymptotic behaviours, the divergent velocity at large $x$ is not of interest and hence we only consider two classes I and II solutions.  Class I solutions are characterized by $P\\propto\\rho^{4/3}$ with the proportional coefficient related to the specific entropy being an arbitrary function of $x$, while for Class II solutions, this proportional coefficient depends on the enclosed mass $M$ in a power-law form. We discuss these two classes separately.  \\subsection{Class I Self-Similar Solutions}  Class I self-similar solutions represent a substantial extension of the special solutions with a constant entropy derived by Goldreich \\& Weber (1980). For an astrophysical system such as stars, the specific entropy is not expected to be a global constant in general. For a stellar interior, this depends on the competition between thermal kinetic energy and Fermi energy as determined by the mass density. Qualitatively speaking, especially for a compact object, the closer to the centre, the closer the material is in a degenerate state; this would correspond to a smaller specific entropy. However, the density is relatively small and the temperature is relatively low in the outer part of a star, perhaps also leading to a lower level of entropy. We do not yet know the exact distribution of specific entropy within a star so far. Thus the case of a constant entropy is the simplest to consider and provides a certain sense for a homologous dynamic process. The model analysis of this paper is more general and allows for a fairly arbitrary distribution of specific entropy along streamlines. Meanwhile, the radial velocity profile remains always equal to $2r/(3t)$. For a given time $t$, the radial velocity increases linearly with increasing radius $r$. Hence, this solution can be valid within a finite radial extent. It turns out that the mass density vanishes at some place referred as the outer boundary of the flow system.  According to the model analysis of Goldreich \\& Weber (1980), a pre-collapse progenitor star of a static configuration may evolve into a homologous core collapsing phase (see Figure \\ref{GWf}), when the pressure suddenly decreases by a fraction within a range of $\\sim 2.9\\%$. Early simulations of Bethe et al. (1979) indicated a substantially larger pressure reduction of $26\\%$ is needed in order to initiate collapse in supernova explosions. The much smaller fraction change of pressure reduction for a homologous core collapse given by Goldreich \\& Weber (1980) is actually related to the assumption of a constant specific entropy (i.e., their constant $\\kappa$) in space and time. Requiring specific entropy conservation along streamlines and allowing the specific entropy to be a function of space and time, it is possible to have a homologous core collapse for a much larger fractional change of pressure reduction. In our more general analysis and notations, we find that other forms of $g(x)$ instead of $g(x)=1$ can give rise to a fractional change of $26\\%$ or larger for a pressure reduction. Physically, this corresponds to different distributions of specific entropy along streamlines.  To illustrate this case specifically, we choose $g(x)=1/(1+\\epsilon x)$ where $\\epsilon>0$ is an adjustable parameter to gradually modify the shape of $g(x)$. When $x$ is sufficiently small, $g(x)$ is nearly equal to $1$, analogous to the $g(x)=1$ case. Globally $g(x)$ is a decreasing function with increasing $x$. When $\\epsilon$ is small, $g(x)$ decreases slowly and only deviates from $g(x)=1$ case when $x$ is sufficiently large. For large $\\epsilon>0$, the result will differ considerably from that of Goldreich \\& Weber (1980). We carry out such a $g(x)$ experiment numerically.  Substituting dimensional quantities into the dimensionless state function $p=g(x)\\alpha^{4/3}$, the dimensional equation of state can be written explicitly as \\begin{eqnarray} P=\\frac{g(x)}{A^2}(4\\pi G)^{1/3}\\rho^{4/3}\\ , \\end{eqnarray} where for a given time $t$, $g(x)$ corresponds to a radial distribution of specific entropy. Parameter $A$ also varies for different values of chosen $f(0)$ for a certain system in which the total enclosed mass $M$ is conserved and thus the value of $A$ parameter can be deduced from equation (\\ref{GWM}). The variation in $A$ value actually corresponds to the variation range for fractional change in pressure denoted by $r_p$; in Goldreich \\& Weber (1980), this $r_p$ is $2.9\\%$ as $f(0)$ increases from $f_c$ to infinity. For any given form of $g(x)$, the limiting case is the Lane-Emden equation (e.g., Chandrasekhar 1939) as long as $g(x)\\rightarrow1$ as $x\\rightarrow 0^+$. Of course, this applies to our chosen form of $g(x)=1/(1+\\epsilon x)$ as $x\\rightarrow 0^+$.  \\begin{table} \\begin{center} \\caption{When $g(x)$ takes the form of $1/(1+\\epsilon x)$, the results are shown in Figure \\ref{GWexf}. As $f(0)$ is sufficiently large, the outer boundary at a finite $x$ where $f(x)=0$ is extremely small and thus $g(x)$ can be almost treated as a constant; these results vary little as compared with those of Goldreich \\& Weber (1980) for $f(0)\\rightarrow+\\infty$ corresponding to the Lane-Emden equation. Results in Figure \\ref{GWexf} mainly focus on the limiting case of $f(0)\\rightarrow f_c^+$. For possible $f(x)$ solutions of a homologous core collapse, we use $r_p$ to denote the range for fractional change of pressure variation as $f(0)$ increases from $f_c$ to infinity. }\\label{GWext} \\begin{tabular}{c|c|c|c} $\\epsilon$ & $f_c$ & $\\bar{\\rho}/\\rho_c$ at $f_c$ & $r_p$ \\\\ \\hline $0$ & $4.67047$ & $0.00655$ & $2.9\\%$ \\\\ $0.01$ & $4.58642$ & $0.00693$ & $3.6\\%$ \\\\ $0.05$ & $4.28755$ & $0.00848$ & $6.4\\%$ \\\\ $0.1$ & $3.98354$ & $0.0106$ & $10\\%$ \\\\ $0.3$ & $3.23130$ & $0.0198$ & $26\\%$ \\\\ $0.5$ & $2.83696$ & $0.0293$ & $42\\%$ \\\\ \\hline \\end{tabular} \\end{center} \\end{table}  \\begin{figure} \\includegraphics[width=0.5\\textwidth]{GWex.eps} \\caption{Results of the normalized $f(x)/f(0)$ is displayed for $g(x)=1/(1+0.1x)$ and $\\epsilon=0.1$. In this case, the minimum of $f(0)$ is $f_c=3.9835$ for physical solutions. }\\label{GWexf} \\end{figure}  Table \\ref{GWext} and Figure \\ref{GWexf} show major results for a range of $\\epsilon>0$ values. Numerical experiments indicate that as $\\epsilon$ increases from $0$, the limiting value $f_c$ of $f(0)$ decreases while the density ratio $\\bar\\rho/\\rho_c$ increases. More importantly, the range of fractional change $r_p$ by which the pressure can be reduced for a homologous core collapse becomes larger and larger. Table \\ref{GWext} shows that for $\\epsilon=0.3$, we have $r_p\\cong 26\\%$; in other words, for such a large fractional change in pressure reduction (e.g., Bethe et al. 1979) in order to initiate supernovae, it is still possible for a homologous core collapse prior to the development of a rebound shock. It is conceptually important that a different specific entropy distribution of $g(x)$ from a constant value can lead to a better agreement with numerical simulations; this appears more effective than the inclusion of a less massive core in the centre as mentioned in Goldreich \\& Weber(1980).  \\subsection{Class II Self-Similar Solutions}  In addition to extensions of Goldreich \\& Weber (1980) discussed in the above subsection, we also substantially generalize the self-similar solution space for $\\gamma=4/3$ by adjusting the scaling index $a$. In contrast to $\\gamma\\neq 4/3$, a straightforward analysis with $\\gamma=4/3$ leads to an exact value of $q=2/3$ that is independent of $a$. The dimensional equation of state then takes the form of $P\\propto M^{2/3}\\rho^{4/3}$ with a constant proportional coefficient. We can compare the thermal energy $k_BT$, where $k_B$ is the Boltzmann constant, with the Fermi energy $\\varepsilon_F$. Neglecting the rest mass in the relativistic regime, the relationship between total energy $\\varepsilon$ and momentum $\\mathbbm{p}$ for a single particle can be written as $\\varepsilon=c\\mathbbm{p}$ where $c$ is the speed of light, leading to $\\varepsilon_F\\propto\\rho^{1/3}$. In our model, the state function gives $k_BT\\propto P/\\rho\\propto M^{2/3}\\rho^{1/3}$. It follows that \\begin{eqnarray} \\frac{k_BT}{\\varepsilon_F}\\propto M^{2/3}\\ . \\end{eqnarray} The enclosed mass $M$ is a non-decreasing function in radius $r$. At a given time $t$, we see from this relation that, at small $r$, the enclosed mass is small and hence this ratio is also small. Physically in the inner core of a star, where materials are highly condensed and may be close to a degenerate state, the specific entropy is low.  A modified self-similar transformation is introduced for $\\gamma=4/3$. In the self-similar transformation for a conventional polytropic gas (i.e., $P=\\kappa\\rho^{\\gamma}$ with globally constant $\\kappa$ at all times), the sound speed appears either explicitly or implicitly. In contrast, we here use an integration constant $C_0$ relating to the sound speed and transformation (\\ref{eq07}) does not involve the sound speed because of the uniqueness of the $\\gamma=4/3$ case; this $C_0$ coefficient is in fact allowed by the transformation and is an adjustable parameter in our analysis for astrophysical applications. At a deeper level, we realize that as the special self-similar transformation (\\ref{eq07}) for $\\gamma=4/3$ does not involve the sound speed, we have a scaling invariance (\\ref{scaleinvariance}) which simplifies our theoretical analysis considerably.  By comparisons and analogies of solutions known for $\\gamma\\neq4/3$, we try our best to find the counterpart solutions and to discover new solutions for $\\gamma=4/3$. Global analytic solutions, i.e., static SPS solution (\\ref{static}) and (\\ref{staticcond}) and Einstein-de Sitter expansion solution (\\ref{eqq}), still exist for $\\gamma=4/3$ with some modifications. Analytic asymptotic solutions of various kinds are also derived for both large and small $x$. However, the LP-type solution no longer exists for $\\gamma=4/3$ (thus $q=2/3$) except for rare situations, while other counterpart solutions are readily found. In addition, a new type asymptotic solution (\\ref{new1}) is discovered in the regime of small $x$. It seems that this type of asymptotic polytropic solutions only exists for $\\gamma=4/3$.  We have also examined properties of the sonic singular surface and the sonic critical line of coupled nonlinear ODEs (\\ref{eq11}). A salient feature of $\\gamma=4/3$ case is that all sonic critical lines are straight and pass through the origin $x=0$ and $v=0$ in the $-v(x)$ versus $x$ presentation; while first revealed by extensive numerical experiments, these remarkable results can be proven analytically. It is also fairly straightforward to derive two eigensolutions to smoothly cross the straight sonic critical line. In later analyses, we realize that the sonic critical line cannot fully represents the behaviours of the sonic singular surface, especially for constructing shocks. We hence use $v_c(x)$ defined by equation (\\ref{defvc}) for each solution, which is another curve on the sonic singular surface and tightly relates to the current solution, to show the interrelation between the current solution and the sonic singular surface. According to definition (\\ref{defvc}), the solution meets the singular surface if and only if the solution and the corresponding $v_c$ intersects; meanwhile, a shock solution jumping across the sonic singular surface also jumps across the corresponding $v_c$ in $v(x)$ versus $x$ plane. % The standard Runge-Kutta scheme (e.g., Press et al. 1986) is used to numerically integrate coupled nonlinear ODEs (\\ref{eq11}) and (\\ref{eq12}) to connect various asymptotic solutions and eigensolutions across the sonic critical line. We also construct possible solutions with and without shocks for potential astrophysical applications. From the behaviours of these semi-complete solutions, we can see that many such solutions have similar behaviours as those for $\\gamma\\neq 4/3$ in a qualitative manner. And our solutions can be sensibly regarded as limits of $\\gamma\\rightarrow4/3$ when a polytropic gas becomes relativistically hot or degenerate. Our analysis for $\\gamma=4/3$ does share a certain common characteristics with cases for $\\gamma\\neq4/3$.  We can readily solve for eigensolutions crossing the straight sonic critical line. Once the slope $k$ of the sonic critical line is positive, we can construct a new type of self-similar solution characterized by an expanding central spherical void within which the enclosed mass is zero or negligible. Around the edge of such an expanding void, there exists an overdensed shell where the density variation becomes rather steep. Diffusion processes are expected to smooth out such relatively steep gradients locally. To consider properties of spherical void boundary, we first note that such a void expands with a radial speed \\begin{eqnarray} \\dot{r}_e=-ar_e/t\\qquad\\quad\\Rightarrow \\qquad\\quad r_e\\propto t^{-a}\\ , \\end{eqnarray} where $r_e$ stands for the radius of the spherical void edge. The spherical void edge evolves as a power law of time $t$ with a scaling index $-a$. One also notes by equation (\\ref{voidlimit}) that the density gradient approaches a negative infinity near the void edge. We expect that in a narrow region near the spherical void edge, materials are actually diffused instead of being so sharply distributed as shown by our solution mathematically. Within this narrow region, the local evolution does not behave self-similarly and may not be spherically symmetric, while the overall self-similar profile remains on large scales. In the outer part, the mass density scales as $\\rho\\propto r^{2/a}$ and the radial flow velocity remains finite with a wind. In fact, it is also possible to construct various shock solutions with a central void.  At this stage, we may outline a physical scenario in the context of a supernova explosion. During the core collapse of a progenitor, neutrons are formed in abundance and neutrinos of relativistic energies are released. In a high-density environment, neutrino opacity is extremely high so that neutrino pressure, radiation pressure and gas pressure work together to drive the central core expansion. In the relativistic regime, we may ignore tiny neutrino masses and regard the neutrino gas as polytropic with an index $\\gamma_{n}=4/3$. Similarly, the radiation pressure resulting from the photon gas trapped in the stellar interior can also be regarded as polytropic with an index $\\gamma_{\\nu}=4/3$. In the hot stellar core of high temperatures, we may approximate the thermal gas pressure as polytropic with an index $\\gamma_p\\cong 4/3$. It might be conceivable that under certain situations, the neutrino pressure is so overwhelming such that a central void may start to form. As the outer part expands and density drops, neutrinos escape while the radiation and thermal gas pressures continue to drive the expansion. It should be emphasized that in real situations, a grossly spherical void may still encompass materials here and there but the mass density inside is substantially lower than that of surroundings.  In this context, we note the model work of Fillmore \\& Goldreich (1984b) who considered a collection of collisionless particles in an expanding universe of the Einstein-de Sitter form. There is no pressure effect from particles in their model. In essence, the background Einstein-de Sitter expansion prescribed is similar to the rapid expansion driven by the thermal pressure force in our model, both providing the tendency for particles to move outwards in competition with the inward self-gravity. The key physical difference is that the Einstein-de Sitter expansion of the universe is homogeneous (presumably driven by the ubiquitous dark energy) while our gas expansion is driven by the thermal gas pressure closely related to gas mass density and temperature. Not only in the case $\\gamma=4/3$, self-similar void solutions can also be constructed for $\\gamma\\neq4/3$ and isothermal gas which we shall investigate more thoroughly in separate papers.  Besides certain similarities with previous polytropic model analysis with $\\gamma\\neq 4/3$, the case of $\\gamma=4/3$ carries its own unique features. First, because of scaling invariance (\\ref{scaleinvariance}), various self-similar solutions can be readily classified, especially for the two eigensolutions across the sonic critical line. Once solution properties at a chosen point $x$ have been examined completely, other points will have the same solution characteristics by scaling invariance (\\ref{scaleinvariance}). This simplifies the analysis to a considerable extent. Fundamentally, the cause of this scale invariance (\\ref{scaleinvariance}) is due to the fact that the sound speed is not involved in self-similar transformation (\\ref{eq07}). In various solutions, the case of $\\gamma=4/3$ also shows some differences: (i) the LP type solution does not exist, except for rare situations (see Appendix \\ref{LP}); (ii) when discussing the quasi-static solution at small $x$, two sensible roots may be found for $\\gamma<4/3$ (see Lou \\& Wang 2006), while one of the two roots is always unacceptable for $\\gamma=4/3$ with only one sensible root being available in our model calculations; (iii) it is no longer possible for a quasi-static solution at small $x$ to show a vibration behaviour here. The solution quickly converges to an outer asymptotic solution; and (iv) the sonic critical lines with constant density are straight lines emanating from the origin in the $-v(x)$ versus $x$ presentation.  \\subsection{Conclusions}  To sum up this paper, we have explored possible self-similar solutions, both analytical and numerical, for a generalized polytropic gas with $\\gamma=4/3$. The classical analysis of Goldreich \\& Weber (1980) for a constant specific entropy everywhere at all times is substantially extended by specific entropy conservation along streamlines with specific entropy dependent on time and space. This differs from what Yahil (1983) did in this context. In addition to counterparts of various previously known types of polytropic solutions with $\\gamma\\neq 4/3$, we find two new asymptotic solutions. One notable feature is that sonic critical lines are all straight lines emanating from the origin in the $-v(x)$ versus $x$ presentation with constant densities. Using all asymptotic solutions available, two eigensolutions across the sonic critical line and self-similar shock conditions, global semi-complete solutions are constructed numerically. Two classes of self-similar solutions are investigated separately according to the value of the scaling index $a$. For Class I solutions with $a=-2/3$ precisely, we can simulate the homologous evolution of a flow system once the distribution of specific entropy is prescribed. These more general solutions for a homologous core collapse (Goldreich \\& Weber 1980) may be utilized to model the dynamic formation of an inner compact core from a pre-collapse stellar interior. Collapsing solutions of Class II with $a<-2/3$ may also explain the formation of compact objects and other similar flow systems, while expansion solutions with shocks can be utilized to model supernova explosions (e.g., Lou \\& Wang 2006, 2007).  By equations (\\ref{eq1})$-$(\\ref{eq5}), we emphasize several aspects of solutions for $(3a+2)=0$ and for $(3a+2)\\rightarrow 0$. Based on our analysis, the Einstein-de Sitter solution with $\\gamma=4/3$ exists for $(3a+2)\\neq 0$, $(3a+2)\\rightarrow 0$ and $(3a+2)=0$. Except for this special Einstein-de Sitter solution, Class I solutions valid for $(3a+2)=0$ cannot be obtained by taking the limit of $(3a+2)\\rightarrow 0$ for Class II solutions. In other words, Class I and II solutions are qualitatively different solutions and we need to consider them separately. By equations (\\ref{eq1}) and (\\ref{meq}) during the limiting process of $(3a+2)\\rightarrow 0$, we must require $(ax+v)\\rightarrow 0$ in order to have a finite reduced mass $m(x)$. Only the Einstein-de Sitter solution and Class I solutions with $(3a+2)=0$ bear this unique feature for the reduced flow speed $v(x)$ while all other Class II solutions are excluded by this limiting procedure.  In the course of investigation, we realize the possibility of constructing self-similar solutions for dynamic evolution of central spherical void in a flow system involving self-gravity and thermal pressure. Here, the thermal pressure force drives the gas expansion sufficiently fast and creates a central spherical void by pushing materials outwards. By specific examples, we now prove by analytical and numerical calculations that a spherical void can indeed form in astrophysical flow systems under the joint action of thermal pressure force and self-gravity. We expect that such processes could happen in association with supernova explosions and evolution of supernova remnants.  At the beginning of our model formulation, several physical effects, such as nuclear reactions, radiation pressure, neutrino transport, general relativistic effects, rotational effects and magnetic field, are not taken into account. Under various situations, these effects can be very important in real astrophysical systems. Given these approximations and idealizations of our model, it is still hoped that this simple theoretical model framework may catch certain essential characteristics or features of flow phenomena of relevant scenarios and interpretations."
    },
    "0711/0711.3801_arXiv.txt": {
        "abstract": "A simple model of the coronal magnetic field prior to the CME eruption on May 12 1997 is developed. First, the magnetic field is constructed by superimposing a large-scale background field and a localized bipolar field to model the active region (AR) in the current-free approximation. The background field is determined from the normal component of the observed photospheric magnetic field averaged over the longitude of the Sun. The AR field is modeled with the help of a subphotospheric dipole whose strength, location, and orientation are optimized to fit the magnetic field obtained from an MDI magnetogram. Second, this potential configuration is quasi-statically sheared by photospheric vortex motions applied to two flux concentrations of the AR. Third, the resulting force-free field is then evolved by canceling the photospheric magnetic flux with the help of an appropriate tangential electric field applied to the central part of the AR.  To understand the structure of the modeled configuration, we use the field line mapping technique by generalizing it to spherical geometry. It is demonstrated that the initial potential configuration contains a hyperbolic flux tube (HFT) which is a union of two intersecting quasi-separatrix layers. This HFT provides a partition of the closed magnetic flux between the AR and the global solar magnetic field. Such a partition is approximate since the entire flux distribution is perfectly continuous. The vortex motions applied to the AR interlock the field lines in the coronal volume to form additionally two new HFTs pinched into thin current layers. Reconnection in these current layers helps to redistribute the magnetic flux and current within the AR in the flux-cancellation phase. In this phase, a magnetic flux rope is formed together with a bald patch separatrix surface wrapping around the rope. Other important implications of the identified structural features of the modeled configuration are also discussed. ",
        "introduction": "The May 12 1997 CME occurred during solar minimum in an isolated active region near disk center. Since the solar magnetic field during this period had minimum complexity, this event is especially favorable for detailed study. Therefore, it was selected by several working groups (SHINE, CISM and MURI) for modeling and analysis. Several observational features of the event have been reported \\citep{ Hudson1998, Plunkett1998, Sterling2001, Thompson1998, Webb2000, Gopalswamy2002, Lundquist2003, Arge2004, Li2004, Liu2004, Attrill2006, Attrill2007}, which has provided a good base for creating a magnetohydrodynamic (MHD) model of the event. For the latter, we have used the following three-step approach.  First, a potential magnetic field is determined from the observed photospheric distribution of the normal field component. Second, the configuration is sheared quasi-statically by photospheric vortex motions applied to the opposite polarities of the active region (AR). Third, an appropriate tangential electric field is applied then to the central part of the AR to cancel half of the photospheric magnetic flux. This third step is motivated by the observed flux-canceling   properties of the photospheric motions in the event \\citep{Li2004, Ambastha2001}. Theoretically, the third step is intended first to form a flux rope and then to destabilize it by detaching the rope field from the photospheric one, which has eventually to produce its eruption into the corona and heliosphere. This complete modeling program has begun by our group and it will be described in forthcoming papers. Preliminary results for these simulations have been presented by \\citet{LinkeretalFAGU2006}. In this first paper, we consider a simplified model of the magnetic field and its structure at each of the above three steps in order to identify the generic structural features of the configuration prior to its eruption. By the term  ``generic\" we mean those features of the structure that are likely to remain in more accurate models of the AR magnetic field  (AR8038 at N21W08).  Observationally, this field  had an apparently simple bipolar structure \\citep{Liu2004} dominating over a more complicated but distant and weak surrounding field. As a starting point, it is numerically extrapolated in the potential approximation from the synthetic magnetogram obtained by merging an MDI magnetogram with the corresponding synoptic map. From the methodological point of view, it is sensible, however, to begin with a simpler model by approximating the numerical potential field by a sum of two components, an azimuthally averaged large-scale magnetic field plus a bipolar field of the active region itself. This model, comprising essentially two bipolar field components with different scales and orientations, has a deceptive simplicity. It turns out that the structure of such a field is nontrivial even at the initial current-free state, and it becomes even more complicated at the shearing and flux-canceling phases, features that have important implications for understanding the event.   The analysis of this structure prior to eruption is a major goal of the present paper. The paper is organized in the following way: the construction of the field model is described in sections \\ref{s:ipmf} and \\ref{s:mhd_ev}; in section \\ref{s:flcsg}, the so-called squashing factor is generalized to the case of spherical geometry in order to investigate in section \\ref{s:sif} the structure of the modeled field. Section \\ref{s:imp} describes the relationship between field line connectivity and the current distribution. Our results are summarized in section \\ref{s:s}. Some details of the initial potential field model are provided in the Appendix. ",
        "conclusions": "\\label{s:s}  We have developed a simplified model of the coronal magnetic field prior to the CME on May 12 1997 in the following three steps. First, this field is constructed by superimposing a global background field ${\\bm B}_{\\sun}$ and a localized bipolar field ${\\bm B}_{\\rm AR}$ to model the active region (AR) in the potential approximation. The background field ${\\bm B}_{\\sun}$ is determined from the observed photospheric normal field averaged over the longitude of the Sun. The AR field is modeled by a subphotospheric dipole whose parameters were optimized to fit the magnetic field obtained from an MDI magnetogram. To verify the impact of the solar wind on the field structure, ${\\bm B}_{\\sun}$ is calculated for two different sets of boundary conditions. In the first case, ${\\bm B}_{\\sun}$ is modeled as a closed potential field vanishing at infinity. In the second case, ${\\bm B}_{\\sun}$ is constrained to be strictly radial and therefore open at the sphere of radius $r=2.5\\, R_{\\sun}$. The resulting field structures \\onecolumn \\begin{figure}[htbp] \\epsscale{0.8} \\plotone{f11} \\caption{Distributions of $\\log Q$ [(a), (d), (g), (l)] vs. $\\alpha=j_{\\|}/B$ [(b), (e), (h), (m)] in the vicinity of flux concentrations at the end of shear/twist [(a)--(f)] and flux-cancellation [(g)--(n)] phases ($\\alpha$ is normalized to $1/R_{\\sun}$); panels (a)--(c) and (g)--(k) show these distributions at the photospheric $(\\phi,\\theta)$-plane, while panels (d)--(f) and (l)--(k) show them at the $(\\theta,r)$-plane at $\\phi \\approx 3.017$ (yellow dashed line in the photospheric distributions). Superposition of  the corresponding semi-transparent $\\log Q$ distributions and $\\alpha$ distributions is depicted by panels (c), (f), (k), (n). The PILs and contours of $\\left. B_{r} \\right|_{r=R_{\\sun}}$ on the photospheric distributions are the same as in Figure \\ref{f8}. \\label{f11} } \\end{figure} \\twocolumn {\\parindent=0pt in the AR and its neighborhood are quite similar for both choices of boundary conditions, therefore only the first (closed) potential configuration has been used in the next step of MHD modeling. In this second step, the potential field is quasi-statically sheared by photospheric vortex motions applied to the two flux concentrations of the AR. Finally, the resulting force-free field is evolved by canceling the photospheric flux with the help of an appropriate tangential electric field applied to the central part of the AR.}  To investigate the structure of the modeled configurations, the field line mapping technique has been generalized to the case of spherical geometry. In our potential field configuration, this technique reveals a large-scale hyperbolic flux tube (HFT) characterized by very large values of the squashing factor. As in quadrupole-type configurations, such an HFT consists of two intersecting quasi-separatrix layers (QSLs) that provide an approximate partition of magnetic fluxes between four photospheric field sources. In our case, the role of such sources is played by two flux concentrations of the AR and two photospheric polar regions.  It is important that no other structural features, such as magnetic null points or ``bald patches'', are found in the AR and its surrounding area of the initial potential configuration. In this respect, our initial configuration drastically differs from that used in the breakout model \\citep{Antiochos1999}, where the null points and separator are key topological elements of the structure. On the other hand, our large-scale HFT is a generalized analogue of the separator field line \\citep{Titov2002} and is therefore a favorable site for current layer (CL) formation \\citep{Titov2003, Galsgaard2003} and magnetic reconnection \\citep{Hornig2003a, Aulanier2005, Pontin2005}. The accumulation of current density observed in the vicinity of this HFT during the shearing/twisting phase of our simulations is in agreement with this point of view. The mid X-type cross-section of the HFT pinches in the process of this accumulation into a \\Ht-like shape with a strong CL at the center.  This structure is preserved during the cancellation phase when the flux rope is formed. Had the eruption conditions been met in this phase, the flux rope that would form would most likely detach from the dense photosphere and rise by pushing the overlying field upward. This would cause an additional pinching of the CL to subsequently trigger reconnection within it. Such a reconnecting CL is a basic element of the breakout model \\citep{Antiochos1999} and we have demonstrated that it appears in the flux-cancellation model as well if the interaction of the AR with the global background field is included. With this generalization, the only difference between the models is in their interpretation of the drivers of the eruption. The breakout model assumes that the driver is reconnection in the CL, allowing for an explosive expansion of the sheared arcade of the AR. The flux-cancellation model assumes instead that the driver is a rising flux rope unleashed from the strapping field of the AR. While the reconnecting CL is important in our generalized model, the latter mechanism appears to be more likely to drive the eruption.  The following estimate of the magnetic fluxes in the interacting parts of our configuration suggests a picture of how this occurs. The flux associated with the background field is eight times larger than the initial flux in our AR. However, only certain fractions of these fluxes interact with each other. These fractions can be estimated from the geometrical properties of the quasi-separator shown in Figure \\ref{f7}. Since the longitudinal extension of the quasi-separator is about 0.55 radians (Fig. \\ref{f5}b), the fraction of the background magnetic flux overlying the quasi-separator is equal to $8\\times 0.55/(2\\pi)\\approx 0.7$. We must take into account also that only field components perpendicular to the quasi-separator can participate in reconnection (reducing the fraction of the overlying flux from 0.7 to 0.4), and that the flux of the AR is reduced to half of its initial value at the end of the flux-cancellation phase. With these corrections, we obtain the interacting fractions of the background and AR fluxes expressed by the ratio $\\Phi_{{\\rm i}\\sun} : \\Phi_{{\\rm i\\, AR}} = 0.4 : 0.5 $. In other words, the interacting fluxes of the background and AR field are comparable. It implies that a substantial part of $\\Phi_{{\\rm i\\, AR}}$ must be converted into the magnetic flux associated with the rising unstable flux rope in order to allow it to break through the overlying background field. Such a conversion is likely provided by magnetic reconnection in the ``vertical'' CL below the rope, as observed in previous simulations by \\citet{Amari2000} and \\citet{Kliem2004}. This process ought to replenish the magnetic flux in the rope that is gradually destroyed by reconnection in the ``horizontal\" CL above the rising rope.   We have also demonstrated that two small-scale HFTs pinched into thin CLs are formed by nonuniform vortex motions applied to the flux spots of the AR. The reconnection in these CLs helps redistribute the magnetic flux between the AR and the surrounding field during the flux-cancellation phase. Other possible implications of their presence should be explored in a more accurate model of the magnetic field evolution in this event. However, beyond its manifestation in this particular event, the discovered process of current layer formation by rotation of the flux spots has a broader significance. This process seems to be very important for understanding other eruptive/flaring events driven by rotating sunspots \\citep{Brown2003}. Our example shows that their rotation can serve as a mechanism of both the storage of free magnetic energy and its release via gradual formation of the above CLs in the solar corona. More detailed studies are certainly required in this direction.  We have also confirmed that the photospheric flux-canceling process builds a magnetic flux rope at the polarity inversion line \\citep{vanBall1989}. This flux rope appears as a well-defined element of the structure owing to the bald patch separatrix surface that wraps around the rope. The magnetic structure near the rope is topologically equivalent to the twisted configuration described by \\citet{Titov1999}. In particular, the bald patch separatrix surface has a characteristic inverse-S-like shape. This type of sigmoid corresponds to the negative sign of the current helicity in the flux rope, which is in accordance with observations \\citep{Liu2004, Pevtsov1997}."
    },
    "0711/0711.2420_arXiv.txt": {
        "abstract": " ",
        "introduction": "A key challenge of current astrophysical research is to obtain information about the equation of state (EoS) of the ultra-dense nuclear matter making up neutron stars (NSs). The observational constraints on the internal structure of NSs are weak: the observed range of NS masses is $M\\sim 1.1 - 2.2 M_\\odot$ \\citep{Lattimer}, and there is no current method to directly measure the radius. Some estimates using data from X-ray spectroscopy exist, but those are highly model-dependent  (e. g. \\citet{astroph}). Different theoretical models for the NS internal structure predict, for a neutron star of mass $M\\sim 1.4 M_\\odot$, a central density in the range of $ \\rho_c\\sim 2-8 \\times 10^{14}{\\rm g}{\\rm cm}^{-3}$ and a radius in the range of $R\\sim 7-16$km \\citep{Lattimer}. Potential observations of pulsars rotating at frequencies above $1400$Hz could be used to constrain the EoS if the pulsar's mass could also be measured (e. g. \\citet{Bejger}).    Direct and model-independent constraints on the EoS of NSs could be obtained from gravitational wave observations. Coalescing binary neutron stars are one of the most important sources for ground-based gravitational wave detectors \\citep{Cutler:2002me}. LIGO observations have established upper limits on the coalescence rate per comoving volume \\citep{Abbott:2007xi}, and at design sensitivity LIGO II is expected to detect inspirals at a rate of $\\sim 2/$day \\citep{2004ApJ...601...L179}.  In the early, low frequency part of the inspiral ($f\\leq 100$Hz, where $f$ is the gravitational wave frequency), the waveform's phase evolution is dominated by the point-mass dynamics and finite-size effects are only a small correction. Toward the end of the inspiral the internal degrees of freedom of the bodies start to appreciably influence the signal, and there have been many investigations of how well the EoS can be constrained using the last several orbits and merger, including constraints from the gravitational wave energy spectrum \\citep{faber} and from the NS tidal disruption signal for NS-black hole binaries \\citep{vallisneri}. Several numerical simulations of the hydrodynamics of NS-NS mergers have studied the dependence of the gravitational wave spectrum on the radius and EoS (see, e.g. \\citet{baumgarte} and references therein). However, trying to extract EoS information from this late time regime presents several difficulties: (i) the highly complex behavior requires solving the full nonlinear equations of general relativity together with relativistic hydrodynamics; (ii) the signal depends on unknown quantities such as the spins and angular momentum distribution inside the stars, and (iii) the signals from the hydrodynamic merger are outside of LIGO's most sensitive band.   During the early regime of the inspiral the signal is very clean and the influence of tidal effects is only a small correction to the waveform's phase. However, signal detection is based on matched filtering, i. e. integrating the measured waveform against theoretical templates, where the requirement on the templates is that the phasing remain accurate to $\\sim 1$ cycle over the inspiral. If the accumulated phase shift due to the tidal corrections becomes of order unity or larger, it could corrupt the detection of NS-NS signals or alternatively, detecting a phase perturbation could give information about the NS structure. This has motivated several analytical and numerical investigations of tidal effects in NS binaries \\citep{bc,ks,kochanek,ts,mw,shibata,gualteri,pons,berti}. The influence of the internal structure on the gravitational wave phase in this early regime of the inspiral is characterized by a single parameter, namely the ratio $\\lambda$ of the induced quadrupole to the perturbing tidal field. This ratio $\\lambda$ is related to the star's tidal Love number $k_2$ by $k_2 = 3G \\lambda R^{-5}/2$, where $R$ is the star's radius. \\citet{tidal} have shown that for an inspiral of two non-spinning $1.4M_{\\odot}$ NSs at a distance of 50 Mpc, LIGO II detectors will be able to constrain $\\lambda$ to $\\lambda \\leqslant 2.01 \\times 10^{37} {\\rm g}{\\,}{\\rm cm}^2{\\rm s}^2 $ with $90\\%$ confidence. This number is an upper limit on $\\lambda$ in the case that no tidal phase shift is observed. The corresponding constraint on radius would be $R \\leqslant 13.6 {\\,}{\\rm km} {\\;}\\left(15.3{\\,} {\\rm km} \\right)$ for a $n=0.5$ $\\left(n=1.0\\right)$ fully relativistic polytrope, for $1.4M_{\\odot}$ NSs \\citep{tidal}.   Because neutron stars are compact objects with strong internal gravity, their Love numbers could be very different from those for Newtonian stars that have been computed previously, e. g. by \\citet{bo}.  Knowledge of Love number values could also be useful for comparing different numerical simulations of NS binary inspiral by focusing on models with the same masses and values of $\\lambda$.     In \\citet{tidal}, the $l=2$ tidal Love numbers for fully relativistic neutron star models with polytropic pressure-density relation $P=K\\rho^{1+1/n}$, where $K$ and $n$ are constants, were computed for the first time. The present paper will give details of this computation. Using polytropes allows us to explore a wide range of stellar models, since most realistic models can be reasonably approximated as a polytrope with an effective index in the range $n\\sim 0.5-1.0$ \\cite{Lattimer}. Our prescription for computing $\\lambda$ is valid for an arbitrary pressure-density relation and not restricted to polytropes. In Sec. \\ref{sec:lovedef}, we start by defining $\\lambda$ in the fully relativistic context in terms of coefficients in an asymptotic expansion of the metric in the star's local asymptotic rest frame and discuss the extent to which it is uniquely defined. In Sec. \\ref{sec:lovecalc}, we discuss our method of calculating $\\lambda$, which is based on static linearized perturbations of the equilibrium configuration in the Regge-Wheeler gauge as in \\citet{tc}. Section \\ref{sec:results} contains the results of the numerical computations together with a discussion. Unless otherwise specified, we use units in which $c=G=1$. ",
        "conclusions": "\\label{sec:results}  The range of dimensionless Love numbers $k_2$ obtained by numerical integration of Eq. (\\ref{k2expr}) is shown in Fig. \\ref{dimlam} as a function of $M/R$ and $n$ for a variety of different neutron star models, and representative values are given in Table \\ref{tbl-1}. These values can be approximated to an accuracy of $\\sim 6\\%$ in the range $0.5\\leq n\\leq 1.0$ and $0.1\\leq (M/R)\\leq 0.24$ by the fitting formula \\begin{equation} k_2\\approx \\frac{3}{2} \\left(-0.41+\\frac{0.56}{n^{0.33}}\\right)\\left(\\frac{M}{R}\\right)^{-0.003}.\\end{equation}  Both Fig. \\ref{dimlam} and Table \\ref{tbl-1} illustrate that the dimensionless Love numbers $k_2$ depend more strongly on the polytropic index $n$ than on the compactness $C=M/R$. \\footnote{Note, however, that LIGO measurements will yield the combination $k_2 R^5$ and therefore will be more sensitive to the compactness than the polytropic index.} This is expected since the weak field, Newtonian values $k_2^N$ given by Eq. (\\ref{newtk2}) just depend on $n$ (through the dependence on $y$). The additional dependence on the compactness for the Love numbers $k_2$ in Eq. (\\ref{k2expr}) is a relativistic correction to this. For $M/R\\sim 10^{-5}$ our results for $k_2$ agree well with the Newtonian results of \\citet{bo}. Figure \\ref{diff} shows the percent difference $(k_2^N-k_2)/k_2$ between the relativistic and Newtonian dimensionless Love numbers. As can be seen from the figure, the relativistic values are lower than the Newtonian ones for higher values of $n$. This can be explained by the fact that the Love number encodes information about the degree of central condensation of the star. Stars with a higher the polytropic index $n$ are more centrally condensed and therefore have a smaller response to a tidal field, resulting in a smaller Love number.      Some estimates of the masses and radii of neutron stars, given in Table \\ref{tbl-2}, have been inferred from X-ray observations \\citep{nature,astroph} using the information from three measured quantities: the Eddington luminosity, the surface redshift of spectral lines, and the quiescent X-ray flux. The range of the numbers $\\lambda$ for these stars is shown in Fig. \\ref{xray}. LIGO II detectors will be able to establish a $90\\%$ confidence upper limit of $\\lambda \\leqslant 2.01 \\times 10^{37} {\\rm g}{\\,}{\\rm cm}^2{\\rm s}^2 $ for an inspiral of two nonspinning $1.4 M_\\odot$ NSs at a distance of 50 Mpc in the case that no tidal phase shift is observed \\citep{tidal}."
    },
    "0711/0711.2085_arXiv.txt": {
        "abstract": "{ We report the results of our study of magnetic fields in a sample of 15~Be stars using spectropolarimetric data obtained at the European Southern Observatory with the multi-mode instrument FORS\\,1 installed at the 8\\,m Kueyen telescope. We detect weak photospheric magnetic fields in four stars, HD\\,56014, HD\\,148184, HD\\,155806, and HD\\,181615. We note that for HD\\,181615 the evolutionary status is not obvious due to the fact that it is a binary system currently observed in the initial rapid phase of mass exchange between the two components. Further, we notify the possible presence of distinct circular polarisation features in the circumstellar components  of Ca\\,{\\sc ii} H\\&K in three stars, HD\\,58011, HD\\,117357, and HD\\,181615, hinting at a probable presence of magnetic fields in the circumstellar mass loss disks of these stars. We emphasize the need for future spectropolarimetric observations of Be stars with detected magnetic fields to study the temporal evolution of their magnetic fields and the correlation of magnetic field properties with dynamical phenomena taking place in the gaseous circumstellar disks of these stars. } ",
        "introduction": "\\label{sect:intro}  Be stars are defined as rapidly rotating main sequence stars showing normal B-type spectra with superposed Balmer emissions. Further, these stars are characterized by episodic dissipation and formation of a new circumstellar (CS) disk-like environment, non-radial pulsations and photometric and spectroscopic variability. A number of physical processes in classical Be stars (e.g., angular momentum transfer to a CS disk, channeling stellar wind matter, accumulation of material in an equatorial disk, etc.) are more easily explainable if magnetic fields are invoked. In recent years this conclusion was strengthened by numerous theoretical works (e.g., Cassinelli et al.\\ \\cite{2002ApJ...578..951C}, Maheswaran \\cite{2003ApJ...592.1156M}, Brown et al.\\ \\cite{2004MNRAS.352.1061B}). A detailed discussion concerning the mechanisms involved in formation and evolution of disks around Be stars can be found in a recent publication by Owocki (\\cite{2006ASPC..355..219O}).   To date we know only three Be stars with magnetic field detections: $\\omega$\\,Ori (80$\\pm$40\\,G -- Neiner et al.\\ \\cite{2003A&A...409..275N}), $\\beta$\\,Cep (95$\\pm$8\\,G -- Donati et al.\\ \\cite{2001MNRAS.326.1265D}), and 16\\,Peg ($-$156$\\pm$31\\,G -- Hubrig et al.\\ \\cite{2006MNRAS.369L..61H}). These results are cited in the literature as observational evidence for the existence of magnetic fields in Be stars (e.g.\\ Brown et al.\\ \\cite{2004MNRAS.352.1061B}). However, $\\beta$\\,Cep has been shown to be a binary star, where the H$\\alpha$ emission line and the magnetic field originate in two different components (e.g.\\ Schnerr et al.\\ \\cite{2006A&A...459L..21S}), and the magnetic field measured in $\\omega$\\,Ori is below a 3$\\sigma$ threshold. A longitudinal magnetic field at a level larger than 3$\\sigma$ has been diagnosed for the Be star 16\\,Peg (Hubrig et al.\\ \\cite{2006MNRAS.369L..61H}). This star has \\vsini{}=104~km/s and was classified as a Be star by Merrill \\& Burwell (\\cite{1943ApJ....98..153M}) due to the detection of double emission in H$\\alpha$. However, the emission was not confirmed by subsequent observations, and so the question of the presence of magnetic fields in classical Be stars remains open.   \\begin{table*} \\begin{center} \\caption{The sample of stars observed and the resulting measurements of magnetic fields.} \\label{tab:kd} \\begin{tabular}{rcccccccrc} \\hline \\hline \\multicolumn{1}{c}{HD}& Other & V & MJD & Sp.\\ Type & $v\\,\\sin\\,i$ & \\teff{} & \\logg {} & \\multicolumn{1}{c}{$\\langle$$B_z$$\\rangle$} & Circumstellar \\\\ & identifier & & & & [km/s] & [K] & & \\multicolumn{1}{c}{[G]} & Zeeman features \\\\ \\hline 53367  & V750 Mon       & 7.0 & 53503.496 & B0IVe     & 86  & 28000 & 3.9 & $-9\\pm$30           & --- \\\\ 56014  & EW CMa         & 4.7 & 53512.469 & B3IIIp+sh & 144 & 20000 & 3.8 & {\\bf $-$146$\\pm$32} & --- \\\\ 56139  & $\\omega$ CMa   & 4.0 & 53502.547 & B2.5Ve    & 109 & 22000 & 4.0 & $-31\\pm$24          & --- \\\\ 57219  & NW Pup         & 5.1 & 53630.820 & B2Vne     & 117 & 23000 & 4.0 & $+80\\pm$32          & --- \\\\ 58011  & NN CMa         & 7.2 & 53502.529 & B1Vnep    & 18  & 25000 & 3.9 & $+135\\pm$48         & H+K \\\\ 58050  & OT Gem         & 6.4 & 53629.914 & B2Ve      & 123 & 23000 & 4.0 & $-101\\pm$35         & --- \\\\ 79351  & a Car          & 3.4 & 53454.590 & B2.5V     & 35  & 22000 & 4.0 & $-7\\pm$59           & --- \\\\ 105435 & $\\delta$ Cen   & 2.6 & 53474.622 & B2IVne    & 207 & 23000 & 3.9 & $-30\\pm$30          & --- \\\\ 117357 & CD-61 3819     & 9.1 & 53508.491 & B0.5IIIne & 78  & 27000 & 3.9 & $+7\\pm$59           & H+K \\\\ 137432 & HR 5636        & 5.4 & 53531.660 & B4Ve      & 127 & 18000 & 4.1 & $-58\\pm$24          & --- \\\\ 148184 & $\\chi$ Oph     & 4.4 & 53531.722 & B1.5Vpe   & 139 & 24000 & 4.0 & {\\bf +83$\\pm$21}    & --- \\\\ 148184 & $\\chi$ Oph     & 4.4 & 53862.380 & B1.5Vpe   & 139 & 24000 & 4.0 & {\\bf +136$\\pm$16}    & --- \\\\ 148259 & OZ Nor         & 7.4 & 53571.601 & B2Ve      & 86  & 23000 & 4.0 & $+116\\pm$49         & --- \\\\ 153261 & V828 Ara       & 6.3 & 53531.750 & B2IVe     & 184 & 23000 & 3.9 & $-33\\pm$41          & --- \\\\ 155806 & V1075 Sco      & 5.6 & 53531.775 & O7.5 IIIe & 116 & 34000 & 3.8 & {\\bf $-$115$\\pm$37} & ---  \\\\ 181615 & $\\upsilon$ Sgr & 4.6 & 53519.910 & B2Vpe     & 69  & 23000 & 4.0 & {\\bf +38$\\pm$10}    & H+K \\\\ \\hline \\end{tabular} \\end{center} \\end{table*}  It is quite possible that the difficulty of detecting magnetic fields in Be stars is quite similar to that discussed for young Herbig Ae/Be stars with accretion disks by Hubrig et al.\\ (\\cite{2007A&A...463.1039H}). However, in contrast to Herbig stars showing evidence of magnetically mediated disk accretion, the Be disks originate from ejection or decretion of mass from the rapidly rotating B stars. The observed spectral lines may form over a relatively large volume, and the magnetic field topology is likely rather complex. The presence of a mixture of photospheric and CS magnetic fields could drive the net line-of-sight magnetic flux to near zero values. Also, even if there are very strong, small scale magnetic fields distributed over the surface of the star, these could go undetected in our measurements, as the measured mean longitudinal magnetic field $\\langle$$B_z$$\\rangle$ is the average over the stellar hemisphere visible at the time of observation of the component of the magnetic field parallel to the line of sight, weighted by the local emergent spectral line intensity.  Below, we report our results of the study of magnetic fields in a sample of 15~Be stars carried out with the multi-mode instrument FORS\\,1 installed at the 8\\,m Kueyen telescope at the VLT. ",
        "conclusions": "The detected magnetic fields in the studied Be stars are rather weak, with the largest longitudinal magnetic field of $-$146\\,G measured in the Be star HD\\,56014. These results suggest that strong large scale organized magnetic fields are not common among the group of Be type stars. Further, we noticed the possible presence of  distinct circular polarisation features in the CS components  of Ca\\,{\\sc ii} H\\&K in two Be stars, HD\\,58011 and HD\\,117357, as well as in HD\\,181615, indicating the possible presence of magnetic fields in the CS mass loss disks. Interestingly, similar type CS components in Ca\\,{\\sc ii} H\\&K lines have recently been discovered by us from FORS\\,1 observations of Herbig Ae/Be stars (Hubrig et al.\\ \\cite{2007A&A...463.1039H}). The  CS Ca\\,{\\sc ii} H\\&K line profiles in Stokes~I spectra of Herbig Ae/Be stars are frequently quite complex and consist of several components which  are usually assumed to be formed at the base of the stellar wind and in the accretion gaseous flow. A future careful high resolution high signal-to-noise spectropolarimetric study of the temporal behaviour of the Zeeman features in the Stokes~V spectra will allow to get the highly desirable insight into the nature of gaseous disks of these two Be stars.  During our observations the exposure times for each target were rather short, of the order of 15--20~min. Thus, not much can be inferred with respect to the variability and  evolution of their magnetic fields and the correlation of magnetic field properties with dynamical phenomena taking place in these stars. To constrain the structure of magnetic fields in Be stars and to probe the presence of localized transient magnetic fields suggested by several previous studies (e.g., Mathys \\& Smith \\cite{2000ASPC..214..316M} and references therein), the magnetic field  measurements should be carried out over short (minutes) and long (rotational periods) time scales. The existence of small scale, i.e.\\ highly non-dipole magnetic fields has been suggested by X-ray observations of flaring events (e.g. Smith, Robinson \\& Corbet \\cite{1998ApJ...503..877S})."
    },
    "0711/0711.0563_arXiv.txt": {
        "abstract": "To complement the IceCube neutrino telescope currently under construction at the South Pole, the three Mediterranean neutrino telescope projects ANTARES, NEMO and NESTOR have joined forces to develop, construct and operate a km$^3$-scale neutrino telescope in the Mediterranean Sea. Since February 2006, the technical specifications and performance of such a detector are studied in the framework of a 3-year EU-funded Design Study. In 2009 a technical design report will be released laying the foundations for the construction of the detector. In the following, the current status of the Design Study is presented and examples of solutions for the technical challenges are discussed. ",
        "introduction": "The detection of high energy neutrinos from astrophysical sources would be a major breakthrough in our understanding of origin and production mechanisms of cosmic rays. Unfortunately, despite intense search for these neutrinos during recent years no such neutrino has been identified up to now. Recent calculations \\cite{apj:656:870,pr:d74:063007} indicate that detectors with at least a km$^3$ of instrumented volume are required for this task where detectors of the first generation like AMANDA or ANTARES have typical volumes of 0.01\\,km$^3$.  At the South Pole the IceCube detector with an instrumented volume of 1\\,km$^3$ is currently being build as the successor to the AMANDA neutrino telescope. However, due to the large atmospheric muon background for upward observations with neutrino telescopes, the sensitivity of the detector to sources in the southern sky which includes most of the Galactic Plane and the Galactic Centre is greatly reduced. These regions harbour many potential high energy neutrino sources like supernova remnants, pulsar wind nebulae, microquasars and other binary systems, but also unidentified sites of high energy gamma-ray emissions. In order to be able to observe these sources, a km$^3$-scale neutrino telescope in the Northern Hemisphere is required.  Building on the experience gained in the pilot projects ANTARES, NEMO and NESTOR, the three collaborations have joined forces to develop, construct and operate such a km$^3$-scale neutrino telescope, KM3NeT, in the Mediterranean Sea at the beginning of the next decade. At the time of this article the consortium consists of 37 institutes from 10 European countries (Cyprus, France, Germany, Greece, Ireland, Italy, Netherlands, Romania, Spain, UK). Further, also non-European institutions are welcome to join the collaboration. KM3NeT is envisioned as a multidisciplinary research infrastructure with a permanent deep-sea access for marine sciences (such as oceanology, marine biology, environmental sciences, geology and geophysics) through an \\emph{associated science node}. ",
        "conclusions": "The next generation neutrino telescope in the Mediterranean Sea, KM3NeT, is well on its way. With the expertise from the construction of the first generation of deep sea neutrino telescopes an EU funded Design Study is currently conducted that will provide a TDR by 2009. This TDR will contain the technical specifications for the future KM3NeT infrastructure consisting of the neutrino telescope and an associated science node. Subject to the convergence of the accompanying political process, construction is planned to begin shortly afterwards and to be finished within 3 years.  The first half of the Design Study has been dedicated to the exploration of different detector layouts, their physics sensitivity and technical concepts for their realisation. In preparation of the TDR a conceptual design report (CDR) will be compiled at the end of this year containing the results from this first phase. Based on this information the final detector concept will be elaborated in detail in the second phase of the Design Study and published in the TDR."
    },
    "0711/0711.4389_arXiv.txt": {
        "abstract": "Evidence is accumulating suggesting that collisionless shocks in supernova remnants (SNRs) can amplify the interstellar magnetic field to hundreds of microgauss or even milli-gauss levels, as recently claimed for \\SNRJ. If these fields exist, they are almost certainly created by magnetic field amplification (MFA) associated with the efficient production of cosmic rays by diffusive shock acceleration (DSA) and their existence strengthens the case for SNRs being the primary source of galactic cosmic ray ions to the `knee' and beyond.  However, the high magnetic field values in SNRs are obtained exclusively from the interpretation of observations of radiation from relativistic electrons and if MFA via \\NL\\ DSA produces these fields the magnetic field that determines the maximum ion energy will be substantially less than the field that determines the maximum electron energy. We use results of a steady-state Monte Carlo simulation to show how nonlinear effects from efficient cosmic ray production and MFA reduce the maximum energy of protons relative to what would be expected from test-particle acceleration. ",
        "introduction": "Magnetic fields with strengths well above what can be expected from shock compression of the interstellar medium (ISM) field, $\\BISM$, are inferred in supernova remnants (SNRs) from four types of observations: (1) sharp X-ray edges or filaments \\citep[e.g.,][]{BambaEtal2003,WarrenEtal2005,CassamEtal2007}, (2) broadband fits of emission from radio to \\gamray\\ energies \\citep[e.g.,][]{Cowsik80,VL2003,VBK2005a}, (3) spectral curvature in radio emission \\citep[][]{RE92} and, (4) with the recently reported {\\it Chandra} observations of bright filaments in \\SNRJ\\ \\citep{Uchiyama07}, rapid time variations of nonthermal X-ray emission. The most likely explanation for these fields is that $\\BISM$ is amplified as part of the efficient production of cosmic rays by nonlinear diffusive shock acceleration (DSA) \\citep[e.g.,][]{BL2001,AB2006,VEB2006}.  Since the magnetic field largely determines the maximum particle energy a shock can produce, \\MFA\\ (MFA) is critical for interpreting GeV-TeV emission and determining if SNRs produce cosmic rays to the `knee' and above. Equally important is the fact that if strong MFA is an intrinsic part of nonlinear DSA, unexpectedly strong magnetic fields may be present in other systems undergoing efficient shock acceleration such as radio jets and gamma-ray burst (GRB) afterglows. The detailed modeling of broadband observations of SNRs is the only practical way to determine the importance of \\NL\\ effects and MFA in high Mach number collisionless shocks.  Ironically, the evidence for large magnetic fields and, therefore, nonlinear MFA is obtained exclusively from radiation emitted by \\rel\\ electrons, while the nonlinear processes responsible for MFA are driven by the efficient acceleration of \\rel\\ ions, mainly protons. Relativistic protons can only be directly inferred from \\pion\\ emission at GeV-TeV energies and the unambiguous interpretation of these observations is difficult and remains uncertain. In any case, even if \\pion\\ emission is clearly present, these observations must be combined with \\syn\\ emission from electrons to estimate the magnetic field strength.  A recent report by \\citet{Uchiyama07} of rapid time variations ($\\sim 1$\\,yr) in the nonthermal X-ray emission from bright filaments in \\SNRJ\\ provides new clues to the electron acceleration process in SNRs. As emphasized by \\citet{Uchiyama07}, the rapid time variations suggest two conclusions. First, if the variations are determined by the \\syn\\ radiation losses of TeV electrons, the magnetic field can indeed be amplified to milligauss-levels at a SNR shock and, second, the MFA process is tightly bound with the production of \\rel\\ particles (i.e., CRs). An important consequence of these conclusions is that the fraction of the shock ram kinetic energy that is put into accelerated {\\it ions} must be large, e.g., tens of percent.  Therefore, DSA must be in an efficient, nonlinear regime where the feedback of accelerated particles on the shock structure is significant. A simple estimate confirming this consequence is equation~(15) in \\citet{BL2001}.  Here we address a single question: Can, as asserted by Uchiyama et al. (2007), the large amplified fields inferred for {\\it electrons} from radiation losses in a \\NL\\ shock also determine the maximum {\\it proton} energy produced in the SNR shock?  We find the answer to be no because the inevitable \\NL\\ shock modification (due to efficient DSA) and the magnetic field variation in the shock precursor (due to MFA) make the maximum proton energy smaller than what is expected without accounting for these effects.  Using results from \\citet{VEB2006} for DSA limited by the finite size of the shock, we estimate that the maximum proton energy is at least an order of magnitude less than that predicted assuming an \\Unmod\\ shock with a large, nearly constant field inferred from electron radiation losses.  Our result is similar to that found by \\citet{BAC2007} in a time-dependent calculation of DSA where the acceleration is limited by the age of the shock rather than the size, an indication that the \\NL\\ effects we discuss are robust. ",
        "conclusions": "the precursor field must still range from $B_0$ to $B_2$ and the proton maximum energy will be determined by a smaller effective field than that estimated from \\syn\\ losses for electrons.  As a final comment we emphasize a point also made by \\citet{BAC2007}. If MFA is occurring and the system is highly NL, it may not be possible to explain temporal variations in nonthermal X-ray emission simply as a radiation loss time. There cannot be variations in X-ray emission on short time scales unless the accelerator changes in some fashion on these time scales, otherwise the radiation would be steady, or varying on the shock dynamic timescale, regardless of how short the radiation loss time was. Since the injection and acceleration of protons and electrons is nonlinearly connected to the amplified magnetic field, changes in the electron particle distribution and changes in the field producing the \\syn\\ emission, will go together and it may be difficult to unambiguously determine the field strength from temporal variations."
    },
    "0711/0711.1109_arXiv.txt": {
        "abstract": "We present near-infrared (1.15--2.50 microns) medium-resolution (R = 1700) spectroscopy of a sample of 23 brown dwarf candidates in the young Upper Sco association. We confirm membership of 21 brown dwarfs based on their spectral shape, comparison with field dwarfs, and presence of weak gravity-sensitive features. Their spectral types range from M8 to L2 with an uncertainty of a subclass, suggesting effective temperatures between 2700 and 1800 K with an uncertainty up to 300 K and masses in the 30--8 M$_{\\rm Jup}$ range. Among the non-members, we have uncovered a field L2 dwarf at a distance of 120--140 pc, assuming that it is single. The success rate of our photometric selection based on five photometric passbands and complemented partly by proper motion is over 90\\%, a very promising result for future studies of the low-mass star and brown dwarf populations in young open clusters by the UKIDSS Galactic Cluster Survey. We observe a large dispersion in the magnitude versus spectral-type relation which is likely the result of the combination of several effects including age dispersion, extent and depth of the association, a high degree of multiplicity and the occurrence of disks. ",
        "introduction": "Since the discovery of the first genuine brown dwarfs \\citep{rebolo95,nakajima95}, several hundreds of brown dwarfs have been found isolated in the field \\citep[see review by][]{kirkpatrick05}, and as companions \\citep[see review by][]{burgasser07a}. The number of substellar objects unearthed in clusters has increased considerably due to the advent of larger detectors and more efficient instrumentation, allowing for an efficient photometric selection of candidates and then efficient spectroscopic confirmation of the presence of low-gravity features for young (age $\\leq$ 10 Myr) brown dwarfs in regions like Orion \\citep{hillenbrand97,muench02,slesnick04}, $\\sigma$\\,Orionis \\citep{zapatero00,bejar01,caballero07d}, IC\\,348 \\citep{luhman03b,muench03}, Taurus \\citep{briceno02,luhman04a,guieu06}, $\\rho$\\,Ophiuchus \\citep{luhman99b,wilking05}, Chameleon \\citep{comeron04,luhman04b}, and Upper Sco \\citep{slesnick06,lodieu06,lodieu07a}.  The classification of old field L-type dwarfs, a class of objects cooler than M dwarfs with effective temperatures (T$_{\\rm eff}$) between $\\sim$2200--1500 K \\citep{basri00,leggett00} in the field, is now well established at optical wavelengths \\citep{kirkpatrick99,martin99a}. An infrared classification, based on the strength of absorption bands and atomic lines observed at low resolution (R $\\sim$ 50--600), was later set up to match the optical spectral types \\citep{reid01a,testi01,geballe02}. In addition, \\citet[][hereafter M03]{mclean03} produced a homogeneous set of intermediate resolution (R $\\sim$ 2000) spectra obtained with the Near-Infrared Spectrograph (NIRSPEC) on Keck II \\citep{mclean98} and developed a set of spectral indices to derive a classification based on this higher resolution data. Likewise, \\citet{cushing05} generated a set of spectra at similar resolution but over a larger wavelength range (0.6--4.1 $\\mu$m).  However, the current spectral classification does not take into account two important parameters: gravity and metallicity \\citep{kirkpatrick05}. Indeed, the number of spectroscopic L dwarfs reported to date in young clusters and star forming regions \\citep{metchev06,allers07,mohanty07} is insufficient to develop a spectral classification for young brown dwarfs. The tools currently available comprise several methods. Firstly, the spectral shape and the changes observed in the strength of gravity-sensitive features is often employed to assess the youth of brown dwarfs when compared to known field dwarfs \\citep{lucas01a,gorlova03,slesnick04,mcgovern04,wilking04,meeus05,allers07}. Secondly, \\citet{luhman99a} proposed a T$_{\\rm eff}$ vs spectral type relation for young dwarfs, that relation later being revised by \\citet{luhman03b} and extended by \\citet{allers07}. Finally, the direct comparison of observed spectra with theoretical models \\citep{allard00} allows an estimate of T$_{\\rm eff}$, gravity, and masses \\citep{lucas01a,meeus05}.  The Upper Sco (hereafter USco) OB association is part of the Scorpius Centaurus complex \\citep{deGeus89}, located at a distance of 145$\\pm$2 pc from the Sun \\citep{deBruijne97,deZeeuw99}. The age of USco is 5 Myr with little scatter \\citep{preibisch99}. The region is relatively free of extinction (Av $\\leq$ 2 mag) and star formation has already ended \\citep{walter94}. The association has been targeted at multiple wavelengths over the past decade starting with X-rays \\citep{walter94,preibisch98,kunkel99,preibisch99} but also with Hipparcos \\citep{deBruijne97} as well as in the optical \\citep{preibisch01,preibisch02,ardila00,martin04,slesnick06} and in the near-infrared \\citep{lodieu06,lodieu07a}. Several tens of low-mass stars and brown dwarfs with spectral types later than M6 have been confirmed as genuine spectroscopic substellar members in USco \\citep{ardila00,martin04,slesnick06,lodieu06}.  Recently, we have found a large number of new BD candidates with spectral types later than M6 and estimated masses below 0.05 M$_{\\odot}$  from several square degrees imaged by the UKIRT Infrared Deep Sky Survey \\citep[UKIDSS;][]{lawrence07} Galactic Cluster Survey (GCS). Additionally, we have extended the mass function down to 0.01 M$_{\\odot}$ with proper motion confirmation for sources more massive than 0.02 M$_{\\odot}$ \\citep{lodieu06,lodieu07a}. In this paper we present the near-infrared cross-dispersed spectra of the 23 faintest BDs in Upper Sco extracted from the UKIDSS GCS as well as two known late-M dwarfs used as spectral templates \\citep{martin04,slesnick06}. In Section \\ref{USco:obs_dr} we detail the observations and the data reduction of the cross-dispersed spectra obtained with the Gemini Near-InfraRed Spectrograph \\citep[GNIRS;][]{elias06a}. In Section \\ref{USco:memb} we assess the membership of all sources using the spectral shape, gravity-sensitive features, and direct comparison with near-infrared spectra of field dwarfs at similar resolution. In Section \\ref{USco:discussion} we discuss the implications of our spectroscopic follow-up and the current uncertainties on the determination of the effective temperatures and masses of these young brown dwarfs.   \\begin{table*} \\caption{Photometric candidates in Upper Sco with near-infrared spectra. We list the IAU name, coordinates (J2000), WFCAM $JHK$ photometry, the number of exposures and the on-source integration times in seconds (numbers in brackets represent the number of frames effectively used during the data reduction process), the spectral type (SpT) and the membership (Y$\\equiv$Member; N$\\equiv$Non-member). Typical uncertainties on the photometry is better than 0.05 mag in $J$, 0.02 mag in $H$, and 0.01 mag in $K$ \\citep[full photometry is available in][]{lodieu07a}. Proper motions are accurate to 10 mas/yr. Two known members, also observed with GNIRS, are added to this list: DENIS1611 \\citep[M9;][]{martin04} and SCH1625 \\citep[M8;][]{slesnick06}. We list the 2MASS photometry for these two sources because they lie outside the current GCS coverage. } \\label{tab_USco:list_cand} \\begin{tabular}{c c c c c c c c c c c} \\hline IAU Name                   &    R.A.\\    &      dec      &  $J$  &   $H$   &  $K$   & $\\mu_{\\alpha}\\cos\\delta$ & $\\mu_{\\delta}$ & ExpT & SpT & Memb \\cr \\hline USco J154722.82$-$213914.3 & 15:47:22.82 & $-$21:39:14.3 & 15.64 & 14.83 & 14.18 &  $-$13 &  $-$31 &  8 (8)$\\times$150 & L0 & Y \\cr  % USco J155419.99$-$213543.1 & 15:54:19.99 & $-$21:35:43.1 & 14.93 & 14.28 & 13.71 &  $-$14 &  $-$18 &  8 (4)$\\times$150 & M8 & Y \\cr  % USco J160603.75$-$221930.0 & 16:06:03.75 & $-$22:19:30.0 & 15.85 & 15.10 & 14.44 &    --- &    --- &  4 (4)$\\times$300 & L2 & Y \\cr  % USco J160606.29$-$233513.3 & 16:06:06.29 & $-$23:35:13.3 & 16.20 & 15.54 & 14.97 &    --- &    --- &  4 (4)$\\times$300 & L0 & Y \\cr  % USco J160648.18$-$223040.1 & 16:06:48.18 & $-$22:30:40.1 & 14.93 & 14.35 & 13.84 &  $-$23 &  $-$15 &  8 (8)$\\times$150 & M8 & Y \\cr  % USco J160714.79$-$232101.2 & 16:07:14.79 & $-$23:21:01.2 & 16.56 & 15.83 & 15.07 &    --- &    --- &  4 (4)$\\times$300 & L0 & Y \\cr  % USco J160723.82$-$221102.0 & 16:07:23.82 & $-$22:11:02.0 & 15.20 & 14.56 & 14.01 &  $-$11 &  $-$31 &  8 (8)$\\times$150 & L1 & Y \\cr  % USco J160727.82$-$223904.0 & 16:07:27.82 & $-$22:39:04.0 & 16.81 & 16.09 & 15.39 &    --- &    --- &  6 (6)$\\times$300 & L1 & Y \\cr  % USco J160737.99$-$224247.0 & 16:07:37.99 & $-$22:42:47.0 & 16.76 & 16.00 & 15.33 &    --- &    --- &  8 (8)$\\times$300 & L0 & Y \\cr  % USco J160818.43$-$223225.0 & 16:08:18.43 & $-$22:32:25.0 & 16.01 & 15.44 & 14.70 &    --- &    --- &  8 (8)$\\times$300 & L0 & Y \\cr  % USco J160828.47$-$231510.4 & 16:08:28.47 & $-$23:15:10.4 & 15.45 & 14.78 & 14.16 &  $-$12 &  $-$13 &  8 (4)$\\times$150 & L1 & Y \\cr  % USco J160830.49$-$233511.0 & 16:08:30.49 & $-$23:35:11.0 & 14.88 & 14.29 & 13.76 &  $-$5  &  $-$12 &  8 (8)$\\times$150 & M9 & Y \\cr  % USco J160847.44$-$223547.9 & 16:08:47.44 & $-$22:35:47.9 & 15.69 & 15.09 & 14.53 &     0  &  $-$20 &  4 (4)$\\times$300 & M9 & Y \\cr  % USco J160843.44$-$224516.0 & 16:08:43.44 & $-$22:45:16.0 & 18.58 & 17.22 & 16.26 &    --- &    --- & 12 (8)$\\times$300 & L1 & Y \\cr  % USco J160918.69$-$222923.7 & 16:09:18.69 & $-$22:29:23.7 & 18.08 & 17.06 & 16.16 &    --- &    --- &  8 (8)$\\times$300 & L1 & Y \\cr  % USco J160956.34$-$222245.5 & 16:09:56.34 & $-$22:22:45.5 & 17.83 & 16.99 & 16.29 &    --- &    --- &  8 (8)$\\times$300 & dL2 & N \\cr  % USco J161047.13$-$223949.4 & 16:10:47.13 & $-$22:39:49.4 & 15.26 & 14.57 & 14.01 &  $-$15 &  $-$24 &  8 (8)$\\times$150 & M9 & Y \\cr  % DENIS 161103.60$-$242642.9 & 16:11:03.60 & $-$24:26:42.9 & 14.86 & 14.14 & 13.70 &    --- &    --- &  8 (8)$\\times$150 & M9 & Y \\cr  % USco J161227.64$-$215640.8 & 16:12:27.64 & $-$21:56:40.8 & 17.14 & 16.37 & 15.78 &    --- &    --- &  8 (4)$\\times$300 & L0 & ? \\cr  % USco J161228.95$-$215936.1 & 16:12:28.95 & $-$21:59:36.1 & 16.41 & 15.56 & 14.79 &    --- &    --- &  4 (4)$\\times$300 & L1 & Y \\cr  % USco J161302.32$-$212428.4 & 16:13:02.32 & $-$21:24:28.4 & 17.17 & 16.37 & 15.65 &    --- &    --- &  8 (8)$\\times$300 & L0 & Y \\cr  % USco J161421.44$-$233914.8 & 16:14:21.44 & $-$23:39:14.8 & 14.97 & 14.44 & 13.94 &     0  &  $-$18 &  8 (8)$\\times$150 & A0 & N \\cr  % USco J161441.68$-$235105.9 & 16:14:41.68 & $-$23:51:05.9 & 16.07 & 15.34 & 14.62 &    --- &    --- &  4 (4)$\\times$300 & L1 & Y \\cr  % SCH 162528.62$-$165850.6   & 16:25:28.62 & $-$16:58:50.6 & 13.68 & 13.01 & 12.63 &    --- &    --- &  8 (8)$\\times$150 & M8 & Y \\cr  % USco J163919.15$-$253409.9 & 16:39:19.15 & $-$25:34:09.9 & 17.20 & 16.39 & 15.61 &  $-$1  &  $-$17 &  8 (8)$\\times$300 & L1 & Y \\cr  % \\hline \\end{tabular} \\end{table*} ",
        "conclusions": "\\label{USco:conclusions}  We have presented the results of a near-infrared spectroscopic follow-up study aimed at assessing the membership of a sample of brown dwarf candidates in the young Upper Sco association. We have confirmed a total of 21 out of 23 candidates, implying that the luminosity function derived form the photometric survey remains unchanged. By the same token, the power law index $\\alpha$=0.6 derived from our photometric survey \\citep{lodieu07a} remains valid after spectroscopic confirmation of over 90\\% of the photometric candidates below 30 M$_{\\rm Jup}$. The spectral types of the newly confirmed brown dwarfs range from M8 to L2, yielding effective temperatures in the 2700--1800 K range with uncertainties up to 300 K, following the scale developed by \\citet{luhman03b} for young brown dwarfs. The interval in mass probed by our study spans 30--8 M$_{\\rm Jup}$ according to theoretical models and is consistent with the masses inferred from our photometric survey. The dispersion observed in the spectral type-magnitude relation is large, in particular for the latest spectral types. The reasons for this dispersion might be the consequence of several effects, including reddening (although unlikely for our sample), errors on the age (small contribution), uncertainty on the distance, binarity, and the presence of disks (affecting at least 20\\% of the sources).  To investigate the binary frequency in USco, several follow-up observations are required, including adaptive optics and high-resolution spectroscopy. Several studies pointed out the possibility of a large fraction of close and wide low-mass binaries in USco \\citep{kraus05,bouy06b} compared to other clusters although no brown dwarf companions were found (in particular in the mass range discussed here). If binarity is the main effect for the dispersion of spectroscopic members observed in colour-magnitude diagrams, those systems would have separations smaller than th limits of high-resolution imaging surveys ($\\sim$5 au for HST, 7--10 au for adaptive optics). The cause for the difference between clusters remains however unclear: does it depend on the environment as proposed by \\citet{kraus05} and \\citet{bouy06b} or is it a consequence of the disruption of binaries with age? The occurrence of disks around low-mass brown dwarfs and planetary-mass objects and the influence of disks on the association sequence could be addressed with a $L'$-band imaging survey \\citep{jayawardhana03b,liu03} or Spitzer follow-up \\citep{carpenter06}. Thus additional photometry at some of these wavelengths for our USco brown dwarfs could be made to indicate the presence of disks, estimate their effect on the $J-K$ colours, and provide some indication about the disk properties of our sample."
    },
    "0711/0711.4815_arXiv.txt": {
        "abstract": "Jet-induced supernovae (SNe) have been suggested to occur in gamma-ray bursts (GRBs) and highly-energetic SNe (hypernovae). I investigate hydrodynamical and nucleosynthetic properties of the jet-induced explosion of a population III $40\\Msun$ star with a two-dimensional special relativistic hydrodynamical code. The abundance distribution after the explosion and the angular dependence of the yield are obtained for the models with high and low energy deposition rates $\\Ed=120\\times10^{51}$~\\ergs\\ and $1.5\\times10^{51}$~\\ergs. The ejection of Fe-peak products and the fallback of unprocessed materials in the jet-induced SNe account for the abundance patterns of the extremely metal-poor (EMP) stars. It is also found that the peculiar abundance pattern of a Si-deficient metal-poor star HE~1424--0241 is reproduced by the angle-delimited yield for $\\theta=30^\\circ-35^\\circ$ of the model with $\\Ed=120\\times10^{51}$~\\ergs.  Furthermore, I compare the yield of the jet-induced explosion with that of the spherical explosion and confirm the ejection and fallback in the jet-induced explosion is almost equivalent to the ``mixing-fallback'' in spherical explosions. In contrast to the spherical models, however, the high-entropy environment is realized in the jet-induced explosion and thus [(Sc, Ti, V, Cr, Co, Zn)/Fe] are enhanced. The enhancements of [Sc/Fe] and [Ti/Fe] improve agreements with the abundance patterns of the EMP stars. ",
        "introduction": "\\label{sec:2Drelintroduction}  Gamma-ray bursts (GRBs) are phenomena emitting $\\gamma$-ray for short periods followed by a power-law decaying afterglow. The origin had been covered for a long while, but it has become clear that long GRBs are associated with supernovae (SNe). Three GRB-associated SNe have been observed so far: GRB~980425/SN~1998bw \\citep{gal98}, GRB~030329/SN~2003dh \\citep{sta03,hjo03}, and GRB~031203/SN~2003lw \\citep{mal04}. They are all energetic explosions of massive stars, called hypernovae (\\eg \\citealt{nom06,nom07} for a review).  Although the explosion mechanism of GRBs and GRB-associated SNe is still uncovered, the following photometric and spectroscopic observations indicate that they are aspherical explosions with jet(s). (1) The light curve of the GRB afterglow has shown a polychromatic break in the power-law decay. The break is called a ``jet break'' explained by a deceleration of a relativistic jet and relativistic beaming of light (\\eg \\citealt{fra01,pir05}). (2) The nebular spectra of SN~1998bw show narrower [\\ion{O}{1}] lines than [\\ion{Fe}{2}] lines \\citep{pat01} indicating that O locates in the inner and lower-velocity region than Fe. This is realized in an aspherical explosion but not in a spherical explosion \\citep{mae02,mae06a}.  The aspherical explosions are indirectly suggested from the abundance patterns of extremely metal-poor (EMP) stars with [Fe/H] $<-3$.\\footnote{Here [A/B] $=\\log_{10}(N_{\\rm A}/N_{\\rm B})-\\log_{10}(N_{\\rm A}/N_{\\rm B})_\\odot$, where the subscript $\\odot$ refers to the solar value and $N_{\\rm A}$ and $N_{\\rm B}$ are the abundances of elements A and B, respectively.} Such EMP stars are likely to show the nucleosynthesis yields of a single or a few core-collapse SN/SNe \\citep{aud95,bee05}. The abundance patterns are reproduced by the ``mixing-fallback'' models that assume the extensive mixing of the shocked material before the fallback in a spherical SN model \\citep{ume02,ume03,ume05,iwa05,tom07b}. In particular, the abundance patterns of EMP stars with high [Co/Fe] and [Zn/Fe] ($\\sim0.5$) are explained by energetic explosion models because the high ratios require explosive nucleosynthesis under high entropy (\\eg \\citealt{ume05,tom07b}). Since the fallback doesnot take place in energetic spherical or quasi-spherical explosions (\\eg \\citealt{woo95,iwa05}), the mixing and fallback are interpreted as a consequence of the aspherical explosion. Indeed, the abundance patterns of the EMP stars are reproduced by the jet-induced explosion model (\\citealt{tom07a}).  Since the jet-induced explosion contributes to many astronomical phenomena, it is important to make a quantitative prediction on the nucleosynthesis outcome of the jet-induced explosions. In order to follow the jet propagation, fallback, and nucleosynthesis on the site and compare the jet-induced explosion models with the EMP stars, GRBs, and SNe, it is required to calculate multi-dimensional hydrodynamics and nucleosynthesis with the gravity and the special relativity.\\footnote{In order to follow the core collapse, jet formation, and nucleosynthesis on the site, it is also required to involve the neutrino, the general relativity, and possibly the magnetic field (\\eg \\citealt{pru05,fro06,jan07}).} Although there are many studies on the GRB jet propagation of the stellar mantle using the special relativistic hydrodynamics (\\eg \\citealt{alo00,zha04,miz06}), they did not include the gravity or calculate nucleosynthesis.  To investigate the yields of the jet-induced SNe, it is crucial to include the gravity because the fallback plays an important role on the nucleosynthesis yields. The studies by the use of Newtonian calculations have concluded that the energy deposition rate ($\\Ed$) sensitively affects SN nucleosynthesis \\citep{mae03b,nag06}. This result has been confirmed for the special relativistic cases \\citep{tom07a}. In particular, \\cite{tom07a} have shown that the jet-induced explosions with various $\\Ed$ can explain both the variations of GRB-associated SNe and the EMP, carbon-enhanced EMP (CEMP, [C/Fe]~$>1$), and hyper metal-poor (HMP, [Fe/H]~$<-5$) stars.  The previous studies have proved mostly the angle-integrated yields and shown that the abundance patterns of the EMP stars are reproduced by the angle-integrated yield. However, the abundance distribution of the jet-induced explosion depends on the direction (\\eg \\citealt{mae03b}). Thus, the abundance patterns of the next-generation stars might depend on the direction. I calculate aspherical stellar explosions induced by highly relativistic jets and obtain hydrodynamical and nucleosynthetic structures of such jet-induced explosion models. In particular, I investigate the angular dependence of the yield to compare the yields with the abundance patterns of the metal-poor stars. Furthermore, I compare the jet-induced explosion with the spherical SN model applied the mixing-fallback model and connect properties of the jet-induced explosion to the mixing-fallback model.  In \\S~\\ref{sec:model}, the applied models are described. In \\S~\\ref{sec:hyd}, I present the hydrodynamical and nucleosynthetic structures of the jet-induced explosion model, investigate the angular dependence of the yields, and compared the jet-induced explosion model with the spherical SN model. In \\S~\\ref{sec:summary}, the conclusion and discussion are presented. In Appendixes, the hydrodynamics and nucleosynthesis code is described and tested. ",
        "conclusions": "\\label{sec:summary}  I perform two-dimensional hydrodynamical and nucleosynthesis calculations of the jet-induced explosions of a Pop III $40\\Msun$ star. Here I test three jet-induced explosion models A, B, and C as summarized in Table~\\ref{tab:2Drelmodel} and conclude as follows.  (1) {\\bf Fallback}: The dynamics and the abundance distributions depend sensitively on the energy deposition rate $\\Ed$. The explosion with lower $\\Ed$ leads to a larger amount of fallback, and consequently smaller $\\Mfe$ and higher [C/O], [C/Mg], and [C/Fe]. Such dependences of [C/Fe] and $\\Mfe$ on $\\Ed$ predict that higher [C/Fe] tends to be realized for lower [Fe/H]. This is consistent with the observations (\\eg \\citealt{bee05}). Note, however, the formation of star with low [C/Fe] and [Fe/H] is possible because [Fe/H] depends on the swept-up H mass, \\ie the interaction between the SN ejecta and ISM (\\eg \\citealt{cio88}).  (2) {\\bf Abundance distribution and angular dependence}: I present the aspherical abundance distributions and investigate the angular dependence of the yield. After the explosion, the elements ejected even along the equatorial plane locate in the all directions, while the elements ejected only along the jet axis expand laterally but tend to be distributed along the jet axis. The abundance distributions in the jet-induced SN ejecta could be examined by spatially-resolved observations of supernova remnants (\\eg \\citealt{hwa04,keo07}).  The angle-delimited yield could reproduce the extremely peculiar abundance pattern of HE~1424--0241. However, the angle-delimited yields of model A have a large scatter that may be inconsistent with the relatively small scatter in the abundance ratios of the EMP stars. This implies that the angular dependence of the yield in most SNe is diluted by the strong mixing of the SN ejecta. On the other hand, the angle-delimited yields of model B show high [C/Fe] that is the same feature as the angle-integrated yield.  The angle-delimited yield strongly depends on which mass elements are included into the integration. This would be determined by the abundance mixing in the SN ejecta and by the region where the next-generation star takes in the metal-enriched gas. To investigate this issue further, it is required to calculate three-dimensional evolution of the supernova remnant in the ISM.  (3) {\\bf Comparison with the spherical explosion}: The angle-integrated yield of the jet-induced explosion is well reproduced by a spherical SN model applied the mixing-fallback model. This confirms that the mixing-fallback model approximates the jet-induced explosion well and that the mixing and fallback in hypernovae assumed in the mixing-fallback model are actually achieved in aspherical explosions. The abundance ratios between elements synthesized in different regions (\\eg C, O, Mg, and Fe) depend on the hydrodynamical structure of the explosion, \\eg the fallback. Thus, such macroscopic properties of the jet-induced explosion are represented by the mixing-fallback model. In particular, $\\Min$, $f_{\\rm Si}$, and $\\Mbh$ in the jet-induced explosion model are represented by $M_{\\rm cut}{\\rm (ini)}$, $f$, and $\\Mbh$ in the mixing-fallback model, respectively.  On the other hand, the ratios between the explosively-synthesized elements depend on the thermodynamical properties of the explosion. In particular, [Sc/Fe], [Ti/Fe], [V/Fe], [Cr/Fe], [Co/Fe] and [Zn/Fe] are enhanced by the high-entropy environment in the jet-induced explosion, thus showing differences from the mixing-fallback model. The enhancement of [Sc/Fe] and [Ti/Fe] improve the agreement with the observations. Note, the enhancement of [Ti/Fe] relative to the spherical SN model in this paper is larger than in \\cite{mae03b}. This might be because the relativistic jet suppresses the lateral expansion and enhances the energy concentration.  The enhancement of [Sc/Fe] is also suggested to be obtained by nucleosynthesis in the $p$-rich ejecta expelled from the innermost region \\citep{pru04,pru05,fro06,iwa06}. However, it seems difficult to enhance [Zn/Fe] ($\\sim0.5$) and [Sc/Fe] ($\\sim0$) only with nucleosynthesis in the $p$-rich ejecta (see \\S~5.2 in \\citealt{tom07b}). On the other hand, a jet-induced explosion enhances [Sc/Fe], [Ti/Fe], and [Zn/Fe] simultaneously and realizes the mixing and fallback in an energetic explosion. Therefore, I propose that the jet-induced explosion is a presumable origin of the enhancement of [Sc/Fe]. The origin would be concluded by future quantitative investigations including various elements.  Although the models with same $\\thj$ are shown in this paper, the asphericity is likely to be different for each SN. For example, the observations of supernova remnants show various aspherical structures (\\eg Cassiopeia A, \\citealt{hwa04,fes06}; and W49B, \\citealt{mic06,keo07}). Furthermore, recent observations of the nebular spectra of SNe suggested that all core-collapse SNe are aspherical explosions \\citep{mae08,mod08}. The degree of asphericity in SNe in the early universe will be revealed by detailed comparisons of abundance ratios between the aspherical SN models with different $\\thj$ and the metal-poor stars which represent the hydrodynamical and/or thermodynamical properties."
    },
    "0711/0711.3604_arXiv.txt": {
        "abstract": "% This review discusses recent results on the astrochemistry of (mostly high-mass) star-forming regions. After an introduction on the use of chemistry in astrophysics and some basic concepts of astrochemistry, specific results are presented.  Highlighted areas are the use of chemistry in the search for massive circumstellar disks, the interaction of molecular clouds with cosmic rays, and the feedback effects of protostellar irradiation on the parent molecular cloud.  The review concludes with a discussion of future observational opportunities. ",
        "introduction": "\\label{sec:use}  There are several ways in which knowledge of chemistry is helpful to gain a better understanding of the Universe. The prime area of interest is the formation of stars and planets, which occurs in cold dark clouds which require observation at long (infrared and radio) wavelengths. The line radiation of molecules is an essential part of this effort, because it is the only probe of the kinematics of these clouds, and because it is the major way to determine their temperatures, volume densities, and other conditions. Successful use of molecular lines to derive physical parameters requires some understanding of chemistry to predict which molecules may be abundant under which conditions. This use may be called `passive' astrochemistry.  A more active way to use chemistry in astrophysics makes use of the dependence of the molecular composition of the gas on parameters which are otherwise hard to estimate. This use involves the construction of chemical models typically containing thousands of reactions; \\citet{wakelam:errors} have studied the accuracy of such models.  First, the chemical composition of the gas in such models usually depends on the time, so that observations of molecular lines may be used to estimate the ages of star-forming regions \\citep{doty:hotcores}. Second, since the chemistry needs time to respond to changing conditions, molecular abundances often contain some memory of the source history. An example are the deuterium bearing molecules seen in hot molecular cores, which must be remnants of a previous cold phase. Third, the chemical composition of star-forming matter may give clues to the presence of (energetic) radiation which is difficult or impossible to observe directly. It is clear that chemistry is a significant help in the understanding of astrophysical processes.  More than 130 molecules are known to exist in interstellar space, of which 36 are known outside our Galaxy and 10 are known in the solid state \\citep{gibb:ice}. The eight new discoveries of 2007 (see www.cdms.de for the latest updates) may be grouped in a few `threads': (i) Complex organics (by which astronomers usually mean 4$^+$-atomic carbon chains), which probe gas-grain interactions and serve as a link to pre-biotic molecules; the most recent addition is C$_3$H$_6$ \\citep{marcelino:propylene}. (ii) Fluorine and phosphorus compounds (PO, HCP, CF$^+$), which help to determine elemental abundances and the composition of dust grains in dense clouds (\\citealt{neufeld:cf+}; \\citealt{agundez:hcp}; \\citealt{tenenbaum:po}). (iii) Negative ions, which are useful to measure the electron fraction of molecular gas (\\S~\\ref{sec:zeta}). (iv) Deuterated molecules, which are important as tracers of very early (cold and dense) phases of star formation (\\S~\\ref{sec:filters}).  Due to space limitations, this review is not comprehensive, but biased toward the interests of its author. For a recent general review of astrochemistry see the books edited by \\citet{lis:asilomar} and \\citet{combes:paris}. The physics and chemistry of regions of low-mass star formation are reviewed by \\citet{bergin:review} and \\citet{ceccarelli:ppv}. An overview of methods to derive molecular abundances from spectral line data can be found in \\cite{vdtak:alma}. The formation of high-mass stars is reviewed by \\citet{zinnecker:araa} and \\citet{beuther:ppv}, and is of course the topic of the current volume. ",
        "conclusions": ""
    },
    "0711/0711.2572_arXiv.txt": {
        "abstract": "{ We have identified three possible ways in which future \\xmm\\, observations can provide significant constraints on the equation of state of neutron stars. First, using a long observation of the neutron star X-ray transient Cen~X--4 in quiescence one can use the RGS spectrum to constrain the interstellar extinction to the source. This removes this parameter from the X-ray spectral fitting of the pn and MOS spectra and allows us to investigate whether the variability observed in the quiescent X--ray spectrum of this source is due to variations in the soft thermal spectral component or variations in the power law spectral component coupled with variations in ${\\rm N_H}$. This will test whether the soft thermal spectral component can indeed be due to the hot thermal glow of the neutron star. Potentially such an observation could also reveal redshifted spectral lines from the neutron star surface. Second, \\xmm\\, observations of radius expansion type~I X--ray bursts might reveal redshifted absorption lines from the surface of the neutron star. Third, \\xmm\\, observations of eclipsing quiescent low--mass X--ray binaries provide the eclipse duration. With this the system inclination can be determined accurately. The inclination determined from the X--ray eclipse duration in quiescence, the rotational velocity of the companion star and the semi--amplitude of the radial velocity curve determined through optical spectroscopy, yield the neutron star mass. } ",
        "introduction": "Low--mass X--ray binaries are binary systems in which a $\\approxlt 1 M_{\\odot}$ star transfers matter to a neutron star or a black hole. A large fraction of the low--mass X--ray binaries is found to be transient -- the so called soft X--ray transients (here, we will just refer to them as X--ray transients, see e.g.~Chen, Shrader, \\& Livio 1997). Before the arrival of the \\xmm\\, and {\\it Chandra} satellites only a few neutron star X--ray transients could be studied in quiescence (e.g.~the system Cen~X--4 and Aql~X--1; Van Paradijs et al.~1987; McClintock, Horne, \\& Remillard 1995; Campana et al.~1997). Using the \\xmm\\, and {\\it Chandra} satellites many more systems have been studied in quiescence (see e.g.~Rutledge et al.~2001, Wijnands et al.~2001, 2002, Campana et al.~2003, Jonker et al.~2003, Tomsick et al.~2004, Heinke et al.~2007 to name but a few references) and Cen~X--4 and Aql~X--1 were studied in much more detail than was possible before (e.g.~Campana et al.~2004, Rutledge et al.~2002). As we will explain below, these observations can have a profound impact on an important area of astrophysics: determining the neutron star equation of state (EoS). This is one of the ultimate goals of the study of neutron stars.  Recent theoretical progress provides the framework for the interpretation of X--ray observations of quiescent low--mass X--ray binaries (e.g.~Brown, Bildsten, \\& Rutledge 1998, Colpi et al.~2001, Zavlin, Pavlov, \\& Shibanov~1996, Gaensicke, Braje, \\& Romani 2002, Heinke et al.~2006). Theoretically, one expects the neutron star to emit X--rays even after accretion has stopped. This emission can be modelled by a neutron star atmosphere (NSA) model. The NSA models have four free parameters. The neutron star distance, mass, radius and temperature. Due to the large heat capacity the temperature of the neutron star core is set on timescales of tens of thousands of years (Colpi et al.~2001) and depends on the equilibrium between the heating and the cooling rate of the neutron star. The heating rate depends on the total amount of accreted baryons and the pycnonuclear reactions taking place a few hundred meters deep in the crust (Haensel \\& Zdunik, 1990; Brown, Bildsten, \\& Rutledge 1998; Colpi et al.~2001). The time--averaged mass accretion rate in neutron star transients can be derived from binary evolution models if the orbital period is known (Kraft, Mathews, Greenstein 1962). The pycnonuclear reactions taking place in the neutron star crust are described in Salpeter \\& Van Horn (1969) and  Kitamura (2000). The balance between the heating and cooling rates sets the neutron star core temperature (for a review of the cooling properties see e.g.~Yakovlev \\& Pethick 2004). In theory, a NSA--fit provides means to measure the mass and radius of the neutron star and thus constrain the equation of state (EoS) of matter at supranuclear densities.  In practice, spectra of neutron star transients in quiescence where the flux is high enough to allow for a spectral study, can indeed be well--fit by a neutron star atmosphere model (NSA). Sometimes, additional emission is present at energies above a few keV, often  quantified by a power--law. When systems are selected for which the distance is well known (e.g.~sources in globular clusters), neutron star masses and radii can be determined accurately.  Canonical neutron star values were found (e.g.~Heinke et al.~2003, Webb \\& Barret 2007), rendering support for the interpretation that the soft X-ray spectral component is due to the NSA. Besides fitting for the neutron star radius and mass in the NSA modelling, there is another way to derive a constraint on the neutron star mass. Namely, by investigating the neutron star temperature. The stringent upper limit on the NSA spectral component to the luminosity of SAX J1808.4--3658 ($\\approxlt$10\\%; Campana et al. 2002, Heinke et al. 2007) and the stringent limit on the luminosity and thus neutron star temperature in 1H~1905+000 (Jonker et al.~2006, 2007) hint at massive neutron stars in these two transients. The reasoning behind this is the following: the upper limit on the thermal spectral component implies low core temperatures, which implies a rapid release  via enhanced neutrino emission of the energy produced in the crust. This enhanced neutrino emission can only occur if the neutron star mass is larger than 1.6--1.7 $M_\\odot$ (Yakovlev \\& Pethick ~2004). Neutron stars with masses well above $1.4\\,M_\\odot$ cannot exist for so--called soft equations of state (EoS), in which matter at high densities is relatively compressible (e.g., due to a meson condensate or a transition between the hadron and quark--gluon phases). An important caveat is that the heating accretion history of the last several tens of thousands of years has to be determined from binary evolution models.  In this Manuscript we identify three possible ways in which future \\xmm\\, observations can provide constraints on the neutron star EoS. ",
        "conclusions": "We have identified three possible ways in which \\xmm\\, can either lead or have an important role in determining the neutron star EoS. First, a long observation of a neutron star low--mass X--ray binary in quiescence with a flux high enough to allow an RGS spectrum to be extracted is necessary to investigate if indeed the soft thermal spectral component is stationary. If so, and if the distance to the system is accurately known the neutron star mass and radius can be determined from a NSA model fit to the EPIC pn and MOS spectra. Currently, the best source for such a study is Cen~X--4. Secondly, one can potentially use radius expansion bursts to search for redshifted absorption lines from the neutron star surface. Again nearby systems are preferred to maximise the source flux. Finally, \\xmm\\, observations of eclipsing low--mass X--ray binaries yield the eclipse duration, which together with a measurement of the mass ratio yields an accurate determination of the inclination angle that enters the mass function. Using \\xmm\\, observations mass measurements of these eclipsing sources become much more accurate."
    },
    "0711/0711.2091_arXiv.txt": {
        "abstract": "The evolution of the first generations of stars at zero or extremly low metallicity, and especially some crucial properties like the primary \\nvi\\ production, is charactarized by convective-reactive mixing events that are mostly absent from similar evolution phases at solar-like metallicity. These episodes occur when unprocessed H-rich material is mixed accross a convective boundary into \\czw\\ rich He-burning material, as for example in He-shell flashes of extremely-low metallicity AGB stars. In this paper we describe the astrophysical context of such convective-reactive events, including the difficulty of current one-dimensional stellar evolution models to correctly simulate these evolutionary phases. We then describe the requirements and current state of modeling convective-reactive processes in the first stars environment. We demonstrate some of the new concepts that we are applying to this problem, i.e. the highly accurate PPB advection scheme in the framework of PPM hydrodynamic simulations of mixing accross a very stiff convective boundary. We show initial results of such simulations that address the first non-reactive step of this problem, which is the entrainment of H at the top boundary of the He-shell flash convection zone. ",
        "introduction": "The formation and evolution of the first generations of stars traces the formation and evolution of the first structure on cosmological scales (e.g.\\ O'Shea, this volume), and starts the continuing process of chemical evolution in galaxies. The nuclear production in these stars provides an important tracer of events that took place in the early Universe \\citep[e.g.][]{frebel:07}. An increasing number of very metal-poor stars ($\\mem{[Fe/H]} < -2$) are now discovered in large spectroscopic surveys \\citep{beers:05} reaching down to metalicities below $\\mem{[Fe/H]}=-5$ \\citep{aoki:06a}. These low-mass stars preserve the signature of the nuclear processes in the first generations of stars, of both low mass as well as from massive stars. Some important findings from the observational study of stars that carry the signature of the first generations of stars include the following.  The fraction of C-enhanced stars ($\\mem{[C/Fe]}>1$ ) is between about $10\\%$ \\citep{cohen:05} and $20\\%$ \\citep{lucatello:06} for stars with $\\mem{[Fe/H]} < -2$, and increases with decreasing metallicity. Below $\\mem{[Fe/H]} < -3$ there are now about $300$ C-rich stars, below $\\mem{[Fe/H]} < -3.5$ $40\\%$ are C-rich. $80\\%$ of these CEMP stars are enhanced with s-process elements \\citep{aoki:07,beers:07a}. Many of these CEMP-s stars show evidence for binarity and \\citet{lucatello:05} find that present data is statistically consistent with all CEMP-s stars beeing in binaries. The abundance signature of this large group of CEMP-s stars has been interpreted as beeing the result of mass transfer from an Asymptotic Giant Branch (AGB) star, that is now a white dwarf. An additional important finding is that CEMP-s stars have C to N ratios that are slightly lower than predicted by non hot-bottom burning, low-mass, third dredge-up AGB stars ($0 \\leq \\mem{[C/N]} \\leq 1.5$), but about 2 orders of magnitude larger than any hot-bottom burning AGB models \\citep{johnson:07}.   As far as heavy elements are concerned one of the most remarkable findings is the discovery of the solar-scaled, apparently universal r-process element distribution for for the heavy neutron capture elements, $56\\leq \\mem{Z} < 83$ \\citep{cowan:06}. However, for lower elemental mass down to Sr the observed element abundance distribution in several extremely metal poor stars can not be only accounted for by the standard solar-scaled r-process. Since at very low metallicity the average contribution from the weak and main s process is negligable, an extra contribution from a primary s process called light element primary process (LEPP) has been proposed \\citep{travaglio:04}. Recently, \\citet{montes:07} showed that the LEPP signature at very low metallicity could be reproduced by a process with neutron densities of less than $10^{13} \\mem{cm^{-3}}$, which are intermediate between the typical s- (much lower) and r-process (much higher). Concerning the light elements, the observations of extremely metal poor stars of any variety point at a primary production of \\nvi\\ that is currently not yet well understood \\citep[e.g.][]{chiappini:06}.  These observational findings of the first generations of stars, that have been discussed at length at this First Stars III conference, have one thing in common. They are sensitively related to the physics of mixing, and in some cases the hydrodynamic interaction of mixing and nuclear burning. The present 1-D spherically symmetric stellar evolution models of these low-metallicity stars do not account properly for this physics. While this is true also for solar-like metallicity, the uncertainties are amplified by the flash-like burning of H mixed into \\czw-rich He-burning layers. As we will describe below these convective-reactive events are peculiar to extremely low-metallicity stellar conditions. Uncertainty in mixing is not limited to convection-induced mixing of course. Rotation must play an important role. The uncertainty of present rotating stellar models is evident from their inability to properly meet constraints from pulsars \\citep{heger:00,hirschi:05} in the case of massive stars models, or the inhibiting effect of rotation in simulations of the s process in low-mass stars \\citep{herwig:02a}. However, some areas related to the comparatively fast and dynamic evolution of convective and convective-reactive events in the stellar interior are open to detailed hydrodynamic simulations. In the following sections we will discuss the discussed observational properties of the extremely metal poor stars in the framework of one-dimensional stellar evolution simulations with the goal to identify the most promising, high-gain situations in first stars that can be investigated by hydrodynamical simulations now. We will discuss the one-dimensional spherically symmetric assumption in the convective-reactive cases. We will then present some new simulations of one of these important convective boundaries, showcasing the PPB advection scheme at work inside the most recent version of the LCSE PPM stellar hydodynamics code. Finally we report high-resolution simulations of the top convective boundary of He-shell flash convection. Mixing at this boundary determines the properties if the H-ingestion flash in EMP AGB stars. ",
        "conclusions": ""
    },
    "0711/0711.5022_arXiv.txt": {
        "abstract": "A model of non-local dilaton interactions, motivated by string duality symmetries, is applied to a scenario of ``coupled quintessence\" in which the dilaton dark energy is non-locally coupled to the dark-matter sources. It is shown that the non-local effects tend to generate a backreaction which -- for strong enough coupling -- can automatically compensate the acceleration due to the negative pressure of the dilaton potential, thus asymptotically restoring the standard (dust-dominated) decelerated regime. This result is illustrated by analytical computations and numerical examples. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.2015.txt": {
        "abstract": "To date, two planetary systems have been discovered with close-in, terrestrial-mass planets ($\\la 5-10 \\mearth$).  Many more such discoveries are anticipated in the coming years with radial velocity and transit searches. Here we investigate the different mechanisms that could form ``hot Earths'' and their observable predictions.  Models include: 1) {\\it in situ} accretion; 2) formation at larger orbital distance followed by inward ``type 1'' migration; 3) formation from material being ``shepherded'' inward by a migrating gas giant planet; 4) formation from material being shepherded by moving secular resonances during dispersal of the protoplanetary disk; 5) tidal circularization of eccentric terrestrial planets with close-in perihelion distances; and 6) photo-evaporative mass loss of a close-in giant planet.  Models 1-4 have been validated in previous work.  We show that tidal circularization can form hot Earths, but only for relatively massive planets ($\\ga 5 \\mearth$) with very close-in perihelion distances ($\\la$ 0.025 AU), and even then the net inward movement in orbital distance is at most only 0.1-0.15 AU.  For planets of less than $\\sim 70 \\mearth$, photo-evaporation can remove the planet's envelope and leave behind the solid core on a Gyr timescale, but only for planets inside 0.025-0.05 AU.  Using two quantities that are observable by current and upcoming missions, we show that these models each produce unique signatures, and can be observationally distinguished.  These observables are the planetary system architecture (detectable with radial velocities, transits and transit-timing) and the bulk composition of transiting close-in terrestrial planets (measured by transits via the planet's radius). ",
        "introduction": "Both radial velocity (RV) and transit searches are biased toward finding large/massive planets at small orbital distances (e.g., Marcy \\& Butler 1998; Charbonneau \\etal 2007).  Given the increased sensitivity of new instruments, ever-smaller close-in planets are being detected.  Currently, two systems are thought to contain close-in planets of less than 10 Earth masses ($\\mearth$): GJ 876 (Rivera \\etal 2005) and GJ 581 (Udry \\etal 2007).  Transit missions {\\it CoRoT} (Baglin 2003; Aigrain \\etal 2007) and {\\it Kepler} (Basri \\etal 2005) expect to find perhaps a few hundred close-in planets with masses less than 5-10 $\\mearth$.  In this paper we focus on these ``hot Earth'' planets, which we assume to have masses in the range $0.1 < m_p < 10 M_\\oplus$, and semi-major axes $a \\la 0.2$ AU.  We propose that it is possible to determine the formation history of a given hot Earth planetary system with two observable quantities: the architecture of the inner planetary system, and the bulk composition of the hot Earth(s). The planetary system architecture can be detected by a combination of RV and transit measurements, as well as additional analysis of transit signals (e.g., timing variations: TTV; Agol \\etal 2005, Holman \\& Murray 2005).  The composition of a transiting terrestrial planet can be determined by its physical size, i.e. the transit depth.  Structure models indicate that very water-rich planets ($\\ga$10\\% water by mass) have detectably larger radii than dry, rocky planets or iron-dominated planets (Valencia \\etal 2007a, 2007b; Fortney \\etal 2007; Sotin \\etal 2007; Seager \\etal 2007), although a massive H/He envelope can also inflate the observed planetary radius (Adams \\etal 2007).  Several mechanisms for the formation of close-in terrestrial planets have been proposed (Zhou \\etal 2005; Gaidos \\etal 2007).  In Section 2 we describe the observable quantities that can distinguish between models.  In Section 3, we summarize four known models, and test two unproven models: a) tidal circularization of terrestrial planets on eccentric orbits, and b) photo-evaporation of hot Neptunes or hot Jupiters.  We have tried to include all reasonable models, which include various combinations of accretionary growth, planet migration, and evaporative loss.  Table 1 summarizes the observable differences between these models.  In section 4, we apply these models to the two known hot Earth systems.  Section 5 concludes the paper with a discussion of whether the mechanism for giant planet formation -- core-accretion or gravitational instability -- can affect the abundance of hot Earths, as claimed by Zhou \\etal (2005).  \\section {Observable Quantities}  The observables considered in this paper are the architecture of the inner planetary system and the bulk planetary composition.  The planetary system architecture, i.e. the co-existence (or lack) of additional planets in hot Earth systems, can provide strong circumstantial evidence for or against certain formation models, as described below.  In particular, certain characteristic planetary configurations are smoking guns (see Table 1). Determining the bulk composition of a planet requires transit measurements. Thus, our analysis applies only to systems with at least one transiting planet.  In most cases, but not all, the transiting planet must be a hot Earth.  The architecture of hot Earth planetary systems may be determined in three primary ways: ($i$) via the detection of transits of multiple planets; ($ii$) via radial velocity (RV) monitoring of the host star; and ($iii$) by analysis of transit timing variations (TTV).  Other techniques such as astrometry may be used in conjunction with these techniques, but note that astrometry is not optimal for detecting close-in planets (Black \\& Scargle 1982).  Detection of multiple transiting planets in the same system requires extremely low mutual inclinations between planetary orbits, which are thought to be rare (e.g., Levison \\etal 1998).  The RV technique has discovered several planets with minimum masses less than Neptune, including the two known systems with hot Earths (Rivera \\etal 2005; Udry \\etal 2007).  There exist several currently-operational instruments capable of RV followup for {\\it CoRoT} and {\\it Kepler} targets, such as Keck HIRES (Vogt \\etal 1994), the Hobby-Eberly Telescope's HRS spectrograph (Cochran \\etal 2004), and the HARPS instrument at ESO (Mayor \\etal 2003).  In addition, the HARPS-North spectrograph is being built specifically to do RV followup of {\\it Kepler} candidate transiting planets (Latham 2007).  However, given that many of the target stars will be very faint, RV followup of a large number of stars may not be possible.  For those that can be followed up, the $\\la m \\, s^{-1}$ sensitivity of current RV instruments should be able to detect close-in, $\\la 5-10 \\mearth$ planets and to probe the inner regions of {\\it CoRoT} and {\\it Kepler}-detected targets.  %\\begin{deluxetable}{p{4.5cm}|p{5cm}|p{5cm}} %\\tablewidth{0pt} %\\tablecaption{Observable Predictions of Hot Earth Formation Models} %\\tabletypesize{\\normalsize} %\\tablecolumns{5} %\\renewcommand{\\arraystretch}{.6} %\\tablehead{ %\\colhead{Model} & %\\colhead{System Architecture} & %\\colhead{Planet Composition}} %\\startdata  \\begin{table*} \\centering \\begin{minipage}{150mm} \\caption{Observable Predictions of Hot Earth Formation Models} \\begin{tabular}{p{5cm} || p{5cm} | p{5cm}}   %{@{}llrrrrlrlr@{}} \\hline \\bf \\large Model & \\bf \\large System Architecture & \\bf \\large Planet Composition\\\\  \\\\ \\hline  \\bf {\\it \\bf In Situ} Accretion & Several hot Earths, spaced by $\\sim$20-60 mutual Hill radii & Relatively dry for Solar-type stars.  Up to 0.1-1 percent water for low-mass stars \\\\ \\hline  \\bf Type 1 Migration & Chain of many terrestrial planets, close to mutual mean motion resonances & Icy or Rocky, depending on formation zone.  Most likely to be icy ($\\ga$ 10 \\% water by mass) . \\\\ \\hline  \\bf Giant Planet Migration Shepherding & Co-existence of hot Earths and close-in giant planets near (but not in) strong mean motion resonances & Rocky with moderate water content: a few percent water by mass at time of formation. \\\\ \\hline  \\bf Secular Resonance Shepherding during Disk Dissipation & Co-existence of hot Earths and at least two, interacting giant planets & Depends on the details of the giant planet's orbital history.  Rocky if formed mainly {\\it in situ} \\\\   \\hline  \\bf Tidal Circularization of Eccentric Planets & Single hot Earth, with possible distant companion (giant planet or stellar binary) to explain high eccentricity & Depends on formation zone of planet -- rocky unless migrated inward. \\\\ \\hline  \\bf Photo-evaporation of hot Neptunes & Hot Earth inside 0.025-0.05 AU.  Likely chain of several planets, as for type 1 migration.  Correlation between hot Earth vs. hot Neptune frequency and stellar age.  & Icy, assuming what remains is a giant planet core.\\\\ \\hline  \\end{tabular} \\end{minipage} \\end{table*}  %\\enddata %\\end{deluxetable}  Transit timing variations (TTV) analysis measures the deviation of a series of transits from a perfect chronometer, representing a deviation of the transiting planet's orbit from a perfect Keplerian ellipse due to perturbations from one or more additional planets (Holman \\& Murray 2005; Agol \\etal 2005).  The TTV signal scales with the transiting planet's orbital period, and increases for more massive and closer perturbing planets.  For sufficiently accurate transit timing data, TTV analysis can either derive the mass and orbit of a perturbing planet or place constraints on the existence of nearby perturbers (Steffen \\& Agol 2005; Agol \\& Steffen 2007).  TTV is especially sensitive to planets that lie in or close to mean motion resonances with the transiting planet, which is convenient given that several formation models predict near-resonant planetary configurations (see $\\S$ 3 below).  The bulk composition of a planet determines its density and therefore its physical size: ice-planets are far larger than iron-planets.  Recently, several studies have calculated mass-radius relations for planets with different compositions (Valencia \\etal 2006, 2007a, 2007b; Fortney \\etal 2007; Sotin \\etal 2007; Seager \\etal 2007).  For a fixed mass, there exists a roughly 40\\% difference in radius between pure ice planets and pure rock planets, and a similar 40\\% difference between pure rock and pure iron planets; these ratios of sizes are independent of planet mass.  In addition, there is a $\\sim$ 35\\% difference in size between Earth-like planets (2/3 rock, 1/3 iron) and ocean planets (1/2 rock, 1/2 water; Fortney \\etal 2007).\\footnote{Note that water contents of a few percent by mass, though $\\ga$10-20 times larger than the Earth's estimated water budget (L\\'ecuyer \\etal 1998), would have a negligible effect on the planetary radius compared with the observational uncertainties.}  Estimates of both the planetary mass and radius are needed to derive a bulk composition (Selsis \\etal 2007).  For transiting planets, errors in stellar masses and radii (Ford \\etal 1999; Cody \\& Sasselov 2002; Fischer \\& Valenti 2005; Sozzetti \\etal 2007) are likely to lead to errors in planetary radii on the order of 2-10\\% (see section 6 of Seager \\etal 2007).  With a determination of the planetary mass to within 5\\% and radius to within 10\\%, it may be possible to differentiate between mostly rocky (Earth-like) planets and icy planets with $\\geq$ 10\\% water by mass (Valencia \\etal 2007b).  For transiting planets with very precise mass and radius measurements (better than 2\\%), more detailed compositions may be derived (Seager \\etal 2007).  Atmospheric envelops of H/He can inflate the observed radii of solid planets by tens of percent (Adams \\etal 2007).  For a given radius measurement, solutions for the planetary structure become degenerate with respect to water content and envelope mass: small radii are clear signatures of rocky planets, but larger radii are ambiguous.  Theoretical models suggest that the ability of a planet to accrete a gaseous envelope depends on the planet's mass, and is not sensitive to the orbital distance (e.g., Ikoma \\etal 2001; Ida \\& Lin 2004).\\footnote{In addition, as Adams \\etal (2007) point out, significant envelopes of hydrogen may be formed as a result of outgassing from the planetary interior.}  Thus, additional information about the presence and thickness of the planet's atmosphere is needed to determine whether the planet is water-rich or rocky.   For example, more information could be gathered from a situation in which an atmosphere has almost certainly photo-evaporated away (old stellar age plus very close-in planet).  Another favorable case would be a situation for which some spectral information about the planet's atmosphere could potentially be obtained.  Observational limitations are such that certain planets will have no mass estimates, given the faintness of their host stars and the consequent difficulty of RV followup.  For such cases, it is possible to place mass limits based on maximum and minimum radius estimates, i.e., by assuming the planet to be made of pure iron or pure water (or pure hydrogen for gaseous planets).  In addition, the bulk planetary composition may not be determined in many cases because of observational limitations (Selsis \\etal 2007) or degeneracy between model parameters (Adams \\etal 2007).  With no composition information it becomes more difficult to differentiate between formation models.  Nonetheless, several cases can be distinguished if the inner planetary system architecture is known.  Thus, current and future programs have the sensitivity to determine the orbits, masses, radii and companions of a large number of hot Earths. Although it will not be feasible in all cases, information about other planets in hot Earth systems will be determined via RV, transits, and transit timing analysis.  In this paper we focus on systems in which both the inner planetary architecture and the composition of a hot Earth (rocky vs. $>$10\\% water) can be determined (i.e., the brightest {\\it CoRoT} and {\\it Kepler} targets; see FIg.~6 of Selsis \\etal 2007). As explained below and summarized in Table 1, analysis of these data may be able to identify the formation mechanism of such planets.  \\section {Models for hot Earth formation}  Here we investigate six mechanisms for hot Earth formation, including proven and previously untested mechanisms.  The six models are: ($\\S3.1$) {\\it in situ} accretion; ($\\S3.2$) type 1 migration; ($\\S3.3$) shepherding during giant planet migration; ($\\S3.4$) shepherding via secular resonance sweeping; ($\\S3.5$) tidal circularization of eccentric planets; and ($\\S3.6$) photo-evaporation of close-in giant planets.  For each model, we discuss the state of the inner planetary system, as well as the likely composition of the hot Earth(s). Table 1 summarizes the differences between models.  The first four models listed have been demonstrated in previous work.  We introduce two additional models for hot Earth formation, and test them quantitatively below.  \\subsection{{\\it In situ} formation}  If protoplanetary disks contain a substantial mass in solids close to their stars, then perhaps hot Earths can form from local material.  This depends critically on the condensation temperatures of grains (Pollack \\etal 1994; Lodders 2003), disks' inner truncation radii (e.g., Eisner \\etal 2005; Akeson \\etal 2005), and the surface density profile of solids (Weidenschilling 1977; Hayashi 1981; Davis 2005; Raymond \\etal 2005).  If hot Earths form {\\it in situ}, then their growth would be similar to that of Solar System's terrestrial planets (Wetherill 1990, 1996; Chambers \\& Wetherill 1998; Agnor \\etal 1999; Morbidelli \\etal 2000; Chambers 2001; Kenyon \\& Bromley 2006), but minus the dynamical effects of Jupiter and Saturn (although giant planets may co-exist with some hot Earths, e.g. Gliese 876; Rivera \\etal 2005).  If there is sufficient mass to form one hot Earth {\\it in situ}, then we expect a population of several hot Earths to form, with masses determined by the local disk mass and spacings similar to those in the Solar System and in accretion simulations (roughly 20-80 mutual Hill radii: $R_{H,m} = 0.5 (a_1 + a_2) (M_1+M_2/3 M_\\star)^{1/3}$; $a_1$ and $a_2$ are the orbital radii and $M_1$ and $M_2$ the masses of two adjacent planets).  The surface density distribution of protoplanetary disks, $\\Sigma$, is thought to scale with radial distance $r$ as $\\Sigma \\sim f r^{-\\alpha}$, where $f$ is a scale factor to account for the large variability in observed disk masses (Andre \\& Montmerle 1994; Eisner \\& Carpenter 2003; Andrews \\& Williams 2005; Scholz \\etal 2006), and the value of $\\alpha$ lies between 0.5 and 2 (Weidenschilling 1977; Hayashi 1981; Kuchner 2004; Davis 2005; Dullemond \\etal 2007; Andrews \\& Williams 2007; Garaud \\& Lin 2007).  Accretion models suggest that planets close to their stars are generally smaller than those farther out for $\\alpha < 2$ (Lissauer 1987; Kokubo \\& Ida 2002; Raymond \\etal 2005, 2007; Kokubo \\etal 2006).  For solar-type stars, hot Earths that form {\\it in situ} are likely to be dry because of the low efficiency of water delivery from both comets (Levison \\etal 2000) and asteroids (Raymond \\etal 2004).  Because of the very hot local temperatures, these planets would be mainly composed of refractory materials such as iron and rock (Pollack \\etal 1994; Lodders 2003).  However, for the case of low-mass stars, the snow line is located very close-in (as is the habitable zone -- Kasting \\etal 1993).  Water delivery to hot Earths may therefore be more favorable around low-mass stars, although impact speeds are high and formation times fast compared with Earth's formation zone (Lissauer 2007; Raymond \\etal 2007), and the snow line moves significantly in the disk lifetime (Sasselov \\& Lecar 2000; Kennedy \\etal 2006).  Figure~\\ref{fig:insitu} shows snapshots of {\\it in situ} accretion of terrestrial material close to a 0.31 $\\msun$ star from Raymond \\etal (2007, in preparation), designed to examine the GJ 581 system.  The simulation started from a disk of 57 planetary embryos (initially separated by 3-6 Hill radii, as in Raymond \\etal 2006a) and 500 planetesimals in a very massive disk totaling 40 $\\mearth$ between 0.03 and 0.5 AU.  The disk's surface density decreased with orbital distance $r$ as $r^{-1}$ (i.e., $\\alpha = 1$), and was roughly 30 times more massive than the minimum-mass solar nebula model (Weidenschilling 1977; Hayashi 1981; Davis 2005).\\footnote{Note that this disks as massive as the one in Fig.~\\ref{fig:insitu} (30 times the minimum-mass disk) are likely to be quite rare (e.g., Andrews \\& Williams 2005).  Thus, {\\it in situ} accretion of close-in planets of several Earth masses can probably occur only for a small fraction of protoplanetary disks (see $\\S$ 4).}  The three planets that formed in this simulation have masses of 6.6 $\\mearth$ (at 0.06 AU), 10.9 $\\mearth$ (0.12 AU), and 10.6 $\\mearth$ (0.30 AU).  Each has a substantial water content, acquired via collisions with material originating beyond the \"water line\" at 0.29 AU, but note that the effects of water depletion during impacts (Genda \\& Abe 2005; Canup \\& Pierazzo 2006) and hydrodynamic escape (Matsui \\& Abe 1986; Kasting 1988) have not been accounted for.  \\begin{figure*} \\begin{center} \\leavevmode \\epsfxsize=15cm\\epsfbox{fig1.eps} \\caption[]{Snapshots of {\\it in situ} accretion of a system of hot Earths. Each object's size is proportional to its mass$^{1/3}$, and the color corresponds to its water content, from 10$^{-5}$ water by mass (red) to 5\\% water by mass (darkest blue; see color bar).  The dark circle in the center of each body refers to the relative size of its iron core (see Raymond \\etal (2005) for details).  In this case, the central star is 0.31 $\\msun$, the same as Gliese 581.}  \\label{fig:insitu} \\end{center} \\end{figure*}  Thus, if hot Earths form {\\it in situ} then we expect systems of several hot Earths to form, with spacings comparable to the solar system terrestrial planets.  These planets will contain mainly local, dry material, although for low-mass stars they may contain up to perhaps 0.1-1\\% water by mass (but see Raymond \\etal 2007, Lissauer 2007).  \\subsection {Inward Type 1 Migration}  Earth- to $>$Neptune-mass planets excite density waves in the disk (Goldreich \\& Tremaine 1979).  The back reaction of these waves on the planet causes inward orbital migration on a timescale of 10$^{4-6}$ years (Goldreich \\& Tremaine 1980; Ward 1986, 1997; D'Angelo \\etal 2003; Masset \\etal 2006a). However, type 1 migration may be stopped or even reversed in the very inner regions of the disk, because of net disk torque changes at disk edges (Masset \\etal 2006b) or in the optically thick inner disk (Paardekooper \\& Mellema 2006).  Thus, hot Earths could form far from their stars and migrate in to the location in the disk where the disk torques cancel out and migration is stopped.  Presumably, if there is enough solid and gas mass to enable type 1 migration of one planet, then others should follow.  Co-migrating planets may end up trapped in mean motion resonances (e.g., Lee \\& Peale 2002), and can form chains of many planets in paired resonances.  These orbital chains of planets can survive for long times, as the outward-directed torques on the inner planets may be balanced by inward-directed torques on the outer planets (Terquem \\& Papaloizou 2007).  Surviving planets do not remain on strictly resonant orbits, and collisions between planets can occur after the disk dissipates.  The amount of solid material in the disk is thought to increase by a factor of 2-4 or perhaps more beyond the snow line (Hayashi 1981, Stevenson \\& Lunine 1988; Lodders 2003).  In addition, most disk surface density profiles contain far more mass in their outer regions.  Thus, it seems reasonable to assume that, in this model, most hot Earths must form beyond the snow line and are therefore icy in composition rather than rocky.  Indeed, transits of the hot Neptune GJ 436 b have been interpreted as an indication that it may be largely composed of water and may therefore have formed beyond the snow line and migrated inward (Gillon \\etal 2007). Note, however, that inferring a detailed planetary composition from a radius measurement is ambiguous because different combinations of rock, ice and H/He envelopes can form planets with the same mass and radius (see Fig.~3 from Adams \\etal 2007).  In addition, formation and migration models suggest that in a system of several hot Earths it is possible for the innermost hot Earth to be rocky (Alibert \\etal 2006).  The main consequence of the type 1 migration model is simply that hot Earths didn't form locally but farther out in the disk, probably in the water-rich icy regions.  Therefore, hot Earths should contain a large quantity of ice and have measurably larger radii.  In addition, the migration process favors the formation of a chain of resonant or near-resonant planets (Terquem \\& Papaloizou 2007).  \\subsection{Shepherding by giant planet migration}  Giant planets more massive than a critical value carve an annular gap in the protoplanetary disk and are thus locked to the disk's viscous evolution (Lin \\& Papaloizou 1986; Takeuchi \\etal 1996; Bryden \\etal 1999; Rafikov 2002; Crida \\etal 2006).  These planet subsequently ``type 2'' migrate, usually inward, on a $\\sim 10^5 - 10^6$ year timescale, depending on the disk's viscosity (Lin \\& Papaloizou 1986, Lin \\etal 1996; Ward 1997; D'Angelo \\etal 2003).  Such a planet migrates through a disk composed of both gas and solids in the form of km-sized planetesimals and Moon- to Mars-sized planetary embryos, which formed in series of dynamical steps from micron-sized dust grains (as in the {\\it in situ} formation model; see $\\S$ 3.1 or Chambers 2004, Papaloizou \\& Terquem 2006 for reviews).  As the giant planet migrates inward, it shepherds material in front of strong mean motion resonances (MMRs).  The evolution of a typical planetary embryo in the inner disk proceeds as follows.  As the giant planet approaches the embryo, the embryo's eccentricity is increased by an MMR (usually the 2:1 or 3:2, but higher-order resonances are stronger for more eccentric giant planets; Murray \\& Dermott 1999).  Gas drag and dynamical friction with nearby planetesimals act to recircularize the embryo's orbit and decrease its energy, thereby reducing its semimajor axis and moving it just interior to the MMR (Adachi \\etal 1976; Tanaka \\& Ida 1999).  As the giant planet continues its migration, the embryo is again excited by the approaching MMR and the cycle continues.  Thus, embryos and planetesimals are shepherded inward by moving MMRs and accrete into planet-sized bodies during giant planet migration\\footnote{The formation timescale of shepherded hot Earths is on the order of the migration timescale (Mandell \\etal 2007).  Thus, hot Earths may form in $\\sim 10^{5-6}$ years, as opposed to the 10$^{7-8}$ year timescale for the Earth calculated from Hf/W isotopic measurements (Kleine \\etal 2002; Jacobsen 2005).  This very short formation timescale for hot Earths could have consequences for their geological evolution.}  (Fogg \\& Nelson 2005, 2007; Zhou \\etal 2005; Raymond \\etal 2006b; Mandell \\etal 2007).  However, during this process, many bodies' eccentricities are damped too slowly to avoid a close encounter with the giant planet.  Such bodies are usually scattered outward, and can form a subsequent generation of exterior terrestrial planets (Raymond \\etal 2006b; Mandell \\etal 2007).  Figure~\\ref{fig:migshep} shows snapshots in time of this shepherding process from a simulation by Mandell \\etal (2007).  It is clear that the 2:1 MMR is responsible for the bulk of the shepherding in this simulation.  The two hot Earths formed are on low-eccentricity orbits immediately interior to strong resonances, as expected.  However, in simulations including weaker gas drag, hot Earths can form on higher eccentricity orbits (Fogg \\& Nelson 2007, Mandell \\etal 2007).  The survival of high eccentricity hot Earths is uncertain, given that tides may act to alter the planets' orbits and possibly lead them into unstable giant planet resonances or drive them into the star (see $\\S$ 3.5 below).  In the giant planet migration shepherding model, hot Earths are a mixture of material that originated interior to the giant planet's orbit.  Both the core-accretion and gravitational collapse models predict that giant planets are likely to form at large orbital distances, beyond the snow line, which itself moves inward in time (Pollack \\etal 1996; Boss 1997; Bodenheimer \\etal 2000; Sasselov \\& Lecar 2000; Mayer \\etal 2002).  In the simulations of Raymond \\etal (2006b) and Mandell \\etal (2007), the hot Earths that formed contained 1-2\\% water by mass, as is the case for the two hot Earths in Fig.~\\ref{fig:migshep}.  The assumed starting water distribution in those cases was similar to that of current-day primitive asteroids (Abe \\etal 2000; Fig.~2 from Raymond \\etal 2004).  Note that water depletion from impacts and hydrodynamic escape was neglected in these calculations.  \\begin{figure*} \\begin{center} \\leavevmode \\epsfxsize=15cm\\epsfbox{fig2.eps} \\caption[]{Snapshots in time of the migration of a Jupiter-mass giant planet through a disk of terrestrial bodies, including the effects of aerodynamic gas drag, from a simulation by Mandell \\etal (2007).  Specific strong resonances with the giant planet are indicated.  Colors refer to bodies' water contents, as in Fig.~\\ref{fig:insitu}.  This version reproduced from Gaidos \\etal (2007).}  \\label{fig:migshep} \\end{center} \\end{figure*}  The giant planet migration shepherding model predicts that hot Earths should lie close to a strong MMR, most likely the 2:1 MMR, interior to a giant planet (Fogg \\& Nelson 2005; Zhou \\etal 2005; Raymond \\etal 2006b; Mandell \\etal 2007).  In this model, hot Earths are formed from a mixture of material that originated interior to the giant planet's orbit.  Giant planets are expected to form just outside the snow line, given the increase in solid material (Hayashi 1981; Stevenson \\& Lunine 1988; Ida \\& Lin 2004).  Thus, hot Earths formed by shepherding are likely to contain up to a few percent water, but probably not more (Mandell \\etal 2007).  Note that planets with water contents of a few percent by mass cannot be distinguished from rocky planets by radius measurements given observational uncertainties (Valencia \\etal 2007b; Fortney \\etal 2007; Sotin \\etal 2007; Seager \\etal 2007).  \\subsection{Shepherding by Sweeping Secular Resonances during Disk Dispersal}  Moving secular resonances (SRs) can shepherd material in a similar way to mean motion resonances (MMRs) if gas drag is present.  An SR occurs when the apsidal precession frequency of two bodies' orbits are commensurate (e.g., Murray \\& Dermott 1999).  In a disk with two or more giant planets, interactions between the planets cause each of their orbital alignments to precess.  In addition, the gravitational potential of the massive gaseous disk affects the precession rates, and therefore the location of SRs with each planet in the disk (Ward 1981; Nagasawa \\etal 2005).  As the disk dissipates, SRs can move progressively (``sweep'') across a given region, increasing the eccentricities of bodies.  In the case of a smooth, inward-sweeping SR, shepherding of material can happen similar to MMR shepherding for migrating giant planets.  In the context of hot Earth formation, the SR shepherding model applies to cases with two or more giant planets that have stopped migrating.  A smooth dissipation of the disk can induce SR sweeping.  Much as in the migration shepherding mechanisms, a sweeping SR excites the eccentricities of nearby protoplanets.  These eccentricities are subsequently damped by gas drag and the body's orbit is moved interior to the resonance.  This process continues for the duration of the SR sweeping, unless a planet gets close enough to the star that its precession rate becomes dominated by general relativistic effects rather than dynamical ones (Zhou \\etal 2005).  Thus, the secular resonance shepherding model involves a complex interaction between two giant planets, the massive gaseous disk, and relatively low-mass terrestrial material.  It requires a monotonic, inward secular resonance sweeping which itself requires a smooth dispersal of the gaseous disk (Ward 1981), which is uncertain given that most stars form in large clusters and may lose disk mass in periodic photo-evaporation events (Lada \\& Lada 2003; Adams \\etal 2004; Hester \\etal 2004).  In the SR shepherding model, a hot Earth system must also contain at least two more distant, interacting giant planets. The compositions of hot Earths in this scenario are a mixture of material from interior to the giant planets' starting orbits.  Estimating the compositions of hot Earths in this model therefore requires a knowledge of the giant planet's formation locations, specifically how far past the snow line they formed.  \\subsection{Tidal Circularization of Eccentric Terrestrial Planets}  The circular orbits of hot Jupiters have been attributed to energy and angular momentum dissipation via tides raised on the planet by the star (Rasio \\etal 1996).  In fact, it has been proposed that close-in giant planets may have been scattered onto high-eccentricity orbits and tidally circularized (Rasio \\& Ford 1996; Weidenschilling \\& Marzari 1996; Mardling \\& Lin 2004; Jackson \\etal 2007).  Could tidal circularization act as a mechanism to transport terrestrial planets inward?  To address this possibility, we integrated the second-order, coupled semimajor axis $a$ and eccentricity $e$ tidal evolution equations (Goldreich \\& Soter 1966; Kaula 1964; Greenberg 1977): \\begin{equation} \\frac{da}{dt} = -\\Big(21\\frac{\\sqrt{GM_*^3}R_p^5k_p}{m_pQ'_p}e^2 + \\frac{9}{2}\\frac{\\sqrt{G/M_\\star}R_\\star^5m_pk_\\star}{Q'_\\star}\\Big)a^{-11/2} \\end{equation}  \\begin{equation} \\frac{de}{dt} = -\\Big(\\frac{21}{2}\\frac{\\sqrt{GM_\\star^3}R_p^5k_p}{m_pQ'_p} + \\frac{171}{16}\\frac{\\sqrt{G/M_\\star}R_\\star^5m_pk_\\star}{Q'_\\star}\\Big)a^{-13/2}e \\end{equation}  \\noindent where $Q'_p$ and $Q'_\\star$ are the tidal dissipation functions of the planet and star, respectively, $k_p$ and $k_*$ are the Love numbers of the planet and star, $m_p$ and $M_\\star$ are the masses of the planet and star, $R_p$ and $R_\\star$ are the radii of the planet and star, and $G$ is the gravitational constant.  Note that solutions with higher order terms in $e$ have been derived (e.g., Hut 1981; Eggleton \\etal 1998). However, such models include, in effect, assumptions about how a body responds to the ever-changing tidal potential, effects that have not been observed. Therefore not enough is known about the actual response of real bodies to evaluate these higher-order effects. As we are only interested in the qualitative differences between planets whose orbits have evolved through tidal decay and those that did not, the second order solution should suffice.  We considered a stellar mass of 0.3, 1, and 3 $\\msun$ with radii determined from Gorda \\& Svechnikov (1999), a planet mass of 1 and 5 $\\mearth$ (assuming $R_p \\propto m_p^{0.27}$, as suggested by Valencia \\etal 2006), a perihelion distance from 0.025 to 0.1 AU, and eccentricities $e$ from 0 to 0.9.  For the planet, we assumed $k_p = 0.3$ and $Q'_p = 21.5$ (Dickey \\etal 1994; Mardling \\& Lin 2004); for the star, $k_\\star = 1.5$ and $Q'_\\star = 10^{5.5}$ (Ogilvie \\& Lin 2007; Jackson \\etal 2007). Each orbit was integrated for 10 Gyr using a 10$^3$ year timestep, which convergence tests showed is three orders of magnitude smaller than necessary to produce reliable results.  Figure~\\ref{fig:tides} shows the evolution of a set of 5 $\\mearth$ planets with the same initial perihelion distance of 0.025 AU, and starting eccentricities ranging from 0.05 ($a$ = 0.026 AU) to 0.9 ($a$ = 0.25 AU).  As expected, evolution proceeds much faster for bodies at smaller orbital distances (in this case, those with lower eccentricities).  For planets with starting eccentricities of 0.6 or less ($a \\leq 0.06$ AU), orbital circularization takes place within $10^8$ years, including an inward drift in semimajor axis of up to 0.01-0.02 AU.  Circularization takes longer for larger $a$ values, but the amount of inward drift is also increased.  For the $e$ = 0.8 planet, circularization requires several Gyr, but the planet moves inward from 0.125 to 0.065 AU.  At still-larger orbital distances (and eccentricities), circularization takes longer than the age of the star.  Note that orbital evolution continues slowly after the planet's orbit becomes circular, via tides raised on the star by the planet.  Mardling \\& Lin (2004) showed that should $a$ become very small ($\\la 0.01$ AU), then the planet is doomed to fall into the star within a few Gyr. We confirm that assessment here. Therefore we expect no planets inside 0.01 AU for any formation scenario.  \\begin{figure*} \\begin{center} \\leavevmode \\epsfxsize=15cm\\epsfbox{fig3.eps} \\caption[]{Orbital semimajor axis $a$ and eccentricity $e$ vs. time for a series of $5 \\mearth$ planets orbiting Sun-like stars with starting perihelion distances of 0.025 AU.  Each curve corresponds to a planet with a given starting eccentricity, from 0.05 ($a$ = 0.026 AU) to 0.9 ($a$ = 0.25 AU). If $a$ drops below $\\sim 0.01$ AU, then the planet will fall into the star due to tidal evolution.}  \\label{fig:tides} \\end{center} \\end{figure*}  The degree to which tidal circularization can move a planet inward clearly depends on its starting orbit, size, and ability to dissipate energy. Figure~\\ref{fig:tides} shows the most tidal evolution of any of the cases we explored for a Solar-mass star; in most cases evolution was slower (except for the 3 $\\msun$ cases, which were faster).  Thus, it appears that inward movement of planets during tidal evolution is relatively small, at most a $\\sim$0.05 AU change in semimajor axis.  However, if the large eccentricity were due to a perturbative event, then the planet's starting aphelion might be representative of its pre-encounter semimajor axis.  In that case, the effective inward movement due to tides is doubled (starting aphelion to final semimajor axis), although it would still be less than about 0.1-0.15 AU on a $>$Gyr timescale.  The composition of the planet depends on its formation history, especially whether it formed locally or migrated inward.  What is the source of the large eccentricity needed to drive tidal circularization? Planet-planet scattering has been invoked to explain the large eccentricities of the known extra solar planets (Rasio \\& Ford 1996; Weidenschilling \\& Marzari 1996; Lin \\& Ida 1997; Ford \\etal 2005; but see Barnes \\& Greenberg 2007).  The strength of a scattering event depends on a combination of the escape speed of the perturber, the encounter velocity, and the escape speed from the system.  For close-in planets, the system escape speed is large, and so only very massive bodies can excite large eccentricities.  Indeed, accretion may be preferred over scattering in these situations (Goldreich \\etal 2004).  One alternative mechanism for eccentricity growth is an instability in nearby giant planets could alter the orbit of a terrestrial planet (Veras \\& Armitage 2006).  In that case, one or more exterior giant planets should exist in the system, on eccentric orbits.  Another possible source of eccentricity could arise if the host star had a binary companion.  If the orbital plane of the planet were significantly inclined with respect to that of the binary, then large eccentricities could be induced via the Kozai mechanism (Kozai 1962). In most of these models, some evidence for an external perturber should be evident.  Thus, tidal circularization can move a highly-eccentric terrestrial planet inward to some extent, although the planet must be relatively massive ($\\ga 5 \\mearth$) and have a very small starting perihelion distance ($\\la$ 0.03 AU). If the planet formed locally, then its composition is likely to be relatively dry ($\\la$ 1\\% water by mass).  A source of high eccentricity may also be evident, such as a binary stellar companion or a distant eccentric giant planet.  \\subsection{Giant Planet Migration and Photo-evaporation}  Baraffe \\etal (2004, 2006) proposed that close-in, Neptune-mass planets might be the remains of larger planets that have been photo-evaporated away.  Such planets would form farther from their parent stars and migrate inward (Ida \\etal 2004, Alibert \\etal 2005), losing a portion of their gaseous envelopes hydrodynamically via irradiative XUV heating (Lammer \\etal 2003, Baraffe \\etal 2004).  Here we investigate the possibility that photo-evaporation could lead to the removal of the entire envelope of a hot Jupiter or hot Neptune, leaving behind a solid planet, i.e., the core of the irradiated giant planet.  Recent estimates derive evaporation rates that are far smaller than those used by Lammer \\etal (2003), and include the effects of two-dimensional layering (Tian \\etal 2005) and improved atmospheric chemistry (Yelle 2004, 2006).  Indeed, these new evaporation rates are closer to those of Watson \\etal (1981).  Perhaps most convincing that the Lammer \\etal evaporation rates are too large is empirical evidence that the mass distribution of highly irradiated extra-solar planets (inside 0.07 AU) is identical to that of more distant planets (Hubbard \\etal 2007a).  A substantial change in the mass function is predicted for evaporation models (i.e., fewer massive planets and more less-massive ones).\\footnote{Fortney \\etal (2007) also showed that if hot Neptunes form via photo-evaporation of hot Jupiters, then their radii should be on the order of one Jupiter radius.  However, if they are not remnants of hot Jupiters, then their radii should be 0.3-0.4 $R_J$.  The first transiting hot Neptune indeed has a radius of $\\sim 0.35-0.4 R_J$ (Gillon \\etal 2007; Deming \\etal 2007).  Note, however, that Gliese 436 is an M dwarf (0.41 $\\msun$) and therefore has low EUV and FUV emission (except during flares), which are key for driving evaporative mass loss (Butler \\etal 2004).} Such an effect may exist at lower masses, but not in the currently-probed sample of planets. Hubbard \\etal (2007b) show that at the minimum orbital radius of known extra-solar planets (0.023 AU), the initial mass must be less than about a Saturn mass to evaporate completely, i.e., to its core.  For more typical hot Jupiter orbits, at 0.05-0.1 AU, this critical mass is smaller still.\\footnote{Tidal evolution models suggest that close-in planets may have originated at somewhat larger orbital distances and slowly evolved inward, concurrently decreasing their semimajor axes and eccentricities (Jackson \\etal 2007).  If close-in giant planets did indeed originate on somewhat more distant, eccentric orbits, then their time-averaged fluxes would likely be reduced.}  The models of Baraffe \\etal (2004) and Hubbard \\etal (2007b) do not account for the presence of a core, which is important once the planet mass is less than $\\sim 100 \\mearth$, such that a 5-10 $\\mearth$ core constitutes a non-negligible fraction of the planet mass.  Note that the Baraffe \\etal (2006) models do incorporate this effect, as we do implicitly by using their internal structure models.  Mass loss due to hydrodynamic escape, limited only by energy deposition, depends critically on the stellar irradiance of the atmosphere, and can be approximated by the relation \\begin{equation} \\dot{M} = \\frac{3\\, \\beta (F_* , a)^3}{G\\rho}\\, \\frac{F_{XUV} + F_{\\alpha}}{a^2} \\end{equation} where $F_{XUV}$ and $F_{\\alpha}$ represent the high-energy radiation incident on the planet, and $\\rho$ and $a$ are the planet density and orbital distance from the star respectively (Lammer \\etal 2003; Baraffe \\etal 2004). The parameter $\\beta$ is the ratio of the irradiated planetary radius to the planet's ``original'', non-irradiated radius for a specific stellar flux and orbital distance (Lammer \\etal 2003, Baraffe \\etal 2004, 2006, and Hubbard \\etal 2007b all assume $\\beta$ = 3 based on atmospheric models of Watson \\etal 1981).  For a constant orbital distance, the mass loss will therefore initially decrease in time as the planet cools and becomes more dense and the star's UV and x-ray flux decreases.  Over time, as hot material escapes from the top of the planetary atmosphere new layers are irradiated and stellar flux is converted to expansion energy, gravitational contraction of the planet slows.  If enough mass is evaporated, expansion surpasses contraction and the planet experiences run-away mass loss, leaving behind only the solid core.  To construct a simplified model of photo-evaporative mass loss, we need to constrain certain parameters.  We assumed evaporation rates 100 times smaller than the energy-limited case from Lammer \\etal (2003).  The radius of an evaporating planet stays relatively constant regardless of mass, such that we extrapolated radii for planets of various masses and heavy-element abundances from Baraffe \\etal (2006, with corrections for semimajor axis from Chabrier \\etal 2004) to find a mass-radius relation for irradiated planets as a function of time. Planets less massive than $\\sim50 \\mearth$ are more likely than larger planets to contain significant concentrations of molecular species, simply because the ratio of core mass to envelope mass is decreased (e.g., Uranus and Neptune).  Molecules such as $\\rm H_2O$ and $\\rm CH_4$ play an important part in the energetics of atmospheric expansion, and therefore affect the mass loss rate (Hubbard \\etal 2007a).  Although our approach does not directly incorporate changes in evaporation rate with chemistry, the mass-radius relations from Baraffe \\etal (2006) are based on Alibert \\etal (2005)\u00d5s values for heavy-element enrichment of the planets' atmospheres and therefore implicitly include a decreased evaporation rate for smaller planets with heavy-element-rich atmospheres since evaporation rates depend on the planet radius.  Our simplified model demonstrates good agreement with the results from more detailed models by Hubbard et al. (2007b) and the reduced-evaporation models of Baraffe et al. (2006) for higher-mass planets.  \\begin{figure*} \\begin{center} \\leavevmode \\epsfxsize=15cm\\epsfbox{fig4.eps} \\caption[]{The evolution of the masses of two highly-irradiated planets due to photo-evaporation.  Our model is based on planetary structure models of irradiated planets from Baraffe \\etal (2006) and Chabrier \\etal (2004) and evaporation rates based on the model of energy-limited hydrodynamic escape from Lammer \\etal (2003), but reduced by a factor of 100 in accordance with the results of Hubbard \\etal (2007). The intermediate-mass planet ($72 \\mearth$) represents the boundary between planets that would undergo type I migration (M$\\la 70 \\mearth$) and planets that would undergo type II migration (M$\\ga 70 \\mearth$; D'Angelo \\etal 2003). The planets are placed at three different orbital radii: 0.025 AU (solid), 0.05 AU (dotted), and 0.1 AU (dashed).}  \\label{fig:photo} \\end{center} \\end{figure*}  Figure~\\ref{fig:photo} shows the masses of two highly-irradiated planets as a function of time, on circular orbits at 0.025, 0.05, and 0.1 AU.  Even on its closest orbit, the more massive planet ($72 \\mearth$) took 5 Gyr to evaporate to its core.  Given that the transition from type 1 to type 2 migration is thought to occur at roughly 70 $\\mearth$ (Ward 1997; D'Angelo \\etal 2003), this effectively rules out the formation of hot Earths by type 2 migration and subsequent photo-evaporation, simply because planets massive enough to type 2 migrate will not lose enough mass by photo-evaporation.  In contrast, the less massive planet ($25 \\mearth$) evaporated to its core in less than 50 Myr at 0.025 AU, but required Gyr to lose mass past 0.05 AU.  Additionally, beyond $\\sim$0.1 AU the mass loss over 5 Gyrs is negligible; planets found beyond this orbital distance maintain their mass over the lifetime of the system.  After the primordial atmosphere has been lost, the core of the planet is exposed.  If the core of the planet is composed primarily of low-temperature condensates (consistent with formation in the cold outer disk -- Pollack \\etal 1996; Boss 1997), the outer layers of the planetary core may continue to vaporize. To determine the original mass of a hot Earth formed from the core of a photo-evaporated massive planet, models of the evolution of intensely irradiated icy bodies must be developed; current models of volatile-dominated Earth-mass planets (Valencia \\etal 2006, 2007a, 2007b; Fortney \\etal 2007; Sotin \\etal 2007; Seager \\etal 2007) have not yet probed these processes.  Thus, photo-evaporation of close-in gaseous planets may remove their atmospheres and leave behind solid cores.  This process is only effective for planets within about 0.05 AU that are below the type 2 migration threshold of $\\sim 70 \\mearth$.  For those cases, ``hot Neptunes\" could become ``hot Earths\" on a $10^{8-10}$ year timescale.  This process can occur in conjunction with the type 1 migration scenario discussed above in $\\S$ 3.2. Indeed, the most likely source of hot Neptunes is the outer disk (Gillon \\etal 2007).  In some cases, a series of $\\ga 10 \\mearth$ planets may form in the cold outer disk and then type 1 migrate inward as described above, into a chain of hot Neptunes.  Depending on their orbits, the innermost planet or two could be photo-evaporated over time into a very water-rich hot Earth.  A diagnostic of photo-evaporation could therefore be a system with 1) a water-rich hot Earth inside 0.05 AU and 2) additional, $\\sim$ Neptune-mass planets exterior to the hot Earth in near-resonant orbits.  However, it would still be difficult to definitely assess the degree of photo-evaporation in such a setting.  More definitive detections of the importance of photo-evaporation would require statistics of a large number of hot Earths and hot Neptunes orbiting stars with a range of ages.  A correlation between the number of hot Earths vs. hot Neptunes and the stellar age would indicate that such planets were losing mass in time, presumably via photo-evaporation.  Alternatively, exploring the mass functions of close-in planets down to lower masses could reveal time-dependent mass loss (as in Hubbard \\etal 2007a). ",
        "conclusions": "We anticipate that a large number of planetary systems containing close-in terrestrial planets, referred to here as ``hot Earths'', will be discovered in the coming years with radial velocity and transit measurements.  In some cases, both an accurate determination of the architecture of the inner planetary system and of the bulk composition of a hot Earth (rocky vs. icy; but see Adams \\etal 2007) will be possible (see $\\S$ 2).  The goal of this paper is to determine whether the formation history of such systems can be unraveled, given the relatively small amount of information available.  In addition to four already-known mechanisms for hot Earth formation, we have shown that tidal circularization of highly eccentric planets can move terrestrial planets' orbits inward, but only by perhaps 0.1 AU, and only for very close-in perihelion distances ($\\la 0.05$ AU).  In addition, our simple model suggests that photo-evaporation can remove a giant planet's atmosphere and leave behind the core.  However, this is only possible for very close-in orbits ($< 0.025-0.05$ AU) and relatively low-mass planets (``hot Neptunes\" with masses below $70 \\mearth$), as suggested by Hubbard \\etal (2007a).  Table 1 summarizes the observable consequences of these six models for hot Earth formation.  There exist several clear differences between the models that should be detectable in the near future.  Given a planetary system with a transiting hot Earth, considerable RV measurements, and perhaps transit timing analysis, Table 1 provides a simple way to determine the formation history of hot Earth planetary system.  Note that in some cases, more than one of these mechanisms can act in concert.  For example, the case of GJ 876 may be explained in a two-step process, via shepherding during migration and then during secular resonance sweeping (see $\\S$ 4.1; Zhou \\etal 2005).  In addition, tides affect the orbits of all hot Earths to some degree, regardless of their formation history.  However, certain mechanisms cannot act together: planets massive enough to type 2 migrate cannot have their envelopes photo-evaporated and become hot Earths (see $\\S$ 3.6).  The formation mechanisms of the two known hot Earth systems are not entirely clear (see $\\S$ 4 above).  However, transit measurements of the hot Earth of either of the known systems would make it far easier to discern between models.  In particular, for the case of GJ 581, a transit measurement of planet c or d would distinguish between {\\it in situ} formation (rocky) and type 1 migration (icy).  Clearly, more work is needed to better characterize and quantify some of these models, and to examine the long-term survival of hot Earths in different systems.  In addition, it is possible that additional mechanisms exist for hot Earth formation that have not yet been considered.  Zhou \\etal (2005) claimed that hot Earths should be numerous if giant planets form via core-accretion (Mizuno 1980; Pollack \\etal 1996; Lissauer \\& Stevenson 2007), but rare if they form via gravitational instability (Boss 1997; Mayer \\etal 2002; Durisen \\etal 2007).  Given the large number of avenues for hot Earth formation, we disagree with Zhou \\etal on this point. Indeed, three of the candidate mechanisms for hot Earth formation -- {\\it in situ} accretion, type 1 migration, and tidal circularization -- do not require a giant planet at all and so are unaffected.  Photo-evaporation of hot Neptunes may be affected because the two giant planet formation models predict different core masses: core-accretion predicts 5-20 $\\mearth$ cores (e.g., Alibert \\etal 2005) while the cores of giant planets formed via disk instability are likely to be smaller (Boss 1998, 2006).  For the other two mechanisms -- giant planet migration shepherding and secular resonance shepherding -- is there a reason that the outcome should depend on the mode of giant planet formation?  The main difference between the two models is the timing of giant planet formation: core-accretion predicts that giant planets form late in the lifetime of the gaseous disk, while gravitational instability forms planets very quickly.  Giant planet migration starts immediately after, or even during, formation (Lufkin \\etal 2004).  Thus, if giant planets form via core-accretion, they migrate through a disk that has undergone at least 1 Myr of accretion, and contains both $\\sim$Moon-sized planetary embryos and km-sized planetesimals (plus $\\sim$ 99\\% gas; e.g., Kokubo \\& Ida 2000; Chambers 2004).  If, however, giant planets form via gravitational instability, then they would migrate through a disk containing predominantly smaller bodies such as planetesimals.  Fogg \\& Nelson (2005, 2007) showed that the prevalence of shepherding vs. scattering during migration is relatively insensitive to the accretion history of the inner disk.  In the secular resonance shepherding model, two giant planets must be on interacting orbits by the late stages of the dispersal of the gaseous disk; the planets' prior orbital histories are not relevant.  Thus, we see no reason that the abundance or rarity of hot Earths should be affected by the mechanism for giant planet formation.  One other interesting difference between the core accretion and gravitational instability models is the expected location of giant planet formation.  In core accretion, there are several reasons to expect giant planets to form just past the snow line: 1) the density of solid building blocks increases by a factor of 2-4 or more (Hayashi 1981; Stevenson \\& Lunine 1988; Lodders 2003), 2) accretion timescales are shorter than anywhere else beyond the snow line (Kokubo \\& Ida 2002; Ida \\& Lin 2004), and 3) the surface density jump at the snow line, if it is steep enough, can trap inward-migrating planetary cores and form a pileup (Masset \\etal 2006b).  If these arguments hold, then core-accretion predicts that material interior to the giant planet is therefore relatively dry.  However, gravitational instability forms planets in the more distant reaches of protoplanetary disks, where the Toomre $Q$ value is lowest (Boss 1997; Mayer \\etal 2002).  Thus, icy material is included interior to the giant planet.  In the giant planet migration shepherding model, hot Earths are a mixture of material interior to the giant planet's starting orbit (Mandell \\etal 2007); they would be rocky for the core accretion model, and icy for the instability model.  Thus, transit measurements of hot Earth in systems formed by giant planet migration shepherding may provide a test to distinguish between the two dominant giant planet formation models.  As observational uncertainties of planetary orbits and masses become smaller, it will become possible to differentiate formation mechanisms based on these observations. We have laid out the qualitative differences between six different mechanisms that may form hot Earths (although some phenomena may operate simultaneously). Determining how hot Earths form is an important step toward understanding planet formation, identifying target stars for future surveys, and searching for habitable planets."
    },
    "0711/0711.4604_arXiv.txt": {
        "abstract": "Based on a new spectroscopic sample observed using the WHT, we examine the kinematic properties of the various emission line regions in narrow line Seyfert 1 galaxies (NLS1s) by modelling their profiles using multiple component fits.  We interpret these results by comparison with velocity components observed for different lines species covered in the same spectrum, and equivalent components measured in the spectra of some broad line Seyfert 1s and a representative Seyfert 2 galaxy.  We find that the fits to the H$\\alpha$ and H$\\beta$ line profiles in NLS1s require an additional broad ($\\gtrsim$3000km/s) component that might correspond to a suppressed broad line region with similar kinematics to those of typical broad line Seyfert 1s.  From the profiles of the forbidden high ionisation lines (FHILs) in NLS1s, we find evidence that they appear to trace an `intermediate' velocity region with kinematics between the standard broad and narrow line regions.  Weaker evidence of this region is also present in the profiles of the permitted Balmer lines.  Finally, we note that despite having similar ionisation potentials, the relative intensities of the highly ionised lines of [Fe X]$\\lambda$6374 and [FeXI]$\\lambda$7892 show considerable dispersion from one galaxy to another. The interpretation of this requires further modelling, but suggests the possibility of using the ratio as a diagnostic to constrain the physical conditions of the FHIL emitting region and possibly the shape of the spectral energy distribution in the vicinity of 200eV.  This spectral region is very difficult to observe directly due to photoelectric absorption both in our Galaxy and intrinsic to the source. ",
        "introduction": "Narrow line Seyfert 1 (NLS1) galaxies are a sub-set of the AGN population which share a number of characteristic properties \\citep{Pogge00}. Superficially, their optical spectra resemble those of type 2 Seyferts in that their permitted emission lines are quite narrow and only marginally wider than their forbidden lines - a property  that has unfortunately led to numerous misidentifications as Sy2s in the literature.    However, \\cite{Osterbrock85} noted that their strong permitted Fe II lines and high H$\\beta$/[O III]$\\lambda$5007 ratios were similar to the traditional type 1 Seyferts.  Although the velocity widths of the permitted lines in AGN spectra are observed to cover a continuous range ($\\sim$10$^{2}$ to $\\sim$10$^{4}$ km/s), \\cite{Goodrich89} placed a somewhat arbitrary limit of to the FWHM of H$\\beta$ of 2000km/s, for an object to be classed as a NLS1. Since then most studies have continued to use this value as the boundary between NLS1s and their broad line counterparts, here referred to as broad line  Seyferts, or BLS1s, in order to distinguish them from NLS1s.  Studies of the continuum emission from NLS1 revealed interesting features across the electromagnetic spectrum, for example, significant soft X-ray excesses, rapid X-ray variability and strong emission from high ionisation forbidden lines (FHILs) (e.g. \\citealt{Stephens89}; \\citealt{Puchnarewicz92}; \\citealt{Pogge00}). Their X-ray properties encouraged searches using X-ray surveys with follow-up confirmation from optical spectroscopy. This showed that around 50\\% of all \\textit{ROSAT} soft X-ray selected AGN were NLS1s (\\citealt{Grupe94}; \\citealt{Pogge00}).  It should be noted, however, that this method introduces serious selection effects, since \\cite{Boller96} found that not all NLS1s display strong soft X-ray excesses.  The features described above, in particular the rapid X-ray variability and lack of significantly broadened permitted emission lines has led to speculation that NLS1s may harbour a lower mass central black hole than is typical in BLS1s (10$^{7-8}$M$_{\\odot}$ and 10$^{8-9}$M$_{\\odot}$, respectively; e.g. \\citealt{Mathur01}).  This, coupled with their bolometric luminosities,  implies a high accretion rate, possibly super-Eddington suggesting that the key parameter in determining the distinctive physical properties of this class is their high ratio of $\\dot{m}/m$.  Indeed high luminosity AGN which do not fulfil the velocity width criteria of NLS1s (i.e. $<$2000km/s) but share many of their other characteristics are also thought to have high $\\dot{m}/m$, e.g. PDS456 (\\citealt{Reeves03}), consistent with the model of \\cite{McHardy06}.  It is partly for this reason that NLS1s have attracted much interest. In addition, we can potentially test the full extent of the relationships between observable parameters that have been established for the more general AGN population, such as relations between black hole mass and stellar bulge luminosity and velocity dispersion (see e.g. \\citealt{Ferrarese00}; \\citealt{Gebhardt00b}).  Also, evidence of some of the more dramatic features of AGN activity, in particular relativistic and non-relativistic outflows, are sometimes  prominent in NLS1 spectra (e.g. \\citealt{Leighly01}; \\citealt{Leighly04}), although it is unclear whether this is an indication of a real propensity toward outflows in NLS1s, or whether other spectral features simply allow these features to be more easily measured in these AGN.  This has encouraged their study as laboratories for some phenomena present in general in AGN, but exhibited in a more extreme form in the NLS1s.  The optical spectra of NLS1 provide a wealth of information. The large number of permitted and forbidden lines give  information on the kinematics and physical properties of the various emitting regions, as well as the shape of the (unobscured) ionising continuum.  For example, \\cite{Dietrich05} compared the low ionisation ($<$100eV) emission lines of NLS1 to show that the far-UV continuum of NLS1 was not significantly different from that of BLS1s, thereby indicating that the soft X-ray excess, seen in many NLS1s, is unlikely to extend to significantly lower energies approaching 100eV.  However the true extent of this continuum feature, which is likely to contribute a significant fraction of the bolometric luminosity in a large proportion of NLS1s, is still unknown.  With ionisation potentials greater than 100eV, FHILs offer the possibility of determining the true shape of the soft X-ray excess. Previous studies of the optical FHILs in NLS1s, most notably [Fe VII]$\\lambda$6087 (ionisation energy: 99eV), [Fe X]$\\lambda$6374 (235eV), [Fe XI]$\\lambda$7892 (262eV) and [Fe XIV]$\\lambda$5303 (361eV) noted that these lines were significantly blueshifted and broadened, with respect to the low-ionisation forbidden lines, suggesting the possibility of outflows between the traditional broad- and narrow- line regions (e.g. \\citealt{Erkens97}; \\citealt{Porquet99}; \\citealt{Nagao00}).  However, despite the fact they  often display broad wings and have prominent profile asymmetries, little effort has been directed towards modelling the FHILs using multiple component fits. This could yield insights into the detailed kinematics of the highly ionised gas close to the AGN.  The information gained from gas kinematics would help us to constrain the location of the FHIL emitting gas with respect to other emitting regions, and also help us obtain a better  understanding of outflows from AGN.  In this paper, we examine the optical emission line spectra of 10 nearby (z $<$ 0.1) AGN at sufficiently high spectral resolution and signal to noise to permit us to model their FHILs using multiple components.  Six of these AGN have previously been classified as NLS1s.  We also observed three BLS1s to facilitate comparison between the Seyfert groups.  In an attempt to determine whether the FHILs are affected by the obscuration of the putative dusty molecular torus, we also observed one Seyfert 2 galaxy, Mrk573, which has been shown to contain a hidden broad line region by \\cite{Nagao04}.  The high quality of our spectra also enable accurate multiple component fits to both the forbidden and permitted emission lines, allowing precise measurement of the kinematics of the broad and narrow emission line regions in NLS1s and BLS1s.  This paper is structured as follows: in \\S2 we discuss the observations and spectral reduction techniques.  We present our main results, including measurement of the FHILs, in \\S3 and in \\S4 we interpret our findings in terms of current AGN models. We conclude with a summary our principal results in the final section. ",
        "conclusions": ""
    },
    "0711/0711.3212_arXiv.txt": {
        "abstract": "Although the existence of large-scale hot gaseous halos around massive disk galaxies have been theorized for a long time, there is yet very little observational evidence. We report the {\\sl Chandra} and {\\sl XMM-Newton} grating spectral detection of \\ovii\\ and \\neix\\ K$\\alpha$ absorption lines along the sight-line of 4U~1957+11. The line absorption is consistent with the interstellar medium in origin. Attributing these line absorptions to the hot gas associated with the Galactic disk, we search for the gaseous halo around the Milky Way by comparing this sight-line with more distant ones (toward LMC X--3 and the AGN Mrk~421). We find that all the line absorptions along the LMC X-3 and Mrk~421 sight-lines are attributable to the hot gas in a thick Galactic disk, as traced by the absorption lines in the spectra of 4U~1957+11 after a Galactic latitude dependent correction. We constrain the \\ovii\\ column density through the halo to be $N_{\\rm OVII}<5\\times10^{15}~{\\rm cm^{-2}}$ (95\\% confidence limit), and conclude that the hot gas contribution to the metal line absorptions, if existing, is negligible. ",
        "introduction": "\\label{sec:intro}  Many semi-analytic calculations and numerical simulations for disk galaxy formation predict the existence of extended hot gaseous halos around massive spirals due to the accretion of the intergalactic medium. For the Milky Way, the gas temperature can be shock-heated to $\\sim10^6$ K at the virial radius ($\\sim250$ kpc; e.g., Birnboim \\& Dekel 2003; Fukugita \\& Peebles 2006), and the total mass contained in the large scale halo can be comparable with or even greater than the total baryonic mass of stars and the interstellar medium in the Galactic disk (e.g., Sommer-Larsen 2006). Clearly, an observational measurement of such an extended halo will provide an important test of galaxy formation theories.  Searching for X-ray emission from large scale halos around nearby disk galaxies has proven unsuccessful. Extraplanar X-ray emissions have indeed been routinely detected around a number of nearby spirals, but only on scales of several kpc (except for the star burst galaxies), and most likely as the result from on-going stellar feedback in galactic disks (e.g., T\\\"ullmann \\etal 2006; Li \\etal 2006). The X-ray surface brightness of a galactic halo must be very weak at large scales, at least partly because of the density square dependence of the emission and the possible low metallicity of the gas. So far, the only claimed detection of an apparent X-ray-emitting halo on a scale of $\\sim20$ kpc is around the quiescent edge-on disk galaxy NGC~5746 (Pedersen \\etal 2006).  The large-scale hot gaseous halo around the Milky Way is indirectly evidenced by the presence of high velocity clouds (HVCs), in particular, the detection of their associated \\ovi\\ line absorption (e.g., Spitzer 1956;  Sembach \\etal 2003). Because the \\ovi-bearing gas is preferentially populated at intermediate temperatures of $\\sim 3 \\times 10^5$ K where thermally unstable cooling occurs, it is believed to be produced at the interfaces between the cold/warm and the hot media. Without knowing the temperature distribution and the metallicity of the hot gas, however, it is very hard to reliably estimate its total mass.  The best way to directly measure the hot medium is through its X-ray absorption line features. Indeed, \\ovii, \\oviii, and/or \\neix\\ absorption lines consistent with zero velocity shifts ($cz\\sim0$) have been detected unambiguously toward several bright extragalactic sources (e.g., LMC~X--3, Mrk~421, and PKS~2155-304; Nicastro \\etal 2002; Fang \\etal 2003; Wang \\etal 2005; Yao \\& Wang 2007a, YW07a hereafter). But these sight-lines pass through the Galactic disk, the extended Galactic halo, and (for those AGNs) the possible warm hot intergalactic medium (WHIM) in the Local Group. Recently, it has been argued that the bulk of these absorptions can {\\sl not} be produced in the WHIM \\citep{fang06, yao07a, bre07}. Indeed, these highly ionized absorption lines have also been detected in spectra of many Galactic sources % (e.g., Yao \\& Wang 2005; Juett \\etal 2006), clearly indicating the existence of the hot gas in the Galactic disk. All existing X-ray absorption and emission data are consistent with the hot gas located in the Galactic disk with a vertical exponential scale height of $\\sim 2$ kpc (Yao \\& Wang 2005, 2007a), comparable to those inferred from the angular distributions of column densities of the \\ovi-bearing gas and pulsar dispersion measures (Savage \\etal 2003; Berkhuijsen \\etal 2006). Clearly, the constraint on the location of the hot gas depends on the modeling. For instance, \\citet{bre07} found that the absorbing gas toward AGNs are also consistent with a uniformly spherical distribution within a radius of 20 kpc. The questions here are, whether the extended Galactic halo (defined as the hot gas in the region of $>$10 kpc beyond the Galactic plane but within the Galactic virial radius) exists, and how much does it contribute to the observed highly ionized X-ray absorptions?  In this Letter, we present a differential study searching for such a large-scale hot gaseous halo around our Galaxy by comparing observations between sight-lines of 4U~1957+11 and Mrk~421, and between those of LMC~X--3 and Mrk~421. \\neix, \\ovii, and/or \\oviii\\ absorption lines at $cz\\sim0$ have been detected in spectra of all three sources observed with {\\sl Chandra} and/or {\\sl XMM-Newton X-ray Observatories}. These comparisons enable us to probe the Galactic halo by examining its contribution to the observed absorption lines, and then to estimate its total mass and/or its metallicity.  Throughout the Letter, we adopt the solar abundances from \\citet{wil00}, quote statistical errors (or upper limits) for single floating parameters at 90\\% (95\\%) confidence levels, and assume that the hot gas in both the Galactic disk and halo is in the collisional ionization equilibrium state and is approximately isothermal. We further assume that the disk gas is distributed following an exponential decay law with a vertical scale height of 2 kpc (e.g., Savage, \\etal 2003; YW07a). All the data analyses are performed with the software package XSPEC (version 11.3.2). ",
        "conclusions": "\\label{sec:dis}  In this study, we have assumed that the absorption lines observed in the spectra of 4U~1957+11 are produced in the ISM. However, the photo-ionized in- or out-flow material intrinsic to 4U~1957+11 could contaminate these lines. From the best fit to the {\\sl Chandra} spectrum in the range of 2--23 \\AA, % we estimate the source luminosity over 0.5--10 keV as $L=1.19D_{\\rm 10kpc}^2\\times10^{37}~{\\rm ergs~s^{-1}}$, where $D_{\\rm 10kpc}$ is the source distance in units of 10 kpc. A curve of growth analysis with a dispersion velocity ($v_b$) of 0 km~s$^{-1}$ gives a firm upper limit of the \\neix\\ column density to be $N_{\\rm NeIX}<6.6\\times10^{17}~{\\rm cm^{-2}}$ (assuming the solar abundance of Ne). This is equivalent to a hydrogen column density $N_{\\rm H}<1.1\\times10^{22}~{\\rm cm^{-2}}$, assuming an ionization fraction of 0.5 for \\neix. For the absorber, the ionization parameter $U=L_x/n(r)r^2$ must satisfy $\\log(U)\\leq2.5$, since the \\nex\\ K$\\alpha$ absorption line has not been detected~\\citep{kal82}. Further assuming a radial gas density distribution $n(r)\\propto(r_w/r)^2$, where $r_w$ is the location where wind launches and $r>r_w$, we obtain $r_w > 50 R_\\odot$, which is much larger than the star separation ($\\lsim6 R_\\odot$) of the system (estimated from its orbital period 9.33 hours and taking an upper limit of the X-ray compact object $M_X<16 M_\\odot$). These estimates make the in-flow scenario very unlikely. On the other hand, if the absorber is out-flow material, the lines are expected to be blue-shifted at velocities larger than the escape velocity of the system (e.g., GRO J1655--40; Miller \\etal 2006), which is 260 ${\\rm km\\ s^{-1}}$ for 4U~1957+11 (taking $M_X=10 M_\\odot$ and $r_w = 6 R_\\odot$). In contrast, the detected \\neix\\ K$\\alpha$ velocity is $0\\pm114~{\\rm km~s^{-1}}$. The above arguments suggest that the absorption lines are more likely produced in the ISM (see also Juett \\etal 2006), although the intrinsic scenario can not be completely ruled out.  We have compared the highly ionized line absorptions of \\ovii, \\oviii, and \\neix\\ observed in the spectra of 4U~1957+11 and Mrk~421, and in the spectra of LMC~X--3 and Mrk~421 to search for the large scale Galactic halo. We found that there is no significant X-ray absorption due to hot gas beyond either 4U~1957+11 or LMC~X--3. These results are consistent with our previous conclusion that bulk of the absorptions observed toward the Mrk~421 sight-line are due to the hot gas around the Galactic disk with a scale height of $\\sim2$ kpc based on a joint analysis of the absorption and emission data (YW07). Note that the upper limits obtained toward LMC~X--3 are slightly lower than those toward 4U~1957+11 (Table~\\ref{tab:results}). Some of the difference are due to the partial sampling of the Galactic disk by the latter sight-line, and the better spectral quality of LMC~X--3 also matters.  Large scale hot gaseous halos are commonly expected in modern disk formation models for massive spirals (e.g. White \\& Frenk 1991). For the Milky Way in particular, the hot gaseous halo can extend up to $r_{\\rm virial}\\sim 250$~kpc and the total mass contained is expected to be $\\sim6\\times10^{10}\\ M_\\odot$, which is currently missing in the baryon mass inventory in the Lambda Cold Dark Matter ($\\Lambda$CDM) cosmology \\citep{deh98, kly02, som06}. Recently, \\citet{mal04} assumed that the fragmentated warm clouds could be formed out of the hot gas halo and proposed a multiphase cooling scenario for the galaxy formation. This model alleviates the so-called ``over-cooling'' problem faced by the standard cooling model (e.g., Klypin \\etal 2002). They then obtained a density profile for the residual hot halo gas as a function of radius (Eq. 21 and Fig. 4 in Maller \\& Bullock 2004). Normalizing their density profile by requiring the total column density integrated from 10 to 250 kpc to be $N_{\\rm H}A_{\\rm O} < 10^{19}\\ {\\rm cm^{-2}}$ (Tabel~\\ref{tab:results}), we obtain the total mass contained in the halo as $M_{\\rm H}<1.2/A_{\\rm O}\\times10^{10}\\ M_\\odot$. Clearly, to match this upper limit with the total missing baryon mass in the MW, $A_{\\rm O}$ should be $<20\\%$ of the solar value. More recently, \\citet{han06} assumed that gaseous halos of galaxies are in a hydrostatic equilibrium state and that the gas density and the total mass density profiles are power laws, and then derived a hot halo gas density profile for the MW-like galaxy (Fig.~2 in Hansen \\& Sommer-Larsen 2006). Again, normalizing their profile we obtain the total hot gas mass as $<2.2/A_{\\rm O}\\times10^9$ M$_\\odot$, requiring $A_{\\rm O}$ to be $<3.7\\%$ of the solar value.  Low metallicity of the hot halo gas is not surprising; the halo gas is believed to be primarily accreted from the intergalactic medium that is expected to be metal poor. This would also suggest that the star formation and the associated stellar feedback via Type II supernovae (SNeII) that occurred in the Galactic disk and bulge are not important in regulating the metal content of the halo. Hot gas resulting from SNeII, if leaking into the halo (e.g., Mac Low \\& Ferrara 1999), may cool quickly enough (via adiabatic expansion and atomic line radiation) to form metal-rich cold clouds in the local environment before homogenizing with the halo gas (Wang 2007). As these clouds return back to the disk plane as Galactic fountains, their interaction with the embedding hot gas could be another mechanism of warm cloud formation as required in multiphase cooling scenarios \\citep{mal04}. The interfaces between these two media may harbor the bulk of HVCs as observed via \\ovi\\ absorptions at large distance (e.g., Sembach \\etal 2003).  In this Letter, we derive the mass and metallicity limits on the putative Galactic halo by comparing the observed absorption lines among the sight-lines of 4U~1957+11, LMC~X--3, and Mrk~421. However, the amount of absorptions observed along different AGN sight-lines could vary substantially (e.g., Bregman \\& Lloyd-Davies 2007), which may be due to additional absorption components (e.g., 3C~273; Yao \\& Wang 2007b) and/or different absorbing gas properties. More comparisons among other sight-lines are therefore needed to confirm the presented results."
    },
    "0711/0711.3481_arXiv.txt": {
        "abstract": "We present a geometrodynamical method for determining distances to orbital streams of HI gas in the Galaxy.  The method makes use of our offset from the Galactic centre and assumes that the gas comprising the stream nearly follows a planar orbit about the Galactic centre. We apply this technique to the Magellanic Stream and determine the distances to all points along it; a consistency check shows that the angular momentum is approximately constant. Applying this technique to the Large Magellanic Cloud itself gives an independent distance which agrees within its accuracy of around 10\\%.  Relaxing the demand for exact conservation of energy and angular momentum at all points along the stream allows for an increase in orbital period between the lagging end and the front end led by the Magellanic Clouds.  Similar methods are applicable to other long streams of high-velocity clouds, provided they also nearly follow planar orbits; these would allow otherwise unknown distances to be determined. ",
        "introduction": "\\label{sec:intro}  High-velocity clouds (HVCs) are defined as HI gas which have anomalous velocities that are incompatible with what is expected from a simple Galactic rotation model.  Ever since their discovery by \\cite{HVCdiscovery}, their role in the formation and evolution of our Galaxy has been a source of speculation.  Amongst the attempts to explain their origin, the two most well-known and contested are the Galactic Fountain model, in which gas is either being thrown out of the Galactic disk from supernovae or being heated and ionized by massive stars, and the alternative model in which the clouds are pockets of neutral hydrogen left over from the formation epoch of the Galaxy, and which are only now being accreted.  Given the vast sky-coverage of HVCs, the two theories are not necessarily contradictory, as it is likely that some clouds belong to each model.  However, with the exception of those HVCs that are obvious components of a stream with a known progenitor, such as the Magellanic \\citep{1974ApJ...190..291M} and Sagittarius streams \\citep{2004ApJ...603L..77P}, the lack of distance measurements has hindered the progress in understanding the role that HVCs play in the Galaxy.  Since newly accreted gas would have a lower metallicity than recycled gas from the Galaxy, metallicity measurements, if available, should be able to distinguish the applicability of each theory on a cloud-by-cloud basis \\citep[e.g.][]{2002ApJ...572..178S,2000ASPC..218..407V,2001ApJS..136..463W}, but it would still be instructive to know exactly where these clouds lie.  Observationally, the difficulty in determining direct distances arises from the fact that most of the HVCs do not host stars.  If, however, there is one star behind and another to the front of the cloud along the line of sight, both of which have known distances, then a distance bracket can be set for the HVC.  The high resolution spectrum of a star directly behind the HVC will exhibit absorption lines at the velocity of that cloud; features that will be absent in the spectrum of a star to the front of the cloud.  Starting with the first such measurement for Complex M in the early 90's \\citep{1993ApJ...416L..29D}, other HVC complexes for which a distance bracket that now exists include Complex A \\citep{2001ApJS..136..463W} and Complex WB \\citep{2006ApJ...638L..97T}.  The high latitudes of much of the neutral hydrogen gas in question means there is a lack of background stars bright enough for high resolution spectroscopy, so this technique has not been without difficulty.  Many detailed simulations of streams such as the Magellanic Stream exist \\citep[e.g.][]{2006MNRAS.371..108C,2005MNRAS.363..509M,1996MNRAS.278..191G, 1985IAUS..106..471F,1976A&A....47..263F}, however, all such simulations have many parameters which must be tweaked to give the best fits to the data.  In this paper, we describe an alternative set of methods with the aim of probing the elusive distances to streams of HI gas comprising the HVCs.  These methods instead make strong hypotheses that enable us to attempt to read the distances out of the data themselves; the studies presented here are neither simulations nor attempts at modelling.  The following section introduces the geometrodynamical methods for determining distances to streams of HVCs and discusses which streams are appropriate for each method.  We developed the methods in response to a challenge laid down by Majewski, who asked whether it was possible to determine distances from radial velocities alone, when those were known along a stream \\citep{2004ASPC..327...63M}.  In Section \\ref{subsec:masterpoints}, we consider distant streams and use heliocentric radial velocities corrected to the Galactic Standard of Rest (GSR) as a lowest order approximation to the Galactocentric radial velocities. Section \\ref{sec:velcor} shows how these radial velocities can be corrected iteratively, and shows that the corrections are negligible for the Magellanic Stream.  The choice of Galactic potential is briefly discussed in Section \\ref{sec:potential}.  On the assumption that the Magellanic Stream is composed of material that was torn off the Magellanic Clouds at an earlier passage through perigalacticon, it seems reasonable to assume an angular momentum gradient to be present along the length of the stream; this topic is addressed in Section \\ref{sec:deltah}.  We present our summary and conclusions in Section \\ref{sec:sum}.  There is a great body of literature on streams and associations of bodies in planes through the Galactic centre.  The Sagittarius stream \\citep{2001ApJ...547L.133I,2003AJ....125.1352P} features prominently in this literature, however it is not relevant to the methods developed here which work off the parallax due to the Sun's displacement from the plane of the orbit.  This displacement is almost zero for the Sagittarius stream.  Some frequently used symbols and their definitions are summarised in Table \\ref{tab:defs} with a schematic guide to a generic point on the stream given by Figure \\ref{fig:schematic}.   \\begin{table*} \\caption{ Definition of commonly used symbols, where hats denote unit vectors. \\label{tab:defs} } \\begin{tabular}{ccl} \\hline Symbol & \\multicolumn{1}{c}{Equation where first used} & \\multicolumn{1}{c}{Description} \\\\ \\hline $\\norms$, $\\normsigma$ & (\\ref{eqn:theta}),  (\\ref{eqn:v}) & Apparent and true direction vectors for a particular point on the stream\\\\ $\\normn$ & (\\ref{eqn:theta}) & Vector normal to the stream plane\\\\ $\\norml$, $\\normr$ & (\\ref{eqn:theta}), (\\ref{eqn:d}) & Unit vectors to a stream point, from the Sun and the  Galactic centre\\\\ $v_l$ & (\\ref{eqn:v}) & Heliocentric line-of-sight velocity corrected to the Galactic System of Rest\\\\ $\\vecrsun$ & (\\ref{eqn:d}) & Vector joining Galactic centre to Sun $=(-\\rsun,0,0)$ \\\\ $d$, $r$ & (\\ref{eqn:d}), (\\ref{eqn:energy}) & Distance to a specific point on the stream from the Sun and from the Galactic centre\\\\ ${\\bf{h}}$, $\\varepsilon$ & (\\ref{eqn:am}), (\\ref{eqn:energy}) & Specific angular momentum and energy associated with the stream \\\\ $D$ & (\\ref{eqn:norm}) & Ratio of distances to the $Z$ point and the Galactic centre from the Sun \\\\ $\\psi$ & (\\ref{eqn:energy}) & Galactic potential \\\\ $v_c$ & (\\ref{eqn:rH}) & Galactic circular velocity at $\\rsun$ (unless specified otherwise)\\\\ \\hline \\end{tabular} \\end{table*}  \\begin{figure} \\hspace{-10pt} \\includegraphics[width=0.5\\textwidth]{schematic.eps} \\caption{An orbital stream lying in a plane through the Galactic centre.  The unit vector ${\\rm {\\bm{\\hat{\\sigma}}}}$ lies along the stream at a general point $P$, which is seen from the Sun in the direction ${\\bf{\\hat{l}}}$ at a distance $d$.  Its observed radial velocity, corrected to the Galactic Standard of Rest, is $v_l$ = ${\\bf{v}}.{\\bm{\\hat{l}}}$; at the position labelled $Z$, $v_l$ = 0. The unit vector $\\norms$ is along the projection of $\\normsigma$ onto the sky, so $\\norms.\\norml = 0$ and $\\norms$ is $\\normsigma - (\\normsigma.\\norml)\\norml$ normalised to unity.  At $Z$, $\\norms$ and $\\normsigma$ coincide.} \\label{fig:schematic} \\end{figure} ",
        "conclusions": "\\label{sec:sum}  We have applied a variety of geometrodynamical methods to streams of high-velocity clouds in an attempt to constrain their distances.  Some of these methods, including the near method, have the nice feature that distance determinations would be possible without the parameter dependency inherent in any simulations.  We find, however, that the near method in particular becomes rather inaccurate when the angle subtended at the stream by $\\vecrsun$ becomes small.  At the expense of becoming more model dependent by assuming a halo form for the Galactic potential, we have improved on this accuracy.  We then find distances to the entire length of the Magellanic Stream by assuming that the gas comprising the stream nearly follows a planar orbit about the Galactic centre.  We find the perigalacticon to be at $45\\kpc$ and the maximum distance at the trailing end to be $75\\kpc$.  The heliocentric distances are not very different from the Galactocentric distances, varying typically by $1\\kpc$ or less.  The centre-of-mass of the Magellanic Clouds would be at $47\\pm5\\kpc$ from us, according to this method; this is consistent with various observationally determined values for the Large Magellanic Cloud (see \\citet{2004NewAR..48..659A} for a review).  Application of our methods to Complex A and C did not yield very promising results under the assumptions that were made.  This indicates that our hypothesis, namely that the streams were nearly following an orbit around the Galactic centre in a well-defined plane, was not valid for these two streams of high-velocity clouds.  In this paper, we have tested the idea that streams of HVCs follow orbits.  However the assumption that the orbit lies in a plane that passes through the Galactic centre is rather restrictive when applied to nearby streams that do not pass close to the Galactic Pole. It is possible to remove the restriction of a planar orbit, provided a model of the Galaxy's potential is given. We find that the transverse velocity of the stream can be deduced from the run of radial velocities as a function of angle along the stream, provided the distance to one point of the stream is assumed. Orbits can then be computed in the given potential. Both directions in the sky and radial velocities are then computed and these are compared with the observations.  Such a method and an example of its application to Complex A and the stellar Orphan stream has already been presented \\citep{2007MNRAS.378L..64J}.  We note that such methods of deriving the transverse velocity component were used in a related problem by \\cite{1977A&A....57..265F}, but probably date back to an earlier century.  We are also exploring the possible existence of stream associations.  Following on from these works, we hope, eventually, to give a definitive answer to the question of whether high-velocity cloud streams follow orbits around the Galaxy.   \\subsection*{Acknowledgements} The authors thank B. Wakker and R.~D.~Davies for the kind use of their data for Complex A, C and the Magellanic Stream and the anonymous referee for helpful comments.  SJ acknowledges financial support from PPARC."
    },
    "0711/0711.1167_arXiv.txt": {
        "abstract": "We report on the detection of reddening toward $z\\sim 0.2$ galaxy clusters. This is measured by correlating the {\\it Sloan Digital Sky Survey} cluster and quasar catalogs and by comparing the photometric and spectroscopic properties of quasars behind the clusters to those in the field. We find mean $E(B-V)$ values of a few$\\times 10^{-3}$\\,mag for sight lines passing $\\sim$Mpc from the clusters' center. The reddening curve is typical of dust but cannot be used to distinguish between different dust types. The radial dependence of the extinction is shallow near the cluster center suggesting that most of the detected dust lies at the outskirts of the clusters.  Gravitational magnification of background $z\\sim 1.7$ sources seen on Mpc (projected) scales around the clusters is found to be of order a few per cent, in qualitative agreement with theoretical predictions.  Contamination by different spectral properties of the lensed quasar population is unlikely but cannot be excluded. ",
        "introduction": "Zwicky (1957) suggested the presence of dust  in the intracluster medium (ICM) to explain galaxy number counts behind the Coma cluster and a search has begun to quantify its properties. Similar analyses were carried out by Karachentsev \\& Lipovetskii (1969), Bogart \\& Wagoner (1973), and Boyle et al. (1988 using quasars)  and yielded B-band extinctions of order 0.2\\,mag.  However, a more recent study by Maoz (1993) using radio-selected quasars behind rich Abell clusters yielded only upper-limits on the reddening [$E(B-V)<0.05$\\,mag] with similar limits [$E(B-V)< 0.02$\\,mag] obtained from the analysis of galaxy colors behind APM clusters (e.g., Nollenberg et al. 2003 see also Ferguson 1993). In contrast, Hu (1992) found $E(B-V)\\sim 0.2$\\,mag toward cooling flow clusters (but see Annis \\& Jewitt 1993).  Dust emission from clusters has been marginally detected in the infrared  resulting in a dust-to-gas ratio as low as $\\sim 10^{-6}$ in a few Abell clusters (Stickel et al.  2002; but see Bai et al. 2007).  Cluster members are known to contain dust (e.g., Edge et al. 1999, Popescu et al.  2002, Boselli et al. 2004) and galaxies are seen to shed gas and dust into their environments (e.g., Heckman et al. 2000; see also Jonsson et al. 2006). However, the survival time for small dust grains mixed with the ICM is much shorter than Hubble time (Draine \\& Salpeter 1979), and would be absent from the cluster core (but could survive on large scales; see Dwek et al. 1990) unless constantly replenished (e.g., Montier \\& Giard 2004) or  shielded from the hot plasma (e.g., Voit \\& Donahue 1995).   Presently, there are few meaningful constraints on the composition and quantity of intracluster dust, an understanding of which would have important implications for cluster evolution and metal enrichment processes, as well as for dust physics. In this {\\it Letter} we provide new answers to 50 years old questions concerning dust in galaxy clusters, by harnessing the power of the {\\it Sloan Digital Sky Survey} (SDSS). ",
        "conclusions": "We deduce finite extinction  for quasars behind clusters  and our results are consistent with a large covering factor of dust toward these objects. The dust mass enclosed within a sphere of radius $b$ is \\begin{equation} M_{\\rm dust} \\sim 10^{9} \\left ( \\frac{b}{{\\rm Mpc}} \\right )^2 \\frac{\\left < E(g-i) \\right >}{10^{-2}\\,{\\rm mag}} ~{\\rm M_\\odot}. \\end{equation} (Here we used the Kramers-Kronig relations and grain density of $2.5\\,{\\rm g~cm^{-3}}$; Kr\\\"ugel et al. 2003). This  implies a dust-to-gas ratio for the ICM which is $<5\\%$ of local inter-stellar medium (ISM; here we used ICM mass of $4\\times 10^{13}\\,{\\rm M_\\odot}$ within the central Mpc; Morandi et al. 2007).  If, as suggested by some theoretical models, galaxies eject their ISM in amounts comparable to their own mass during their formation, then our results imply that either the associated dust is in a clumpy form with a small covering factor or else it is rapidly destroyed and the metals are deposited in the ICM. Large scale distributions of dust could arise from a non-virialized gas component at the outskirts of clusters left from past merger events or related to halos of neighboring galaxies and small groups provided their size is of Mpc scales. Singular isothermal sphere (SIS) models, where dust follows dark matter, seem to over-predict the extinction on small  scales (see Figs. \\ref{photo} \\& \\ref{red_mag}) indicating that the centers of clusters may be relatively devoid of ISM-like dust.  Our results indicate that foreground emission is not important.  This means that the contribution of foreground light that enters a (median) quasar's  point-spread function is $>7$\\,mags fainter in the $z$-band (Fig.1 and appendix A) and requires that most sight lines pass $>10$ (de-Vaucouleur) half-light radii ($\\sim15$\\,kpc; see Fig.  4 in Bernardi et al. 2003) from a median cluster member.  Repeating the entire analysis for quasars whose sight lines pass $>30$\\,kpc from confirmed members  yields similar results. We estimate foreground light contamination by resolved and non-resolved cluster members (using a luminosity function with a slope of $-0.8$ extrapolated down to $\\sim 0.04L^\\star$, and including 25\\% of blue galaxies; Hansen et al. 2007) to be at the per cent level, and hence negligible.  The observed magnification was found to be few per cent on Mpc scales. Correcting for luminosity function effects (see appendix B) the intrinsic magnification is of order 10\\%, which is in qualitative agreement with SIS model predictions for the magnification of  a $z= 1.7$ quasar behind a rich $z= 0.2$ cluster being $\\mu-1\\sim 0.2 (b/{\\rm Mpc})^{-1}$ (e.g., M\\'enard et al. 2007). Given the uncertainties on the dark-matter profile, our approximation for the SDSS selection function and the luminosity function of quasars, we find this agreement to be surprisingly good. It also implies that grey opacity  (by large grains that survived for a Hubble time at the cluster center) is not required.  The above agreement, the fact that the deduced reddening and magnification follow different trends at small $\\tilde{b}$ (Figs. \\ref{photo} \\& \\ref{red_mag}), and the good fit to known reddening curves, all suggest that our results are not significantly contaminated by luminosity-dependent intrinsic color differences in the lensed quasar population. The agreement between the photometric and spectroscopic results, which are prone to different systematics, is also reassuring."
    },
    "0711/0711.0357_arXiv.txt": {
        "abstract": "We present observations of transient active region heating events observed with the Extreme Ultraviolet Imaging Spectrometer (EIS) and X-ray Telescope (XRT) on Hinode. This initial investigation focuses on NOAA active region 10940 as observed by Hinode on February 1, 2007 between 12 and 19 UT.  In these observations we find numerous examples of transient heating events within the active region. The high spatial resolution and broad temperature coverage of these instruments allows us to track the evolution of coronal plasma. The evolution of the emission observed with XRT and EIS during these events is generally consistent with loops that have been heated and are cooling.  We have analyzed the most energetic heating event observed during this period, a small GOES B-class flare, in some detail and present some of the spectral signatures of the event, such as relative Doppler shifts at one of the loop footpoints and enhanced line widths during the rise phase of the event.  While the analysis of these transient events has the potential to yield insights into the coronal heating mechanism, these observations do not rule out the possibility that there is a strong steady heating level in the active region. Detailed statistical analysis will be required to address this question definitively. ",
        "introduction": "Understanding how the solar corona is heated to high temperatures is one of the principal objectives of the Hinode mission \\citep{kosugi2007}. The Extreme Ultraviolet Imaging Spectrometer (EIS) on Hinode has an unprecedented combination of high spatial, spectral, and temporal resolution and provides a wealth of information on the temperatures, densities, velocities, and nonthermal motions in the solar corona. The X-ray Telescope (XRT) on Hinode is one of the highest spatial resolution soft X-ray telescopes ever built, and combines this high spatial resolution with high temporal resolution and a broad temperature coverage. Together, these instruments give us an unparalleled ability to probe the morphology and evolution of coronal plasma (see, \\cite{culhane2007} and \\cite{golub2007}, for more details on the instruments).  One question of particular importance is the time scale for heating in solar active regions. Recent observations have yielded contradictory results. At high temperatures, observations at soft X-ray wavelengths have suggested that the emission is relatively steady and can be described with steady heating models. The analysis of individual loops (e.g., \\cite{porter1995,kano1996}) has shown that the evolution of the intensity is generally much slower than the conductive or radiative cooling times and that the loop emission is consistent with the RTV scaling laws \\citep{rosner1978}. Transient brightenings have been observed in active regions previously (e.g., \\cite{shimizu1992}), but they do not appear to provide sufficient heating (e.g., \\cite{shimizu1995}). The application of steady heating models to entire active regions (e.g., \\cite{warren2006b}), coronal bright points \\citep{brooks2007}, and the full Sun \\citep{schrijver2004} have successfully reproduced the observed high temperature emission.  \\begin{figure*}[t!] \\centerline{% \\FigureFile(6.4in,3.2in){warren_fig1.small.ps}} \\caption{XRT observations of AR 10940 on February 2, 2007. On the left is a composite full disk image. On the right is an example of a $512\\times512$ pixel cutout. The field of view for the EIS context and sit-and-stare observations is indicated. A movie is provided as on-line material.} \\label{fig:1} \\end{figure*}  At lower temperatures, observations of active regions at temperatures near 1\\,MK suggest that coronal plasma is far from equilibrium. The analysis of coronal loops at these temperatures indicate very large densities relative to the RTV scaling laws (e.g., \\cite{winebarger2003,aschwanden2001b,lenz1999}). Full active region modeling shows that while the steady heating models can reproduce the high temperature emission, these models cannot account for the presence of cooler active region loops \\citep{warren2006b}. Because loops cool more rapidly than they drain (e.g., \\cite{cargill1995}), impulsive heating models can account for the bright active region loop emission observed at these temperatures (e.g., \\cite{warren2002b,warren2003}).  Examples of active region loops cooling from high temperatures down through 1\\,MK have been documented (e.g., \\cite{winebarger2005,urra2006}).  The different heating time scales suggested by the observations at different temperatures can be resolved in a number of ways. It may be that there is a different heating mechanism responsible for the plasma at different temperatures. Alternatively, it may be that high temperature emission is more dynamic than previously thought and the modest spatial resolution of previous instruments has obscured this. The purpose of this paper is to make an initial survey of some EIS and XRT active region observations and to explore the available diagnostics. ",
        "conclusions": "We have presented an initial look at active region transient brightenings observed with the EIS and XRT instruments on Hinode. The high spatial and temporal resolution of these instruments indicate that high temperature active region emission can be highly dynamic, suggesting that impulsive heating is more important than previously thought. The fine temperature resolution of EIS allows us to follow the evolution of active region plasma in great detail. We see numerous examples of plasma cooling from high temperatures to lower temperatures and have presented one example in some detail. To our knowledge this is the most complete observation of an active region transient to date. The very broad temperature coverage of EIS allows us to follow the temperature evolution from above 10\\,MK to below 0.1\\,MK. The high spectral resolution of EIS also allows us to identify possible velocity signatures of active region heating in one event. We find what appear to be bi-directional flows at the loop footpoints and excess line widths during the rise phase of the event.  These results, however, are preliminary. They generally come from the analysis of a single event from a single active region. Our impression is that this event is qualitatively similar to the other, smaller events that are observed in this region. Many more systematic studies of active region transient activity with Hinode are clearly required.  While the analysis of transient events observed with XRT and EIS has the potential to yield insights into the coronal heating mechanism, it is also possible that there is a strong steady heating level that contributes to the active region heating. The light curves presented in Figure~\\ref{fig:2} clearly show a high basal level of several thousand DN~s$^{-1}$. Differentiating between steady heating with transient events superimposed and purely impulsive heating will require detailed statistical analysis and modeling of the observed light curves (e.g., \\cite{shimizu1992}).  The heating of the solar corona is undoubtedly related to the evolution of surface magnetic fields and incorporating the analysis of magnetic field data and chromospheric emission from the SOT on Hinode is also very important. We have surveyed the SOT data from this period and we see a wealth of activity at these heights in the solar atmosphere. Much of the photospheric flux, however, closes before it reaches the corona \\citep{close2003} so the challenge is to determine what activity is relevant to the coronal dynamics we observe.  \\vspace{0.1in}  Hinode is a Japanese mission developed and launched by ISAS/JAXA, with NAOJ as domestic partner and NASA and STFC (UK) as international partners. It is operated by these agencies in co-operation with ESA and NSC (Norway). We thank Spiros Patsourakos for preparing the STEREO EUVI data for us. We also thank the referee for a thorough reading of the manuscript and many insightful comments."
    },
    "0711/0711.0482_arXiv.txt": {
        "abstract": "In 2003 a previously unpulsed Einstein and ROSAT source cataloged as soft and dim (L$_X$ of few $\\times$10$^{33}\\ergs$) thermal emitting object, namely \\xte, was identified as the first unambiguous transient Anomalous X-ray Pulsar. Two years later this source was also found to be a bright highly polarized transient radio pulsar, a unique property among both AXPs and radio pulsars. In September 2006 \\swift\\ Burst Alert Telescope (BAT) detected an intense burst from the candidate AXP \\cxou, which entered in an outburst state reaching a peak emission of at least a factor of 300 higher than quiescence. Here, we briefly outline the recent results concerning the outburst phenomena observed in these two AXPs. In particular, \\xte\\ has probed to be a unique laboratory to monitor the timing and spectral properties of a cooling/fading AXP, while new important information have been inferred from X-ray and radio band simultaneous observations. \\cxou\\ has been monitored in X-rays and radio bands since the very beginning of its outbursting state allowing us to cover the first phases of the outburst and to study the timing and spectral behavior during the first months. ",
        "introduction": "At the beginning of the X--ray astronomy era the study of the X-ray variability (both of the flux and timing properties) of Cen X--3 allowed astronomers to unambiguously assess the binary nature of the source and to identify the accretion of mass, flowing from the companion onto a rotating neutron star, as the main mechanism to produce X--rays \\citep{cenx-3a,cenx-3b}. Since then, several classes of high energy sources hosting a compact object have been identified and their variability, observed on different timescales, used to test models and/or to study the emission mechanisms as a function of flux (while keeping fixed other parameters, such as the distance and the geometry of the system). For the class of accreting neutron stars this approach was successfull in confirming, among others, the presence of a centrifugal barrier to accretion, testing the dependence of the period derivative with respect to the source luminosity, and studying the change in the photon propagation direction as a function of the accretion rate by means of the pulse shape changes \\citep[fan  and pencil beam model; as an example see ][]{exo1,exo2,lorella89}. Isolated neutron stars are relatively constant sources, making the above approach unreliable. However, there are two small classes of isolated neutron stars which show spectacular events during which their luminosity may change up to 10 orders of magnitude on timescales down to few milliseconds. These objects are better known as Anomalous X--ray Pulsars (AXPs; 10 objects plus 1 candidate) and Soft $\\gamma$-ray Repeaters \\citep[SGRs; 4 objects plus 2 candidates; for a review see][]{wot06}. It is believed that AXPs and SGRs are linked at some level, owing to their similar timing properties (spin periods in the 2-12\\,s range and  period derivatives \\.P in the  10$^{-13}$--10$^{-11}$ s\\,s$^{-1}$ range). Both classes  have been proposed to host neutron stars whose emission is powered by the decay of their extremely strong magnetic fields \\citep[$>10^{15}$\\,G;][]{DT92, TD95}.  Different types of  X-ray flux variability have been displayed by AXPs. From slow and moderate flux changes (up to a factor of few) on timescales of years (virtually all the object of the class),  to moderate-intense outbursts (flux variations of a factor up to 10) lasting for 1-3 years (\\ea, and \\ee),  to dramatic and intense SGR-like burst activity (fluence of $10^{36}-10^{37}$ ergs) on sub-second timescales \\citep[\\uu, \\xte, \\ea\\ and \\ee; see ][for a review on the X-ray variability]{2007London}. The first notable recorded case of flux variability was the 2002 bursting/outbursting event detected from \\ea, the only known event in which a factor of $\\sim$10 persistent flux enhancement in an AXP was followed (or proceeded) by the onset of a bursting activity phase during which the source displayed more than 80 short bursts \\cite{fotis04,pete04}. The timing and spectral properties of the sources changed significantly and recovered the pre-bursting activity phase values within few days, likely due to the relatively high luminosity DC level ($\\sim 10^{35}\\ergs$). However, it was only in 2003 that the first transient AXP was discovered, namely \\xte, which displayed a factor of $>$100  persistent flux enhancement with respect to the unpulsed pre-outburst quiescent luminosity level \\citep[$\\sim 10^{33}\\ergs$; ][]{alaa04,eric04,ioxtej1810,nanda04}. Unfortunately, the initial phases of the outburst were missed and we do not know whether a bursting activity phase, similar to that of \\ea, occurred also for this source, though four bursts have been detected by RossiXTE between 2003 September and 2004 April and unambiguously associated with \\xte\\ \\cite{pete05}. In 2006 the candidate AXP \\cxou\\ displayed a bursting-outbursting behavior with a maximum flux variability of $>300$, followed by extreme and daily changes both in the spectral and timing properties \\cite{mike06a,mike06b,iocxou}. \\begin{figure}[thb] \\includegraphics[width=1.\\textwidth]{GL-Israel_fig1} \\caption{Evolution of the spectral parameters for the 2BB (left panel) and 3BB (right panel) models together with the pulsed fraction as a function of the observed 1-10\\,keV band flux. Central Panel: schematic view of the assumed scenario with marked the emission regions of the hottest BB components.} \\label{fig:XTE} \\end{figure}  These two sources currently represent our best opportunity in order to study the evolution of the main spectral and timing parameters as a function of flux by keeping fixed other parameters (such as distance and geometrical angles) otherwise difficult to infer (similarly to the pioneering studies on accreting X--ray pulsars). In the following pages we will briefly outline (i) the recent results obtained from the analysis of the latest 4 years \\XMM\\ monitoring observations of \\xte\\ as it approached to quiescence, (ii) the comparison of X-ray and radio emission from \\xte\\ by means of two $\\sim$8hr-long simultaneous campaigns with \\XMM\\ and Parkes, (iii) and the results of the first 6 months monitoring of \\cxou\\ during the first phases of its current outburst by means of \\swift, \\CXO\\ and \\XMM\\ in the X-rays and Parkes in the radio band. ",
        "conclusions": ""
    },
    "0711/0711.2955_arXiv.txt": {
        "abstract": "The contact eclipsing binary system XX Leonis (p = 0.97 days, Sp A9) has been analysed using the PHOEBE programme, based on the Wilson-Devinney code. The BVRI light curves were obtained during spring 2006 using the 20-cm telescope and ST-7 CCD detector. The effective temperature of the primary component determined from the photometric analysis is $T=(7342\\pm14)\\,K$, the inclination of the orbit is $i=(84.83\\pm0.29)^{\\circ}$ and the photometric mass ratio $q=(0.40\\pm0.01)$. Also the third body hypothesis was suggested, based on the period analysis using 57 minimum times and resulting the period of the third body $p_3=(59.66\\pm0.05)\\,yr$, amplitude $A=(0.036\\pm0.028)\\,day$ and zero eccentricity which gives the minimum mass $M_{3,min}=(0.91\\pm0.01)\\,M_{\\odot}$. ",
        "introduction": "The binary system XX Leonis (SV* P 3370, AN 355.1934, BD+14 2177) was discovered to be a variable by \\cite{Tsesevich1954}. It is an eclipsing binary star of $11^{th}$ magnitude, the period is close to one day and the depth of primary minimum is about 0.5 mag. The spectral type of the primary is A9V and of secondary K0. It was included in a catalogue of apsidal motion binaries by \\cite{Heged1988}, but this hypothesis was rejected. ",
        "conclusions": "\\label{sec:concl}  We have performed new analysis of the contact binary of the Beta-Lyrae type XX Leo. New light curves were measured and analysed. Also new period study has been done. The hypothesis of the third body was suggested, according to the third light in the LC solution and also from the light-time effect hypothesis. Both hypothesis are in very good agreement and leads to the third body comparable to the secondary.  \\begin{table}\\def~{\\hphantom{0}} \\begin{center} \\caption{Parameters of the third component.} \\label{tab:LITE} \\begin{tabular}{cc}\\hline Parameter &  Value \\\\\\hline% $p_3 [yr]$ & $59.66 \\pm 0.05$ \\\\ $A [d]$    & $0.036 \\pm 0.028$ \\\\ $\\omega [^\\circ]$ & $19.22 \\pm 12.98$ \\\\ $e$ & 0 (fixed) \\\\ $T_0 [JD]$ & $2437200.62 \\pm 216.63$\\\\\\hline \\end{tabular} \\end{center} \\end{table}"
    },
    "0711/0711.0211_arXiv.txt": {
        "abstract": "We use a simple optical/infrared (IR) photometric selection of high-redshift QSOs that identifies a Lyman Break in the optical photometry and requires a red IR color to distinguish QSOs from common interlopers.  The search yields 100 $z\\sim3$ ($U$-dropout) QSO candidates with $19<r'<22$ over 11.7 deg$^2$ in the ELAIS-N1 (EN1) and ELAIS-N2 (EN2) fields of the {\\it Spitzer} Wide-area Infrared Extragalactic (SWIRE) Legacy Survey.  The $z\\sim3$ selection is reliable, with spectroscopic follow-up of 10 candidates confirming they are all QSOs at $2.83<z<3.44$.  We find that our $z\\sim4$ ($g'$-dropout) sample suffers from both unreliability and incompleteness but present 7 previously unidentified QSOs at $3.50<z<3.89$.  Detailed simulations show our $z\\sim3$ completeness to be $\\sim80-90\\%$ from $3.0<z<3.5$, significantly better than the $\\sim30-80\\%$ completeness of the SDSS at these redshifts.  The resulting luminosity function extends two magnitudes fainter than SDSS and has a faint end slope of $\\beta=-1.42\\pm0.15$, consistent with values measured at lower redshift.  Therefore, we see no evidence for evolution of the faint end slope of the QSO luminosity function.  Including the SDSS QSO sample, we have now directly measured the space density of QSOs responsible for $\\sim70$\\% of the QSO UV luminosity density at $z\\sim3$. We derive a maximum rate of HI photoionization from QSOs at $z\\sim3.2$, $\\Gamma = 4.8\\times10^{-13}$ s$^{-1}$, about half of the total rate inferred through studies of the Ly$\\alpha$ forest.  Therefore, star-forming galaxies and QSOs must contribute comparably to the photoionization of HI in the intergalactic medium at $z\\sim3$. ",
        "introduction": "The QSO luminosity function (QLF) is an observable constraint on models of galaxy formation and the corresponding growth of super-massive black holes (SMBHs) \\citep[eg.][]{small92,kauffmann00,haiman00}.  These models are useful in interpreting observed phenomena, such as the relation between a galaxy's black hole mass and bulge luminosity \\citep{kormendy95,magorrian98}, as well as inferring specifics of QSO activity, such as initial black hole mass functions, QSO light curves, and accretion rates \\citep[see eg.][]{hopkins06}.  In addition to galaxy formation models, the QLF can be used to derive the QSOs' contribution to HI and HeII reionization.   QSOs are responsible for HeII reionization at $z\\sim3$ \\citep{jakobsen94,reimers97,hogan97,sokasian02} and are presumed to have a neglible contribution to the HI reionization at $z\\sim6$ \\citep{madau99}.  However, both of these claims require assumptions about the faint end slope of the QLF, as this has not been measured accurately at $z>2$.  The first QSO luminosity functions demonstrated a rapid increase in space densities toward higher redshift \\citep{schmidt68,schmidt83} .  Deeper surveys, which primarily identified QSOs by their ``UV-excess''\\footnote{QSOs were most easily identified as point sources with $U-B$ colors bluer than most stars.}, found that the faint end of the QLF was shallower than the bright end \\citep{boyle88, heisler88, koo88, hartwick90}.  Recent large surveys, most notably the 2dF Quasar Redshift Survey \\citep[2QZ,][]{boyle00,croom04} have found thousands of $z<2.5$ QSOs.  With these large samples, the QSOs have been placed into smaller bins in both luminosity and redshift, accurately constraining the shape and evolution of the QLF.  The data are typically fit to a broken power-law \\citep{boyle88,pei95}  \\begin{equation} \\Phi(L,z) = \\frac{\\Phi(L^*)/L^*}{(L/L^*)^{-\\alpha}+(L/L^*)^{-\\beta}}, \\label{eqn:qlf_lum} \\end{equation}  {\\noindent} with the break at $L^*$ and a bright end slope, $\\alpha$, steeper than the faint end slope, $\\beta$.  The QLF evolution with redshift is consistent with Pure Luminosity Evolution (PLE) \\citep{marshall83,marshall85,boyle88,hartwick90,boyle00,croom04,richards05}.  That is, the evolution can be parametrized by a shift in the luminosity of the break, $L^*(z)$, without any change in its shape.  Recently, the 2dF-SDSS LRG and QSO Survey \\citep[2SLAQ,][]{richards05} has extended the ``UV-excess'' QSO search one magnitude fainter to more accurately measure the faint end slope of the QLF out to $z=2.1$.  These deeper data, when combined with the bright end slope from the 2QZ and 6QZ \\citep{croom04}, fit a faint-end slope, $\\beta=-1.45$, and bright-end slope, $\\alpha=-3.31$, and demonstrates a $\\sim$40-fold increase in $L^*$ from $z=0$ to $z=2$.  At high redshift, various surveys have been conducted through grism searches for UV emission lines \\citep{schmidt95}, spectroscopic follow-up of point sources with optical colors away from the stellar main sequence \\citep{warren94,kennefick95,fan01}, X-ray \\citep{hasinger05}, searches for radio-loud \\citep{dunlop90} or infrared luminous QSOs \\citep{brown06}.  These surveys all show a precipitous decrease in space densities at $z>3$.  Early data from the Sloan Digital Sky Survey \\citep[SDSS,][]{york00}, \\citet{fan01} show a factor of six decrease in $M<-25.5$ QSOs from $z=3.5$ to $z=6.0$.  Recent results from SDSS suggest that the bright end slope is no longer constant at $z>3$.  Rather, it is getting shallower toward higher redshift (Richards et al. 2006). Unfortunately, these high redshift surveys are shallow ($i \\lesssim 20$) and can only measure the bright end of the QLF at $z>3$.  Therefore, little can be said about the shape of the faint end of the high-z QLF, or its integrated properties (eg. contribution to intergalactic HI, HeII ionizing radiation or black hole growth) without a census of fainter QSOs at $z>3$.  \\citet{hunt04} (hereafter, H04) searched for faint AGN at $z\\sim3$ with deep Keck spectroscopy over 0.43 deg$^2$ and found 11 QSOs.  Though limited by small numbers, the fitted faint-end slope, $\\beta=-1.24 \\pm 0.07$, is substantially shallower than low redshift measurements.  X-ray selected samples suggest a rather modest evolution in AGN space density at high-z \\citep{barger05}, or even a luminosity dependent density evolution (LDDE) where lower luminosity AGN peak in number density at lower redshifts than more luminous QSOs \\citep{ueda03,hasinger05}.  Given the importance of the QSO luminosity function in constraining models of galaxy and black hole formation as well as the contribution of AGN to the ionizing background, these initial indications that the QLF shape is evolving at $z>3$ warrant deeper surveys to better constrain the high redshift QLF.  The SWIRE Legacy Survey \\citep{lonsdale03}, a wide-area infrared survey with deep ground-based optical data, is optimal for searches of faint QSOs at high redshift as it is deep enough to detect sub-$L^*$ QSOs at $z\\leq4$ and covers sufficient area to detect large numbers of them.  In Section \\ref{template} we outline our method for creating a new QSO template spanning far-UV to mid-IR wavelengths.  In Section \\ref{selection}, we present a simple optical/IR color selection for QSOs at $z>2.8$ and identify areas of possible contaminations or incompleteness.  Our selection results are given in Section \\ref{results}.  In Section \\ref{reliability}, the reliability of the sample is assessed through spectroscopic follow-up and analysis of the infrared colors.  In Section \\ref{completeness} we determine, through simulations and comparisons with known samples, our sample completeness as a function of redshift.  In Sections \\ref{qlf} and \\ref{qlf_comp} we present our measurement of the QLF at $z\\sim3$ and compare to previous studies.  In Section \\ref{ion_bkg}, the QSO contribution to photoionization of HI in the IGM is computed and compared to measurements of the {\\it total} photoionization rate.  Although many early studies of high redshift QLFs use an Einstein-DeSitter cosmology with $H_0=50 $km s$^{-1}$ Mpc$^{-1}$, throughout this paper we choose to use a more recent cosmology with $H_0=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m = 0.3$, $\\Omega_{\\Lambda}=0.7$ and correct other measurements accordingly.  All optical magnitudes are Vega magnitudes unless stated otherwise. ",
        "conclusions": "We present our method of finding high redshift, $z>2.8$, QSOs by identifying a Lyman Break in the optical photometry, and ensuring red mid-IR ($[3.6]-[4.5]$) colors indicative of QSOs.  The use of only three optical filters allows a search over larger areas in the SWIRE fields as most of the area does not have coverage in four or more bands.  The use of only IRAC1 and IRAC2 channels is emphasized as these two bands are a factor of seven times more sensitive than IRAC3 and IRAC4.  Spectroscopic follow-up of 10 $z\\sim3$ ($U$-dropout) candidates confirms that all 10 are QSOs between $2.83<z<3.44$.  Spectroscopy of 10 $z\\sim4$ ($g'$-dropout) candidates confirmed 7 QSOs with $3.48<z<3.88$, two galaxies at low redshift ($z=0.354$, $0.390$) and one unconfirmed redshift.  We place reliability estimates on our $z\\sim3$ and $z\\sim4$ samples of 100\\% ($>69$\\% $1\\sigma$) and $70^{+16}_{-26}$\\%, respectively.  Since we have not spectroscopically confirmed all of our candidates, we only use the more reliable $z\\sim3$ sample for determining a luminosity function.  By using detailed models which include variations in number and column density of line-of-sight HI absorbers, UV spectral slope, emission line equivalent width, redshift, observed magnitude, and photometric errors, we assess the completeness of the optical color selection.   Completeness near the center of the redshift range of our $U$-dropout selection is 85-90\\%.  However, our completeness decreases significantly in our faintest magnitudes bins ($\\sim 75$\\%), due to the shallow depth of the $U$-band imaging.  We find 100 $z \\sim 3$ QSO candidates with $r'<22$ over 11.7 deg$^2$.  Through our models of completeness versus redshift, we derive effective volumes for each half-magnitude bin and compute the $z\\sim3$ QSO luminosity function.  When combined with SDSS data, a least-squares fit to a double power-law gives a faint-end slope, $\\beta = -1.62\\pm0.19$, and location of the break at $M^*= -25.6$.  Our binned QLF assumes that all of the QSO candidates are at $z=3.2$, which may skew the fitted parameters because of luminosity function evolution over our redshift range and Eddington Bias \\citep{eddington13} due to large dispersions in the actual absolute magnitude distribution.  Therefore, we have performed a maximum likelihood fit of the {\\it apparent} magnitude distribution of our sample with that inferred from a specific QLF over this redshift range.  Our results are slightly different, with a shallower faint-end slope, $\\beta = -1.43 \\pm 0.15$, and a somewhat fainter break at $M^*= -24.9$.  This fit is more accurate as it does not assume that all of the QSOs are at the same redshift.  The fitted slope is consistent, within the errors, with values measured at low redshift ($0.5<z<2.0$), $\\beta=-1.45$ \\citep{richards05} and therefore does not require evolution in the faint end slope of the luminosity function.  Our QLF predicts significantly more faint QSOs than suggested with initial measurements at $z\\sim3$ \\citep{hunt04}.  Although it is difficult to tell with our limited spectroscopic sample, some of the difference between our faint-end slope and that of \\citet{hunt04} may be attributed to an increasing number of narrow-line, moderately reddened AGN at fainter UV luminosities which were excluded from the H04 sample.  The QLF exhibits some curvature at all magnitudes and, because of this, the parameters for a double power-law fit are degenerate.  That is, the position of the break, ($M^*$,$\\Phi^*$), can be fit at different locations along the binned QLF, with appropriate changes in the bright and faint end slopes ($\\alpha$,$\\beta$).  This is especially true at high redshift, where the difference between the bright and faint end slopes appears to decrease.  Therefore, one must be careful when assigning physical significance to the measured values of these parameters when comparing to models.  The QSOs in our sample span the break in the luminosity function ($0.25L^* < L < 4.0L*$) and thus we measure the space density of QSOs that comprise the majority (55\\%) of the QSO UV luminosity density at these redshifts.  When combined with the SDSS sample this percentage is more than $70$\\%.  Therefore, large extrapolations are not required to estimate the effects of undetected QSOs.  The integrated UV luminosity density at $z \\sim 3.2$ is $\\epsilon_{1450} = 7.3\\times10^{24}$ ergs s$^{-1}$ Hz$^{-1}$ $h$ Mpc$^{-3}$.  Using the scaling relation derived by \\citet{haardt96}, we infer a maximum HI photoionization rate by QSOs, $\\Gamma = 4.5\\times10^{-13}$s$^{-1}$.  This is about 50\\% of the {\\it total} IGM HI photionization rate at $z=3$, requiring comparable ionizing flux from either starburst galaxies or redder AGN that lie outside our color criteria."
    },
    "0711/0711.2214_arXiv.txt": {
        "abstract": "We present a thermodynamical description of the interaction between holographic dark energy and dark matter. If holographic dark energy and dark matter evolve separately, each of them remains in thermodynamic equilibrium. A small interaction between them may be viewed as a stable thermal fluctuation that brings a logarithmic correction to the equilibrium entropy. From this correction we obtain a physical expression for the interaction which is consistent with phenomenological descriptions and passes reasonably well the observational tests. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.2508_arXiv.txt": {
        "abstract": "Planet frequency shows a strong positive correlation with host mass from the hydrogen-burning limit to $M\\sim2\\,M_\\odot$.  No search has yet been conducted for planets of higher-mass hosts because all existing techniques are insensitive to these planets. We show that infrared observations of the white-dwarf (WD) remnants of massive stars $3\\,M_\\odot \\la M\\la 7\\,M_\\odot$ would be sensitive to these planets for reasons that are closely connected to the insensitivity of other methods.  We identify 49 reasonably bright, young, massive WDs from the Palomar-Green survey and discuss methods for detecting planets and for distinguishing between planet and disk explanations for any excess flux observed.  The young, bright, massive WD sample could be expanded by a factor 4--5 by surveying the remainder of the sky for bright UV-excess objects. ",
        "introduction": "Extrasolar planets have been discovered by 4 different techniques orbiting stars ranging from the hydrogen-burning limit to about 2.5 solar masses. As first pointed out by \\citet{laws03}, the fraction of monitored stars that are found to have planets climbs steadily as the host mass increases from $0.75\\,M_\\odot$ to $1.5\\,M_\\odot$.  \\citet{fischer05} found a more significant trend in a larger sample and noted that while this might reflect a higher rate of planet formation around massive stars, it could also be an indirect selection effect due to the known correlation between planet frequency and metal abundance, combined with the fact that their sample shows a strong correlation between stellar mass and metallicity. \\citet{johnson07} extended this correlation to the entire observed range of host masses $0.1\\,M_\\odot< M < 1.9\\,M_\\odot$ (considering only planet masses $m>0.8\\,M_{\\rm jup}$ and semimajor axes $a<2.5\\,$AU) and argued that it was in fact independent of the correlation of frequency with metallicity.  It would be quite interesting to determine whether the correlation of planet frequency with host mass continues to yet higher masses. However, current techniques are extremely challenged in the regime $M\\ga 3\\,M_\\odot$. Radial-velocity (RV) searches are restricted to FGKM stars because of the absence of usable lines in earlier-type stars. RV searches have managed an ``end run'' around this problem for stars of 1--3 $M_\\odot$ by monitoring the {\\it red giant descendants} of these early type stars \\citep{frink02,reffert06,johnson07,lovis07}. Reasonable mass estimates can be obtained from spectroscopic gravities, although these are not quite as accurate as the mass estimates of main-sequence (MS) stars.  Massive red giants still pose significant challenges for the RV technique.  Their low gravities induce atmospheric instabilities, which radically degrade RV precision relative to what is possible for stable MS stars.  So far, this reduced precision has limited searches to relatively massive planets with periods of order a few hundred days, i.e., well inside the ``snow line'' that is believed to be the boundary of massive-planet formation.  Hence, these searches probably probe only migrated planets and so provide a much smaller window than is available for less massive stars.  At yet higher masses, this technique becomes less effective: declining stability of the red-giant atmospheres and increasing stellar mass both tend to decrease the size of orbits that can be probed, while increasing red-giant radii remove close-in planets.  Finally, the declining mass function and declining lifetime both tend to reduce the overall population of hosts. To date, there are no planets detected with host masses $M>2.5\\,M_\\odot$, although \\citet{lovis07} did find a low-mass brown dwarf orbiting a star with $M=3.9\\,M_\\odot$. Figure \\ref{fig:cmd} is a color-magnitude diagram of all hosts detected by the RV, transits, and microlensing as of June 2007.  Figure \\ref{fig:mm} shows the mass of these planets vs.\\ the mass of their host stars.  The transit technique is challenged to detect planets around massive stars for three reasons: first they are big, second they are rare, and third (most critically) there is no way to confirm detections. In principle, the first problem could be overcome by more accurate photometry, which would probe the smaller fractional transits caused by large stellar radii.  Lack of targets is a more fundamental problem, since even very large surveys are detecting relatively few transiting planets even around much more common stars.  So, in practice, selection pressure drives this technique to a relatively narrow range of F and  G stars (see Fig.\\ \\ref{fig:cmd}). But the key problem is that if a planet were detected, the same line-lacunae that plague RV in this mass range would make confirmation extremely difficult.  Perhaps with prodigious effort, late M-dwarf ``Jupiter impersonators'' could be rejected, but the planets could not be reliably distinguished from brown dwarfs, nor could their masses be measured.  Microlensing is often said to be insensitive to host mass, and indeed Figure \\ref{fig:cmd} shows that microlensing host masses trace the MS mass function: low mass hosts dominate.  However, this very fact makes microlensing relatively insensitive to high-mass hosts, just exactly because of their low frequency of birth and early demise. Moreover, microlensing's ``insensitivity to host mass'' is predicated on the fact that all potential GKM hosts are fainter (or at least not much brighter) than the bulge sources (which are typically turnoff stars, with some subgiants and giants).  By contrast, massive stars are substantially brighter than typical microlensed sources, greatly reducing the probability that the underlying microlensing event could be detected at all under the ``glare'' \\citep{einstein36} of the lens.  Of course, 3 planets have been detected around a pulsar \\citep{wolszczan91}, which is a remnant of a massive star, $M>8\\,M_\\odot$.  However, it is difficult to see how planets could survive a supernova explosion without becoming unbound, so these planets are believed to have formed out of an accretion disk generated by fallback from the explosion. Thus, while an extremely interesting system, the pulsar planets probably do not provide any direct information on planets around massive stars. ",
        "conclusions": "Probing WDs for sub-stellar companions began with brown dwarf searches (Probst 1980; Zuckerman \\& Becklin 1987; Farihi et al. 2005). \\citet{winget03} proposed to search for planetary companions from timing residuals of pulsating white dwarfs. \\citet{mullally07a} reported a possible detection of a planet around a pulsating white dwarf, however follow up observations are required to confirm that the observed change in the pulsation periods is not due to the cooling of the star itself. \\citet{burleigh02,burleigh06}, \\citet{debes05,debes07}, Livio et al. (2005), Friedrich et al. (2006), and \\citet{mullally07b} proposed and looked for planetary companions from excess near- and mid-IR emission. \\citet{hansen06} searched for excess IR emission from disks of very massive WDs, $M_{\\rm WD}>1\\,M_\\odot$.  What is specifically new to this paper is the idea of probing the previously unexplored stellar mass regime $3\\,M_\\odot \\la M \\la 7\\,M_\\odot$ by looking for the planets around the WD remnants of these stars. Typical field WDs have masses $M_{\\rm WD}\\sim 0.6\\,M_\\odot$ and so have MS progenitors $M\\sim 2\\,M_\\odot$.  This is a very interesting mass range, currently probed only from red-giant RV studies, which (as noted above) are restricted in sensitivity to relatively close companions $P\\la 2\\,$yr.  Hence, thermal-excess searches could potentially probe new regimes of parameter space for these stars.  The very massive WDs studied by \\citet{hansen06} are believed to be the result of WD-WD mergers.  If planets survived this merger process, it would be both surprising and very interesting. Of course, some of the massive white dwarfs in the PG sample may be the products of WD-WD mergers as well. These mergers would result in rapid rotation rates, which could be measured from the non-LTE cores of the Balmer lines. Massive WDs with planetary companions could be spectroscopically checked for high rotation rates in order to see if they are the products of binary mergers or single star evolution.  The regime we have targeted has not been systematically explored. The \\citet{mullally07b} sample does contain about a dozen WDs in this mass range, but only a few of these have ages less than 1 Gyr, when planets of a few Jupiter masses can still be robustly detected. We have identified a sample of 49 massive WD targets, the great majority of them young, and suggested that 3--4 times as many could be found in other areas of the sky.  There is no guarantee that the observed correlation between stellar mass and planet occurrence will continue to these higher masses. In fact, \\citet{kennedy07} suggest that the frequency function declines for $M>3\\,M_\\odot$.  Only observations can settle this question."
    },
    "0711/0711.1044_arXiv.txt": {
        "abstract": "Recent high-resolution simulations together with theoretical studies of the dynamical evolution of galactic disks have shown that contrary to wide-held beliefs a `live', dynamically responsive, dark halo surrounding a disk does not stabilize the disk against dynamical instabilities. We generalize Toomre's Q stability parameter for a disk-halo system and show that if a disk, which would be otherwise stable, is embedded in a halo, which is too massive and cold, the combined disk-halo system can become locally Jeans unstable. The good news is, on the other hand, that this will not happen in real dark haloes, which are in radial hydrostatic equilibrium. Even very low-mass disks are not prone to such dynamical instabilities. ",
        "introduction": "The classical paradigm is that at given total mass of the system a galactic disk is stabilized against local dynamical Jeans instabilities, if it is embedded in a dark halo. This can be seen for instance from Toomre's $Q_{\\rm T}$ stability index for a stellar disk, \\begin{equation} Q_{\\rm T} = \\frac{\\kappa \\sigma_{\\rm d}}{3.36\\, G\\Sigma_{\\rm d}}\\,, \\label{eq1} \\end{equation} where $\\kappa$ denotes the epicyclic frequency of the orbits of the stars, $\\sigma_{\\rm d}$ is the radial velocity dispersion of the stars, and $\\Sigma_{\\rm d}$ the surface density of the disk. $G$ denotes the constant of gravitation. If all other parameters are kept constant, but $\\Sigma_{\\rm d}$ is lowered, $Q_{\\rm T}$ rises and the disk becomes more stable against Jeans instabilities. The physical reasoning is that the self-gravity of the disk, which has a destabilizing effect, is reduced by the surrounding halo. Similarly the onset of non-axisymmetric coherent large-scale instabilities of the entire disk such as the bar instability was thought to be damped by a surrounding halo. Ostriker \\& Peebles \\cite{OP73} showed in their classical numerical simulations of the dynamical evolution of a self-gravitating disk that the bar instability could be suppressed, if the disk was embedded in a halo potential. However, modern high-resolution simulations in which the surrounding halo is treated as a dynamically responsive system, have shown that actually the opposite is true. Athanassoula \\cite{A02} showed that in her simulations of the bar instability the bar grows stronger, if the disk is embedded in a live dark halo rather than in a static halo potential. This was explained there and, particularly, in Athanassoula \\cite{A03} as due to the effect of the live halo on the angular momentum exchange within the galaxy. First doubts about an entirely passive role of the halo were already raised by Toomre\\cite{T77} . These findings were supported by theoretical studies of the swing amplification of shearing spiral density waves which is also enhanced, if the disk is embedded in a live dark halo instead of a static potential (Fuchs\\cite{F04} , Fuchs \\& Athanassoula \\cite{FA05}).   Here we return to the severe local Jeans instability of a self-gravitating disk and investigate the effect of the presence of a live dark halo. In the next section we demonstrate how Toomre's concept of the $Q_{\\rm T}$ parameter can be generalized in order to take into account the effect of such a halo. In the final section we discuss implications for realistic disk-halo systems (cf.~also Esquivel \\& Fuchs \\cite{E07}). ",
        "conclusions": "As a first application of the stability criterion derived here we test the stability of the Milky Way disk and the surrounding dark halo in the vicinity of the Sun. The local disk and halo parameters listed in Table \\ref{table1} imply $Q_{\\rm T} = 2.8$ and $\\beta = 0.0057$, respectively. If we include in our estimate the cold interstellar gas with a local surface density of $4 M_{\\odot}$/pc$^2$ (Dame \\cite{D93}) and assume a velocity dispersion of the interstellar gas of $\\sigma_{\\rm g}$ = 5 km/s, which leads to a reduced mass weighted effective velocity dispersion of the combined stellar and gaseous disks, the parameter values change to $Q_{\\rm T} = 2.2$ and $\\beta = 0.0078$, respectively. Equation (\\ref{eq7}) implies that $\\lambda_{\\rm J}= 8 \\,\\lambda_{\\rm crit}$ = 39 kpc. Since the instability sets in nominally at $\\lambda = 127 \\,\\lambda_{\\rm crit}$ the Milky Way disk and halo system seems to be very stable.  \\begin{table} \\tbl{Local parameters of the Milky Way} {\\begin{tabular}{ccc} \\toprule $\\Sigma_{\\rm d}$  & 38 $M_{\\odot}/pc^{2}$ & (Holmberg \\& Flynn \\cite{HF04})\\\\ $\\sigma_{\\rm d}$  & 40 km/s & (Jahrei{\\ss} \\& Wielen \\cite{Jawi97})\\\\ $\\kappa$  & $\\sqrt{2}\\cdot220$ km/s/8.5 kpc & (flat rotation curve)\\\\ $\\rho_{\\rm h}$  & 0.01 $M_{\\odot}/pc^{3}$ &  (Bahcall \\& Soneira \\cite{BS80})\\\\ $\\sigma_{\\rm h}$   &  220 km/s/$\\sqrt{2}$ & (isothermal sphere)\\\\ $\\lambda_{\\rm crit}$ & 4.8 kpc & \\\\ $\\lambda_{\\rm J}$ & 39 kpc & ($\\beta$ = 0.0078)\\\\ \\botrule \\end{tabular}} \\label{table1} \\end{table}  In order to explore in what range the $\\beta$-parameter of spiral galaxies is to be expected in general, we consider the model of a Mestel disk with the surface density $\\Sigma_d = \\Sigma_0 \\,R^{-1}$ embedded in a singular isothermal sphere representing the dark halo with the density distribution $\\rho_h = \\rho_0 \\,R^{-2}$. The rotation curve of the model galaxy is given by \\begin{equation} \\upsilon^2_{\\rm c}(R) = \\upsilon^2_{\\rm d}(R) + \\upsilon^2_{\\rm h}(R) \\label{eq8} \\end{equation} with the disk contribution $\\upsilon^2_{\\rm d}(R)=2\\pi G \\Sigma_{\\rm d} R$ = const.~and the halo contribution $\\upsilon^2_{\\rm h}(R)=4\\pi G \\rho_{\\rm h} R^2$ = const.~(Binney \\& Tremaine \\cite{BT87}). Assuming isotropic velocity dispersions of the dark matter particles it follows from the radial Jeans equation, \\begin{equation} \\frac{1}{\\rho_{\\rm h}}\\frac{d\\rho_{\\rm h} \\sigma_{\\rm h}^2}{dR} =-\\frac{v_c^2}{R}\\,, \\label{eq8a} \\end{equation} that their velocity dispersion is given by \\begin{equation} \\sigma^2_{\\rm h} = \\frac{1}{2}(\\upsilon^2_{\\rm d} + \\upsilon^2_{\\rm h})\\,, \\label{eq9} \\end{equation} because the dark matter particles are bound by both the gravitational disk and halo potentials. We find then \\begin{equation} \\beta = \\frac{\\upsilon^2_{\\rm h}}{R^2}\\frac{1} {\\upsilon^2_{\\rm d} + \\upsilon^2_{\\rm h}}\\frac{R^2}{4} \\frac{\\upsilon^4_{\\rm d}}{(\\upsilon^2_{\\rm d} + \\upsilon^2_{\\rm h})^2} = \\frac{1}{4}\\frac{\\upsilon^2_{\\rm h}\\upsilon^4_{\\rm d}} {(\\upsilon^2_{\\rm d} + \\upsilon^2_{\\rm h})^3}\\,, \\label{eq10} \\end{equation} which implies the maximal value \\begin{equation} \\beta \\leq \\beta_{\\rm max} (\\upsilon^2_{\\rm d} = 2\\upsilon^2_{\\rm h}) = 0.037\\,. \\label{eq11} \\end{equation} This means that in realistic halo models its density cannot be increased, on one hand, and the velocity dispersion of the halo particles lowered, on the other hand, indiscriminately because the halo has to stay in radial hydrostatic equilibrium. Equation (\\ref{eq11}) implies $\\lambda_{\\rm J} = 3.7 \\,\\lambda_{\\rm crit}$, whereas the instability sets nominally in at $\\lambda=26\\,\\lambda_{\\rm crit}$. In order to ensure stability at smaller wave lengths the Toomre stability index must be larger than $Q_{\\rm T} \\geq 1.02$.  \\begin{figure} \\centering \\psfig{file=bfdark2007_f2.ps,width=4.5in} \\caption{$\\beta$-parameters of mass models of the low surface brightness galaxy F568-1.} \\label{fig2} \\end{figure}  As a further concrete example we analyse the dynamics of the low surface brightness galaxy F568-1. De Blok et al.~\\cite{db01} have observed its rotation curve and constructed several concurring mass models for the galaxy, which all fit the rotation curve equally well. Using the parameters of the models with isothermal haloes and with either a `realistic' $M/L$-ratio of the disk or a `maximum-disk' we have solved numerically the radial Jeans equation (\\ref{eq8a}) for the dark matter particles, \\begin{equation} \\sigma_{\\rm h}^2(R)= \\frac{1}{\\rho_{\\rm h}(R)}\\int_R^\\infty dR \\, \\frac{\\rho_{\\rm h}(R) v_{\\rm c}^2}{R}. \\label{eq12} \\end{equation} The resulting $\\beta$-parameters of both models are shown in Fig.~\\ref{fig2}. As can be seen from Fig.~\\ref{fig2} the $\\beta$-parameters are of the same small order of magnitude as in the Milky Way or the simple disk -- halo model described above.  We conclude from this discussion that embedded galactic disks are not prone to Jeans instabilities, provided their Toomre stability index is a few percent higher than $Q_{\\rm T}$ = 1. From a practical point of view the destabilizing effect of the surrounding dark halo on the Jeans instability of the embedded galactic disks seems to be negligible."
    },
    "0711/0711.2337_arXiv.txt": {
        "abstract": "We present a search for RR Lyrae variable stars from archival observations of the Southern Edgeworth-Kuiper Belt Object survey. The survey covers 1675 square degrees along the ecliptic to a mean depth of $V$=19.5, i.e. a heliocentric distance of $\\sim50$kpc for RR Lyrae stars. The survey reveals 2016 RR Lyrae candidates. Follow-up photometric monitoring of a subset of these candidates shows $\\sim$24\\% contamination by non-RR Lyrae variables. We derive a map of over-density of RR Lyraes in the halo that reveals a series of structures coincident with the leading and trailing arms of debris from the Sagittarius dwarf galaxy. One of the regions of over-density is found on the trailing arm, 200$^{\\circ}$ from the main body of the Sagittarius dwarf at a distance of $\\sim45$kpc. This distant detection of the stellar population of the outer trailing arm of Sagittarius offers a tight constraint on the motion of the dwarf galaxy. A distinctly separate region of over-density is seen towards the Virgo Over Density. ",
        "introduction": "The stellar halo of our galaxy offers a window to the past. Stars in the halo have long orbital periods and inhabit regions of space where the galactic potential is relatively smooth and slowly evolving. Consequently, structures in the halo suffer reduced spatial and kinematic dissipation and may persist for several Gigayears \\citep{Bullock05}. Such halo structures offer an insight into the process of galaxy formation, arguably one of the most important processes yet to be understood by modern astrophysics.  One schematic for the process of halo formation is that due to \\citet{ELS}, who found that the stars on the most eccentric orbits were also the most metal-poor. This was interpreted as the signature of monolithic collapse of the proto-galactic cloud; the most metal-poor, and hence oldest, stars forming during the collapse. This solitary process of galaxy formation was challenged by \\citet{SZ}, who proposed the large dispersion seen in the metallicity of globular clusters and their lack of radial abundance gradient indicated that the Galaxy was the composite of a large number of sub-galaxy sized pieces.  Recent observational studies have converged on a combination of scenarios \\citep[for a review]{Freeman02}. This is perhaps most clearly demonstrated by the works of \\citet{Kinman07}, \\citet{Carollo07} and \\citet{Miceli07} who find a distinct outer halo population. \\citet{Carollo07} shows that whereas the inner halo (galactocentric radius less than 15kpc) is dominated by highly eccentric, prograde orbits and a metallicity of [Fe/H]$\\sim-1.6$, the outer halo exterior to this possesses a more uniform distribution of eccentricities, includes highly retrograde orbits and [Fe/H]$\\sim-2.2$. \\citet{Kinman07} study of the local halo similarly identifies a component that possesses retrograde rotation and streaming motions and another that exhibits negligible rotation and less signs of streaming. Kinman et al.\\ furthermore show that the horizontal branch morphologies of the two components of the halo are consistent with those seen in the young and old globular clusters, reinforcing the conclusions of \\citet{Lee99} that the inner halo is consistent with rapid collapse and the outer halo with the ongoing accretion of dwarf galaxies.  A halo that formed from \\emph{in situ} star formation would possess, at the current epoch many dynamical times later, very little substructure. Simulations of galaxy formation in the context of $\\Lambda$CDM predict that the majority of the stars in the outer halo formed in progenitor dwarf galaxies and were subsequently accreted \\citep{Abadi06,Bullock05}, with the general conclusion that significant amounts of substructure should survive to this time. However, we must point out that $\\Lambda$CDM predictions on galaxy scales are contentious at this stage. \\citet{Bell07} compare the level of structure evident in the SDSS data with simulations and finds that the Milky Way's halo is well matched to a halo built exclusively from disrupted satellite galaxies.  Numerous studies of the spatial distribution of stars in the halo have shown the importance of ongoing accretion of satellite galaxies on to the Milky Way. The most striking example of ongoing accretion is the Sagittarius dwarf galaxy \\citep[Sgr]{Ibata94}. The ensuing trail of debris has been found to wrap around the sky \\citep{fieldofstreams,Majewski03,Ibata03,Yanny03,Newberg02}. The debris of Sgr is arguably the largest contributor yet found to the halo substructure at radii of over 50kpc. The orbit of Sgr, as traced by its extensive debris tails, has been used as a probe of the mass and shape of the Galaxy \\citep{Ibata01, Helmi01, Helmi04, Johnston05, Law05, Fellhauer06}.  Another halo structure, the Virgo Stellar Stream (VSS) was found as an over-density of RR Lyrae stars (RRLs) \\citep{VivasZinn06,Zinn04} at galactocentric radius $\\sim20$kpc. Later radial velocity measurements \\citep{Duffau06} revealed the grouping shared a common space velocity. \\citet{Newberg02}, \\citet{Juric05} and more recently \\citet{Newberg07}, find a coincident diffuse over-density of F-type main sequence stars spanning some 1000 square degrees of sky called the Virgo Overdensity (VOD). \\citet{Newberg07} proposes that the VSS and VOD are the same structure. On the basis of radial velocity measurements Newberg et al.\\ concludes it is not associated with a confluence of the leading and trailing arms of Sgr as proposed by \\citet{MartinezDelgado07} but is possibly remnant material from another merger event.  The Monoceros Stream \\citep{Newberg02} is believed to be a low surface brightness ring encircling the Galaxy at 15-20kpc \\citep{Ibata03}. The structure was first traced by \\citet{Newberg02} using halo turnoff stars and subsequently by \\citet{Majewski03} using M giants. The stream is not thought to be related to Sgr tidal debris or the VSS/VOD, but rather may result from the disruption of the Canis Major dwarf \\citep{MartinezDelgado07,Dinescu06}. \\citet{Conn07} reports that interpretation of the Monoceros stream as a warp or flare in the galactic disk does not account for the observed stellar densities.  Possibly related to the Monoceros Stream are the Tri-And over-density \\citep{RochaPinto04} at $\\sim20$kpc and a yet more distant stellar structure at $\\sim28$kpc \\citep{Martin07}.  Most recently, \\citet{Belokurov07} reports the detection of a significant stellar over-density that covers some 80$^\\circ$ of  galactic latitude, dubbed the Hercules-Aquila Cloud. This over-density lies at a distance of between 10-20kpc and straddles the Galactic disk offset by 6-15kpc from the Galactic centre. Spectroscopy of probable cloud members shows that the structure possesses a radial velocity offset of $\\sim$180km/s from the halo and thick disk.  The features discussed above are a small fraction of the number that could be expected from simulations of $\\Lambda$CDM hierarchical structure formation \\citep{Bullock01,Bullock05}. The search for relic substructure could help illuminate the outstanding 'missing satellite problem' in galaxy formation in  $\\Lambda$CMD cosmology \\citep{Freeman02}. Here, models predict of order 500 small dark matter halos should surround the Milky Way today, in contrast to the $\\sim40$ known dwarf galaxies in the Local Group \\citep{Wilkinson07}. If the level of substructure does not match that predicted from a large population of dwarf galaxies this may prompt revisions to model assumptions. ",
        "conclusions": "In this paper we have presented a survey for RR Lyrae stars that has revealed a sample of 2016 RR Lyrae candidates from multiple (2-10) epoch data covering 1675 square degrees. Follow up photometric observations of these candidates show 76\\% to be RR Lyrae variables. Our survey reveals a series of significant over densities in the Galactic halo. Many of the over dense regions lie along the leading and trailing arms of debris from the Sagittarius dwarf galaxy extending up to $\\sim200^{\\circ}$ from the main body of the Sagittarius dwarf. Another locality of over density can not be ascribed to the Sagittarius dwarf: namely the region of the Virgo Over Density. Here we find two significant over densities 10$^{\\circ}$ to the south west of the centre previously reported.  Radial velocities of RR Lyrae stars drawn from our over densities will give insight into the dynamics and origins of these systems. In particular, regions coincident with the leading and trailing arms of the Sagittarius dwarf could offer powerful constraints on the orbit and evolution of the Sagittarius dwarf."
    },
    "0711/0711.2101_arXiv.txt": {
        "abstract": "Distribution and kinematics of molecular gas in the central region of the barred spiral galaxy Maffei 2 were investigated using a data set of $^{12}$CO(1--0), $^{12}$CO(2--1), CS(2--1) lines and 103 GHz continuum. We found that the offset ridges along the kpc-scale bar continue to the central spiral structure embedded in the weak oval structure which is regarded as $x$2 orbit in the bar potential. The spiral structure continues toward the center diverging from the oval structure. The size of these structures is less than $R \\sim 100$ pc. The mass concentration within $R =$ 35 pc is estimated to be $2 \\times 10^{8} M_{\\odot}$. The high mass concentration is consistent with theoretical predictions concerning the creation of such a nuclear spiral structure. A comparison with the tracers of dense gas and star-forming region suggests that the dense molecular gas traced by CS(2--1) line is formed at the crossing points of $x$1 and $x$2 orbits and the star-forming region appears after $2 \\times 10^{5}$ yr which is comparable with the free-fall time of the dense gas traced by the CS line ($\\sim 10^{5}$ cm$^{-3}$). ",
        "introduction": "To maintain activities of central starbursts or active galactic nuclei in galaxies, fueling of molecular gas toward the center is necessary. Bar is one of the possible mechanisms, for the gas fueling. Many numerical simulations of gas motion in the bar potential show that bars can accumulate gas toward the central region (e.g., Athanassoula 1992). Observational evidence of gas accumulation toward the central kilo-parsec region has also been shown. For example, Sakamato et al. (1999) and Sheth et al. (2005) showed that barred spiral galaxies have higher central concentration of molecular gas than non-barred spiral galaxies. Furthermore, Kuno et al. (2007) showed that the higher central concentration in barred spirals is made by gas accumulation within a region on the order of the bar length and found a correlation between bar strength and the degree of central concentration.  On the other hand, observations of molecular gas with high angular resolution sufficient to resolve the central structure less than 100 pc are still limited (e.g., Schinnerer et al. 2003; Garc\\'ia-Burillo et al. 2003). Such observations are necessary to understand how the gas flows into the nucleus and relates with central activities. For example, the size of the twin peaks often seen in the central region of barred galaxies is less than several hundred pc (Kenney et al. 1992). Therefore, we need such high angular resolution to investigate the distribution and dynamics of molecular gas in the central region.  Maffei 2 is a barred spiral galaxy that belongs to the IC 342/Maffei group. Since Maffei 2 is located behind the Galactic plane, its optical features are not clear. However, the bar and spiral structures can be seen clearly in a near infrared image (Jarrett et al. 2003). Many CO observations have so far been done (Ishiguro et al. 1989; Kawabe et al. 1991; Israel, Baas 2003; Mason, Wilson 2004; Kuno et al. 2007). The CO map obtained with the Nobeyama 45-m telescope shows structures typical of the distribution of molecular gas in barred spiral galaxies; central peak, offset ridges along the leading side of the large-scale bar, condensations of molecular gas at the end of the bar, and spiral arms from the end of the bar. They are symmetric against the center. The motion of molecular gas in the bar is also consistent with that seen in barred spiral galaxies. Namely, the large velocity change is seen across the bar implying that the gas moves along the offset ridges toward the center. Since Maffei 2 is one of the nearest barred spiral galaxies, we can investigate the distribution and dynamics of molecular gas in the center and the bar in detail. Furthermore, since Maffei 2 has a moderate starburst at the center as suggested from IR and radio observations (Rickard, Harvey 1984; Turner, Ho 1994), observations of Maffei 2 with high angular resolution are helpful to understand the mechanism of the gas fueling toward the galactic center by bars and relation with a starburst phenomenon in barred spiral galaxies.  In this paper, we present the results of $^{12}$CO(1--0), $^{12}$CO(2--1), CS(2--1) lines and 103 GHz continuum observations of Maffei 2 with the Nobeyama Millimeter Array (NMA). We discuss the central structure of molecular gas and star formation in Maffei 2. We adopt 2.8 Mpc as the distance of Mafei 2 throughout this paper (Karachentsev et al. 2003). ",
        "conclusions": "\\subsection{Nuclear spiral structure of molecular gas}  As mentioned in the previous section, the P -- V diagrams along the major axis (Figure 7) shows two components almost overlapping along the line of sight. These structures can be regarded as a spiral structure embedded in a weak oval structure. We made integrated intensity maps of each component of the parallel structure in Figure 7 separately using the channel map. Figure 9 shows the integrated intensity map overlaid on the HST image (F814W) obtained from the HST Archive.  They look like the two-arm spiral structure continues from the offset ridges along the large-scale bar in the CO map obtained with the NRO 45-m telescope. The structure can be seen in both $^{12}$CO(1--0) and $^{12}$CO(2--1). The spiral structure is symmetric against the center determined from the 2MASS data. The northern arm coincides with the dust lane in the HST image very well near the center, while the dust lane along the southern arm is not clear. This is because the northern arm is on the near side (western side against the major axis of the galaxy) near the center.  It is apparent that the oval structure in Figure 7b is not an expanding ring. The part of the oval structure that overlaps with the northern arm is on the near side near the center, as mentioned above. The radial velocity of the part is red-shift from the systemic velocity of Maffei 2 (around C' in figure 7c). On the other hand, if the ring is shrinking, the velocity toward the center must be higher than 50 km s$^{-1}$. In that case, the oval structure collapses into the nucleus in less than $2 \\times 10^{6}$ yr. This seems to be unlikely. The oval structure can be naturally understood as an oval motion of molecular gas. The oval structure connects the spiral structure. The schematic figure of the spiral and oval structures is shown in Figure 10. Here, we assume inclination and position angle derived for the galactic disk. Similar structures with the offset ridges and oval structure are often seen in numerical simulations (e.g., Athanassoula 1992). These structures are made from $x$1 and $x$2 orbits in a bar potential. Similar structure has been found in NGC 4303 (Schinnerer et al. 2002; Koda, Sofue 2006). The complex P--V diagrams can be understood with our interpretation. The offset ridges along the large-scale bar contact with the x2 orbit at B and B' in Figure 7. The constant gradient of the radial velocity B -- C and B' -- C' in Figure 7 is attributed to the flow of the molecular gas along the spiral arms around the center. The spiral arms diverge from the oval structure toward the center. This corresponds to C -- D and C' -- D' in the P -- V diagrams. The deviation from the systemic velocity becomes max there, since the gas stream line closes to the line of sight.  The sizes of the structure shown in Figure 10 were derived from the P -- V diagrams along the major and minor axes and the channel map as follows: (1) The separation along the major axis between the spiral arms derived from Figure 7(b) ($\\sim 70$ pc). (2) The separation along the minor axis between the spiral arms derived from Figure 8 ($\\sim 90$ pc). (3) The separation along the major axis direction between the positions where radial velocity equals to the systemic velocity. If we assume that most of the gas moves along the spiral arms, the spiral arms are perpendicular to the line of sight at these positions. This can be measured in channel map at the systemic velocity of Maffei 2 ($\\sim$ -24 km s$^{-1}$) (Figure 5 and 6). ($\\sim 75$ pc). (4) The separation along the minor axis between the positions where radial velocity equals the systemic velocity derived from Figures 5 and 6. ($\\sim 120$ pc). (5) The length of the oval structure along the major axis derived from Figure 7(b) ($\\sim 200$ pc).  \\subsection{Formation mechanism of the spiral structure}  Martini et al. (2003a; 2003b) show from HST images that galaxies with strong bars often have a grand-design nuclear dust spiral structure, which appears to be the continuation of the dust lanes along the leading edges of a large-scale bar toward the center. On the other hand, observations of such spiral structures of molecular gas are still rare. The distribution of molecular gas in Maffei 2 can be regarded as such a spiral structure. Although Maffei 2 is classified as SAB(rs) in RC3, it seems to be due to the unclear optical image. According to Buta and Block (2001), Maffei 2 should be classified as SB(rs). Therefore, our result is consistent with Martini(2004)'s finding that grand-design spiral structure is only found in SB(s) or SB(rs). Some numerical simulations show that when there is a massive concentration at the center, such as a super massive black-hole, a nuclear spiral structure is formed (Fukuda et al. 1998; Ann, Lee 2004; Maciejewski 2004). This is because the mass concentration removes the rigid rotation part from the rotation curve. As a result, IILR (inner Inner Lindblad Resonance), where a circum-nuclear ring is formed, disappears or another resonance (Nuclear Lindblad Resonance) appears inside the IILR. The mass of the central concentration assumed in numerical simulations to make the spiral structure is 0.1 \\% -- 1 \\% of the total mass of the host galaxies. The observational data of molecular gas has the advantage that it can be used to investigate the kinematics of the central structure. From the P -- V diagram along the major axis (Figure 7), the rotation velocity is estimated to be about 140 km s$^{-1}$ at $R = 35$ pc. The central mass within $R =$ 35 pc is estimated to be $2 \\times 10^{8} M_{\\odot}$ from the rotation velocity assuming spherical distribution of mass. The mass is about 0.5 \\% of the total dynamical mass of Maffei 2 (Mason, Wilson 2004). Even if the rotation velocity is overestimated due to the non-circular motion, it seems to be more than 0.1 \\% of the dynamical mass (Koda, Wada 2002). The mass fraction of the central mass concentration in Maffei 2 is comparable to the prediction of the simulations. The rotation velocity still rising toward the center at $R = 35$ pc. It is desirable to obtain the inner rotation curve with higher spatial resolution and higher sensitivity observations to know the mass distribution in the inner region.  \\subsection{Inflow of Molecular Gas along the Bar}  The molecular gas changes the direction of motion with a large velocity change when it enters the offset ridge along the large-scale bar. Then, the gas flows toward the center along the offset ridge. Since the radial velocity is almost constant along the straight ridges, as seen in Figure 7, the streaming velocity along the ridge seems to be constant. (The component of the pattern speed of the bar is relatively small as compared with the velocity along the bar.) The velocity along the ridge is estimated to be 70-120 km s$^{-1}$ assuming the inclination angle and position angle of the galactic disk. Here, we assumed a pattern speed of the bar of 53.5 km s$^{-1}$ kpc$^{-1}$ which was adopted from Hurt et al. (1996) and made a correction for the difference of the adopted distance. The velocity is comparable with those in NGC 253 (Sorai et al. 2000) and NGC 1530 (Regan et al. 1997). Using the velocity and surface density of molecular gas at the radius of 200 pc in the offset ridge, the mass flow of molecular gas along the ridges is derived to be 7--12 $M_{\\odot} {\\rm yr}^{-1}$. This is comparable with the estimation for NGC 1530 by Regan et al. (1997).  If the molecular gas in the flow is accumulated in the center directly, it takes only $< 10^{8}$ yr to concentrate all of the molecular gas in the disk on the oval structure. However, only a portion of the flowing gas must be accumulated in the central region as estimated by Regan et al. (1997). They evaluated it to be 20 \\% of the flowing gas.   \\subsection{Central starburst in barred spiral galaxies}  There are concentrations of molecular gas where the offset ridges connect to the spiral structure (Figures 7 and 9). The mass of the condensations within 65 pc $\\times$ 94 pc is estimated to be $\\sim 7 \\times 10^{6} M_{\\odot}$ from the $^{12}$CO(1--0) intensity using a conversion factor of $1 \\times 10^{20}$ cm$^{-2}$ [K km s$^{-1}$]$^{-1}$ (Nakai, Kuno 1995). The peaks correspond to $x$1/$x$2 orbit-crowding region.  We compare the position of the peaks with the peaks in the CS(2--1) and 3 mm continuum maps. Since the northern peak is much stronger than the southern peak in the CS(2--1) and 103 GHz maps, we discuss the northern peak here. The peak of the 3mm continuum which is a tracer of star-forming regions is located downstream of the CO peak. The offset between the CO and 3 mm peaks is about 20 pc, while the CS peak coincides with the CO peak. We can estimate the time scale for the gas to move from the CS peak to the 3 mm peak along the ridge using the streaming velocity measured at the offset ridge, 70 -- 120 km s$^{-1}$, assuming that the peaks are near the point where the spiral structure is perpendicular to the line of sight (cf. Figure 10). The time scale is about 1.6 -- 2.8 $\\times 10^{5}$ yr. On the other hand, free fall time of the dense gas traced by CS ($\\sim 10^{5}$ cm$^{-3}$) is estimated to be $1.4 \\times 10^{5}$ yr. This is comparable to the time for the gas to move from the CS peak to the 3 mm peak. From these results, the following scenario can be described: The molecular gas in $x$1 orbit moves along the offset ridge in the large bar. At the $x$1/$x$2 orbit-crowding region, the gas contacts with $x$2 orbit and is compressed by shock and the density rises abruptly there. Once the dense gas is formed, it evolves into stars with free-fall time. It is interesting to compare the structure with that seen in NGC 6951 (Kohno et al. 1999). Kohno et al. (1999) show that the condensations of dense gas traced by HCN in NGC 6951 are located downstream of more diffuse gas traced by CO. They showed that the time scale of gravitational instability and flowing time of the gas along the ridges are comparable ($\\sim 1 \\times 10^{6}$ yr). This result supports the idea that the dense gas is made by gravitational instability during the flow along the ridge. We speculate that the difference of our results from Kohno et al. (1999) may be caused by the difference in the size of the structure. The structure found in NGC 6951 is several times larger than that in Maffei 2. The size of the CO peak is more than hundreds pc, while the condensations in Maffei 2 are comparable to the Galactic GMCs. Therefore, the molecular clouds should be treated as particles in NGC 6951. In that case, since the viscosity of the gas is small, shock may not play an important role in dense gas formation at the $x$1/$x$2 orbit-crowding region, while the gravitational instability does. On the other hand, the clump of molecular gas whose size is comparable with a GMC is compressed at the crossing point in Maffei 2. As a result, strong shock may occur and dense gas is formed at a burst there."
    },
    "0711/0711.0274_arXiv.txt": {
        "abstract": "A follow-up survey using the Submillimetre High-Angular Resolution Camera (SHARC-II) at $350\\,\\mathrm{\\mu m}$ has been carried out to map the regions around several $850\\,\\mathrm{\\mu m}$-selected sources from the Submillimetre HAlf Degree Extragalactic Survey (SHADES). These observations probe the infrared luminosities and hence star-formation rates in the largest existing, most robust sample of submillimetre galaxies (SMGs).  We measure $350\\,\\mathrm{\\mu m}$ flux densities for 24 $850\\,\\mathrm{\\mu m}$ sources, seven of which are detected at $\\geq2.5\\,\\sigma$ within a 10\\,arcsec search radius of the $850\\,\\mathrm{\\mu m}$ positions. When results from the literature are included the total number of $350\\,\\mathrm{\\mu m}$ flux density constraints of SHADES SMGs is 31, with 15 detections. We fit a modified blackbody to the far-infrared (FIR) photometry of each SMG, and confirm that typical SMGs are dust-rich ($M_\\mathrm{dust}\\simeq9\\times10^{8}\\,\\mathrm{M}_{\\odot}$), luminous ($L_\\mathrm{FIR}\\simeq2\\times10^{12}\\,\\mathrm{L}_{\\odot}$) star-forming galaxies with intrinsic dust temperatures of $\\simeq35\\,\\mathrm{K}$ and star-formation rates of $\\simeq400\\,\\mathrm{M}_{\\odot}\\,\\mathrm{yr}^{-1}$. We have measured the temperature distribution of SMGs and find that the underlying distribution is slightly broader than implied by the error bars, and that most SMGs are at 28\\,K with a few hotter.  We also place new constraints on the $350\\,\\mathrm{\\mu m}$ source counts, $N_\\mathrm{350}(>25\\,\\mathrm{mJy})\\sim200$--$500\\,\\mathrm{deg}^{-2}$. ",
        "introduction": "The SCUBA HAlf Degree Extragalactic Survey (SHADES; \\citealt{Mortier}; \\citealt{Coppin06}) mapped $\\simeq0.25\\,\\mathrm{deg^{2}}$ of sky with an RMS of 2\\,mJy at $850\\,\\mathrm{\\mu m}$ with the Submillimetre Common-User Bolometer Array (SCUBA; \\citealt{Holland}).  The area was split approximately evenly between the Lockman Hole (LH) and the Subaru-\\textit{XMM} Deep Field (SXDF).  Using uniform selection criteria, the survey uncovered 120 SMGs with a median deboosted flux density of $\\sim\\!5\\,\\mathrm{mJy}$ \\citep{Coppin06}.  The SHADES programme was designed to study the nature and evolution of high star-formation rate (SFR) submillimetre galaxies (SMGs) via a systematic study of a well-characterised and statistically meaningful sample. The programme includes an effort to identify members of the source list at other wavelengths using deep follow-up data from the radio \\citep{paper3} to X-ray, in order to characterise the SHADES population and to probe the variation in the star-formation and clustering with redshift. The relatively precise positions available from the radio data greatly aid in identifying secure counterparts at other wavelengths, which can then be used to provide spectroscopic or photometric redshifts \\citep{paper4} and to categorise the sources.  Even combined with knowledge of source redshift, SCUBA $850\\,\\mathrm{\\mu m}$ and millimetre wavelength fluxes do not constrain the total dust mass of a galaxy because there is an ambiguity between column density and source temperature.  The submillimetre (submm) spectral energy distribution (SED) of a luminous dusty galaxy arises from the re-emission at far-infrared (FIR) wavelengths of absorbed optical/UV radiation from regions of intense star-formation (see e.g.~\\citealt{SandMir} and references therein).  Typically the dust temperature is within a factor of 2 of $T_\\mathrm{d}\\sim40\\,\\mathrm{K}$ \\citep{Blain}, so the restframe SED peaks in the range 60--$120\\,\\mathrm{\\mu m}$ and the SED is almost a simple power law at the SCUBA wavelengths and longer.  At a redshift of $\\left\\langle z\\right\\rangle\\,{\\sim}\\,2$ or 3, typical of SMGs \\citep{Chapman2005}, the peak of the SED is shifted to be near $350\\,\\mathrm{\\mu m}$, and the Submillimetre High Angular Resolution Camera (SHARC-II; \\citealt{Dowell}) at the Caltech Submillimetre Observatory (CSO) is very well situated to provide the photometry needed to constrain the temperatures, and therefore the luminosities and masses, of the SHADES sources.  We have therefore mapped a subset of the SHADES catalogue with SHARC-II.  In this paper we constrain the FIR SEDs of our sample by fitting modified blackbody curves to the FIR photometry of 31 SHADES SMGs including SHARC-II data at $350\\,\\mathrm{\\mu m}$. Sections~\\ref{obs} and \\ref{dr} describe the observations and data reduction.  Section~\\ref{results} presents the $350\\,\\mathrm{\\mu m}$ flux densities of the SHADES galaxies and the FIR SEDs. Section~\\ref{discussion} provides a discussion of the results. Conclusions and future prospects are discussed in Section~\\ref{fp}.  We adopt the Cosmological parameters from the \\textit{WMAP} fits in \\citet{Spergel}: $\\Omega_\\Lambda=0.73$, $\\Omega_\\mathrm{m}=0.27$, and $H_\\mathrm{0}=71$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$. ",
        "conclusions": "\\label{fp}  We have performed follow-up mapping at $350\\,\\mathrm{\\mu m}$ using SHARC-II of several $850\\,\\mathrm{\\mu m}$-selected sources from SHADES.  In total there are flux density measurements for about 25 per cent of the SHADES SMGs, of which this work has provided 21 new flux density constraints.  The combination of 350 and $850\\,\\mathrm{\\mu m}$, 1.1, and/or 1.2\\,mm photometry and spectroscopic or photometric redshifts have provided estimates of $T_\\mathrm{d}$ and $L_\\mathrm{FIR}$ for each $350\\,\\mathrm{\\mu m}$-observed SHADES SMG by fitting a modified blackbody function to the data.  The SMGs population is confirmed to be dominated by very dusty ($M_\\mathrm{dust}\\simeq9\\times10^{8}\\,\\mathrm{M}_{\\odot}$), luminous ($L_\\mathrm{FIR}\\simeq2\\times10^{12}\\,\\mathrm{L}_{\\odot}$) star-forming ($\\mathrm{SFR}\\simeq400\\,\\mathrm{M}_{\\odot}\\,\\mathrm{yr}^{-1}$) galaxies with dust temperatures of $\\simeq35\\,\\mathrm{K}$.  We have measured the temperature distribution of SMGs and find that the underlying distribution is slightly broader than implied by the error bars, and that most SMGs are at 28\\,K with a few hotter.  With the data in hand, it is very difficult to understand whether we are just seeing a combination of selection effects and the known luminosity evolution of SMGs, or whether the SEDs might also systematically change with redshift.  Larger samples of SMGs with a broader selection function will enable this to be fully investigated, e.g.~with a deep $450\\,\\mathrm{\\mu m}$ survey, a primary goal of the upcoming SCUBA-2 Cosmological Legacy Survey\\footnote{{\\tt http://www.jach.hawaii.edu/JCMT/surveys/Cosmology.html}}.  Using these data, we have also been able to estimate the $350\\,\\mathrm{\\mu m}$ source counts at $200\\,\\mathrm{deg}^{-2}\\lesssim N(\\gtrsim25\\,\\mathrm{mJy})\\lesssim500\\,\\mathrm{deg}^{-2}$.  Large blank-field surveys at 350 and/or $450\\,\\mathrm{\\mu m}$ will allow us to study a statistically significant population of lower redshift and/or hotter dust temperature galaxies, responsible for emission nearer to the peak of the extragalactic background than those selected at $850\\,\\mathrm{\\mu m}$.  These studies should provide a link between galaxies detected at $\\sim\\!1\\,\\mathrm{mm}$ and those selected in the mid-IR with \\textit{Spitzer}.  Such comparisons should allow us to investigate the evolution of FIR emitting galaxies and how they contribute to the background at different wavelengths."
    },
    "0711/0711.3985_arXiv.txt": {
        "abstract": "After our tentative detection of an optical counterpart to \\smc\\ from archival Hubble Space Telescope (HST) imaging, we have followed up with further images in four bands. Unfortunately, the source originally identified is not confirmed. We provide deep photometric limits in four bands and accurate photometry of field stars around the location of the magnetar. ",
        "introduction": "Magnetars are neutron stars that derive their large X-ray luminosity from the decay of a super-strong magnetic field or the order $10^{15}$\\,G (and even greater internally), many orders greater even than normal radio pulsars and accreting X-ray sources (Woods \\& Thompson, 2006). They are found preferentially in the Galactic plane, in accordance with their presumed youth and high-mass progenitors (e.g., Figer et al., 2005). This makes observations in soft X-rays and the optical difficult, due to the large columns of extincting material to each of the objects (Durant \\& van Kerkwijk, 2006a).  For the nearest of the magnetars, 4U 0142+61, an intriguing break was seen in the broad-band optical spectrum between the B and U bands (Hulleman et al., 2004). Due to the faintness and extinction to this source, it has not proved possible so far to further characterize the optical feature.  A source for which the extinction will be much less of an issue is \\smc, a magnetar in the Small Magellanic Cloud. This source was detected by {\\em Chandra} as a slow-spinning X-ray source, with a bright thermal or power-law spectrum (Lamb et al., 2002). Although relatively faint in the first observations by Lamb et al. and Majid et al., in later {\\em Chandra} and {\\em XMM} observations, the object had brightened somewhat towards a similar luminosity to the other Anomalous X-ray Pulsars ($\\sim2\\times10^{35}$\\,erg\\,s$^{-1}$, 0.5--10\\,keV range; McGarry et al., 2005). The AXPs (the more stable type of magnetar) all seem to have the same 2--10\\,keV luminosity ($\\sim 1\\times 10^{35}$\\,erg\\,s$^{-1}$; Durant \\& van Kerkwijk, 2006b).  In Durant \\& van Kerkwijk (2005a, hereafter DvK05), we presented evidence for the detection of an optical counterpart to \\object{CXOU J010043.1-721134} in archival Hubble Space Telescope imaging with the Wide Field/Planetary Camera 2 (WFPC2). The serendipitous detection was based on a single exposure from a survey of the SMC (\\dataset[MAST,mission=hst&dataid=U6744102R]{Tolstoy, 1999}). Although a faint source, the detection parameters and statistics of non-detections in the field pointed to it likely being a true detection.  Here we present deeper HST imaging of the field of \\smc\\ in four bands to attempt to confirm the tentative detection presented in DvK05. ",
        "conclusions": "Given the non-detection probability of $\\sim1.5$\\% estimated in DvK05, it is rather surprising that the deeper observations presented here failed to detect the same source. The estimate was based on the number of good detections in the field which would also have been considered interesting (by color), had they fallen within the positional error circle. Only one new source is detected within the positional error circle, Star 6, and this is red enough in $m_{606}-m_{814}$ to appear as a normal star on the color-magnitude diagram, Figure \\ref{CMD}  Star Y  does not appear to have varied significantly between the two observations in either filter where it was detected, and Star Z  is still consistent with a bright early-type star.  Two distinct possibilities exist for the non-detection of the counterpart to \\smc\\ proposed in DvK05: the original detection was false, or the detection was real, but the source has faded considerably in the intervening time. For the detected object to physically leave the positional uncertainty circle at the distance of the SMC, or even by a foreground white dwarf is not possible, the proper motion required would be too great.  One obvious solution is that the detection in DvK05 was, in fact, a cosmic ray or some other artifact. {\\tt HSTphot} rejects detections as cosmic rays if they are much sharper than typical stars, and this source was, if anything, more diffuse. Some objects were, however, in the same part of the color-magnitude diagram as Star X, so there is a small but non-negligible chance that this was a mere fluke. In this deeper set of observations, there remain a few detections of similar brightness in $m_{606}$ and blueness in $m_{606}-m_{814}$ as the original detection. The large scatter at the bottom of Figure \\ref{CMD} is not significant in this respect, as these sources (including the couple near the error circle) are more poorly measured than Star X appeared to be in DvK05.  Although the detection in DvK05 is clearly in doubt following these observations, it is possible that the object faded considerably in the optical, and became undetectable. The only AXP that has been detected multiple times in the optical, 4U 0142+61, does at times show large variations in flux, and the X-ray to optical flux ratio inferred in DvK05 was much larger for Star X than it had typically been for 4U 0142+61. Assuming that none of the sources detected is the counterpart to \\smc, we derive a limit on the flux ratio $f_X/f_V>114$ (with $V\\approx m_{606}$, $A_V=0.3$ from Hilditch et al., 2005; and X-ray flux in the 2--10\\,keV range from Woods \\& Thompson, 2006), which compares to a typical value of 460 for 4U 0142+61. Thus, the general consistency found between different AXPs (Durant \\& van Kerkwijk, 2005b) could still hold for this source as well.   To summarize, follow-up HST/WFPC2 observations of the field of \\smc\\ have failed to confirm our earlier tentative detection of an optical counterpart. No convincing counterpart is seen, with much better limiting magnitudes than before. The absence of a detection could either mean that the original detection was false, or that the counterpart has faded significantly. If the latter, its X-ray to optical flux ratio could now be the same as for 4U 0142+61."
    },
    "0711/0711.1980_arXiv.txt": {
        "abstract": "{The role of magnetic fields in gamma-ray burst (GRB) flows remains controversial. The study of the early afterglow phases and, in particular, of the reverse shock dynamics and associated emission offers a promising probe of the magnetization of the ejecta. In this paper, we derive the conditions for the existence of a reverse shock in arbitrarily magnetized ejecta that decelerate and interact with the circumburst medium. Both constant and wind-like density profiles are considered. We show, in contrast to previous estimates, that ejecta with magnetization $\\sigma_0\\simmore 1$ are not crossed by a reverse shock for a large fraction of the parameter space relevant to GRB flows.  Allowing for shell spreading, there is always a relativistic or mildly relativistic reverse shock forming in $\\sigma_0\\simless 0.3$ ejecta. From this, we conclude that the paucity of optical flashes, believed to be a distinctive signature of a reverse shock, may be explained by the existence of dynamically important magnetic fields in the ejecta. \\\\  {\\bf Key words:} Gamma rays: bursts -- MHD -- shock waves} ",
        "introduction": "\\label{intro}  The role of magnetic fields in the GRB flow remains uncertain. Neutrino annihilation may be the dominant process that leads to the launching of the GRB flow by forming a fireball, i.e.  a flow initially dominated by its thermal energy (Goodman 1986; Paczynski 1986). Alternatively, if they are powerful enough and develop the appropriate topology, magnetic fields can efficiently extract the rotational energy from the central engine launching a Poynting-flux-dominated flow (hereafter PDF; Usov 1992).  In the fireball model the magnetic fields are dynamically unimportant. The GRB flow during the onset of the afterglow phase is expected to be, at most, weakly magnetized.  Though model dependent, the conversion of magnetic into kinetic  energy in a PDF is partial (Michel 1969; Li et al. 1992; Vlahakis \\& K\\\"onigl 2003; Giannios \\& Spruit 2006). The flow is expected to have a significant fraction of its energy in the form of Poynting flux at large radii where there is substantial interaction with the external medium.  The study of the early phases of the ejecta interaction with the external medium can probe the strength of magnetic fields in flow. The early emission from the forward shock, for example, depends on the magnetization of the ejecta (Lyutikov \\& Blandford 2003; Zhang \\& Kobayashi 2005, hereafter ZK05; Genet et al. 2006).  Modeling of the microphysical parameters of the forward and the reverse shock supports the existence of strong magnetic fields in the ejecta (Fan et al. 2002; Kumar \\& Panaitescu 2003; Zhang et al. 2003). Furthermore, polarization measurements can provide information on the structure of the magnetic fields (Lazzati et al. 2003; Granot \\& Taylor 2005; Mundell et al. 2007).  The early optical afterglow lightcurves depend critically on the very existence of a reverse shock into strongly magnetized ejecta. A reverse shock may also contribute in a dominant way to the observed emission in later afterglow phases (Genet et al. 2007; Uhm \\& Beloborodov 2007). In this paper, we investigate the parameter space of the magnetization of the flow for which there is a reverse shock crossing the ejecta. This parameter space is described by simple analytic expressions. We show that, in contrast to previous claims, if the flow contains similar amounts of kinetic and magnetic energy at the onset of the afterglow, there is no reverse shock forming in a large fraction of GRBs. The observational implications of our findings are also discussed. ",
        "conclusions": "The role of magnetic fields in GRB flows remains uncertain. In fireball models, the GRB flow is launched with dynamically unimportant magnetic fields while MHD energy extraction can lead to a Poynting-flux dominated flow. Observations of the prompt emission alone appear hard to distinguish between fireball or PDF models\\footnote {with the possible exception of multi GeV prompt emission that may show signature of pair cascades in a neutron rich fireball that is absent in PDF flows; Koers \\& Giannios (2007).}. Early afterglow observations can, however, probe the magnetic content of the ejecta. The magnetization can leave its signature in the early forward shock (Lyutikov \\& Blandford 2003; Genet et al. 2006) and the reverse shock emission (Kumar \\& Panaitescu 2003; Zhang et al. 2003; Fan et al. 2004; ZK05).  \\subsection{Magnetization and existence of a reverse shock}  Here, we derive the condition for the existence of a reverse shock in the ejecta that contain a large scale magnetic field, assuming the ideal MHD description is adequate. The ejecta decelerate interacting with the circumburst medium. External medium with both constant and wind-like density profiles are considered. It is shown that, for high enough magnetization, the ejecta interact smoothly with the external medium and decelerate without the formation of a reverse shock. The conditions for the formation of a reverse shock are derived by demanding that the reverse shock propagates supersonically in the ejecta. The conditions are summarized by eqs.~(\\ref{ksi-sigma}) and (\\ref{rscondition}) and Fig.~1. Our results are substantially different from those of ZK05 who do not consider the internal evolution of the ejecta when deriving their criterion for the existence of a reverse shock.  One important finding of the study is that $\\sigma_0 \\simmore 1$ suffices to suppress the formation of a reverse shock for a large parameter space that characterizes GRB flows. This parameter space is larger for a constant density external medium in comparison to the wind case. This means that only a small fraction of PDFs are expected to show reverse shock emission. Even when a reverse shock forms in a high $\\sigma_0$ flow, it is not expected to result in bright emission (at least in the ideal MHD limit discussed here) since the shock dissipates only part of the kinetic energy which is a small fraction of the total energy of the shell (see, e.g., ZK05). On the other hand, allowing for radial spreading of the ejecta, practically all the flows with $\\sigma_0 \\simless 0.3$ are expected to be crossed by a mildly-relativistic or relativistic reverse shock.  This work contributes in understanding the early deceleration dynamics of magnetized ejecta. The complete evolution of the system is not considered here. Important questions related to the timescale over which the magnetic energy of the shell (which may dominate the total energy budget) is passed into the shocked external medium remain unexplored. We are currently attacking these questions through relativistic MHD simulations to be reported elsewhere (Mimica et al. 2007).   \\subsection{Connection to observations}  A clear prediction of the fireball models is that there is a reverse shock crossing the ejecta on the onset of the interaction of the ejecta with the external medium. The reverse shock emission is expected to result in a short lived, bright emission event. Although optical flashes have been observed in a few bursts (e.g. GRB990123, Akerlof et al. 1999; GRB021211, Fox et al. 2003) and are believed to be powered by the reverse shock emission, they are rare events. This may be a result of a more luminous forward shock masking the reverse shock optical emission (McMahon et al. 2006; Mundell et al. 2007b).  Our analysis points to a different possibility, i.e. that the paucity of bursts with clear  reverse shock emission is a result of the frequent reverse shock suppression. It is shown here that this happens for a large GRB parameter space in a strongly magnetized flow. Moreover, in cases where reverse shock emission is identified, modeling of the microphysical parameter connected to the magnetic field strength indicates ejecta with dynamically important magnetic fields (Fan et al. 2002; Kumar \\& Panaitescu 2003; Zhang et al. 2003; ZK05) in agreement with our suggestion.  Polarized emission, result of synchrotron emitting, (reverse) shock-accelerated particles is expected from a coherent magnetic field. Note, however, that the lack of significant optical polarization in the early afterglow does not necessarily show the absence of large scale field in the ejecta but may be result of the absence of reverse shock (suppressed by the strong field). This is particularity likely if there is no evidence for the distinct emission patterns expected from a reverse shock. In light of these considerations the claim of Mundell et al. (2007a) that the low early polarization signal in GRB060418 rules out the presence of large scale magnetic field in the ejecta is premature.  This study has focused in the ideal MHD limit. On the other hand non-ideal MHD effects can lead to additional energy release and emission during the afterglow phases of strongly magnetized flows. This emission is {\\it not} directly related to shock accelerated particles and therefore to the existence or not of a reverse shock. On the other hand, the crossing of the reverse shock can lead to forced reconnection in the ejecta that powers the GRB emission (Thompson 2006). Furthermore, the deceleration of the ejecta can lead to revival of MHD instabilities that lead to magnetic dissipation. Such late-time energy release can power the afterglow flares observed in many early X-ray afterglows (Giannios 2006a)."
    },
    "0711/0711.1387_arXiv.txt": {
        "abstract": "Radiative transfer in a relativistic accretion disk wind is examined under the plane-parallel approximation in the fully special relativistic treatment. For an equilibrium flow, where the flow speed and the source function are constant, the emergent intensity is analytically obtained. In such an equilibrium flow the usual limb-darkening effect does not appear, since the source function is constant. Due to the Doppler and aberration effects associated with the relativistic motion of winds, however, the emergent intensity is strongly enhanced toward the flow direction. This is the {\\it relativistic peaking effect}. We thus carefully treat and estimate the appearance of relativistic winds and jets, when we observe them in an arbitrary direction. ",
        "introduction": "Accretion disks are now widely believed to be energy sources in various active phenomena in the universe (see Kato et al. 2007 for a review). Relating to energetic accretion disks, accretion disk winds have been extensively examined in connection with astrophysical jets and outflows: in bipolar outflows from young stellar objects (YSOs), in outflows from cataclysmic variables (CVs) and supersoft X-ray sources (SSXSs), in relativistic jets from microquasars ($\\mu$QSOs), active galactic nuclei (AGNs), quasars (QSOs), and in gamma-ray bursts (GRBs). In particular, intense radiation fields of luminous supercritical accretion disks may be responsible for relativistic jets from super-Eddington sources, such as luminous $\\mu$QSOs, GRS~1915$+$105 and SS~433, luminous QSOs, 3C~273, and energetic GRBs (see, e.g., Fukue 2004 for references). Furthermore, energetic emissions from relativistic jets have been examined, relating to, e.g., gamma-ray blazars and gamma-ray bursts (Dermer, Schlickeiser 1993, 2002; Dermer 1998; B\\\"ottcher, Dermer 2002; Dermer et al. 2007).   In such circumstances, radiative transfer in accretion disk winds as well as accretion disks becomes more and more important. Radiative transfer in the standard disk has been investigated in relation to the structure of a static disk atmosphere and the spectral energy distribution from the disk surface (e.g., Meyer, Meyer-Hofmeister 1982; Cannizzo, Wheeler 1984). Furthermore, gray and non-gray models of accretion disks were constructed under numerical treatments (K\\v ri\\v z, Hubeny 1986; Shaviv, Wehrse 1986; Adam et al. 1988; Mineshige, Wood 1990; Ross et al. 1992; Shimura, Takahara 1993; Hubeny, Hubeny 1997, 1998; Hubeny et al. 2000, 2001; Davis et al. 2005; Hui et al. 2005) and under analytical ones (Hubeny 1990; Artemova et al. 1996; Fukue, Akizuki 2006a).  Radiative transfer in the accretion disk wind, on the other hand, has not been well considered both in the non-relativistic and relativistic regimes. For example, transformation properties of disk radiation fields in the proper frame of a relativistic jet were examined by, e.g., Dermer and Schlickeiser (2002). In these earlier works, however, the radiation fields are set to be external sources, and the radiation transfer was not considered. Recently, radiative transfer in a moving disk atmosphere was firstly investigated in the subrelativistic regime (Fukue 2005a, 2006a), and in the relativistic regime (Fukue 2005b, 2006b; Fukue, Akizuki 2006b). In these studies, however, only the radiative moments were obtained under the moment formalism, and the specific intensity was not solved. Hence, in Fukue (2007) the specific intensity from an accretion disk wind was obtained in the subrelativistic regime, where the flow speed $v$ is of the order of $(v/c)^1$.   In this paper, we thus extend the previous work. Namely, we examine radiative transfer in the relativistic accretion disk wind, which is assumed to blow off from the luminous disk in the vertical direction (plane-parallel approximation), with the relativistic speed up to the order of $c$.       In the next section we describe the basic equations. In section 3, we show analytical solutions of the specific intensity. The final section is devoted to concluding remarks. ",
        "conclusions": "In this paper we have examined the radiative transfer problem in an accretion disk wind under the plane-parallel approximation in the fully special relativistic regime. For an equilibrium flow, where the radiative quantities and source function are constant, we analytically obtain the specific and emergent intensities. We found that the emergent intensity depends on the flow speed as well as the direction cosine, and exhibits a relativistic peaking effect. As a result, a wind luminosity would be overestimated by a pole-on observer and underestimated by an edge-on observer, when we observe an accretion disk wind (cf. Sumitomo et al. 2007; Nishiyama et al. 2007).   It should be noted that the {\\it apparent} optical depth in the relativistically moving media. Abramowicz et al. (1991) pointed out that the optical depth in the relativistic flow decreases as $\\gamma(1-\\beta\\mu)\\tau$ toward the downstream direction, due to the Doppler and aberration effects. Inspecting equation (\\ref{itransf_pp3}) or solution (\\ref{i_sol13}), we find that in the present case the optical depth $\\tau$ is apparently replaced by $\\gamma (1-\\beta\\mu)\\tau$. This is just consistent with the results by Abramowicz et al. (1991).  In this paper we only examined the frequency-integrated intensity under the plane-parallel and gray approximations. It should be briefly remarked on the frequency dependence; i.e., the frequency-dependent intensity $I_\\nu$. As long as the opacity is gray, the spectral transformation is determined by the relativistic effect in the present situation. Namely, for the frequency-integrated intensity the relativistic invariant is $I/\\nu^4$, and this effect appears in solution (\\ref{i_sol13}) and other equations as a factor of $[\\gamma(1-\\beta\\mu)]^4$. Since the relativistic invariant for the frequency-dependent intensity is $I_\\nu/\\nu^3$, the corresponding factor should be changed as $[\\gamma(1-\\beta\\mu)]^3$. On the other hand, the relativistic modification in the optical depth discussed above is not changed in the frequency-dependent case. Hence, the present results may be valid for the frequency-dependent emergent intensity; e.g., the incident intensity $I_{\\nu 0}$ would be boosted in the polar direction, according to solution (\\ref{i_sol13}), but $[\\gamma(1-\\beta\\mu)]^4$ is replaced by $[\\gamma(1-\\beta\\mu)]^3$.  The present study can be applied to energetic jets in, e.g., gamma-ray blazars and gamma-ray bursts. The effect of the relativistic jets on the emergent spectrum has been studied in several literatures (Dermer, Schlickeiser 1993, 2002; Dermer 1998; B\\\"ottcher, Dermer 2002; Dermer et al. 2007), In these earlier works, however, the radiation transfer was not treated at all. Hence, the present approach may be very usefull in these fields of active phenomena.   The radiative transfer problem investigated in the present paper must be quite {\\it fundamental problems} for accretion disk physics and astrophysical jet formation. In onder to demonstrate the existence of the relativistic peaking effect, we have imposed various assumptions, including a constant flow speed, gray approximation, no heating source, and so on. By relaxing these assumptions and integrating the relativistic transfer equation numerically, we could obtain the emergent intensity and spectra more quantitatively. These are left as future works.    \\vspace*{1pc}  The author would like to thank S. Kato, S. Mineshige, M. Umemura, T. Koizumi, and C. Akizuki for enlightening and stimulating discussions. This work has been supported in part by a Grant-in-Aid for Scientific Research (18540240 J.F.) of the Ministry of Education, Culture, Sports, Science and Technology."
    },
    "0711/0711.4794_arXiv.txt": {
        "abstract": "{\\it Chandra}'s high resolution observations of radio galaxies require a revisit of the relevant electron acceleration processes. Although the diffusive shock particle acceleration model may explain spectra of spatially unresolved sources, it encounters difficulties in explaining the structure and spectral properties of recently discovered {\\it Chandra} X-ray features in several low-power radio sources. We argue that these observations strongly suggest stochastic electron acceleration by magnetized turbulence, and show that the simplest stochastic particle acceleration model with energy independent acceleration and escape timescales can overcome most of these difficulties. We use the bright core of the western hotspot of Pictor A as an example to demonstrate the model characteristics, which may be tested with high energy observations. ",
        "introduction": "In the classical diffusive shock particle acceleration model the observed relativistic electrons in radio galaxies are accelerated at the edges and knots of jets. For relatively more luminous FR II radio galaxies, electrons may also be energized in radio hotspots, where jets terminate. It is usually assumed that the jet power is initially stored in the bulk motion of outflows and is converted into the internal energy of the jet plasma by shocks at the observed radio bright features \\citep{br74}. The model requires that particles repeatedly pass the shock front, characterized by a jump in the density and velocity field, to reach very high energies. Although the physical processes responsible for the scattering and/or diffusion of particles through the shock front in most collisionless astrophysical plasmas are not well established, it has been widely cited due to its simplicity and its achievement in explaining the spectra of spatial unresolved sources \\citep[e.g.][]{hbw01}.  It however is challenged by recent high spatial and spectral resolution observations with {\\it Chandra}. Although nonthermal X-rays from radio galaxies can be produced by relativistic electrons through synchrotron, synchrotron self-Comptonization (SSC), and inverse Compton scattering of the cosmic microwave background radiation (CMBR) or other background photons, X-rays from jets of FR I galaxies and some low-power hotspots of FR II galaxies are generally attributed to synchrotron process \\citep{wbh01, h04}. If the magnetic field is nearly in energy equipartition with the relativistic electrons, the synchrotron cooling timescale of electrons producing the observed X-rays is short, and the electrons need to be accelerated very close to the observed compact X-ray features, which usually are displaced from peaks of the corresponding radio sources and occasionally do not have optical and/or radio counterparts \\citep{hbw01, hck07, m02, wy02}. The diffuse X-ray emission from jets of FR I galaxies, on the other hand, indicates a distributed particle acceleration process. The X-ray spectral hardening and relatively high X-ray flux as compared to the extrapolation from optical fluxes of some compact sources are also puzzling.  The bright core of the western hotspot of Pictor A (WHPA), a famous nearby\\footnote{At a redshift of $z = 0.035$, the luminosity distance to Pictor A is $d_{\\rm L} = 151.9$\\,Mpc, and the conversion scale is $1\\arcsec = 0.687$\\,kpc for the modern cosmology with $\\Omega_{\\rm M} = 0.27$, $\\Omega_{\\Lambda} = 0.73$, and $H_0 = 71$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$ \\citep{s03}.} FR II galaxy, is an excellent candidate to demonstrate difficulties of the shock model. \\citet{myr97} discussed its radio to optical spectrum, and \\citet{w01} reported X-ray observations with {\\it Chandra}. The latter studied in detail several possible X-ray emission processes and found that none of them can explain the broadband spectrum (BBS) naturally. They concluded that the X-rays are most likely synchrotron emission by an electron population distinct from these responsible for the radio to optical spectrum. In this picture the fact that the radio, optical, and X-ray emissions come from the same spatial region would have to be a coincidence.  In this Letter, we show how the simplest stochastic particle acceleration (SA) model % explains the BBS of WHPA. The model also predicts a spectral cutoff near the {\\it Chandra} band and an intrinsic volume of $0\\farcs03^3$ of the particle acceleration region (AR), which may be tested with future observations. In \\S\\ 2, we discuss the major difficulties of the shock model in explaining {\\it Chandra} observations and why the SA by turbulence is favorable. A simple SA model is built and applied to WHPA in \\S\\ 3. In \\S\\ 4 we draw conclusions and discuss how the model may also explain other {\\it Chandra} X-ray features. ",
        "conclusions": "We have shown that SA reproduces the X-ray spectral hardening of the bright core of WHPA and predicts a spectral cutoff near $1$ keV. With few modification, the model can be applied to other {\\it Chandra} X-ray features with similar properties. We emphasize that the shock model has difficulties in explaining these observations since the lifetime of X-ray emitting electrons is too short for the shock to produce the observed high flux levels. This challenges any models with particles accelerated efficiently in very small regions With a very large shock cross section ($>0\\farcs3^2$) and high downstream flow velocity ($>0.1c$), the shock model may fit the BBS of WHPA by assuming a quick adiabatic expansion phase of the downstream flow, similar to the SA model discussed above. However, a power-law spectrum with a photon spectral index of 2.24 is expected beyond the X-ray band except that the acceleration rate of highest energy electrons is fine-tuned to be equal to the cooling rate of X-ray emitting electrons. In this case the model predicted BBS will be indistinguishable from that of the SA model. For Bohm diffusion, the acceleration rate is equal to $eB\\Delta r^2 c/2\\pi$, the emission spectrum cuts off at $\\sim 70$ MeV \\citep{bg70}. Observations above the {\\it Chandra} band (e.g. with {\\it Suzaku}) can readily distinguish these two models.  The SA is less efficient and operates over an extended region with its size determined by the turbulence generation and decay processes. We predict that the observed harder X-ray emission originates from the AR while the radio (optical) emission comes from the CR. The energy dissipation and the accompanying SA may still be triggered by shock waves. In this case, the traditional study of shock dynamics can be used to make predictions on the displacements between peaks in the X-ray and radio (optical) images. These may explain observations of bright knots in FR I galaxies.  The dynamics of jet flows is sophisticated. Some of the energy dissipation regions may not be associated with shock fronts. In this case, one may use MHD simulations to identify regions with strong energy dissipation and apply the SA model to make predictions on the source morphology. MHD simulations may also address the dynamics related to the turbulence and the adiabatic expansion phase introduced to produce the observed spectral hardening.   \\vspace{-.4cm}"
    },
    "0711/0711.4107_arXiv.txt": {
        "abstract": "{The supernova remnant (SNR) Kes 75/PSR J1846-0258 association can be regarded as a certainty due to the accurate location of young PSR J1846-0258 at the center of Kes 75 and the detected bright radio/X-ray synchrotron nebula surrounding the pulsar. We provide a new distance estimate to the SNR/pulsar system by analyzing the HI and $^{13}$CO maps, the HI emission and absorption spectra, and the $^{13}$CO emission spectrum of Kes 75. That there are no absorption features at negative velocities strongly argues against the widely-used large distance of 19 to 21 kpc for Kes 75, and shows that Kes 75 is within the Solar circle, i.e. a distance $d<$13.2 kpc. Kes 75 is likely at a distance of 5.1 to 7.5 kpc because the highest HI absorption velocity is at 95 km/s, and no absorption is associated with a nearby HI emission peak at  102 km/s in the direction of Kes 75. This distance to Kes 75 gives a reasonable luminosity of PSR J1846-0258 and its PWN and also leads to a much smaller radius for Kes 75, so the age of the SNR is consistent with the spin-down age of PSR J1846-0258, confirming this pulsar as the second-youngest in the Galaxy. ",
        "introduction": "The supernova remnant (SNR) Kes 75 (G29.7-0.3) is one of a few examples of a shell-type remnant harboring an extremely young pulsar that powers an extended radio/X-ray synchrotron nebula. Ever since the young energetic X-ray pulsar, PSR J1846-0258, was discovered by Gotthelf et al. (2000) using the Rossi X-Ray Timing Explorer, many observations of this system at different wave-bands have been executed (Sugizaki et al. 2001, Mereghetti et al. 2002, Bock \\& Gaensler 2005, Djannati-Ata\\\"i 2007, McBride et al. 2008). These observations strongly support the SNR/pulsar association because of both the position of PSR J1846-0258 at the geometrical centers of the SNR shell and the bright radio/X-ray nebula within the SNR.  A young pulsar usually favors a small radius of the associated SNR, but Kes 75 had an estimated radius of 10 pc based on the previously estimated 19 kpc distance. Caswell et al. (1975) employed the Parkes HI-line interferometer to find the most distant reliable absorption feature at velocity of $\\sim$ 90 km/s in the direction of Kes 75, and suggested the lower-limit distance of 6.6 kpc (assuming Solar - Galactic center distance R$_{0}$= 10 kpc). This drops to 5.0 kpc using a recent measurement of R$_{0}$= 7.6$\\pm$0.3 kpc (Eisenhauer et al 2005). Subsequent distance estimates of 19 to 21 kpc for Kes 75 were obtained. Becker \\& Helfand (1984) suggested $\\sim$ 21 kpc using VLA (D array) HI absorption spectrum observations with features at negative velocity. Based on d=19 to 21 kpc, the extremely high X-ray luminosities for Kes 75 ($\\sim$ 4$\\times10^{37}$ ergs s$^{-1}$), the pulsar, and the PWN require a highly implausible explosion energy ($\\sim$ 10$^{53}$ erg), leading to several progenitor evolutionary scenarios (Morton et al. 2007, Helfand et al. 2003). The most likely one of the scenarios points to a large error in the distance estimate of $\\sim$ 19 kpc (Morton et al. 2007).  The 21 cm neutral hydrogen absorption measurements to the Galactic objects can generally provide a distance estimate with an uncertainty of $<$1 kpc. Recently we used  HI+$^{13}$CO observations to measure and revise distances to a few SNRs with high precision (Leahy \\& Tian 2008, Tian et al. 2007, Tian \\& Leahy 2008).  In this paper, we significantly revise the distance to Kes 75 using these methods. The radio data come from 1420 MHz continuum plus HI-line observations of the VLA Galactic Plane Survey (VGPS) and  the $^{13}$CO-line (J = 1-0) observations of the Galactic ring survey (Stil et al. 2006, Jackson et al. 2006).   \\begin{figure*} \\vspace{50mm} \\begin{picture}(50,50) \\put(-80,230){\\special{psfile=f1a.eps hoffset=0 voffset=0 hscale=40 vscale=40 angle=-90}} \\put(240,-20){\\special{psfile=f1b.eps hoffset=0 voffset=0 hscale=48 vscale=48 angle=0}} \\end{picture} \\caption{The 1420 MHz continuum image (left) and HI image (right) of Kes 75 from a single channel. The HI map has superimposed contours (28, 40, 60, 100 K) of the 1420 MHz continuum emission to show the SNR.} \\end{figure*}  \\begin{figure*} \\vspace{160mm} \\begin{picture}(160,160) \\put(-55,645){\\special{psfile=f2a.eps hoffset=0 voffset=0 hscale=39 vscale=33 angle=-90}} \\put(230,375){\\special{psfile=f2a1.eps hoffset=0 voffset=0 hscale=46 vscale=41 angle=0}} \\put(-20,308){\\special{psfile=f2b.eps hoffset=0 voffset=0 hscale=40 vscale=31 angle=0}} \\put(230,478){\\special{psfile=f2b1.eps hoffset=0 voffset=0 hscale=31 vscale=30 angle=-90}} \\put(-20,155){\\special{psfile=f2c.eps hoffset=0 voffset=0 hscale=40 vscale=31 angle=0}} \\put(230,325){\\special{psfile=f2c1.eps hoffset=0 voffset=0 hscale=31 vscale=30 angle=-90}} \\put(-20,-10){\\special{psfile=f2d.eps hoffset=0 voffset=0 hscale=40 vscale=31 angle=0}} \\put(200,180){\\special{psfile=f2d1.eps hoffset=0 voffset=0 hscale=40 vscale=34 angle=-90}} \\end{picture} \\caption{Two channel HI images of Kes 75 (top row). The second and third rows show the HI spectra (left) and the CO emission spectra (right), extracted from boxes 1 and 2 shown in Fig. 1. The bottom row shows an additional HI spectra(left) with the source spectrum from box 1 and background spectrum from the area at $\\l$[29.705, 29.725],$b$[-0.285, -0.310], and one-channel CO image (right). Both the HI and CO maps have superimposed the same contours as Fig. 1.} \\end{figure*} ",
        "conclusions": "We find strong evidence about the distance to Kes 75, and if disagrees with what is given by Becker \\& Helfand et al (1984). By testing different background regions, we were not able to obtain similar HI absorption spectra to those shown in that study. The VGPS data used here is superior to their data in sensitivity, and it also includes single-dish observations from the 100 m Green Bank telescope, yielding the large-scale HI and continuum emission in addition to small-scale emission. This is probably one reason we get a different, better, result.  Next we discuss implications of the shorter distance to Kes 75. The new distance gives a much smaller radius for Kes 75: the radio and X-ray sizes nearly match (Morton et al. 2007) and give a radius of 1.60 arcmin. Our new distance range gives a physical radius range of 2.37 to 3.49 pc. We write the distance to Kes 75 as $d=d_6 \\times$6 kpc, so the radius is 2.80$d_6$ pc. As estimate of the age of the SNR, we use the spin-down age of 723 years for PSR J1846-0258 from Gotthelf et al (2000). Livingstone et al (2006) use measurements of braking index of PSR J1846-0258 to put an upper limit on age of 884 yr, confirming the spin-down age estimate. Assuming free expansion, the derived expansion velocity is then 3800$d_6$ km/s with this spin-down age and a radius of 2.80$d_6$ pc. If the SNR is expanding according to Sedov evolution, the current Sedov shock velocity is 0.4 times the average expansion velocity, giving a shock velocity of 1500$d_6$ km/s. For the more realistic case, the initial phases will be dominated by free-expansion, so the current shock velocity will be higher than the Sedov estimate and lower than the free expansion estimate. However, we note that Helfand et al. (2006) estimate the expansion velocity of Kes 75 using the Si X-ray line width to obtain $\\sim$3700 km/s. This indicates that Kes 75 is probably closer to the free-expansion phase.  The electron temperature of the high-temperature component in the X-ray spectrum is 1.4-1.5 keV (Morton et al. 2007), which corresponds to a shock velocity (for electron-ion equilibration) of 1200 km/s. Since the minimum shock velocity is 1500$d_6$ km/s, there must not be full electron-ion equilibration: one finds a maximum electron-to-ion temperature ratio of 0.64. The shock velocity derived from the electron temperature increases in this case. For an electron-to-ion temperature ratio of 0.1, one obtains a shock velocity of 3800 km/s. Thus the electron-to-on temperature ratio is restricted to the range of 0.10 to 0.64 for Kes 75.  Morton et al (2007) give X-ray emitting gas masses in hot (1.4 keV)  and cool (0.23 keV) components. For our new distance, the masses are reduced to reasonable values: $5.3 f_1^{1/2}d_6^{5/2}M_{\\odot}$ for the cool component and $0.84 f_1^{1/2}d_6^{5/2}M_{\\odot}$ for the hot component with $f_1$ the filling factor. They interpret the cool component as shocked interstellar medium (ISM) or circumstellar material (CSM) and the hot component as reverse-shocked ejecta. That two temperature components are seen is clear evidence that the ejecta has not fully entered the Sedov phase. The shocked-ejecta mass is likely to be smaller than the ISM/CSM shocked mass since only a small part of the ejecta has been shocked so far. The expansion should have already slowed down somewhat from the initial ejection velocity, so the current shock velocity and electron-to-ion temperature ratio should be somewhere between the Sedov estimates (1500$d_6$ km/s and 0.64) and the free-expansion values (3800$d_6$ km/s and 0.10).  Using the Sedov model and age of 723 yr, the observed outer shock radius yields an estimate of the explosion energy ($E_0$) divided by interstellar density ($n_0$): $E_0/n_0=0.69d_6^5\\times10^{50}$erg.  The x-ray spectral fitting yields an estimate of the density for the shocked CSM/ISM component of $n=100f_1^{-1/2}d_6^{-1/2}$cm$^{-3}$ and $n=64f_1^{-1/2}d_6^{-1/2}$cm$^{-3}$ for the SE and SW limbs of Kes 75, respectively. From the inferred X-ray post-shock density ($n$) and using a compression rate of 3 (from the observed radio spectral index of 0.7, see Droge et al. 1987), we obtain a preshock density $n_0=n/3\\simeq 27 f_1^{-1/2}d_6^{-1/2}$ cm$^{-3}$ and an explosion energy of $E_0=1.9d_6^{9/2}f_1^{-1/2}\\times10^{51}$erg. This is entirely consistent with what is expected for a core-collapse explosion now, contrary to the problems caused by having a large distance to Kes 75.  For PSR J1846-0258, the pulsar-plus-nebular luminosity in the 3-10 keV band is $7.8 \\times 10^{34} d_6^2$erg s$^{-1}$, so the pulsar-plus-nebula luminosity to spin-down luminosity has a slightly lower ratio (0.6$d_6^2$) than for the Crab, rather than the previously inferred, much higher ratio.  In summary, we find observational evidence that the distance to Kes 75 is about 6 kpc and that it has a normal explosion energy and harbors a Crab-like rotation powered pulsar and pulsar wind nebula. This is unusual in  that the pulsar has a high value for magnetic field ($5\\times10^{13}$ Gauss),  and it is very young. It also exploded in a somewhat high-density environment: $n_0\\simeq$27 cm$^{-3}$.  We can estimate the absorbing column density to Kes 75 $N_{HI}$ from the HI absorption spectrum (Fig. 2) by $N_{HI}$=1.9$\\times$10$^{18}$$\\tau$$\\Delta{\\it{v}}$$T_{s}$ cm$^{-2}$ (Dickey \\& Lockman 1990). Assuming HI spin temperature $T_{s}$ of 50K, we sum over the absorption features to obtain $N_{HI}$ of $\\sim$ $3\\times$10$^{22}$ and $\\sim$ $4\\times$10$^{22}$ cm$^{-2}$ in front of regions 1 and 2, respectively. These values agree with that of $N_{H}$=4$\\times$10$^{22}$ cm$^{-2}$ derived from X-ray observations to Kes 75. Given the high column density to Kes 75, the derived extinction is $A_V \\sim 21$mag, which would explain why no historical SN has been observed associated with Kes 75 around the time of explosion in 1300 A.D."
    },
    "0711/0711.2517_arXiv.txt": {
        "abstract": "Luminous \\hbox{X-ray} outbursts with variability amplitudes as high as $\\sim 1\\,000$ have been detected from a small number of galactic nuclei. These events are likely associated with transient fueling of nuclear supermassive black holes. In this paper, we constrain \\hbox{X-ray} outbursts with harder spectra, higher redshifts, and lower luminosities than have been studied previously. We performed a systematic survey of $24\\,668$ optical galaxies in the \\chandra\\ Deep Fields to search for such \\hbox{X-ray} outbursts; the median redshift of these galaxies is $\\sim 0.8$. The survey spans $798$ days for the \\chandra\\ Deep Field-North, and $1\\,828$ days for the \\chandra\\ Deep Field-South. No outbursts were found, and thus we set upper limits on the rate of such events in the Universe, which depend upon the adopted outburst \\hbox{X-ray} luminosity. For an outburst with \\hbox{X-ray} luminosity $\\ga 10^{43}$~ergs~${\\rm s}^{-1}$ and a duration of 6 months, the upper limit on its event rate is $\\sim 10^{-4}$~galaxy$^{-1}$~yr$^{-1}$, roughly consistent with theoretical predictions. Compared to previous survey results, our harder-band and deeper survey suggests that the outburst rate may increase by a maximum factor of 10 when considering both obscured \\hbox{X-ray} outbursts and redshift evolution from $z\\sim 0$ to $z\\sim 0.8$. Our results also suggest that the \\hbox{X-ray} luminosity function for moderate-luminosity active galactic nuclei is not primarily due to stellar tidal disruptions. ",
        "introduction": "X-ray observations, mainly with \\rosat, have found about six transient, large-amplitude \\hbox{X-ray} outbursts from galactic nuclei (e.g., Donley et~al. 2002; Komossa 2002; Komossa et~al. 2004; Vaughan, Edelson, \\& Warwick 2004; and references therein). These events have variability factors of \\hbox{$\\approx 20$--1\\,000}, peak $0.5-2$~keV \\hbox{X-ray} luminosities comparable to those of local Seyfert galaxies \\hbox{($\\approx 10^{42}$--$10^{44}$~erg~s$^{-1}$)}, decay timescales of months, and soft \\hbox{X-ray} spectra. They have been observed in both active and inactive galaxies (i.e., galaxies with and without a persistently accreting supermassive black hole, respectively). The estimated event rates of these outbursts suffer from large systematic and statistical uncertainties. In nearby inactive galaxies the event rate is $\\sim 10^{-5}$~galaxy$^{-1}$~yr$^{-1}$, while in active galaxies the event rate appears to be $\\sim 100$ times higher (e.g., Donley et~al. 2002). At least some of these outbursts induce accompanying optical nuclear variability (e.g., Brandt, Pounds, \\& Fink 1995; Grupe et~al. 1995). Recently, five new candidate \\hbox{X-ray} outbursts were found in the {\\it XMM-Newton} Slew Survey \\citep{Esquej2007}, and one outburst has been detected in the ultraviolet and is thought to have the same origin as \\hbox{X-ray} outbursts \\citep{Gezari2006}.  The physical origin of nuclear \\hbox{X-ray} outbursts remains mysterious. The most likely explanation is that they are caused by transient fueling events of nuclear supermassive black holes (SMBHs). Since SMBHs are thought to be ubiquitous in nucleated galaxies (e.g., Ferrarese \\& Ford 2005 and references therein), fueling events seem inevitable in crowded galactic centers. Fueling may occur when a star, planet, or gas cloud is tidally disrupted and partially accreted (e.g., Rees 1990; Ulmer 1999; Li, Narayan, \\& Menou 2002). The predicted event rate for stellar tidal disruptions \\citep[e.g.,][]{Magorrian1999,Syer1999} is roughly consistent with the poorly constrained rate of large-amplitude \\hbox{X-ray} outbursts in inactive galaxies. However, \\citet{Wang2004} predicted a rate that was a factor of \\hbox{$\\sim 10$} higher than earlier results, mainly due to downwardly revised black-hole masses. This predicted rate, apparently in excess of the observed rate of \\hbox{X-ray} outbursts, may indicate that a large fraction of tidal-disruption events do not exhibit expected \\hbox{X-ray} outburst characteristics, perhaps due to short durations of the events or \\hbox{X-ray} obscuration. In some cases, transient fueling could be due to accretion-disk instabilities (e.g., Siemiginowska, Czerny, \\& Kostyunin 1996). Some outbursts could also perhaps be explained by \\hbox{X-ray} afterglows of gamma-ray bursts (GRBs), though no simultaneous GRBs were reported for the known \\hbox{X-ray} outbursts \\citep[e.g.,][]{Komossa99}.  Aside from the innate scientific interest of nuclear \\hbox{X-ray} outbursts, determining the rate and properties of such events is relevant to planning future missions such as the {\\it Black Hole Finder Probe (BHFP)\\/}, {\\it Lobster\\/}, and {\\it eROSITA\\/}. The {\\it BHFP\\/}, for example, is likely to observe in the hard \\hbox{X-ray} band, and the large sky coverage of this mission will allow intrinsically rare transient events to be studied (e.g., Grindlay 2005). Facilities such as the {\\it Large Synoptic Survey Telescope\\/}, the {\\it Joint Dark Energy Mission\\/}, and the {\\it Laser Interferometer Space Antenna\\/} should also allow the accompanying optical (e.g., Brandt 2005) and gravitational-wave (e.g., Kobayashi et~al. 2004) outbursts to be studied.  Previous studies of \\hbox{X-ray} outbursts from galactic nuclei have delivered fascinating results but require extension so that this phenomenon can be better understood. For example, advances are required in the following directions:  \\begin{enumerate}  \\item Outbursts with harder \\hbox{X-ray} spectra. As noted above, the outbursts discovered to date have generally had soft \\hbox{X-ray} spectra with effective \\hbox{0.2--2.0~keV} power-law photon indices of \\hbox{$\\Gamma\\approx 3$--5}. However, this may partially be a selection effect owing to the soft \\hbox{X-ray} bandpass of the \\rosat\\ satellite. Transient fueling events of SMBHs should be capable of generating harder \\hbox{X-ray} spectra \\hbox{($\\Gamma\\approx 1.7$--2.2)} via Compton upscattering in the accretion flow; such spectra are observed from most active galactic nuclei (AGNs) as well as from transient Galactic black holes. Indeed, one \\hbox{X-ray} outburst has shown evidence for spectral hardening as it declined (Komossa et~al. 2004). In addition, obscuration can harden the observed \\hbox{X-ray} spectra of an outburst, which would make it more difficult to detect in the current outburst surveys. A fairly small column density of $N_{\\rm H}=5\\times 10^{21}$~cm$^{-2}$ reduces the expected \\hbox{0.2--2.0~keV} flux from a soft-spectrum outburst by a factor of $\\sim 20$.  \\item Higher redshift outbursts. The \\hbox{X-ray} outbursts discovered to date are at low redshift with \\hbox{$z=0.01$--0.15}. Since only a small fraction of cosmic time is spanned by this redshift range and the source statistics are limited, little is known about the redshift evolution of their frequency. It is plausible that the \\hbox{X-ray} outburst rate could show evolution over cosmic time, considering that AGNs and galaxies both show strong evolution. The number densities of comparably \\hbox{X-ray} luminous AGNs evolve upward by factors of $\\approx 10$--30 out to $z\\sim 1$ (e.g.,  Brandt \\& Hasinger 2005 and references therein), and Milosavljevi\\'c, Merritt, \\& Ho (2006) have suggested that a significant portion of the \\hbox{X-ray} luminosity function of such AGNs may be comprised of sources powered by tidal disruptions.  \\item Lower luminosity outbursts. Transient phenomena in complex natural systems often follow power-law distributions in frequency of occurrence, such that low-amplitude events are more common than high-amplitude events (e.g., earthquakes, avalanches, solar flares, and ecological extinctions). It is plausible that such a distribution applies for \\hbox{X-ray} outbursts in galactic nuclei. Low-mass fueling events, such as partial tidal disruptions or the accretion of brown dwarfs, planets, or small gas clouds, would then be more common than high-mass fueling events.  \\end{enumerate}  \\noindent In this paper, we utilize data from the \\chandra\\ Deep Fields (see Brandt \\& Hasinger 2005 for a review) to constrain the rate of \\hbox{X-ray} outbursts with harder \\hbox{X-ray} spectra, higher redshifts, and lower luminosities than those studied to date. Observations of the \\chandra\\ Deep Fields were made in several well-separated and individually sensitive ``epochs,'' allowing effective constraints to be placed upon \\hbox{X-ray} outbursts that evolve on timescales of months in the rest frame. Some variability work has been performed on the \\chandra\\ Deep Field sources. For example, Paolillo et~al. (2004) studied the \\hbox{X-ray} variability of $\\sim 350$ sources detected in the \\chandra\\ Deep Field-South (\\hbox{CDF-S}), but this work has not been optimized to detect and to constrain systematically \\hbox{X-ray} outbursts of the type relevant here.  We adopt $H_0=70$~\\hbox{km s$^{-1}$ Mpc$^{-1}$}, $\\Omega_{\\rm M}$ = 0.3, and $\\Omega_{\\Lambda}$ = 0.7 throughout this paper. All coordinates are J2000. The relevant Galactic column densities are $1.3\\times 10^{20}$~cm$^{-2}$ for the \\chandra\\ Deep Field-North (\\hbox{CDF-N}; Lockman 2004) and $8.8\\times 10^{19}$~cm$^{-2}$ for the \\hbox{CDF-S} (Stark et~al. 1992). All \\hbox{X-ray} fluxes and luminosities quoted throughout this paper have been corrected for Galactic absorption using these column densities. ",
        "conclusions": "In summary, we constructed a sample of $24\\,668$ optical galaxies in the \\hbox{CDF-N} and \\hbox{CDF-S} with redshifts obtained from the literature; the median redshift is $\\sim 0.8$. We analyzed exceptionally sensitive \\chandra\\ observations of these galaxies, which span $798$ days for the \\hbox{CDF-N} and $1\\,828$ days for the \\hbox{CDF-S}. We searched for \\hbox{X-ray} outbursts with the criterion that the count rate varies by a minimum factor of 20 in one of three standard bands. No outbursts were found, and thus we set upper limits on the rate of such events in the Universe, which depend on the \\hbox{X-ray} luminosity of outbursts. If we only consider those galaxies hosting SMBHs, and those with SMBHs not massive enough to swallow a whole star without disruption, we derive an upper limit on the rate of an outburst with $L_{\\rm X, burst}\\ga 10^{43}$~ergs~${\\rm s}^{-1}$ and $T_{\\rm burst}=6$ months to be $\\sim 10^{-4}$~galaxy$^{-1}$~yr$^{-1}$ (Figure \\ref{ever}), without any other assumptions about the physical model producing \\hbox{X-ray} outbursts. Compared to the survey by Donley et~al. (2002), our survey probes both higher redshifts and harder \\hbox{X-ray} energies. The outburst rate may increase by a maximum factor of 10 when taking into account both obscured \\hbox{X-ray} outbursts and redshift evolution from $z\\sim 0$ to $z\\sim 0.8$. We are able to set constraints on low-luminosity events down to $10^{41}$~ergs~s$^{-1}$, though at this luminosity the event rate is limited by sensitivity.  We also compared our constraints to the predicted tidal-disruption rates of \\citet{Wang2004}, exploring several possibilities in physical parameter space. If the outburst luminosity is limited by the Eddington luminosity with $f_{\\rm Edd}\\la 0.1$, or the predicted rate is computed from a modified relation which does not require isothermal stellar density distributions in galactic nuclei, our results are roughly consistent with theoretical predictions. Otherwise, our constraints are tighter than the predictions. Moreover, if stellar tidal disruptions are largely responsible for the AGN \\hbox{X-ray} luminosity function at luminosities below \\hbox{$10^{43}-10^{44}$}~erg~s$^{-1}$ as proposed by \\citet{Milosavljevic2006}, the predicted rate should be scaled upward by $A(z=0.8)\\approx 10$ to account for the median redshift of our sample. This scaling leads to significant discrepancies with our observational constraints, suggesting that the \\hbox{X-ray} luminosity function at low luminosities is not likely to be dominated by stellar tidal disruptions.   The constraints set by this study could be significantly improved by further deep-field surveys and new missions. For example, the 1 Ms of additional \\hbox{CDF-S} exposure starting in 2007 September will increase the monitoring baseline significantly and provide $\\sim 13$ new observations which could be grouped into several additional epochs for outburst searching and other variability studies. Future missions such as {\\it Lobster\\/} and {\\it eROSITA\\/} in the soft \\hbox{X-ray} band and the {\\it BHFP\\/} in the hard \\hbox{X-ray} band will also be capable of detecting and studying outbursts, benefiting from their large sky coverage. \\citet{Grindlay2004} predicted that the {\\it Energetic \\hbox{X-ray} Imaging Survey Telescope (EXIST)} implementation of the {\\it BHFP\\/} would detect \\hbox{X-ray} outbursts out to $\\sim 100$~Mpc at a rate of $\\sim 30$~yr$^{-1}$, adopting the rates of \\citet{Wang2004}. Such missions will put better constraints on the outburst rate, the fraction of obscured outbursts, and redshift evolution, enhancing our knowledge about the nature of \\hbox{X-ray} outbursts."
    },
    "0711/0711.0038_arXiv.txt": {
        "abstract": " ",
        "introduction": "Gamma-ray burst GRB050922C was detected with satellites (Swift, HETE) and ground-based instruments. 22 Sep 2005 at 19:55:50 UT Swift-BAT triggered and located GRB050922C (Trigger = 156467). The light curve in Fig 1 shows the intense broad peak starting from T - 3 to T + 3 s with two sub-peaks on top. Ground-based instruments were able to measure an afterglow during a few hours both photometrically and spectrally. The Nordic Optical Telescope (La Palma) has obtained spectra of the afterglow [1]. It found several absorption features, including strong Lyman-alpha, OI+SiII, CII, SiIV, CIV, AlII and AlIII, corresponding to a redshift of $z = 2.17 \\pm 0.03$. Assuming z = 2.17 and a standard cosmology model the isotropic energy release is $E_{iso} \\sim 8\\cdot 10^{52}$ erg, the maximum luminosity is $(L_{iso})_{max} \\sim 1.6\\cdot 10^{53}$ erg/s. ",
        "conclusions": "\\begin{itemize} \\item GRB050922C may result from tidal disruption of a white dwarf star by a massive black hole. \\item We may conclude that the GRB050922C optical afterglow confirms the existence of an intermediate mass black hole, of about one thousand solar mass. \\item The periodicity in the light curve can be identified with relativistic precession of an accretion disc. \\item Both the energy release ($E_{iso} \\sim 8\\cdot 10^{52}$ erg) and the   estimate of the burst event dynamical time scale (some seconds) agree   closely with the model of the tidal disruption of a white dwarf star by a massive black hole.   \\end{itemize}"
    },
    "0711/0711.2983_arXiv.txt": {
        "abstract": "{ We discuss correlations in azimuthal angle as well as correlations in two-dimensional space of transverse momenta of two jets as well as photon and jet. Some $k_t$-factorization subprocesses are included for the first time in the literature. Different unintegrated gluon/parton distributions are used in the $k_t$-factorization approach. The results depend on UGDF/UPDF used. The collinear NLO $2 \\to 3$ contributions dominate over $k_t$-factorization cross section at small relative azimuthal angles as well as for asymmetric transverse momentum configurations. } ",
        "introduction": "\\label{intro}  The subject of jet correlations is interesting in the context of recent detailed studies of hadron-hadron correlations in nucleus-nucleus \\cite{RHIC_correlations_nucleus_nucleus} and proton-proton \\cite{RHIC_correlations_proton_proton} collisions. Effects of geometrical jet structure were discussed recently in Ref.\\cite{Levai}. No QCD calculation of parton radiation was performed up to now in this context. Before going into hadron-hadron correlations it seems indispensable to better understand correlations between jets induced by the QCD radiation. Here we discuss the case of elementary hadronic collisions. Our analysis is a first step towards the nuclear case.  In leading-order collinear-factorization approach jets are produced back-to-back. These leading-order jets are therefore not included into correlation function, although they contribute a big ($\\sim \\frac{1}{2}$) fraction to the inclusive cross section. The truly internal momentum distribution of partons in hadrons due to Fermi motion (usually neglected in the literature) and/or any soft emission would lead to a decorrelation from the simple kinematical configuration.  In the fixed-order collinear approach only next-to-leading order terms lead to nonvanishing cross sections at $\\phi \\ne \\pi$ and/or $p_{1,t} \\ne p_{2,t}$ (moduli of transverse momenta of outgoing partons). In the $k_t$-factorization approach, where transverse momenta of gluons entering the hard process are included explicitly, the decorrelations come naturally in a relatively easy to calculate way. In Fig.\\ref{fig:kt_factorization_dijets_diagrams} we show diagrams included in our calculations which illustrate the physics situation. The soft emissions, not explicit in our calculation, are hidden in model unintegrated parton (gluon) distribution functions (UPDF,UGDF). In our calculation UGDFs or UPDFs are assumed to be given and are taken from the literature.  The $k_t$-factorization was originally proposed for heavy quark production \\cite{original_kt_factorization}. In recent years it was used to describe several high-energy processes, such as total cross section in virtual photon - proton scattering \\cite{UGDF_HERA}, heavy quark inclusive production \\cite{BS00,LSZ02}, heavy quark -- heavy antiquark correlations \\cite{LS04,LS06}, inclusive photon production \\cite{LZ05_photon,PS07_inclusive}, inclusive pion production \\cite{szczurek03,CS05}, Higgs boson \\cite{Higgs} or gauge boson \\cite{KS04} production and dijet correlations in photoproduction \\cite{SNSS01} and hadroproduction \\cite{LO00}.    Up to now no theoretical calculation for photon-jet were presented in the literature, even for elementary collisions. In leading-order collinear-factorization approach the photon and the associated jet are produced back-to-back. If transverse momenta of partons entering the hard process are included, the transverse momenta of the photon and the jet are no longer balanced and finite (non-zero) correlations in a broad range of relative azimuthal angle and/or in lengths of transverse momenta of the photon and the jet are obtained. The finite correlations can be also obtained in higher-order collinear-factorization approach \\cite{Berends}.  Here we discuss the region of relatively semi-hard jets/photons, i.e. the region related to the recently measured hadron-hadron correlations at RHIC and photon-hadron correlations being analysed \\cite{JanRak}. Here the resummation effects may be expected to be important. The resummation physics is addressed in our case through the $k_t$-factorization approach.  This presentation is based on recent publications of the authors \\cite{SRS07,PS07}. Here we discuss only the main idea and present some representative results. More details can be found in Refs. \\cite{SRS07,PS07}. ",
        "conclusions": "\\label{conclusions}   Motivated by the recent experimental results of hadron-hadron correlations at RHIC we have discussed jet-jet and photon-jet correlations.  In comparison to recent works on dijet production in the framework of $k_t$-factorization approach, we have included two new mechanisms based on $gq \\to gq$ and $qg \\to qg$ hard subprocesses. This was done using the Kwieci\\'nski unintegrated parton distributions. We find that the new terms give significant contribution at RHIC energies. In general, the results of the $k_t$-factorization approach depend on UGDFs/UPDFs used, i.e. on approximation and assumptions made in their derivation.  An interesting observation has been made for azimuthal angle correlations. At relatively small transverse momenta ($p_t \\sim$ 5--10 GeV) the $2 \\to 2$ subprocesses, not contributing to the correlation function in the collinear approach, dominate over $2 \\to 3$ components. The latter dominate only at larger transverse momenta, i.e. in the traditional jet region.  The results obtained in the standard NLO approach depend significantly whether we consider correlations of any jets or correlations of only leading jets. In the NLO approach one obtains $\\frac{d \\sigma}{d \\phi_{-}}$ = 0 if $\\phi_{-} < \\frac{2}{3} \\pi$ for leading jets as a result of a kinematical constraint. Similarly $\\frac{d \\sigma}{d p_{1,t} d p_{2,t}}$ = 0 if $p_{1,t} > 2 p_{2,t}$ or $p_{2,t} > 2 p_{1,t}$. In this presentation we have discussed explicitly only a similar case of photon-jet correlations.  There is no such a constraint in the $k_t$-factorization approach which gives a nonvanishing cross section at small relative azimuthal angles between leading jets and transverse-momentum asymmetric configurations. We conclude that in these regions the $k_t$-factorization approach is a good and efficient tool for the description of leading-jet correlations. Rather different results are obtained with different UGDFs which opens a possibility to verify them experimentally.  On the contrary, in the case of correlations of any unrestricted jets (all possible dijet combinations) the NLO cross section exceeds the cross section obtained in the $k_t$-factorization approach with different UGDFs. This is therefore a domain of the standard fixed-order pQCD. We recommend such an analysis as an alternative to study leading-jet correlations.  Consequences for particle-particle correlations, measured recently at RHIC, require a separate dedicated analysis.  We have discussed also photon-jet correlation observables. Up to now such correlations have not been studied experimentally. As for the dijet case we have concentrated on the region of small transverse momenta (semi-hard region) where the $k_t$-factorization approach seems to be the most efficient and theoretically justified tool. We have calculated correlation observables for different unintegrated parton distributions from the literature. Our previous analysis of inclusive spectra of direct photons suggests that the Kwieci\\'nski distributions give the best description at low and intermediate energies. We have discussed the role of the evolution scale of the Kwieci\\'nski UPDFs on the azimuthal correlations. In general, the bigger the scale the bigger decorrelation in azimuth is observed. When the scale $\\mu^2 \\sim p_t^2$(photon) $\\sim p_t^2$(associated jet) (for the  kinematics chosen $\\mu^2 \\sim$ 100 GeV$^2$) is assumed, much bigger decorrelations can be observed than from the standard Gaussian smearing prescription often used in phenomenological studies.  The correlation function depends strongly on whether it is the correlation of the photon and any jet or the correlation of the photon and the leading-jet which is considered. In the last case there are regions in azimuth and/or in the two-dimensional ($p_{1,t}, p_{2,t}$) space which cannot be populated in the standard next-to-leading order approach. In the latter case the $k_t$-factorization seems to be a useful and efficient tool.  At RHIC one can measure jet-hadron correlations only for not too high transverse momenta of the trigger photon and of the associated hadron. This is precisely the semihard region discussed here. In this case the theoretical calculations would require inclusion of the fragmentation process. This can be done easily assuming independent parton fragmentation method.  \\vskip 0.5cm  {\\bf Acknowledgments.} This presentation is devoted to the memory of Jozs\\'o Zim\\'anyi, fantastic physicist, extraordinary man and very good friend of the Polish Nuclear and Particle Physics Community. A.S. congratulates the Budapest group for organizing the fantastic memorial workshop. We are also indebted to Jan Rak from the PHENIX collaboration for the discussion of recent correlation results from RHIC. This work was partially supported by the grant of the Polish Ministry of Scientific Research and Information Technology number 1 P03B 028 28."
    },
    "0711/0711.0992_arXiv.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "\\vspace{-6pt}  The calculation presented here is the first calculation of the photon-graviton vacuum polarization in a constant electromagnetic field, and also the first state-of-the-art application of the ``string-inspired'' worldline formalism to an amplitude involving gravitons. Although it was not possible here to go into detail, it should be emphasized that in this formalism this calculation is only moderately more difficult than the photon-photon polarization in the field. Moreover, we have also shown that the properties of the photon-graviton polarization tensor are very similar to the ones of the photon-photon one. We expect that even the graviton-graviton case will be quite feasible in this formalism. In a future sequel, we intend to analyze this case at the same level of the photon-graviton one, and to study the complete one-loop photon-graviton dispersion relations (\\ref{disploop})."
    },
    "0711/0711.0395_arXiv.txt": {
        "abstract": "We have performed uniform analysis of a sample of 54 nearby, early-type galaxies observed with \\emph{Chandra}. In this work we present the spectral results for both the diffuse Interstellar Medium, ISM, and low-mass X-ray binaries, LMXBs.  We determine the metallicity of the hot ISM in the 22 brightest galaxies and find a narrow range of abundance ratios relative to iron. The average iron metallicity of these bright galaxies is $0.96\\pm0.33$ relative to solar. By assuming these enrichment patterns continue to fainter galaxies, we are able to extend accurate ISM modeling down to the faintest galaxies in the sample. The sample of galaxies span 4.5 orders of magnitude in X-ray luminosity starting at $L_X=10^{38}\\, erg \\cdot s^{-1}$ and the average gas temperature in the sample is $0.58\\pm0.24\\, keV$.  We present the X-ray properties of these galaxies scaled to one effective radius as well as radial variations of gas and stellar binary luminosities, and radial variations of gas temperature, metallicity, mass, entropy and density. ",
        "introduction": "Ever since the discovery that most early-type galaxies host hot, X-ray emitting gas halos \\citep{Forman_Einstein}, attempts have been made to understand the origin and evolutionary fate of this gas and how this relates to the formation history of the galaxy itself. This gas provides a probe of the gravitational potential, and signatures of recent tidal or cluster interactions can be detected by examining the mass and temperature of the gas. Although great progress has been made, after \\emph{Einstein}, \\emph{ROSAT}, and \\emph{ASCA}, and the first years of \\emph{Chandra} and \\emph{XMM}, questions remain and the discord between the properties of the interstellar medium (ISM) detected in X-rays and the predicted characteristics from other waveband data has not been adequately resolved. For example, the metallicity of the ISM is expected to be at least solar since the stars shedding their envelopes into the ISM have optically measured metallicities near solar (e.g. \\citet{Trager_2,Trager_1,Forbes_Ages}), but \\emph{ASCA} data were unable to discriminate between extremely sub-solar metallicity \\citep{ASCA_NGC4636_Mushotzky}, solar metallicity \\citep{ASCA_NGC4636} or even two plasmas with different temperatures and higher metallicities \\citep{Buote_2-Temps_1,Buote_2-Temps_2}. Irrefutable metallicity determinations from \\emph{XMM's} Reflection Grating Spectrometer (RGS) have shown the metallicity of several, nearby and bright galaxies to be near solar metallicity, and with non-solar metal-to-metal ratios \\citep{XMMRGS_NGC4636,XMMRGS_NGC5044}, but these results have not yet been reproduced from the \\emph{Chandra} imaging data, which does a better job of isolating the diffuse gas than \\emph{XMM's} slitless RGS. It is also not clear if the two galaxies examined are representative of all nearby early-type galaxies.   The greatest disparity between the inferred properties of the gas from X-ray observations and other waveband observables concerns cooling flows. While there is some debate about the metallicity of the gas, the density of the gas is well determined, typically  $\\sim10^{-2}n_{e}cm^{-3}$ in the center of an early-type galaxy. It is a straightforward observation to determine the density as the emission measure varies as $n_{e}^{2}$, although, this implies a cooling time of $10^{7}yrs$ which is at odds with the observed equilibrium of the gas and a wealth of other waveband observations (e.g. the lack of star formation seen in normal early-type galaxies). Recently, the \\emph{FUSE} observatory has been used to detect cooling gas \\citep{Joel_OVI_FUSE}, but at much lower rates than implied by the densities observed in the X-rays. Many of these galaxies have been surveyed for {[}O II{]} emission (e.g. \\citet{Caon_paper_I}) and FIR emission (e.g. \\citet{Goudfrooij_2}) in efforts attempting to detect hidden star formation, but there are no strong detections. This non-detection of cooling likely means an additional heating source, but all the likely candidates (SNe and/or AGN) have inconsistent rates, energetics, spatial, or spectral distributions from the needed missing heat source.  Observations of individual galaxies provide important contributions to the field, especially when new satellites are launched, however, the baseline of knowledge of the hot gas in early-type galaxies is derived from large (sample size $\\gtrsim$ 25) studies (See Table \\ref{Large Surveys} for large X-ray studies of early-type galaxies). In fact, it was through the large studies that the existence of the hot gas was first discovered as a universal phenomenon in early-type galaxies \\citep{Einstein_2}. These large X-ray studies seek to address one or more of the following goals: 1) infer the physical source(s) of the X-ray emission, 2) calculate accurate global plasma temperatures, densities and metallicities and measure radial variations of these quantities, 3) characterize the morphology of the hot gas via surface brightness profiles, 4) determine the scaling relations between global galaxy properties: X-ray luminosity, temperature, stellar velocity dispersion and optical stellar luminosity, 5) detect any influences the environment may have on the X-ray properties. Depending on the technical capabilities of the X-ray mission employed, the large surveys had differing levels of success in addressing each of these goals.  \\citet{Trevor_ROSAT_study} notes that an order of magnitude disparity exists in the methods for statistically removing the point source contribution and isolating the diffuse emission from the low-mass, X-ray binaries, LMXBs. The spatial resolution of \\emph{Chandra} allows for the isolation of the point sources from the diffuse emission which has been the major source of uncertainty in previous studies of the hot ISM in fainter galaxies.  A large study with the \\emph{Chandra} satellite provides significant advances in addressing the large study goals due to its superb spatial resolution and moderate spectral resolution. As of 2003, no large study of nearby, early-type galaxies have been published with \\emph{Chandra} or \\emph{XMM-Newton} observatories despite the maturity of the archives. Although \\emph{XMM} has a larger collecting area, it is the fine spatial resolution of \\emph{Chandra} that will provide the next critical step in understanding X-ray emission from early-type galaxy as the majority of the ``containmenating'' stellar binary component can finally be removed from gas studies.  Motivated by the scientific returns of a large study of early-type galaxies with \\emph{Chandra} and the availability of over five dozen galaxies in the public archive, we conducted an extensive program investigating all components of the X-ray emission in early-type galaxies. In Section \\ref{Chandra Sample} the sample is defined and limitations of a \\emph{Chandra} archival study are discussed. In Section \\ref{Chandra Data} the data reduction process is described, including efforts taken to correct the known calibration issues. Section \\ref{Analysis} discusses the spectral analysis method and the success of its application to the data. Section \\ref{Chandra Results Spectral} presents the results of the spectral fitting of the data and, where sufficient counts are available, radial variations of the physical parameters fitted in the models. In Section \\ref{Chandra Results Cumulative} global X-ray properties of the sample are presented. We compared the luminosities and temperatures derived with \\emph{Chandra} to \\emph{ROSAT} in Section \\ref{ROSAT Comparison}.   ~  ~ ",
        "conclusions": "We present spectral results from uniformly processed and analyzed \\emph{Chandra} ACIS-S3 observations of 54 nearby galaxies. This sample paves the way for complete sample studies and provides a solid basis for planning future early-type galaxy observations. In the forthcoming papers we will explore some of the implications of this archival study."
    },
    "0711/0711.0506_arXiv.txt": {
        "abstract": "Following a review of our present knowledge about blue subdwarf stars in globular clusters, we present an overview of the results of searches for close binaries among these stars, including results previously published in the literature and reporting recent and preliminary results of new data. Previous investigations revealed a lack of close systems in NGC\\,6752, which we confirm with new, more extensive observations. Our estimate of the close binary fraction in this cluster is only 4\\%. From a review of the relevant literature, there are indications that a low close binary fraction among extreme horizontal branch (EHB) stars is a common feature in globular star clusters. On the other hand, the field EHB population shows evidence of a remarkably high binary fraction. Such a difference among globular cluster and field populations, although not yet explained in detail on the basis of theoretical models, must necessarily be related to different formation histories for EHB stars in the field and in clusters. In this framework, preliminary results indicate that close systems could be relatively common in the peculiar globular cluster NGC\\,2808, although the sample of studied stars is still small. This would imply that not all clusters share the same behavior, as far as EHB star formation is concerned. We briefly explore possible explanations for these results. ",
        "introduction": "The horizontal branch (HB) is a well-known feature in the color-magnitude diagram (CMD) of Galactic globular clusters (GCs), first discovered by \\citet{Bruggencate27} and later identified as a population of stars which, after the evolution along the red giant branch (RGB), eventually ignite He-burning in the core \\citep{Hoyle55}.  \\citet{Greenstein71} and \\citet{Caloi72} identified the hotter ($T_\\mathrm{eff}\\geq \\,20,\\!000$~K) HB extension with the already known field B-type subdwarf (sdB) stars, which as a consequence are also referred to as extreme horizontal branch (EHB) stars.  Because of their intrinsic faintness at optical wavelengths, sdBs in GCs were hardly observed at first \\citep{Cox61}, and only later did it become evident that they were present in many GCs, even at high metallicities \\citep[e.g.][]{Rich97}. Thus they became the most prominent evidence of the well-known {\\it second parameter problem}, i.e. the large degree of variation in HB morphology partly explained by metallicity \\citep[the ``first parameter'';][]{Sandage60} but not completely accounted for by it \\citep{Sandage67,VanDenBergh67}. \\citet{Rood73} pointed out that the spread in temperature along the HB implies a corresponding spread in envelope mass, bluer stars being less massive, and EHB stars must have lost almost their entire envelope during their evolution. The retained envelope is so thin (about $0.02\\,M_\\odot$) that it is insufficient to sustain later evolution along the asymptotic giant branch (AGB), and after He exhaustion in the core EHB stars are predicted to evolve directly to the white dwarf cooling sequence \\citep[see][for a recent review of low-mass stellar evolution]{Catelan07d}. Investigations have accordingly mainly been focused on finding the mechanism(s) responsible for the heavy RGB mass loss required for EHB formation.  \\subsection{Models for EHB Star Formation} \\label{introformation}  Many canonical and non-canonical models have been proposed over the past several decades, either considering EHB stars as distinct populations or just as the hotter components of a continuous HB. Some very recent results allow us to focus our attention on some of these scenarios \\citep[see also][for a discussion in the context of the ultraviolet upturn phenomenon in galaxies]{Catelan07a}.  The recent discovery of a planet orbiting a field sdB star \\citep{Silvotti07} adds new emphasis on the hypothesis that a sub-stellar body interacting with the expanding envelope of an RGB star can cause heavy mass loss, thus leading to the formation of an sdB star \\citep{Soker98}. Searches for planets around main sequence (MS) stars in GCs have been fruitless thus far \\citep{Gilliland00,Weldrake07}; indeed, \\citet{Soker07} argue that similar such searches are hopeless, because of the long predicted orbital periods.  A super-solar He abundance in the stellar envelope can enhance the mass loss along the RGB by increasing the RGB tip luminosity, thus leading to the formation of an sdB star. Non-canonical phenomena were invoked including He mixing driven by internal rotation \\citep{Sweigart79,Sweigart97} and dredge-up induced by H-shell instabilities \\citep{VonRudloff88}, but the recent discovery of He-enriched sub-populations in $\\omega$~Cen = NGC\\,5139 \\citep{Bedin04,Piotto05} and probably NGC\\,2808 \\citep[as suggested by][]{Piotto07,DAntona05} suggests that the He enhancement could be primordial. Such He enhancement would also affect the HB morphology \\citep[e.g.,][]{Lee05,DAntona05}.  Unfortunately, a He enhancement would have little consequence on the stellar surface parameters in this high-temperature regime, independently of its origin, and observed He abundances are altered by atmospheric phenomena like gravitational settling.  Therefore, spectroscopic measurements cannot help decide between the EHB progeny of He-enriched or He-normal stars \\citep[e.g.,][]{Moehler00,Moni07}. Note that there do exist even hotter stars along the extension of the HB, the so-called ``blue hook stars'' \\citep[e.g.,][]{DCruz00}, which will be discussed below. To the best of our knowledge, the crucial phenomenon suspected to be responsible for their formation has yet to be properly incorporated into the CMD simulations that are usually carried out in the framework of the increased He scenario.  In the dense environment of GCs strong dynamical interactions are expected, and these have often been proposed as the cause of sdB formation through stellar collisions, merging, or encounters involving binary stars, which as a consequence can harden \\citep{Heggie75} and eventually merge \\citep{Bailyn92}.  In support of this hypothesis, \\citet{Fusi93} and \\citet{Buonanno97} showed that more concentrated or denser clusters tend to have bluer HB types and longer blue HB ``tails'', while \\citet{Djorgovski92} observed a relation between the dynamical history of GCs and their UV flux, a signature of a hot population that could include sdBs. In some clusters color gradients, with bluer centers, have been identified but, mainly due to low number statistics because the light is dominated by very few bright giants, it is not clear if this feature is a general property of GCs with HB blue tails \\citep{Piotto88,Djorgovski01}  Nevertheless, other observational constraints put this picture in doubt. In fact, one would expect to find EHB stars more concentrated toward cluster centers where interactions are stronger, but no radial gradient has been found in GCs with the most strongly populated EHBs, such as NGC\\,2808 \\citep{Bedin00}, $\\omega$~Cen \\citep{DCruz00}, NGC\\,6388 and NGC\\,6441 \\citep{Rich97}.  On the other hand, it should be noted that blue stragglers, which are also supposed to be formed through stellar encounters, are indeed more concentrated in the central regions of all GCs \\citep{Piotto04}. Moreover, the presence of pairs of clusters that are dynamically similar and yet have very different HB morphology \\citep[e.g., M\\,3/M\\,13][]{Ferraro97b}, or similar in HB morphology but dynamically very different \\citep[e.g., M\\,15/NGC\\,288][]{Crocker88} also argue against dynamical interactions as an important channel for EHB star formation.  \\subsection{Results on EHB stars} \\label{introbservations}  As previously noted, there is a large difference between the apparent magnitudes of field sdB stars and their GC counterparts. The nearest sdBs in the Galactic field reach a visual magnitude $V=10-11$, whereas in GCs they are much farther away, and are never brighter than about $V=17$. To date, the study of EHB stars in GCs has been limited to their photometric properties or low-resolution spectroscopy in few GCs, and many results achieved on field sdBs have not been feasible on EHB cluster stars yet. The most frequent targets of investigations are the nearby GC NGC\\,6752, and the very massive and peculiar (due to their multimodal HB morphologies and multiple MS) GCs $\\omega$~Cen and NGC\\,2808. As a matter of fact, $\\omega$~Cen is not a typical GC at all, and it has been repeatedly proposed as the nucleus of a merged dwarf galaxy \\citep*[][and references therein]{Altmann05}, resembling in this sense the case of M\\,54 (NGC\\,6715) in the Sagittarius dwarf spheroidal galaxy \\citep[][and references therein]{Layden00}.  One important feature in the CMD of many GCs is the often underpopulated region separating sdBs from cooler HB stars. \\citet{Ferraro98} indicated the presence of a gap at about $T_\\mathrm{eff}=18,\\!000$~K, while \\citet{Piotto99} argued for a gap at constant stellar mass, with $T_\\mathrm{eff}$ varying with the metallicity of the cluster \\citep[see][for a more recent discussion]{Momany04}.  The presence of this gap suggested to many authors that sdB stars could be a distinct population with a different formation history compared to the other HB stars  \\citet{Momany02} discovered that EHB stars hotter than $T_\\mathrm{eff}\\approx 23,\\!000$~K deviate from canonical tracks in the CMD of NGC\\,6752, appearing brighter than expected in $U$. They argue that this feature could be due to the onset of radiative levitation after the disappearance of the He\\,{\\sc ii} convective layer, and this so-called ``Momany jump'' was later identified also in other clusters \\citep{Momany04}.  Unfortunately, no high-resolution spectroscopic analysis has probed this phenomenon so far, and in general there is still a lack of atmospheric abundance measurements for EHB stars in GCs. Support for their hypothesis comes from low-resolution observations by \\citet{Moehler00}, who measured a very low He abundance for some targets hotter than $23,\\!000$~K ($\\log{\\mathrm{[He/H]}}\\leq -3$), and a slight depletion for cooler EHB stars. On the contrary, \\citet{Moni07} failed to reproduce these results, finding an almost constant He underabundance regardless of temperature ($\\log{\\mathrm{[He/H]}}\\approx -2$).  On the other hand, among stars hotter than the Momany jump they find puzzling evidence of two distinct groups of stars, with different photometric and spectroscopic properties: the brighter stars are well reproduced by canonical theoretical models, while the remainder show unphysically high masses, which these authors attribute to a mismatch between models and real stars.  Two classes of rapidly pulsating stars are known among field sdB stars \\citep{Kilkenny97,Green03}, and they are currently the subject of extensive studies. Asteroseismological techniques are successfully used to analyze their oscillations and derive their physical parameters \\citep{Randall05}. The detection of such pulsators in GCs is a technically challenging issue, which however holds the promise to unveil the physical mechanism(s) that lead to the production of the EHB (and blue hook) stars. Surveys have proven fruitless until now \\citep{Reed06}, but the search is not hopeless and new projects have recently been undertaken \\citep{Catelan07c}.  \\subsection{Blue Hook Stars} \\label{introbluehook}  Over the past few years, the discovery of very faint EHB stars in the UV CMDs of $\\omega$~Cen \\citep{Whitney98} and NGC\\,2808 \\citep{Brown01} has drawn considerable attentio.  These very hot stars form a so-called ``blue hook'' at the blue end of the HB, and are hotter than the canonical theoretical limit for EHB stars. They were proposed as the progeny of stars which -- due to extreme mass loss -- did not ignite He at the tip of the RGB, but only later along the WD cooling sequence \\citep[``late hot flashers'', e.g.,][]{Castellani93,DCruz96,Brown01}. Spectroscopically measured temperatures and He abundances are consistent with the outlined picture \\citep{Moehler02,Moehler04}. Detailed calculations showed that EHB stars formed through a late He flash should show higher He and C abundances \\citep{Cassisi03,Lanz04} as a consequence of mixing events during the flash, at variance with the primordial He enhancement scenario for which no C overabundance is expected and He should not exceed the primordial (although enhanced) value.  Detailed spectroscopic observations should therefore be able to tell between the competing hypotheses. Very recent results by \\citet{Moehler07} strongly support the late-flash scenario. Blue hook stars have also been found in NGC\\,6388 and NGC\\,6441 \\citep{Busso07}, NGC\\,2419 \\citep{Ripepi07}, and M\\,54 \\citep{Rosenberg04,Momany04,Siegel07}, and a star spectroscopically studied in M\\,15 \\citep{Moehler95} is now recognized as a possible blue hook candidate on the basis of its temperature and He abundance; however, blue hooks are completely lacking in many other clusters.  In NGC\\,6752, for example, the EHB ends at temperatures adequately predicted by standard models, and no blue hook extension is observed. If these stars are the progeny of late flashers, it is unclear why they occur only in the most massive GCs \\citep[the clusters named above are among the most massive in the Galaxy; see also] [for recent evidence pointing to cluster mass as a strong second-parameter candidate]{Recio06}.  \\subsection{Binary sdB Stars} \\label{introbinaries}  Binary interactions are supposed to play a major role in the formation of sdB stars, as first proposed by \\citet{Mengel76}. The most recent models of binary evolution provide a satisfactory explanation for both the origin and observed properties of these stars \\citep{Han02,Han03}. Models in which sdBs arise from binary interaction have also been used to successfully model the observed UV upturn in external galaxies \\citep{Han07}, which is characterized by an increase in UV flux beyond the predictions of canonical population synthesis models, and which is likely to originate from a sdB star population \\citep[see for example][ and Catelan 2007a for a recent review]{Greggio90}.  On the observational side, numerous surveys have pointed to the presence of a large population of binaries among field sdB stars, although the results often disagree on the exact binary fraction, probably because of selection effects that affect the samples to different extents. Early investigations focused on the IR flux excess arising from a MS companion unseen at optical wavelengths, and estimated that 50-66\\% of field sdBs reside in sdB+MS systems \\citep{Ferguson84,Allard94,Ulla98,Aznar01}. These surveys brought no information about binary periods, were blind to compact companions, and limited to a lower limit in MS companion mass for the detection. More recently \\citet{Reed04} found a much lower fraction for sdB+MS systems (20\\%), and they argued that previous results were strongly affected by biased samples.  Later studies relied on radial velocity (RV) measurements for binary detection.  They are more sensitive to short-period systems and higher-mass companions, but searches for wide binaries are ongoing \\citep{MoralesRueda06}.  \\citet{Maxted01} found many RV variables among their sdB targets, mainly with unseen white dwarf companions (sdB+WD binaries), and estimated that nearly 70\\% of sdBs are close systems.  \\citet{Napiwotzki04} found a much lower but still fairly high fraction (about 40\\%), and many investigations of orbital parameters pointed out that close binaries with period $P\\leq 10$~d are very common among field sdB stars \\citep{Moran99,Saffer98,Heber02,MoralesRueda04}: in fact, such systems are suspected to constitute nearly half of the total Galactic sdB population.  The success of theoretical models and the results of observational investigations have led to the generally accepted conclusion that sdB stars are strictly related to binary interactions, with close systems playing a major role amongst them. It is only over the past few years that similar investigations have started being carried out in globular star clusters as well. These have led to some dramatic, and indeed quite unexpected, results. ",
        "conclusions": "With our new investigation we confirm that close binaries are indeed very rare among EHB stars in NGC\\,6752.  Our estimated upper limit for the close binary fraction agrees well with previous results, and the most probable value is small (4\\%).  The difference between field sdB stars and their analogs in this cluster is evident. Still nothing can be said about other types of binaries found among field sdBs, i.e. wide binaries and short-period sdBs with very low-mass companions. We are working to fill this gap, but investigations of more clusters are also needed.  There are indications that close EHB systems are rare also in two other GCs investigated so far, namely M\\,80 and NGC\\,5986. This points to a difference between the formation of sdB stars inside and outside GCs, and any successful model must explain this observational constraint.  Close binary sdB stars are considered the progeny of systems that underwent one or two common envelope (CE) phases \\citep{Han02}. Accordingly, this mechanism should not be efficient inside GCs, contrary to what is observed in the field. Our present knowledge still cannot clarify the role of other binary channels that are theoretically expected to produce sdBs, such as the stable Roche Lobe Overflow (RLOF) channel (which produces wide-orbit binaries) and the merger channel (whose progeny are single sdB stars).  Moni Bidin et al. \\citetext{2007, submitted} propose that among different sdB populations an $f$-age relation should exist, because the CE channel should be less efficient with increasing age. Detailed searches for long-period EHB binaries could test their hypothesis, but wide binaries could also be lacking, if either the merger channel or single-star evolution are the main sdB formation mechanisms inside GCs. In fact, the merger channel should be more efficient for old populations \\citep[][their Fig.~7]{Han07}, and it could also be enhanced by stellar interactions \\citep{Bailyn89}. Wide binaries, on the other hand, can be easily disrupted in dense environments \\citep{Heggie75}. In any case, \\citet{Moni06a} point out that the lack of a radial gradient for EHB stars in GCs (as discussed in \\S1) puts their binary origin in doubt, because typical systems as observed in the field should exceed the turnoff mass (about $0.7-0.8\\, M_\\odot$), and thus migrate toward the central regions due to mass segregation effects. Moreover, we would expect that members of disrupted systems or sdBs formed through stellar encounters and collisions would be a kinematically hotter population, but in our data we find no evidence of a difference with respect to other HB stars.  It is well known that even MS and RGB binaries are much more frequent in the field than in GCs \\citep[see for example][for a recent study]{Sollima07}, but this would not suffice to account for the difference in EHB binary fraction, as it would imply that sdB formation and binarity are unrelated~-- at variance with our present understanding of these stars. Moreover, the binary fraction among MS field stars could even be noticeably lower than has been assumed so far \\citep{Lada06}.  NGC\\,2808 could be a remarkably different case. If our preliminary results are confirmed, this cluster would be relatively rich in close EHB binaries, in contrast to the other clusters studied so far. It is tantalizing to relate this discrepancy to other well-known peculiarities of the cluster, as its multimodal HB \\citep{Sosin97} or its triple MS \\citep{Piotto07}, but our study is still in too early a stage and the sample analyzed is too small to allow definitive conclusions."
    },
    "0711/0711.2818_arXiv.txt": {
        "abstract": "The linear stability of electrically driven flow of liquid metal in circular channel in the presence of vertical magnetic field is studied. It is shown that the instability threshold of such flow is determined by magnetorotational instability of non-axisymmetric modes ($m\\neq0$) and does not depend on the type of the fluid if magnetic Prandtl number is small $\\Pr\\ll1$. Our numerical results are found to be in a good agreement with available experimental data from Grenoble High Magnetic Field Laboratory, France [P. Moresco and T. Alboussi\\`{e}re, J. Fluid Mech. \\textbf{504}, 167 (2004)]. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.3079.txt": {
        "abstract": "We compare our latest single and binary stellar model results from the Cambridge STARS code to several sets of observations. We examine four stellar population ratios, the number of blue to red supergiants, the number of  Wolf-Rayet stars to O supergiants, the number of red supergiants to Wolf-Rayet stars and the relative number of Wolf-Rayet subtypes, WC to WN stars. These four ratios provide a quantitative measure of nuclear burning lifetimes and the importance of mass loss during various stages of the stars' lifetimes. In addition we compare our models to the relative rate of type Ib/c to type II supernovae to measure the amount of mass lost over the entire lives of all stars. We find reasonable agreement between the observationally inferred values and our predicted values by mixing single and binary star populations. However there is evidence that extra mass loss is required to improve the agreement further, to reduce the number of red supergiants and increase the number of Wolf-Rayet stars. ",
        "introduction": "There are many problems outstanding in the field of stellar evolution. The most massive stars, those that end their lives as a supernova (SN), the main questions are `At what rate do stars lose mass?', `Which stars give rise to which SNe?', `What physical processes occur in binary systems?' and `How important are rotation and magnetic fields?' One way to gain insight into these questions is to observe populations of stars and SNe. Massive stars are classified in a range of possible spectroscopic stellar types. The main categories are blue supergiants (BSGs), red supergiants (RSGs) and Wolf-Rayet (WR) stars. If we count the numbers of such stars in a single system (e.g. a stellar cluster or galaxy) we can test the accuracy of our stellar models by comparing to model predictions.  When such comparisons are performed the agreement can be poor. For example the ratio of the BSG to RSG populations has been an outstanding problem for some time \\citep{red2blue,mm2001,mo2003}. The ratio of the number of RSGs to the number of WR stars is also not correctly predicted by stellar models \\citep{massey}. However the ratios of WR star subtypes WN and WC stars are predicted by modern stellar evolution \\citep{mm2005,evink}.  Massive single stars end their lives as core-collapse SNe but the relation between the star and SN type is not straight forward. A list of core-collapse SN types includes IIP, IIL, IIb, IIn, IIpec, Ib and Ic\\footnote{Type Ia SNe are thought to be thermonuclear events from explosive carbon burning in a degenerate carbon-oxygen white dwarf and are not considered in this work.} \\citep{snclassy}. The difference between type II and type I SNe is the presence or absence of hydrogen in the SN spectrum. It is easy to identify which stellar models contain hydrogen or not but the subtype classifications are fairly arbitrary \\citep{popov93,DT03,H03,ETsne,iilafterglow}. The most direct method to link supernovae to stellar end point models is to search for the progenitor stars of SNe in pre-explosion images \\citep{VD03,SJM03}.  If we only split the SN types into II and Ib/c another comparison can be made between models and observations. The relative rates of these two SN types and how they vary with metallicity have been estimated observationally \\citep{snevsZ}. \\citet{mm2005} have shown that standard non-rotating single star models do not agree with these observed rates and that mass loss must be enhanced, in their case by rotation, to find an agreement between predicted SN ratios and those inferred from observations.  The studies mentioned so far only compare single star models to the observations. This ignores the observed fact that many stars are in binaries. Binary evolution has been studied widely, (e.g. \\citealt{podsibin1,langerbin,belcz,izzysne,vanbevwr,vaz07,DT07}). The methods employed to estimate the effect of binaries is based either on rapid population synthesis or a small number of detailed models. Population synthesis employs formulae or tables based on detailed models to follow the evolution, (e.g. \\citealt{HPT00} and the evolution of a star can be calculated in a fraction of a second rather than many minutes taken by a detailed code. The speed of calculations means that the full range of parameters that govern the outcome of binary evolution can be studied and their relative effects evaluated. The disadvantage is that in the complex phases of evolution, such as when the envelope is close to being stripped, the results of population synthesis can be quite spurious.  We have constructed a set of binary stellar models covering a wide range of binary parameter space using a detailed stellar evolution code \\citep{E03}. With these models we have investigated the effect binary evolution has on the lifetimes of the various phases of stellar evolution and therefore the effect it has on the relative populations of massive stars and the relative numbers of different types of SNe. While the binary parameter-space resolution is low compared to rapid population synthesis studies it is still one of the highest resolution studies undertaken with a detailed evolution code.  We note that we do not consider here the effect of rotation or magnetic fields. The effect of rotation on the structure and evolution of stars is a complicated process and has been studied by \\citet{hegrot1} and \\citet{hegrot2} and the extensive study beginning in \\citet{mm1997} and continuing up to most recently in \\citet{mmgrb}. The effect of magnetic fields on stellar structure and evolution has also been investigated by \\citet{mm2003}, \\citet{mm2004}, \\citet{hegermag}. Our code does not include rotation therefore our study reveals the importance of duplicity alone. We discuss how rotation may affect our results in Section 6.  In this paper we start by providing details of our stellar model construction. We then specify our definitions of the different phases of stellar evolution and discuss how binary evolution affects the relative lifetimes of stars in various phases of stellar evolution. Using these lifetimes we predict how the BSG to RSG, the WR to O supergiant, the RSG to WR and WC to WN ratios vary with metallicity. We then predict how the relative rates of type Ib/c and type II SNe vary with metallicity. With each of these predictions we compare the values to the ratios inferred from observations. Finally we discuss the implications of our results. ",
        "conclusions": "We have compared our stellar models of massive stars to several observations and find that by including binary stars we achieve better agreement. This is because duplicity increases the chance of a star losing its hydrogen envelope which leads to fewer RSGs, more WR stars and more type Ib/c SNe than single star models predict as required by observations.  Problems remain and the agreement is not perfect. The BSG/RSG ratio and RSG/WR ratio indicate that we still predict too many RSGs or not enough WR stars. Also the binary models predict to many WN stars and too few WC stars. Physics we have not included in our models might resolve these problems. The main culprit is rotation. It has two effects on our predictions. First it introduces mixing that extends the time stars spend on the main sequence and therefore increases the number of BSGs. Second, if rotation is rapid it enhances mass loss by reducing the depth of the potential from which mass must escape and therefore increases the population of WR stars at the expense of the RSG population. Rotationally enhanced mass loss affects the WC/WN ratio by shortening the WN lifetimes. However the effect of rotation can be much more complex then these simple effects. For example \\citet{mm2001} find that the BSG to RSG ratio is decreased rather than increased by rotation owing to a change of the internal helium gradient at the hydrogen-burning shell. In addition tides in a binary may boost the importance of rotation during important phases of evolution and complicate the situation further \\citep{petrovic2,petrovic}.  Another limitation of our models is that we do not include mass losing eruptions that are common in the luminous-blue variable (LBV) stars. These most massive stars experience giant outbursts, losing large amounts of mass in one short event. This has dramatic implications for the evolution of the object \\citep{solbv}. There is growing evidence that SN progenitors can experience these outbursts prior to the SN explosion \\citep{kvlbv,06jc}. Such outbursts would have a similar effect on stellar populations as rotation. They reduce the time stars spend as RSGs and could reduce the time taken to remove the helium envelopes for WN stars. Therefore a better understanding of these outbursts may also improve agreement between stellar models and observations.  In summary, binary stars must be considered when comparing stellar evolution models to observations. However, binary evolution alone cannot explain all the observations. Fine tuning of stellar models is still required and the effect of enhanced mass loss owing to rotation or LBV-like eruptions must be further considered.  %Finally find that a binary companion can have an important effect on the mass lost by a star and affects the eventual SN type. This is only the case if the star is below the initial single-star mass-limit for a type Ib/c SN ($\\la$ 27M$_{\\odot}$ at solar metallicity). If the star is massive enough that stellar winds can remove the hydrogen envelope, these stars will always result in a type Ib/c SN. The stars below the initial mass-limit for a type Ib/c SN that would normally retain their hydrogen envelopes as single stars are therefore the ones that will give rise to the other SN types, with the final type of SN determined by the strength of the binary interaction during the evolution."
    },
    "0711/0711.1736_arXiv.txt": {
        "abstract": "For two bound-state equations derived as simplified forms of the Bethe--Salpeter equation with confining interaction, stability of all solutions is rigorously shown. ",
        "introduction": "The Salpeter equation is a frequently applied three-dimensional reduction of the Bethe--Salpeter formalism describing bound states within quantum field theory, derived by assuming all interactions to be instantaneous. For given interactions, encoded in its {\\em integral kernel\\/} $K(\\mbox{\\boldmath{$p$}},\\mbox{\\boldmath{$q$}})$ depending on relative three-momenta $\\mbox{\\boldmath{$p$}},\\mbox{\\boldmath{$q$}}$ of the bound-state constituents, it can be regarded as an eigenvalue equation for the Salpeter amplitude $\\Phi(\\mbox{\\boldmath{$p$}}),$ with the mass $M$ of the bound state as eigenvalue. For confining interactions, however, its solutions exhibit in {\\em numerical\\/} studies\\cite{StabSal} certain instabilities, possibly related to Klein's paradox, causing states to decay.  In view of this highly unsatisfactory state of the art, we began a systematic rigorous analysis of the conditions for stability of the energy levels predicted,~for confining interactions, within this framework. We regard a bound state as {\\em stable\\/} if its mass eigenvalue $M$ belongs to a {\\em real discrete\\/} spectrum {\\em bounded from below}.  On energetic grounds, any instabilities should manifest themselves first for pseudoscalar bound states; accordingly, we focus to fermion--antifermion bound states characterized by spin-parity-charge conjugation assignment $J^{PC}=0^{-+}.$ This allows to take advantage of experience gained in earlier studies\\cite{Lucha00:IBSEm=0,Lucha00:IBSE-C4,Lucha00:IBSEnzm,Lucha01:IBSEIAS}.  We analyze three-dimensional reductions of the Bethe--Salpeter formalism for increasing complexity of the problem: the {\\em reduced\\/} Salpeter equation\\cite{Lucha07:HORSE,Lucha07:StabOSS-QCD@Work07}, a generalization thereof, proposed in Ref.~[8] and applied in Ref.~[9], towards~exact propagators of the bound-state constituents\\cite{Lucha07:REPEHI}, and the {\\em full\\/} Salpeter equation.  With Dirac couplings $\\Gamma,$ assumed to be identical for both constituents,~and related potential functions $V_\\Gamma(\\mbox{\\boldmath{$p$}},\\mbox{\\boldmath{$q$}}),$ the action of some kernel $K(\\mbox{\\boldmath{$p$}},\\mbox{\\boldmath{$q$}})$ on $\\Phi(\\mbox{\\boldmath{$p$}})$~is $$[K(\\mbox{\\boldmath{$p$}},\\mbox{\\boldmath{$q$}})\\,\\Phi(\\mbox{\\boldmath{$q$}})]=\\sum_\\Gamma V_\\Gamma(\\mbox{\\boldmath{$p$}},\\mbox{\\boldmath{$q$}})\\,\\Gamma\\,\\Phi(\\mbox{\\boldmath{$q$}})\\,\\Gamma\\ .$$For all interactions of harmonic-oscillator type in configuration space, the above bound-state equations simplify to ordinary differential equations, which may be cast into the form of ``tractable'' eigenvalue equations for Schr\\\"odinger operators, at least in the case of the {\\em reduced\\/} bound-state equations studied in Secs.~\\ref{Sec:RSE} and~\\ref{Sec:IBSEWEP}.  Our primary tool is a (well-known) {\\em theorem\\/} which states that the {\\em spectrum\\/} of a Schr\\\"odinger Hamiltonian operator with a locally bounded positive potential $V$ rising to infinity in all directions, $V(x)\\to+\\infty$ for $|x|\\to\\infty,$ is purely discrete. ",
        "conclusions": ""
    },
    "0711/0711.4922_arXiv.txt": {
        "abstract": "{ We present a large set of theoretical isochrones, whose distinctive features mostly reside on the greatly improved treatment of the thermally pulsing asymptotic giant branch (TP-AGB) phase. Essentially, we have coupled the TP-AGB tracks described in Paper I, at their stages of pre-flash quiescent H-shell burning, with the evolutionary tracks for the previous evolutionary phases from Girardi et al. (2000). Theoretical isochrones for any intermediate value of age and metallicity are then derived by interpolation in the grids. We take care that the isochrones keep, to a good level of detail, the several peculiarities present in these TP-AGB tracks -- e.g. the cool tails of C-type stars owing to the use of proper molecular opacities as convective dredge up occurs along the TP-AGB, the bell-shaped sequences in the Hertzsprung-Russell (HR) diagram for stars with hot-bottom burning, the changes of pulsation mode between fundamental and first overtone, the sudden changes of mean mass-loss rates as the surface chemistry changes from M- to C-type, etc. Theoretical isochrones are then converted to about 20 different photometric systems -- including traditional ground-based systems, and those of recent major wide-field surveys such as SDSS, OGLE, DENIS, 2MASS, UKIDSS, etc., -- by means of synthetic photometry applied to an updated library of stellar spectra, suitably extended to include C-type stars. Finally, we correct the predicted photometry by the effect of circumstellar dust during the mass-losing stages of the AGB evolution, which allows us to improve the results for the optical-to-infrared systems, and to simulate mid- and far-IR systems such as those of Spitzer and AKARI. We illustrate the most striking properties of these isochrones by means of basic comparisons with observational data for the Milky Way disk and the Magellanic Clouds. Access to the data is provided both via a web repository of static tables ({\\tt http://stev.oapd.inaf.it/dustyAGB07} and CDS), and via an interactive web interface ({\\tt http://stev.oapd.inaf.it/cmd}) that provides tables for any intermediate value of age and metallicity, for several photometric systems, and for different choices of dust properties. }  \\authorrunning{Marigo et al.} \\titlerunning{Optical to far-IR isochrones} ",
        "introduction": "\\label{intro}  Theoretical isochrones find a wide range of applications in modern astrophysics, going from the simple age-dating of star clusters, to more complex problems like the derivation of star formation histories of resolved galaxies, and, since Charlot \\& Bruzual (1991), the modelling of their spectrophotometric evolution across the Hubble time.  In fact, many different groups are publishing large grids of stellar isochrones, covering wide enough a range of stellar parameters (age, mass, metal content, alpha-enhancement, etc.) to become useful in the stellar population synthesis of galaxies. These isochrones normally comprehend a detailed description of all evolutionary phases from the zero-age main sequence (ZAMS) up to either carbon ignition or to the onset of the thermally-pulsing asymptotic giant branch phase (TP-AGB). The coverage of the TP-AGB evolution, however, has always been far from detailed and physically accurate, as we will illustrate below.  The poor modelling of this phase greatly contrasts with its importance to the integrated light of stellar populations with ages from a few $10^8$~yr up to the Hubble time. According to Frogel et al. (1990), the TP-AGB contribution to the total luminosity of single-burst stellar populations reaches a maximum of about 40 percent at ages from 1 to 3 Gyr, falling to less than 10 percent at 10 Gyr. Obviously, TP-AGB stars have a great impact on the rest-frame near-infrared spectra of gala\\-xies. Moreover, TP-AGB stars account for most of the bright-infrared (IR) objects in resolved galaxies, as clearly demonstrated by DENIS, 2MASS, SAGE and S$^3$MC data for the Magellanic Clouds (Cioni et al. 1999; Nikolaev \\& Weinberg 2000; Blum et al. 2006; Bolatto et al. 2007).  \\begin{table*} \\small \\caption{Available stellar isochrones including the TP-AGB phase.} \\label{tab_isoc} \\begin{tabular}{lcccl} \\hline Reference & $3^{\\rm rd}$DU and HBB$^{(\\rm a)}$ & $T_{\\rm eff}$$^{(\\rm b)}$ & Dust$^{(\\rm c)}$ &  public www address for retrieval \\\\ \\hline Bertelli et al. (1994)\\comment{, A\\&AS, 106, 275}    &  NO  &  solar-scaled comp.  &  NO  &   http://pleiadi.pd.astro.it \\\\ Bressan et al. (1998)\\comment{, A\\&A, 332, 135}  & NO  &  solar scaled comp. &   YES &               upon request \\\\ Girardi et al. (2000)\\comment{, A\\&AS, 141, 371}  &     NO  &  solar-scaled  comp. &  NO    &        http://pleiadi.pd.astro.it\\\\ Marigo \\& Girardi (2001)\\comment{, A\\&A, 377, 132}   &   YES &  solar-scaled  comp. &  NO   &         http://pleiadi.pd.astro.it \\\\ Mouhcine (2002)\\comment{, A\\&A, 394, 125}  & YES    &  solar scaled comp. &    YES  &               upon request \\\\ Piovan et al. (2003)\\comment{, A\\&A, 408, 559} &  NO  &  solar scaled comp. &   YES  &                upon request \\\\ Marigo et al. (2003a)\\comment{, A\\&A, 403, 225}  & YES  & O-/C-rich comp. &    NO  &          upon request  \\\\ Cioni et al. (2006)\\comment{, A\\&A, 448, 77}  & YES  & O-/C-rich comp. &    NO  &          http://pleiadi.pd.astro.it  \\\\ Cordier et al. (2007)\\comment{, AJ, 133,468}    &  NO & solar scaled comp.  &   NO    & http://www.oa-teramo.inaf.it/BASTI \\\\ This work    &  YES & O-/C-rich comp.  &   YES    & http://stev.oapd.inaf.it/cmd \\\\ \\hline \\end{tabular} \\\\ $^{(\\rm a)}$ Explicit inclusion or not of AGB nucleosynthesis, i.e. the third dredge-up and hot-bottom burning, in the underlying AGB models \\\\ $^{(\\rm b)}$ Choice of low-temperature opacities affecting the predicted effective temperature, i.e. for solar-scaled or variable (O- or C-rich) chemical mixtures \\\\ $^{(\\rm c)}$ Inclusion or not of circumstellar dust in the computation of the emitted spectrum \\\\ \\end{table*}  The situation regarding the TP-AGB phase in published sets of isochrones is the following: Whereas early isochrones sets do not even include the TP-AGB phase, more recent ones do it, but still in a very approximate way, ignoring crucial aspects of the TP-AGB evolution.  A summary table of available isochrones is provided in Marigo (2007) and updated in Table~\\ref{tab_isoc} herewith. In short, Bertelli et al. (1994) and Girardi et al. (2000; hereafter Gi00) adopt a very simple TP-AGB synthetic model that does not include the third dredge-up and the over-luminosity above the core mass--luminosity relation caused by hot bottom burning (HBB). A luminosity correction for stars undergoing HBB is instead included by Cordier et al. (2007), but in this case the third dredge up -- together with its most remarkable consequence, i.e. the formation of C stars -- and the HBB nucleosynthesis are completely missing. Marigo \\& Girardi (2001) present the first isochrones based on tracks (Marigo et al. 1999) with the third dredge up included, and moreover its efficiency is calibrated to reproduce the C-star luminosity functions in both Magellanic Clouds (MC)\\footnote{Many of the improvements in Marigo et al. (1999) TP-AGB tracks were already introduced by Groenewegen \\& de Jong (1993, 1994abc), who computed extended sets of tracks and performed population synthesis of resolved AGB stellar populations. Groenewegen \\& de Jong, however, do not provide theoretical isochrones from their tracks.}. Another distinctive aspect of Marigo et al. (1999) tracks, is that the HBB efficiency, together with the model effective temperatures, are derived from numerical integrations of the envelope structure.  Besides the above mentioned works, the literature abounds in descriptions of evolutionary population synthesis models which have included the complete TP-AGB evolution (e.g. Charlot \\& Bruzual 2003; Raimondo et al. 2005; V\\'azquez \\& Leitherer 2005; only to mention a few). However, in all these cases, critical processes of TP-AGB evolution -- again the third dredge-up, HBB, or superwind-like mass loss -- are missing in the models. In some cases (e.g. V\\'azquez \\& Leitherer 2005; Charlot \\& Bruzual 2003), the tracks for the TP-AGB and previous evolutionary phases come from different sources; this implies that basic stellar properties like the core mass and envelope chemical composition are discontinuous along the composite tracks. As demonstrated by Marigo \\& Girardi (2001), discontinuities in these quantities also produce an energetic inconsistency, in the sense that they violate the  strict coupling between stellar integrated light and chemical yields implied by the condition of energy conservation.   A few authors have computed isochrones aimed at the modelling of spectral evolution of galaxies which consider the effect of circumstellar dust around AGB stars (Bressan et al. 1998; Mouhcine 2002; Piovan et al. 2003), but without improving the treatment of the underlying TP-AGB evolution, with respect to those already mentioned.  In all the previously-mentioned cases, moreover, the effect of variable opacities along the TP-AGB evolution is completely ignored: The basic stellar properties, and in particular the effective temperature \\Teff, are derived using opacities valid for scaled-solar compositions, which are highly inadequate to describe the chemically evolved envelopes of TP-AGB stars. The proper consideration of variable opacities is now known to be crucial (see Marigo 2002; Marigo et al. 2003a), especially for stars undergoing the transition to the C-rich phase -- i.e. comprising, for MC metallicities, all stars with initial masses between about 1.5 and 3.5~\\Msun, and all isochrones with ages between about 0.5 and 5 Gyr. This situation is partially remedied by Cioni et al. (2006), who has made publicly available the first isochrone sets in which variable molecular opacities were considered along the TP-AGB evolution\\footnote{Marigo et al. (2003a) also computed isochrones containing essentially the same kind of TP-AGB tracks as in Cioni et al. (2006), but for a more limited range of metallicities, and for exploratory choices of model parameters. These isochrones were provided a few times upon request, and were superseded first by Cioni et al.'s and now by our new isochrones.}. In fact, the C-star phase in these isochrones is characterized by much lower temperatures (reaching \\Teff\\ as low as 2300~K) than former ones. Those specific TP-AGB models, however, are previous to the extensive calibration of dredge up efficiency  performed in Marigo \\& Girardi (2007, hereafter Paper~I). In this latter work, the dredge up parameters are tuned so that the models reproduce the observed the C-star luminosity functions in the Magellanic Clouds, and the TP-AGB lifetimes of M and C stars in Magellanic Cloud star clusters (Girardi \\& Marigo 2007a).  {\\em The goal of the present paper is to remedy this highly unsatisfactory situation, providing the first set of isochrones in which the TP-AGB evolution is treated with a satisfactory level of completeness and detail -- i.e. considering the crucial effects of third dredge up, hot bottom burning, and variable molecular opacities -- and calibrated at the same time -- i.e. able to reproduce the basic observables of AGB stars in the MCs. Moreover, we wish to provide these results in many photometric systems and including the reprocessing of radiation by circumstellar dust in mass-losing stars, which is crucial for the modelling of the infrared photometry.}  The properties of our stellar evolutionary tracks are briefly recalled in Sect.~\\ref{sec_tracks}, which discusses the basic aspects of constructing isochrones using the new TP-AGB tracks of Paper~I -- including new variables like the C/O ratio, pulsation periods, pulsation modes, etc. Section~\\ref{sec_bc} deals with the transformation of the isochrones to many photometric systems, which require updating and extending the libraries of stellar spectra previously in use. The inclusion of dust reprocessing in the predicted photometry is illustrated in Sect.~\\ref{sec_dust}. A preliminary comparison with observational data is done in Sect.~\\ref{sec_obs}. Section~\\ref{sec_conclu} describes how to access the data tables -- containing isochrones, bolometric corrections, dust attenuation curves, integrated magnitudes of single-burst populations, etc. -- and summarizes both the main novelties and caveats of our new isochrones. ",
        "conclusions": ""
    },
    "0711/0711.4113_arXiv.txt": {
        "abstract": "Measuring the response of the intergalactic medium to a blast of ionizing radiation allows one to infer the physical properties of the medium and, in principle, the lifetime and isotropy of the radiating source. The most sensitive such measurements can be made if the source of radiation is near the line of sight to a bright background QSO. We present results based on deep Keck/HIRES observations of the QSO triplet KP76, KP77 and KP78 at $z \\simeq 2.5$, with separations of 2-3\\arcm\\ on the plane of the sky. Using accurate systemic redshifts of the QSOs from near-IR spectroscopy, we quantify the state of the IGM gas in the proximity regions where the expected ionizing flux from the foreground QSOs exceeds that of the metagalactic background by factors $10-200$, assuming constant and isotropic emission. Based on the unusual ionization properties of the absorption systems with detected \\hi, \\civ\\ and \\ovi, we conclude that the gas has been significantly affected by the UV radiation from the nearby QSOs. Aided by observations of the galaxy density near the foreground QSOs, we discuss several effects that may explain why the transverse proximity effect has eluded most previous attempts to detect it. Our observations suggest that the luminosities of KP76 and KP77 have remained comparable to current values over timescales of, respectively, $\\Delta t > 25$ Myr and 16 Myr $< \\Delta t < 33$ Myr - consistent with typical QSO lifetimes estimated from independent, less-direct methods. There is no evidence that the UV radiation from either QSO was significantly anisotropic during these intervals. ",
        "introduction": "The QSO ``proximity effect'' refers to the observation that the incidence of Lyman $\\alpha$ (Ly$\\alpha$) absorption from the IGM decreases at redshifts close to that of the QSO. If this deficit is interpreted as being due to the influence of the radiation field of the QSO on the nearby IGM, then measurements of the \\hi\\ density near QSOs of known luminosity may be used to determine the intensity of the metagalactic ionizing background \\citep{murdoch86,bajtlik88,scott00}. The statistical significance of the QSO proximity effect is high enough that there is little doubt of its reality. However, the measurement of the intensity of the metagalactic radiation field from observations of the Ly$\\alpha$ forest is subject to a number of systematic uncertainties, in particular the uncertainty in the QSO systemic redshift (e.g., \\citealt{scott00,espey89}), the possibly over-dense large-scale environment in which QSOs may be located (e.g., \\citealt{faucher07}), and the non-linear response of the observed H\\,{\\sc i} optical depth to changes in the radiation field intensity, such that only \\hi\\ systems with relatively low \\hi\\ optical depths ($\\tau_{\\rm H\\,I} \\simlt 10$) will be significantly altered. All of these effects would lead to an overestimate of the background intensity through an underestimate of the sphere of influence of radiation from the QSO: QSO redshifts determined from the broad UV emission lines tend to be lower than the true systemic redshifts, so that the true ``proximity region'' where the QSO radiation field dominates over the background is actually larger than inferred. Similarly, if QSOs are preferentially found in overdense regions, then the incidence of \\hi\\ with sufficient optical depth to remain relatively unscathed by an enhanced radiation field will be higher than at an average location in the universe, suggesting that the QSO's effect on the local radiation field is smaller than it really is. Clearly, there are also statistical (as opposed to systematic) issues associated with the proximity effect, since one expects relatively large sample variance in the \\hi\\ content between the relatively small volumes probed by individual QSOs.  The so-called ``transverse proximity effect'' (hereinafter, TPE), on the other hand, involves searching for the influence of foreground sources of UV radiation on the IGM absorption observed in the spectra of background objects. Unlike the line-of-sight proximity effect, measurements of the transverse proximity effect can in principle constrain the radiative lifetime {\\it and} the isotropy of the sources (e.g., \\citealt{schirber04}). Line of sight measurements of the proximity effect are sensitive only to fluctuations in the source intensity on timescales of $\\simlt 10^4$ yr (e.g., \\citealt{martini04}), whereas transverse measurements can sample longer timescales---determined by the light travel time between the foreground source and the observed line of sight---that are more interesting in the context of QSO lifetimes. The sphere of influence of the radiation field from a source of UV photons, as for the line of sight proximity effect, depends on luminosity of the source. For a typical bright QSO ($V \\sim 18$) at $z \\sim 2.5$, the sphere within which an isotropically radiating QSO significantly dominates over the metagalactic background has a physical radius of $r \\sim 5-10$\\,Mpc, corresponding to a line of sight velocity range of $\\Delta v_{\\rm r} \\simeq 1250-2500$\\,\\kms\\ for material moving with the Hubble expansion. The corresponding light travel time is $15-30$\\,Myr. A number of quasi-independent arguments have suggested that bright QSO typically have lifetimes of $10^6-10^8$\\,yr (see, e.g., \\citealt{haehnelt98,richstone98,martini01,hosokawa02,yu02,steidel02}) so that the potential sphere of influence of the UV radiation from bright QSOs and the timescale of QSO `events' may fortuitously  be of the same order. The relevant angular scales are $\\theta \\simlt 10$\\arcm\\ on the sky at $z \\sim 2-3$, with the amplitude of the expected effects varying as $1/\\theta^2$ and proportional to the (far-UV) luminosity of the sources.  A number of authors have used QSOs with small angular separations on the sky, but at different redshifts, to search for the TPE via observations of the Ly$\\alpha$ forest (e.g., \\citealt{fern-soto95,crotts98,croft04,schirber04}). To date, as far as we are aware, no evidence for a reduction in the number of Ly$\\alpha$ forest lines at redshifts near foreground QSOs has been found\\footnote{\\citet{gallerani07} observe a ``transmission gap'' in the Lyman $\\alpha$ forest of a background $z=6.42$ QSO near the redshift of a foreground QSO at $z=5.65$, although the statistical significance of the region of reduced Ly $\\alpha$ opacity is difficult to evaluate.}. In fact, all of the above-referenced works except \\citet{fern-soto95} reported instead an {\\it excess} of Ly$\\alpha$ absorption systems at the redshifts where TPE deficits would have been expected. The common explanations for the lack of the expected signal include some combination of the QSO large-scale environment, short or intermittent radiative lifetimes of the QSOs, and anisotropy of the QSO radiation. More recently, \\citet{hennawi06} found that high column density H\\,{\\sc i} systems at the redshifts of foreground QSOs are much more common than associated systems of similarly high column density seen in the spectra of QSOs at their own redshifts. These authors also estimated that the clustering of the high column density \\hi\\ systems with the QSOs may be even stronger than that of star-forming galaxies at the same redshifts. Such observations have been interpreted as evidence that H\\,{\\sc i} gas is distributed anisotropically around QSOs, possibly because of QSO variability and/or anisotropic UV radiation. As a counter-argument, however, observations of the He\\,{\\sc ii}, rather than H\\,{\\sc i}, Ly$\\alpha$ forest seem to show that foreground QSOs do in fact have a measurable effect on the ionization state of gas within volumes corresponding to light travel times of $\\sim 10^6-10^7$ Myr \\citep{jakobsen03,worseck06,worseck07}. This may suggest that H\\,{\\sc i} is not the most sensitive barometer of changes in the radiation field near QSOs. The key advantage of the He\\,{\\sc ii} transition is that it is sensitive to the intensity of high-energy photons that can plausibly only be produced by QSOs or AGN, so that one can sense changes in both the intensity and the shape of the radiation field for local sources.  A phenomenon closely related to the TPE is the fluorescence of high column density \\hi\\ gas in the vicinity of bright sources of ionizing photons (see, e.g., \\citealt{gould96,cantalupo05}). Searches for fluorescence at the redshifts of QSOs have so far yielded mixed results. \\citet{adelberger06} identified what appears to be a damped Ly$\\alpha$ system (DLA) at $z=2.84$ that is producing Ly$\\alpha$ emission in response to the ionizing radiation from a very bright QSO located $\\sim 380$\\,kpc away, implying that the QSO luminosity has remained approximately constant over the last $\\sim 10^6$\\,yrs (and that the QSO is radiating isotropically). However, in another case, the lack of fluorescence around a QSO at a similar redshift was interpreted by \\citet{francis04} as evidence for anisotropic emission from the QSO.  Thus, the observational situation concerning the TPE is currently ambiguous. QSO lifetimes of $>10^6$\\,yrs and largely isotropic QSO radiation are supported by some observations, but challenged by others. For this reason, we have elected to investigate the TPE using a different approach from that of most previous attempts. Rather than counting Ly$\\alpha$ lines or evaluating the average Ly$\\alpha$ opacity near foreground QSOs, we focus on the metal line systems. Our strategy is to closely examine the ionization state of metal line systems within the expected spheres of influence of foreground QSOs, using the combination of very high quality optical spectra of the background QSOs and precise determinations of the systemic redshifts of the foreground QSOs from near-infrared (near-IR) spectroscopy. Our aim is to determine the strength of the ionizing radiation field seen by gas at various distances from the foreground QSOs, independently of the statistical limitations and systematic uncertainties inherent in line counting or ``mean flux'' techniques. The hope is that the new data will help to explain why previous results have been so mixed, as well as constrain better the radiative lifetime and solid angle of QSO emission.  In the first stage of this program reported here, we focus on the QSO triplet KP76, KP77, KP78 (also known as Q1623+268). This grouping is among the best studied in the literature, having figured prominently in earlier work aimed at measuring the coherence scale and clustering of the Ly$\\alpha$ forest (e.g., \\citealt{sargent82,crotts89,crotts98,crotts97}). The relative configuration of the three QSOs is illustrated in Figures~\\ref{triplet} and \\ref{scheme}. \\citet{crotts98} searched for the TPE in the Ly$\\alpha$ forests of this triplet and concluded---based on counts of Ly$\\alpha$ absorption lines---that there is no evidence that the foreground QSOs (KP76 and KP77) reduce the density of Ly$\\alpha$ lines in the expected proximity regions. Rather,  the line density appears to be higher than at an average location in the forest at comparable redshifts.   \\begin{figure}[htb] \\plotone{f1.eps} \\caption{The Q1623-KP76, KP77, KP78 triplet on the plane of the sky. The transverse physical distances between the QSO sightlines (evaluated at the redshift of the lower redshift QSO in each case) are indicated. The QSO systemic redshifts were measured from accurate near-IR spectroscopy (\\S 2.2). } \\label{triplet} \\end{figure}   \\begin{figure*}[htb] \\plotone{f2.eps} \\caption{Schematic view of the QSO triplet. The relative positions of the metal line absorption systems studied in this paper, labeled A--J, are indicated within dotted circles representing the proximity regions of the foreground QSOs KP76 and KP77. Within these volumes---spheres of physical radius $\\sim 5$\\,Mpc--- the ionizing flux from the nearby QSO is expected to exceed that due to the metagalactic background by more than one order of magnitude, if the QSOs radiate isotropically and their luminosities have remained approximately constant over the last 20--30\\,Myr. } \\label{scheme} \\end{figure*}  In this paper, we present the highest quality high resolution spectra of these QSOs to date, obtained with the recently upgraded High Resolution Echelle Spectrograph (HIRES) on the Keck~I telescope. We have also obtained, as part of an ongoing survey of galaxies and the IGM at redshifts $1.8 \\simlt z \\simlt 3$ (see \\citealt{steidel04, adelberger05}), spectra of approximately 300 other galaxies, QSOs, and AGN in this redshift range within a $\\sim 15$\\arcm\\ field centered on the QSO triplet, providing unprecedented information on the surrounding large scale structure. We examine the detailed properties of absorption line systems within the expected proximity regions of the two foreground QSOs, and we interpret our findings in the context of our knowledge of the distribution of galaxies in the same regions. ",
        "conclusions": "\\label{sec:discussion}  Using high resolution spectra of a triplet of QSOs with transverse separations of $\\sim 1$\\,Mpc at $z \\simeq 2.5$, together with accurate determinations of the QSO systemic redshifts from rest-frame optical emission lines and extensive spectroscopic observations of galaxies and AGN in the same field, we have conducted the most detailed investigation to date of the transverse proximity effect. By focusing on the regions of the IGM where the QSOs should overwhelm the metagalactic radiation field by factors between  $\\sim10$ and $\\sim200$---if they have radiated isotropically and with similar luminosities over the past 0.2--30\\,Myr---we have examined the details of the ionization state of individual metal absorption systems, rather than counting \\lya\\ lines above a given equivalent width threshold, as has generally been done in previous attempts to measure the TPE. We have shown that the 10 metal line systems within the proximity regions of the two foreground QSOs have properties that are more easily explained if they are being illuminated by a UV radiation field significantly more intense than the metagalactic background. Using photoionization models, we have shown that most of the observed absorption systems have properties consistent with normal (i.e. non-proximate) intermediate column density ($\\log N$(H\\,{\\sc i})\\,$\\sim 14.5$--16.5) metal line systems that have been ionized by QSO continuum radiation with intensity consistent with that inferred from the observed fluxes of the two foreground QSOs.  We have placed the observed QSOs in the context of the large scale distribution of galaxies in the same field, showing that all three lie in regions of moderate galaxy overdensity---conditions typical of high redshift QSOs and AGN. Even with accurate QSO redshifts (which correct the published values for the same QSOs by 1000--2000\\,km~s$^{-1}$), the naive expectation that there should be a dearth of \\lya\\ absorption systems in the proximity regions of the foreground QSOs is not supported by the data. We argue that the environments of the QSOs and an enhancement of moderate column density \\hi\\ absorption compared to average locations in the IGM can easily mask the effects of the TPE if one relies on counting statistics rather than examining the details of the gas-phase physical conditions.  \\subsection{QSO Lifetimes and Isotropy} \\label{sec:lifetimes}  In principle, it should be possible to use information on the distribution of the gas that is clearly affected by QSO radiation relative to the positions and redshifts of the foreground QSOs to measure, or set limits on, both the lifetime and the isotropy of the QSOs radiation field. Each absorption system listed in Table~\\ref{gfactortab} and indicated in Figure~\\ref{deltaz} samples a different time interval $\\Delta t$ and boost factor $g$ under the simple hypothesis that the QSOs shine isotropically and at constant luminosity over their radiative lifetimes. By comparing our sample to ``proximate'' and ``non-proximate'' absorption systems from the literature, as in Figures~\\ref{figsimcoe_civ} and  \\ref{figsimcoe_ovi}, and from arguments based on cloud density and size, we have shown that most of the systems in our sample---systems C, D, E, F, G, H and J---are probably being over-ionized by the nearby QSO, while systems A and B are merely consistent with that hypothesis. Taken at face value, the spread of values of $\\Delta t$  in Table~\\ref{gfactortab} then implies that the minimum QSO lifetime is $\\sim 25$\\,Myr for KP76 (from system E) and $\\sim 16$\\,Myr for KP77 (from system H). As discussed above, system I, in the proximity region for KP77, is the only system whose properties may favor ionization by the metagalactic field alone, without local enhancement; again, taken at face value, this would indicate a radiative lifetime for KP77  in the range 16\\,Myr\\,$<\\Delta t < 33$\\,Myr. Radiative lifetimes of 20--30\\,Myr are entirely consistent with He\\,{\\sc ii} TPE measurements by \\cite{jakobsen03}, and with numerous estimates based on QSO duty cycle arguments (e.g., \\citealt{steidel02}) or the local $M_{\\rm bh} - \\sigma$ relation between supermassive black holes and their host galaxies (e.g., \\citealt{martini01}).  In reality, there are many other variables which potentially could modify such conclusions, even accepting the evidence that at least some of the absorption systems are over-ionized due to their proximity to a bright QSO. For example, even if the lifetime of a QSO event is a well-defined quantity,  it is unlikely that an accreting supermassive black hole maintains a constant luminosity throughout its active phase; rather, its accretion rate and UV output could vary intermittently,  or grow (or decay) exponentially. Because each absorption system samples a different time in the QSO's history, it is entirely possible that systems whose distances from the QSO are larger (so that the $g$ factor is smaller) could have experienced a more intense radiation field in the past, or that a system at some intermediate value of $\\Delta t$ happened to coincide with a dormant period in the QSOs radiative history. We have also seen that the effects of an enhanced radiation field can be quite subtle. It is not always possible to say with confidence, even with high quality data, whether or not a given absorber is experiencing an enhanced ionizing radiation field, because we have no knowledge of its environment or physical conditions prior to the time when the QSO began radiating at its present luminosity.  Given all these {\\it caveats}, the strongest statement we can make about the isotropy of the QSO radiation is that there is no evidence for anisotropy in the present data. If the QSOs' radiation were significantly beamed, one might reasonably expect to find absorption systems well within the proximity zones with properties that are inconsistent with the assumption of a radiation field intensity significantly boosted over the metagalactic  background. Within the two proximity regions considered, the only absorption system that seems marginally inconsistent with the assumption of an enhanced radiation field (system I) also has the largest value of $\\Delta t$ and a relatively small boost factor $g$.  As part of our survey of the Q1623+268 field, we obtained deep \\textit{Spitzer} IRAC and MIPS images which include the QSO triplet. We find that KP76, KP77 and KP78 have nearly identical spectral energy distributions between 0.35 and $24\\,\\mu$m, with flat spectra ($f_\\nu \\simeq$\\,constant) between 0.35\\,$\\mu$m and 4.5\\,$\\mu$m, and with $\\nu f_{\\nu}(0.36\\mu{\\rm m})/\\nu f_{\\nu}(24\\mu{\\rm m}) \\simeq 4$. QSOs that are heavily obscured over a large solid angle would be expected to be very bright in the thermal IR due to emission from heated dust. While modeling of the QSO spectral energy distributions is beyond the scope of this paper, the data suggest that the two foreground QSOs (as well as KP78, though it does not matter for the present purposes) are not heavily obscured over a large fraction of 4$\\pi$ steradians, given the relatively weak 24$\\mu$m luminosity. This provides independent support for the hypothesis that these three QSOs would be seen as UV-bright over a large fraction of a 4$\\pi$ steradian solid angle.  The results we have reported provide a counter-example to recent claims of an excess of high \\nhi\\ absorption systems near the redshifts of foreground QSOs in the spectra of background QSOs, compared to line-of-sight proximate absorbers with the same characteristics (e.g., \\citealt{bowen06,hennawi06}). These authors interpret such an excess as evidence for anisotropy of the QSO radiation which presumably (over)ionizes the clouds in line-of-sight to Earth but is not seen (with the same intensity) by gas at transverse distances. One difference that must be borne in mind is that the analysis by Hennawi et al. (2006) refers to absorption systems with column densities $\\log N$(H\\,{\\sc i})\\,$> 19$,  more than 3.5 orders of magnitude higher than any in the sample considered here. Similarly, the work by Bowen et al. (2006) targeted strong Mg\\,{\\sc ii} absorbers which are likely to be Lyman limit systems with $\\log N$(H\\,{\\sc i})\\,$> 17.5$. While our data appear to be inconsistent with significant beaming of the ionizing radiation from KP76 and KP77, it is certainly possible that QSOs may differ in the solid angle over which they radiate, and that the apparent differences between our study and those referenced above is due to such variations from QSO to QSO. Finally, while we cannot rule out the possibility that there may be an anisotropic distribution of gas (as opposed to anisotropic ionization) surrounding many QSOs, it would be hard to understand if such anisotropy extended over physical scales of $\\sim 1$\\,Mpc as would be required to explain some of the observations.  A lingering concern is the uncertainty in the relevant distances introduced by the systematic errors in the QSO systemic redshifts. For example, it is debatable whether it is surprising or not to find high $N$(H\\,{\\sc i}) absorbers near the redshift of a foreground QSO. Even in cases where the projected sightlines pass within tens of kpc of one another, if the uncertainty in the redshift of the foreground QSO is $\\sim 1500$\\,km~s$^{-1}$  (equivalent to a distance of $\\pm 6$\\,Mpc at $z = 2.5$) the relevant $g$ factor could be uncertain by up to three orders of magnitude. Many such ambiguities could be addressed with more accurate redshifts for the QSOs, as well as more accurate values of $N$(H\\,{\\sc i}) and of other indicators of the physical conditions in the gas.   \\subsection{The Elusive Transverse Proximity Effect}  In this paper, we have used a particularly well-observed triplet of QSOs in a concerted effort to test for the presence of the transverse proximity effect. We have shown that in spite of the absence of an obvious ``clearing'' in the Ly$\\alpha$ forest near the redshifts of bright foreground QSOs, significant evidence for local enhancements in the ionizing radiation field is present when one examines the detailed physical conditions of metal line systems within 5\\,Mpc of the foreground QSOs. Crucial to this analysis are:  (a) the ability to measure accurate H\\,{\\sc i} column densities (for which echelle spectra extending to at least Ly$\\beta$ are required); (b) knowing as precisely as possible where to expect the influence of the foreground QSOs (for which the forbidden line spectroscopy of the QSOs was essential);   (c) the ability to detect weak lines of highly ionized metallic species, in this case \\civ\\ and \\ovi; and (d) knowledge of the large-scale environment inhabited by the foreground QSOs, since galaxy overdensities are likely to be accompanied by the presence of appreciable \\hi\\ absorption.\\footnote{Alternatively, as shown by \\citet{jakobsen03} and \\citet{worseck06}, using the He\\,{\\sc ii} transitions in concert with H\\,{\\sc i} is a very powerful technique because the effects of an enhancement in the radiation field by a nearby QSO are then much more evident than if one has access to H\\,{\\sc i} lines only. However, the downside of this approach is that, given the short rest-frame wavelengths of the He\\,{\\sc ii} Lyman series, there is only a handful of sightlines known at present where this technique can be applied in practice---in most cases intervening absorbers optically thick in the H\\,{\\sc i} Lyman continuum prevent measurement of the He\\,{\\sc ii} lines. }  Essentially all previous searches for the TPE have been missing most or all of these ingredients. As discussed above, reliance on relatively crude line counting or mean flux measurements in the forest, the likelihood that the redshifts of the foreground QSOs, even when taken from carefully compiled catalogs such as SDSS, are incorrect by as much as 2000\\,km~s$^{-1}$ (an error $\\delta z > 1000$\\,km~s$^{-1}$ seems quite typical), and density enhancements of galaxies and intergalactic gas local to the foreground QSOs, could together conspire to mask the TPE. In view of our results, these effects seem at least as plausible as anisotropic emission, short-timescale variability, or very short QSO lifetimes in explaining the difficulties experienced so far in detecting the TPE. Based on the case investigated in this paper, the null hypothesis that bright QSOs radiate isotropically over characteristic timescales of a few $10^7$ yr (timescales suggested by many, less direct, arguments) is consistent with the observations. Whether this is the rule rather than the exception could be established using similar observations of other multiple-QSO-sightline fields. Somewhat farther in the future, the technique can be improved and extended using background \\textit{galaxies} which would provide much finer spatial and temporal sampling of the response of the IGM to radiation from QSOs \\citep{adelberger04}."
    },
    "0711/0711.4263_arXiv.txt": {
        "abstract": "\\noindent Inflationary observables, like the power spectrum, computed at one- and higher-order loop level seem to be plagued by large infra-red corrections. In this short note, we point out  that these large infra-red corrections appear only in quantities which are not directly observable. This is in agreement with general expectations concerning infra-red effects. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.5029_arXiv.txt": {
        "abstract": "Nonlinear evolution is sometimes modeled by assuming there is a deterministic mapping from initial to final values of the locally smoothed overdensity. However, if an underdense region is embedded in a denser one, then it is possible that its evolution is determined by its surroundings, so the mapping between initial and final overdensities is not as `local' as one might have assumed.  If this source of nonlocality is not accounted for, then it appears as stochasticity in the mapping between initial and final densities. Perturbation theory methods ignore this `cloud-in-cloud' effect, whereas methods based on the excursion set approach do account for it; as a result, one may expect the two approaches to provide different estimates of the shape of the nonlinear counts in cells distribution.  We show that, on scales where the rms fluctuation is small, this source of nonlocality has only a small effect, so the predictions of the two approaches differ only on the small scales on which perturbation theory is no longer expected to be valid anyway. We illustrate our results by comparing the predictions of these approaches when the initial-final mapping is given by the spherical collapse model.  Both are in reasonably good agreement with measurements in numerical simulations on scales where the rms fluctuation is of order unity or smaller.  If the deterministic mapping from initial conditions to final density depends on quantities other than the initial density, then this will also manifest as stochasticity in the mapping from initial density to final.  For example, the Zeldovich approximation and the Ellipsoidal Collapse model both assume that the initial shear field plays an important role in determining the evolution. We compare the predictions of these approximations with simulations, both before and after accounting for the `cloud-in-cloud' effect. Our analysis accounts approximately for the fact that the shape of a cell at the present time is different from its initial shape; ignoring this makes a noticable difference on scales where the rms fluctuation in a cell is of order unity or larger.  On scales where the rms fluctuation is 2 or less, methods based on the spherical model are sufficiently accurate to permit a rather accurate reconstruction of the shape of the initial distribution from the nonlinear one.  This can be used as the basis for a method for constraining the statistical properties of the initial fluctuation field from the present day field, under the hypothesis that the evolution was purely gravitational.  We illustrate by showing how the highly non-Gaussian nonlinear density field in a numerical simulation can be transformed to provide an accurate estimate of the initial Gaussian distribution from which it evolved. ",
        "introduction": "The present work is primarily concerned with the probability distribution function (hereafter pdf; some authors prefer to call this the probability density function), which describes the probability that a randomly placed cell of specified shape and volume contains a certain amount of mass.  In the best studied cosmological models, the pdf has a Gaussian form initially, but becomes increasingly positively skewed at later times. There are two methods for estimating the evolution of the dark matter pdf.  One is based on perturbation theory \\cite{b94,ps97,gc99,sg01,ptreview} and the other is based on excursion set methods \\cite{rks98}. The perturbation theory based calculation is not expected to be accurate on scales where the variance is of order unity or larger; in hierarchical models, this means that perturbation theory is not expected to be accurate on small scales. Hyperextended Perturbation Theory (HEPT) \\cite{hept} is expected to be valid in the larger variance (small cell) regime where standard perturbation theory breaks down; it provides explicit expressions for the moments of the pdf rather than a closed form expression. For the special case of clustering from white-noise Gaussian initial conditions, the excursion set method predicts exactly the same pdf as does HEPT, despite the fact that the methods used by the two approaches are very different. Motivated by this coincidence, the purpose of the present note is twofold:  first, to show that, for more general initial conditions, the excursion set approach actually makes rather similar predictions to those of perturbation theory on scales where the variance is small; the second is to understand why.  In the analysis which follows, we distinguish between two steps in the calculation of the pdf.  The first is the approximation for nonlinear evolution which we will call the dynamics.  The second is how this approximation is used to translate from the initial pdf to an evolved one, which we call the statistics.  In the first half of this paper, we study approximations for the dynamics which are based on the assumption of spherical symmetry. In this case, the perturbation theory method provides what is, in effect, a monotonic, deterministic mapping between the initial and final overdensities.  Because the final overdensity at a specified position in space is determined solely by the initial value at that position, this is sometimes also called a `local' mapping, since values of the initial fluctuation field at other positions are assumed to not affect the mapping.  The excursion set approach accounts for the fact that the evolution of a given region may actually be determined by less local surroundings. For example, consider the evolution of an underdense region which is surrounded by a dense shell.  If the shell is sufficiently dense, then it will eventually collapse, crushing the smaller region within it.  The local approximation would have predicted expansion rather than collapse for the smaller underdense region. Sheth \\& van de Weygaert (2004) call this the void-in-cloud problem, although it is clear that this is an extreme example of a more general `cloud-in-cloud' problem.  Clearly, in such cases, the mapping between initial and final overdensities is not as `local' as perturbation theory assumes, and accounting for this `cloud-in-cloud' problem is likely to be more important for small `clouds'. If not accounted for, this effect will manifest both as stochasticity (since the same initial overdensity may map to many different final densities depending on the surroundings) and, perhaps, as a bias. The excursion set approach provides an algorithm which accounts for this source of non-locality; it assumes that, once the correct large scale has been chosen, the mapping is deterministic.   However, there is another source of non-locality which spherical evolution models ignore.  This source plays a crucial role in models which account for the influence of the external shear field on the evolution.  The simplest of these more sophisticated models for the dynamics is the Zeldovich approximation \\cite{zel70}.  Here, the nonlinear density is determined not just by the initial overdensity, but by two other quantities which are related to the surrounding shear field.  These quantities also play an important role in extensions of the Zeldovich approximation \\cite{mkc01} as well as in ellipsoidal collapse models \\cite{ws79,bm96}. In all these models, the nonlinear density is a deterministic function of three quantities associated with the initial fluctuation field. In the context of perturbation theory models for the pdf, the mapping from initial density to final density will appear to be stochastic if the influence of the two other variables is not accounted for. In the excursion set approach, this stochasticity is in addition to that which derives from the cloud-in-cloud problem which is now associated with all three variables.  To explore the consequences of this additional source of stochasticity we show the result of inserting the Ellipsoidal collapse model into the perturbation theory and excursion set calculations.  We do this in two steps:  by showing the predictions when expanding the dynamics to first, and then to second order.  To first order, the Ellipsoidal collapse model reduces to the Zeldovich approximation; hence, our perturbation theory discussion of the first order approximation extends previous work on the Zeldovich approximation \\cite{kbg94,hks00}. However, our analysis also accounts for another subtlety associated with non-spherical collapse models---that the final shape of a patch is different from its initial shape \\cite{br91,pk93,brlc02}. We then present the results of the simplest excursion set treatment of this problem which accounts for the evolution in the volumes but not the shapes of cells.  The perturbation theory and excursion set predictions associated with spherical dynamics are presented in Section~\\ref{models}. This section also contains a discussion of the Lognormal model. It includes a comparison of these predictions with measurements in simulations. The perturbation and excursion set treatments of the Zeldovich and ellipsoidal collapse models are in Section~\\ref{stochastic}. A final section summarizes our findings, and includes a discussion of the fact that local deterministic mapping models of nonlinear evolution motivate a simple method for reconstructing the shape of the initial pdf from that at late times.  We illustrate the method using the spherical evolution model; we also show that, for the present purposes, the spherical model is a good enough approximation to the second order ellipsoidal collapse dynamics. ",
        "conclusions": "We discussed two approaches (perturbation theory and excursion set) for calculating the dark matter pdf using two different models for the dynamics (spherical collapse and the Zeldovich Approximation) in each case.  Although both approaches are deterministic, in the sense that the nonlinear evolution is determined by the initial conditions locally, the excursion set is slightly less `local'.  We showed that both are expected to give similar predictions whenever the dynamics is approximated by local deterministic mappings (equation~\\ref{eqn:PTpdf}), of which the spherical evolution model (equation~\\ref{eqn:SCapprox}) is a special case (Figure~\\ref{fig:besselk}).  The agreement is important, because the excursion set calculation allows one to connect discussions of the dark matter halo distribution with discussions of the pdf \\cite{rks98}.  Both the perturbation and excursion set approaches, when combined with the spherical evolution model, provide good descriptions of the nonlinear pdf seen in simulations (Figures~\\ref{fig:lgrho2Prho} and ~\\ref{fig:lgrho2Prho1-4}). This agreement is slightly unexpected, in the sense that the spherical evolution model ignores the fact that the nonlinear evolution of a region may be determined by quantities other than its initial density.  The ellipsoidal collapse model is a specific example of this; we studied its first (Zeldovich Approximation) and second order expansions in some detail.  In this model, the evolution of a patch is determined by its overdensity as well as the surrounding shear field (only the overdensity matters for the spherical collapse model).  As a result, nonlinear evolution changes the shape as well as the volume of a patch (in the spherical collapse model, only the volume changes). We found that including the evolution of the shapes is important when the density is high, so this is particularly important when the cell size is small. This shape dependence also complicates the excursion set calculation (Section~\\ref{zel}).  Despite its increased complexity compared to the spherical evolution model, we found that the Zeldovich Approximation resulted in significantly worse agreement with the simulations (Figure~\\ref{fig:lgrho2PrhoZA}).  Going to second order in the ellipsoidal collapse model resulted in more accurate predictions, except in the lowest density regions. The relatively large discrepancies in underdense regions, in both models, may be related to the following fact.  A Taylor series in $\\delta$ will converge more rapidly when $\\delta>0$ than when $\\delta<0$, since the latter will be an oscillating series. Thus, it may be that, to accurately represent the evolution of underdense regions, one must go beyond second order in the dynamics. Presumably, going to even higher order would further improve agreement with the simulations; this is the subject of work in progress.  Although our perturbation theory calculation accounts for the evolution of shapes, our excursion set calculation does not. This too is the subject of on-going work. The excursion set calculation is particularly interesting in view of the fact that the marriage of ellipsoidal collapse with the excursion set approach appears to provide a substantially improved prediction of dark halo abundances \\cite{smt01}. Finally, we note that our study of the real space pdf can be extended to redshift space; we are in the process of extending previous work on this problem \\cite{ps97,hks00,ecpdf}.  \\begin{figure*} \\begin{center} \\includegraphics[scale=0.75]{LamShethf6.ps} \\caption{Comparison of the reconstructed linear pdf of $\\delta/\\sigma$ (histogram) with the expected zero-mean unit-rms Gaussian form (solid line).  Different panels show results for reconstructions from different smoothing scales.  The method works well for large cells (top panels), but becomes increasingly inaccurate for smaller cells (bottom panels).  When the nonlinear rms fluctuation is larger than 2 (bottom right), then the reconstructed pdf is rather different from the Gaussian form, as one might expect from Figure~\\ref{fig:lgrho2Prho1-4}. Dotted curves show the bias which results from ignoring the extra weighting factor associated with the transformation from Eulerian to Lagrangian statistics.} \\label{fig:lineardm} \\end{center} \\end{figure*}  \\subsection{An extension} The accuracy of the spherical collapse based predictions, and the fact that the associated deterministic mapping from initial to final overdensity also provides a good description of the mean mapping in the ellipsoidal collapse model (Figure~\\ref{fig:ECmap}), suggests that one might be able to estimate the shape of the initial pdf from a measurement of the evolved one as follows. Starting from equation~\\ref{eqn:SCapprox} and the normalization method in section~\\ref{norm}, there is a mapping from the final density to its initial value. For each measured $M = \\bar\\rho V(1+\\delta_{\\rm NL})$ set \\begin{equation} \\nu = \\frac{1 - [(1+\\delta_{\\rm NL})/N_{\\rm sc}]^{-1/\\delta_{\\rm c}}} {\\sigma_{\\rm L}(M/N_{\\rm sc})/\\delta_{\\rm c}}, \\label{rescalenu} \\end{equation} where $\\delta_{\\rm c}$ is the critical linear density associated with collapse in the spherical model, and $N_{\\rm sc}$ is the corresponding normalization factor (Section~\\ref{norm}). Then make a histogram of $\\nu$, but rather than having each cell count equally, weight each by its value of $(1+\\delta_{\\rm NL})/N_{\\rm sc}$. The resulting distribution of $\\nu$ should provide a good estimate of the shape of the initial pdf of $(\\delta_{\\rm L}/\\sigma_{\\rm L})$.  Figure~\\ref{fig:lineardm} shows the results of this algorithm for a number of different scales.  Reconstruction of the initial pdf from the nonlinear pdf measured on scales larger than $8h^{-1}$Mpc works rather well (top panels).   The reconstructed distributions trace the same Gaussian shape very well; for reference, the smooth black curve shows a Gaussian with zero mean and unit variance. This mapping works well even when the cell size is as small as $4h^{-1}$Mpc (bottom left), but it fails at $1h^{-1}$Mpc (bottom right). The Figure also shows that it is important to account for the factor of $1+\\delta_{\\rm NL}$ which relates Eulerian and Lagrangian statistics: weighting each cell in the nonlinear distribution equally would have lead one to conclude incorrectly that the reconstruction is biased even on very large scales.  Studying whether or not this is relevant for the MAK reconstruction method \\cite{mak} is the subject of work in progress.  In effect, the spherical evolution model, when combined with the nonlinear pdf, is able to provide a good description of the initial pdf for scales where the rms fluctuation is smaller than about 2. In principle, this might be used as the basis for a test of the Gaussianity of the initial conditions because the algorithm makes no explicit assumption about the form of the distribution of $\\nu$.  The fact that the nonlinear pdfs from a range of different scales all map back to the same zero mean, unit variance Gaussian curve suggests that this method is a good test of Gaussianity. Although this is not the first method to reconstruct the shape of the initial one-point density field, it is simple and actually fares rather well compared to a number of previous methods (e.g. Kofman et al. 1994; Narayanan \\& Croft 1999).  This procedure is not as good a test of non-Gaussianity, in the sense that it does not yield a simple estimate of the shape of the initial distribution if it were non-Gaussian. This is because the reconstruction requires an estimate of the normalization factor $N_{\\rm sc}$, and to calculate it, one must first specify the form of the initial distribution.  So the algorithm above is really a self-consistency check:  it checks if the assumed form of the initial pdf is consistent with that reconstructed from the measured nonlinear pdf, under the assumption that spherical evolution is a good model whatever the initial fluctuation field.  Of course, rescaling to $\\nu$ involved a rescaling and transformation of the measured $M$, but it also involved $\\sigma(M)$; this is a consequence of the fact that, for Gaussian fluctuations, information about the volume enters only through the scale dependence of the rms value. For more general distributions this may not remain true, in which case the mapping to the initial variable $\\delta_{\\rm L}$ may be more complicated than equation~(\\ref{rescalenu}). Nevertheless, Figure~\\ref{fig:lineardm} suggests that this may be an interesting avenue to explore in future work.  We are in the process of extending this method to include redshift space effects, as well as galaxy bias.  We believe it may be interesting to merge it with the method recently proposed by Eisenstein et al. (2007) for reconstructing the Baryon Acoustic Oscillation feature in the galaxy distribution."
    },
    "0711/0711.2603.txt": {
        "abstract": "We have investigated the photoneutron cross section of the isotopes $^{148,150}$Nd, $^{154}$Sm, and $^{154,160}$Gd close to the neutron emission threshold in photoactivation experiments at the Darmstadt superconducting electron linear accelerator S-DALINAC. Naturally composed targets were activated with a high-intensity bremsstrahlung beam at various energies and the reaction yields have been determined by measuring the activity of the produced radioactive isotopes with HPGe detectors. The results are compared to two different statistical model calculations. ",
        "introduction": "The nucleosynthesis of the elements heavier than iron is dominated by the slow neutron capture process (\\emph{s}~process) and the rapid neutron capture process (\\emph{r}~process). However, 35 proton-rich stable isotopes cannot be produced in either of these two processes and are believed to be mainly synthesized by a combination of photodisintegration reactions, i.~e.\\ $(\\gamma,n)$, $(\\gamma,p)$ and $(\\gamma, \\alpha)$ reactions, in the explosive scenario of the \\emph{p} process \\cite{Lamb92,Arno03}. An accurate description of the nucleosynthesis within these processes demands reliable input from nuclear physics.\\\\ Calculations of astrophysical reaction networks within the \\emph{p} process have to account for a huge number of reaction rates. Many of the isotopes involved in the network are off the valley of stability and, therefore, are not accessible by experiments in many cases. But even for the stable nuclei experimental data for photodisintegration reactions in the astrophysically relevant energy region, i.~e.\\ close to the particle emission threshold, are rare. Hence, the reaction rates mainly need to be adopted from theoretical predictions. It is mandatory to test the predictive power of the calculations for those isotopes, which can be studied in the laboratory, and to prove the reliability of the predictions for the extrapolation to unstable nuclei.\\\\ The situation is much better for the \\emph{s}~process. Since for the unstable isotopes within the \\emph{s}-process reaction network the $\\beta$-decay rate usually highly exceeds the neutron capture rate, the reaction path follows the valley of stability and, hence, mainly stable isotopes are involved. Therefore, extensive experimental studies have been carried out in neutron capture experiments in the last decades and neutron capture cross sections for a wide range of isotopes have become available with uncertainties of only a few percent \\cite{Bao00,KADONIS}. With this enormous amount of nuclear physics input the understanding of \\emph{s}-process nucleosynthesis and its astrophysical sites has largely improved. However, for some unstable isotopes along the reaction path, the so-called \\emph{branching points}, $\\beta$-decay and neutron capture become competetive. Since the branching ratio of $\\beta$-decay and neutron capture is highly sensitive to the physical conditions, e.~g.\\ temperature and neutron density, in the stellar environment during the \\emph{s} process, these branching points serve as an excellent test for \\emph{s}-process nucleosynthesis models. Thus, it is essential to determine the neutron capture cross sections of these isotopes with high accuracy. Unfortunately, due to their short half-life neutron capture experiments can hardly be performed and, hence, further improvements of nucleosynthesis models are still hampered by the relatively large uncertainties of the cross section predictions for these isotopes \\cite{Kaep06}. Nevertheless, information for the neutron capture cross sections of branching points can be derived from studying the $(\\gamma, n)$ reaction of the stable neighbouring nucleus. The idea to determine $(n,\\gamma)$ reaction rates via photodisintegration experiments has already been presented in \\cite{Sonn03a}.\\\\ Experiments with real photons provide a well-suited tool to study photodisintegration cross sections and ground-state reaction rates. In previous experiments we have concentrated our investigations on isotopes of mass $A\\geq 186$ and confirmed the validity of several theoretical predictions in this mass region \\cite{Vogt01b, Vogt02,Sonn03a,Sonn04b,Muel06}. Average deviations between experiment and theory were typically less than 30\\%.\\\\ In this paper we want to address the photodisintegration reactions in the rare earth region (see Fig.~\\ref{fig:process_flow}). This region is of particular interest for the \\emph{s} process, because the isotopic abundance pattern is shaped by several branchings at the unstable isotopes of Nd, Pm, Sm, Eu and Gd. Since the relative abundances of the involved isotopes are known with an uncertainty of better than two percent \\cite{Ande89}, this mass region is exceptionally suited to test the stellar nucleosynthesis models with high accuracy. To provide precise nuclear physics input, neutron capture cross sections of several branching points have been measured with high accuracy in the last years \\cite{Jaag95,Reif03, Wiss06,Marr06}. But still one has to rely on theoretical predictions of the cross sections of short-lived branching points, e.~g.\\ $^{147}$Nd and $^{153}$Gd. Therefore, we have investigated the inverse photoneutron reactions of $^{148}$Nd and $^{154}$Gd in order to improve the reliability of these predictions.\\\\ Furthermore, the abundances of some isotopes in this region receive large contributions from the \\emph{p} process, e.~g.\\ a 33\\% contribution is predicted for $^{152}$Gd \\cite{Ande89}. For a comprehensive description of the nucleosynthesis in this mass region these contributions cannot be neglected and need to be studied in detail. Hence, the aim of this work is also to provide reliable data for the photoneutron reaction rates of some selected isotopes to prove the predictive power of the theoretical calculations in this mass region.\\\\  \\begin{figure}[htbp] % Requires \\usepackage{graphicx} \\includegraphics[width=1.\\columnwidth]{Fig1_A150Region}\\\\ \\caption{The \\emph{s}-process flow (solid lines) in the rare earth region. The thickness of the solid lines indicates the strength of the reaction flow. The reaction path is influenced by several branching points (dashed boxes). Unstable isotopes are shaded light grey ($\\beta^-$) and dark grey ($\\beta^+,\\epsilon$). Most isotopes also have nucleosynthesis contributions from the \\emph{r} and \\emph{p} process (dotted lines), respectively. The photoneutron reactions of the hatched isotopes have been studied in this work.}\\label{fig:process_flow} \\end{figure}  We have performed photoactivation experiments to study the $(\\gamma,n)$ reactions of the isotopes $^{148,150}$Nd, $^{154}$Sm and $^{154,160}$Gd in the astrophysically relevant energy region close to the neutron emission threshold at the superconducting Darmstadt electron linear accelerator S-DALINAC \\cite{Rich96}. Our experimental setup is discussed in detail in Sec.~\\ref{sec:Setup}. We explain the data analysis and the approach to derive astrophysically relevant ground-state reaction rates directly from the experiment by approximating a thermal Planck spectrum by a superposition of different bremsstrahlung spectra in Sec.~\\ref{sec:DataAnalysis}. The results are presented in Sec.~\\ref{sec:Results} and compared to two theoretical calculations based on the framework of the Hauser-Feshbach theory. ",
        "conclusions": "\\label{sec:Summary} In the last years extensive studies of photoneutron reactions with astrophysical implication for isotopes in the mass region $A\\geq 186$ have been carried out, but no experimental data have been provided for the rare earth isotopes so far. However, reliable experimental data for a wide range of isotopes in various mass regions are mandatory in order to constrain and to improve the theoretical calculations required for astrophysical reaction networks. The intention of this work was to test the theoretical predictions of two different statistical model codes for the photoneutron cross sections and ground-state reaction rates in the rare earth region. Although a variety of photoneutron reactions in this mass region was studied in the 70s, these experiments mainly focused on the investigation of the Giant Dipole Resonance. Hence, data points stemming from these experiments usually underlie large statistical and systematical uncertainties close to the neutron threshold and do not serve as a reliable test for theoretical predictions in the astrophysically relevant energy region. Therefore, we determined normalization factors for the calculations of the photoneutron cross sections of the isotopes $^{148,150}$Nd, $^{154}$Sm, and $^{154,160}$Gd in photoactivation experiments close to the neutron separation energy. Moreover, we derived $(\\gamma,n)$ ground-state reaction rates in the astrophysically relevant energy region for the \\emph{p} process using a superposition of bremsstrahlung spectra of various energies.\\\\ The need of experimental data of photoneutron reactions for astrophysical reaction networks is twofold. First, for \\emph{p}-process studies reliable predictions of reaction rates for a wide range of isotopes need to be provided. Since most of the isotopes involved in the \\emph{p}-process reaction network are unstable, the nuclear physics input for these isotopes is rare and, hence, theoretical calculations usually cannot be adjusted locally. Therefore, the reliability of the theoretical models should be tested using only global input parameters. In this work we found that in the rare earth region the predicted cross sections and the reaction rates for the selected isotopes agree with the experimental data within a factor of two. On average, a mean uncertainty of less than 30\\% has been determined. This is consistent with the results found in previous experiments in the mass region $A\\geq 186$. In addition, no systematic dependence of the deviation with atomic mass has been observed. On the one hand, this proves the reliability of the statistical model codes for a wide range of isotopes along the valley of stability. But it also indicates an appropriate predictive power for the reaction rates of unstable isotopes, since only global input parameters were used in the calculations . However, for almost all studied isotopes the cross sections were found to be overestimated by theory, which calls for further investigations. Moreover, it is planned to study the isotopic chain of the cerium isotopes in the near future to test the validity of the statistical model in this mass region for isotopes close to and at the neutron shell closure. Although the predictions agree fairly well with the experimental data of this work, it needs to be emphasized that experimental studies cannot account for the stellar enhancement of cross sections and reaction rates due to the thermal population of low-lying levels in the parent nuclei under stellar conditions. This so-called \\emph{stellar enhancement factor} can only be derived from theory.\\\\ The second astrophysical application of experimental studies of photoneutron reactions arises in the investigation of branching points within the \\emph{s}~process. In this context the photoneutron reactions of $^{148}$Nd and $^{154}$Gd are of particular interest among the studied isotopes of this work. Since the neighbouring branching points $^{147}$Nd and $^{153}$Gd cannot be studied in neutron capture experiments due to their short half-life, the aim is to derive the neutron capture cross sections of these branching points from the photoneutron reaction of the isotopes $^{148}$Nd and $^{154}$Gd, respectively. It was shown in this work, that the two adopted theoretical predictions overestimate the photoneutron cross sections by up to a factor of two. This might indicate that also the predicted neutron capture cross section of the inverse reactions need to be adjusted accordingly when being calculated within the same statistical model code. For astrophysical studies, a direct correlation between the stellar neutron capture and photodisintegration rate is described by the so-called \\emph{principle of detailed balance} \\cite{Fowl67,Raus00a}. This correlation only holds under the assumption that the low-lying levels in both the parent and the residual nucleus are thermally populated to a significant amount. This condition is fulfilled under sufficiently hot temperatures in a stellar environment. However, experiments in the laboratory only address transitions starting from the ground state of a certain nucleus and, therefore, the principle of detailed balance is not applicable any more for a direct comparison of the laboratory neutron capture and photodisintegration rate. Hence, without any further assumptions, conclusions for the neutron capture rate can hardly be drawn from the reaction rates derived in photodisintegration experiments.\\\\ Instead of directly deriving the neutron capture reaction rate from experiment a promising approach is to improve the nuclear physics input of the statistical model codes in order to increase the reliability of the theoretical predictions. To obtain the most accurate calculations for a single nucleus the input parameters should be adjusted locally. Therefore, the aim of future studies should be to provide improved nuclear physics input for $^{148}$Nd and $^{154}$Gd, while the experimental results of this work serve as a reliable test for any new complete set of input parameters. It can be assumed that a set of input parameters yielding improved predictions for the photoneutron reaction, will then also enhance the calculations of the neutron capture cross sections of the branching points $^{147}$Nd and $^{153}$Gd. At this point, however, we would like to make clear that the experimental data provided by this work set constraints on the absolute value of the photoneutron cross section, but do not allow for separately adjusting the various nuclear physics parameters such as the neutron optical potential, nuclear level densities and the $\\gamma$-ray strength function. Therefore, further experimental investigations with direct access to these parameters are mandatory in the future."
    },
    "0711/0711.1320_arXiv.txt": {
        "abstract": "{The Carina region is an excellent astrophysical laboratory for studying the feedback mechanisms of newly born, very massive stars within their natal giant molecular clouds (GMCs) at only 2.35~kpc distance.} {We use a clumpy PDR model to analyse the observed intensities of atomic carbon and CO and to derive the excitation conditions of the gas.} {The NANTEN2-4m submillimeter telescope was used to map the \\CI\\ $^3P_1-^3P_0$, $^3P_2-^3P_1$ and CO 4--3, 7--6 lines in two $4'\\times4'$ regions of Carina where molecular material interfaces with radiation from the massive star clusters. One region is the northern molecular cloud near the compact OB cluster Tr\\,14, and the second region is in the molecular cloud south of $\\eta$Car and Tr\\,16. These data were combined with $^{13}$CO SEST spectra, HIRES/IRAS $60\\,\\mu$m and $100\\,\\mu$m maps of the FIR continuum, and maps of 8$\\,\\mu$m IRAC/Spitzer and MSX emission.} {We used the HIRES far-infrared dust data to create a map of the FUV field heating the gas. The northern region shows an FUV field of a few $10^3$ in Draine units while the field of the southern region is about a factor 10 weaker. While the IRAC 8$\\,\\mu$m emission lights up at the edges of the molecular clouds, CO and also \\CI\\ appear to trace the H$_2$ gas column density.  The northern region shows a complex velocity and spatial structure, while the southern region shows an edge-on PDR with a single Gaussian velocity component.  We constructed models consisting of an ensemble of small spherically symmetric PDR clumps within the $38''$ beam (0.43~pc), which follow canonical power-law mass and mass-size distributions. We find that an average local clump density of $2\\,10^5$\\,cm$^{-3}$ is needed to reproduce the observed line emission at two selected interface positions. } {Stationary, clumpy PDR models reproduce the observed cooling lines of atomic carbon and CO at two positions in the Carina Nebula. } ",
        "introduction": "%  \\subsection{The importance of atomic carbon}  The fine structure lines of atomic carbon are strong coolants of the interstellar medium (ISM) throughout the universe: in Galactic clouds \\citep[e.g.][]{jakob2007}, in the Milky Way as a whole \\citep{fixsen1999}, as well as in external galaxies \\citep{bayet2006,kramer2005} up to high redshifts \\citep[][]{weiss2003}. Large-scale observations of Galactic clouds show \\CI\\ emission coincident with that of CO. This holds for giant molecular clouds (GMCs) \\citep{sakai2006}, as well as for rather quiescent dark clouds \\citep{tatematsu1999}. It has therefore been suggested that \\CI\\ is a more reliable tracer of the molecular cloud masses in external galaxies than CO \\citep{papadopoulos2004}. However, all steady-state models of photon-dominated regions (PDRs) predict that atomic carbon is formed in a surface layer by recombination of C$^+$ and dissociation of CO. This appears to be in contradiction with atomic carbon emission which is observed even at large distances from UV sources. The dynamics of the ISM may play a decisive role, allowing to explain the distribution of \\CI\\ with time dependent chemical models \\citep{oka2004,stoerzer1997}.  However, stationary, but clumpy PDR models have been successful in explaining e.g.  extended \\CI\\ $^3P_1-^3P_0$ and $^{13}$CO\\,$2-1$ emission in S\\,140 \\citep{spaans1996,spaans1997}.  Inside a clumpy PDR, the UV radiation can penetrate much deeper than it would be possible in a homogeneous layer \\citep{boisse1990}. \\citet{meixner1993} showed that the line intensities of all lines are enhanced compared to a homogeneous model with the same average density. \\cite{juvela2001} presented a model to apply Monte Carlo radiative transfer calculations to the results of magnetohydrodynamic calculations to account for a clumpy and turbulent structure. They also found that local density condensations lead to an increased cooling efficiency and reduced photon trapping. Another approach is to describe both the turbulent velocity field and the density structure by their statistical properties. \\cite{hegmann2007} calculate CO cooling rates based on a stochastic radiative transfer model, which accounts for density and velocity fluctuations with a finite correlation length.  Here, we study the distribution of atomic carbon and warm CO in the Carina star forming region and test clumpy models using the KOSMA-$\\tau$ PDR model \\citep{stoerzer1996,roellig2007} to derive the excitation conditions of the gas. Spaans et al. model the inhomogeneous structure with a Monte-Carlo code, assuming a fixed clump/interclump density ratio of 10. In contrast, we assume that all emission analyzed stems from the clumps. The influence of the interclump medium or a halo leading to pre-shielding of CO has been discussed by \\citet{bensch2006} and \\citet{bensch2003}, who also used KOSMA-$\\tau$, but is not considered here.  To describe the clumpy structure, we assume that the clumps are distributed following the canonical clump mass and mass-size distributions which have been found in molecular clouds.  \\begin{figure*}[ht] \\centering \\includegraphics[angle=-90,width=18cm]{7815fig1.eps} \\caption{Overview of the Carina clouds near $\\eta$Car lying at 0/0 10:45:03.59 -59:41:04.3 (J2000) at the center of the Tr\\,16 cluster.  {\\bf Left:} Estimate of the FUV flux $\\chi$ in Draine units  at a resolution of $\\sim1'$ and derived from HIRES/IRAS. Dashed contours range from 400 to 800 in steps of 200. Solid contours range from $10^3$ to $5\\,10^3$ in steps of $10^3$. {\\bf Right:} MSX emission at 8\\,$\\mu$m at a resolution of $\\sim20''$.  Contours are 1 to 10 in steps of $1~10^{-5}$ Wm$^{-2}$sr$^{-1}$ \\citep[cf.  Figs.\\,7, 11 in][]{rathborne2002}. Note the diffraction artefacts near $\\eta$Car.  The center position of the OB cluster Tr\\,14 is marked by a triangle.  The two $4'\\times4'$ regions mapped with NANTEN2 are delineated by boxes.  Crosses in the left panel mark two cuts discussed in the text.  } \\label{fig-overview} \\end{figure*}  \\begin{center} \\begin{table*}[ht] \\caption[]{\\label{tab-carn-wco} {\\small Physical parameters derived from the observed integrated intensities $I$ in K\\,kms$^{-1}$ and intensity ratios $R^{74}_{\\rm CO}$=$I$(CO 7--6)/$I$(CO 4--3), $R^{21}_{\\rm C}$=$I$(\\CI 2--1)/$I$(\\CI 1--0), $R^{14}_{\\rm CCO}$=$I$(\\CI 1--0)/$I$(CO 4--3) from Gaussian fits to the line profiles at two interface positions. The calibration error of the integrated intensities is estimated to be 15\\%. }} \\begin{tabular}{lllllllllrrrrrr} \\hline \\hline $\\Delta\\alpha/\\Delta\\delta$ & $^{13}$CO & \\multicolumn{2}{c}{CO} & \\multicolumn{2}{c}{\\CI} & $R^{74}_{\\rm CO}$ & $R^{21}_{\\rm C}$ & $R^{14}_{\\rm CCO}$ & $T_{\\rm ex}$ & $N$(CO) & $N$(C) & $N$(H$_2$) & M & C/CO \\\\ $(',')$ & 2--1 & 4--3 & 7--6 & 1--0 & 2--1 & & & & K & & & & \\msol \\\\ \\hline $-11.7/6.8$  & 47 & 227 & 86 & 37 & 27 & 0.38 & 0.72 & 0.16 &  50 &  23$^{30}_{18}$ &   4.9$^{5.8}_{4.1}$ &  28$^{35}_{21}$ &  100 & 0.21$^{0.32}_{0.13}$ \\\\ $+0.8/-7.2$    & 40  & 90 & 24 & 31 & 19 & 0.27 & 0.6 & 0.34 &  30 &  15$^{18}_{22}$ &   3.8$^{4.5}_{3.3}$ &  18$^{22}_{14}$ &  63 & 0.25$^{0.38}_{0.18}$ \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\end{tabular} \\end{table*} \\end{center}  \\subsection{The Carina nebula}  The Carina nebula is a very prominent southern star-forming region at a distance of 2.35~kpc \\citep{smith2006b}.  This giant molecular cloud was first observed by \\citet{grabelsky1988} as part of the Columbia CO 1--0 survey of the Milky Way. They derived a total mass of $6.7\\,10^5$\\,\\msun\\ and an effective radius of 66\\,pc.  In total 65 O-type stars, including 6 of the 11 O3-type stars known in the Milky Way \\citep{smith2006}, produce a strong UV field and strong stellar winds which interact vigorously with the material in the surrounding molecular clouds.  This high concentration of the earliest-type stars is unique in the Galaxy and may serve as a template for more extreme but more distant regions containing stellar super clusters like the central regions of NGC\\,253 or the Antennae \\citep{bayet2006,schulz2007}.  The Carina nebula extends more than 4 square degrees on the sky.  It exhibits a peculiar morphological structure on all scales.  Massive stars are being born in several regions within molecular condensations, located mainly to the south-east and north-west of the cluster Trumpler\\,16, which is centered at $\\eta$Car, a luminous blue variable star.  $\\eta$Car is embedded in the Homunculus nebula \\citep[e.g.][]{smith2006} and its luminosity has changed drastically and repeatedly in the past centuries, leading to strong variations of the FUV field impinging on the surrounding clouds. Its current luminosity of $5~10^6$\\,\\lsun\\ makes it one of the brightest celestial infrared objects \\citep{cox1995}. In the northern part of the Carina nebula, the stellar population is dominated by the cluster Tr\\,14. Recently, \\citet{oberst2006} reported the detection of the \\NII\\ line at $205\\,\\mu$m at the radio peak of the Carina II \\HII\\ region using SPIFI at AST/RO. They compared their data with \\CII\\ data \\citep{mizutani2004,mizutani2002} and concluded that about a third of the \\CII\\ emission stems from a low-density ionized medium. For a recent overview of the Carina region, see \\citet{smith_brooks2007}.      Carina was mapped at $\\sim3'$ resolution in the CO 4--3 \\newline and \\CI\\ $^3P_1-^3P_0$ lines by \\citet{zhang2001} using the AST/RO 1.7~m telescope. These AST/RO data show large-scale \\CI\\ emission, coextensive with CO 4--3.   Here, we have used the 4~m NANTEN2 telescope to map $^{12}$CO in the rotational transitions $J=$4--3, 7--6, and \\CI\\ in the fine structure transitions $^3P_1-^3P_0$, and \\newline $^3P_2-^3P_1$ (henceforth 1--0 and 2--1) at sub-arcminute resolutions in two $4'\\times4'$ regions of Carina (Fig.\\,\\ref{fig-overview}).  One region to the north of $\\eta$Car covers the cloud interface with Tr\\,14.  The second region covers another cloud interface to the south of $\\eta$Car and Tr\\,16.   The northern molecular cloud to the west of Tr\\,14 (Fig\\,\\ref{fig-overview}) contains several sites of massive star formation.  Diffuse PAH $3.29\\,\\mu$m emission has been found by \\citet{rathborne2002} to trace the sharp edge of the cloud.  Peak intensities at the edge towards the Car\\,I \\HII\\ region are almost 10 times higher than in the other regions of Carina.  Car\\,I-E is interacting with the front face of the GMC and creating a PDR seen edge-on \\citep{brooks2003,brooks2001}.  Car\\,I-W manifests a second PDR seen face-on (cf. Fig.\\,\\ref{fig-carina-n-spitzer}).  The CO 2--1 map of \\citet{cox1995} shows a sharp rim of gas delineating the southern molecular cloud, located south of $\\eta$Car and the Tr\\,16 cluster.  Several IRAS sources are located along the edge. A prominent peak in PAH 3.29~$\\mu$m emission is the ultra compact \\HII\\ region IRAS\\,$10430-5931$ \\citep{rathborne2002}. ",
        "conclusions": "%  We have mapped the emission of atomic carbon and CO in two fields of the GMCs surrounding $\\eta$Car. Combining CO 4--3 and \\CI\\ 1--0 data with CO 7--6 and \\CI\\ 2--1 data allows us to study the excitation conditions of both gas tracers.  The northern $4'\\times4'$ field lies adjacent to the compact OB cluster Tr\\,14 and is associated with the \\HII\\ regions Car\\,I-E and Car\\,I-W. The spectral lines show a rich kinematical structure, with several velocity components along the lines of sight spread between $-10$ and $-30\\,$kms$^{-1}$.  The southern $4'\\times4'$ field lies on a molecular ridge south of $\\eta$Car. Here, spectral lines show only one rather narrow Gaussian shaped velocity component. The FUV field derived from HIRES/IRAS far-infrared data is $\\chi\\sim10^{2.5}$, about a factor 10 weaker than in the northern field. IRAC/Spitzer maps at 8$\\mu$m show the detailed structure of the FUV illuminated dust. The edge of the southern region is clearly delineated by the $8\\mu$m emission, while the northern region shows a more complex spatial structure.  In both regions, CO and \\CI\\ emission is co-extensive. Both species appear to trace the bulk of molecular gas, rather than the interface regions. This coincidence has been found in many of the previous large-scale studies of Galactic clouds conducted with the Mt.Fuji, KOSMA, and other telescopes. \\citet{papadopoulos2004} stress the importance of cosmic rays in raising the C/CO abundance ratio and argue that dynamic and non-equilibrium chemistry processes explain the ubiquity of \\CI\\ within molecular clouds as well as at PDR interfaces. Here, we argue that steady state, clumpy PDR models provide an alternative explanation.  The northern and southern Carina regions both show a slight brightening of the CO 7--6 and \\CI\\ 2--1 transitions relative to the 4--3 and 1--0 transitions near the interfaces. In the south-east part of the northern region, the upper transitions brighten towards the interface (Fig.\\,\\ref{fig-carina-n-co43-ci10}).  These transitions also show increased intensities near the embedded IRAS point source in the southern region (Fig.\\,\\ref{fig-carina-s-co43-ci10}).  We selected two cuts from the cloud cores through the interface regions, roughly pointing towards the illuminating sources.  The southern cut lies to the east of IRAS\\,10430-5931, avoiding its complicated structure and additional heating sources. The northern cut runs from the peak of \\CII\\ emission \\citep{brooks2003} through the \\HII\\ regions Car\\,I-W and I-E across the edge-on interface towards the center of the Tr\\,14 cluster.  The peak $^{12}$CO line temperatures along the northern cut indicate kinetic gas temperatures of at least 50\\,K, while a lower limit of 30\\,K is found along the southern cut.  We selected two positions at the northern and southern interfaces for a detailed analysis.  The $^{12}$CO 7--6/4--3 and \\CI\\ 2--1/1--0 intensity ratios at these positions are 0.3--0.4 and 0.6--0.7, respectively.  We use PDR models to interpret the observed line intensities. Our models assume that all observed emission stems from an ensemble of spherically symmetric PDR clumps and there is no emission from an interclump medium or a diffuse halo surrounding the clouds. PDR clumps are modelled using the stationary KOSMA-$\\tau$ code \\citep{roellig2006}.  The clump distributions follow the canonical mass and radius distributions found for molecular clouds. At the two positions analyzed in detail, we find that clumpy PDR models are consistent with the observed absolute intensities of the $^{12}$CO and \\CI\\ lines to within 20\\%, i.e. at about the calibration accuracy.  Since the line ratios observed at the two interface positions are fairly typical for the entire observed regions, we conclude that stationary, clumpy PDR models can simultaneously reproduce the observed \\CI\\ and CO emission of the lower transitions and the upper CO 7--6 and \\CI\\ 2--1 transitions.  However, the observed $^{13}$CO 2--1 intensities are only about half of the modelled intensities at the northern and southern positions, respectively. Interestingly, \\citet{pineda2007} report a similar finding. They analyzed $^{12}$CO and \\CI\\ emission observed in the dark cloud globule B\\,68. This emission can be reproduced using a single KOSMA-$\\tau$ spherically symmetric PDR model, while the observed $^{13}$CO intensities are a factor $\\sim2$ too low.  To study the gas under a broad range of conditions, we plan to extend the present \\CI\\ and CO maps to fully cover the GMCs surrounding $\\eta$Car using the SMART 492/810\\,GHz multipixel array receiver at NANTEN2 \\citep{graf2003}.  Future observations with APEX and Herschel will tell in how far the emission of other species tracing the detailed photo-chemical network is also consistent with stationary, clumpy PDR models."
    },
    "0711/0711.3922_arXiv.txt": {
        "abstract": "The atmospheres of Jupiter and Saturn exhibit strong and stable zonal winds. How deep the winds penetrate unabated into each planet is unknown. Our investigation favors shallow winds. It consists of two parts.  The first part makes use of an Ohmic constraint; Ohmic dissipation associated with the planet's magnetic field cannot exceed the planet's net luminosity. Application to Jupiter (J) and Saturn (S) shows that the observed zonal winds cannot penetrate below a depth at which the electrical conductivity is about six orders of magnitude smaller than its value at the molecular-metallic transition. Measured values of the electrical conductivity of molecular hydrogen yield radii of maximum penetration of $0.96R_J$ and $0.86R_S$, with uncertainties of a few percent of $R$. At these radii, the magnetic Reynolds number based on the zonal wind velocity and the scale height of the magnetic diffusivity is of order unity. These limits are insensitive to difficulties in modeling turbulent convection. They permit complete penetration along cylinders of the equatorial jets observed in the atmospheres of Jupiter and Saturn.  The second part investigates how deep the observed zonal winds actually do penetrate. As it applies heuristic models of turbulent convection, its conclusions must be regarded as tentative. Truncation of the winds in the planet's convective envelope would involve breaking the Taylor-Proudman constraint on cylindrical flow. This would require a suitable nonpotential acceleration which none of the obvious candidates appears able to provide. Accelerations arising from entropy gradients, magnetic stresses, and Reynolds stresses appear to be much too weak.  These considerations suggest that strong zonal winds are confined to shallow, stably stratified layers, with equatorial jets being the possible exception. ",
        "introduction": "\\label{intro}  Jupiter and Saturn, are composed primarily of hydrogen and helium with small additions of heavier elements. Their atmospheres exhibit strong, stable zonal winds composed of multiple jets associated with azimuthal cloud bands \\citep{ingersoll90}. Zonal winds peak in the equatorial region reaching $\\sim 100 \\m \\s^{-1}$ on Jupiter and $\\sim 400 \\m \\s^{-1}$ on Saturn.\\footnote{Wind speeds on Jupiter are determined relative to System III coordinates which rotate with the angular speed of the planet's magnetic field \\citep{dessler83}. Only differences among wind speeds on Saturn are known because the planet's internal rotaton rate is uncertain.} The latitudes of Jupiter's jets have not changed for at least 80 years \\citep{smith76} and their velocities have been constant within $10\\%$ over $25$ years \\citep{porco03}.  \\\\  The depth of the zonal winds is unknown. Both deep and shallow flow models have been proposed.  Wind speeds measured by the Galileo probe at $7.4^{\\circ}N$ on Jupiter rose from $90\\m\\s^{-1}$ at $0.4 \\mathbar$ to $180\\m \\s^{-1}$ at $\\sim 5 \\mathbar$ and then remained nearly constant until $22 \\mathbar$ \\citep{atkinson97,atkinson98}. It is important to bear in mind that these measurements only sample the winds in the outer $1 \\%$ of the planet's radius where the electrical conductivity is low. In regions of high electrical conductivity, the magnetic field lines are frozen into the fluid. Winds in these regions would cause changes in the external magnetic field. By comparing Galileo and Pioneer/Voyager data, \\citet{russell01a,russell01b} find that increases of $0.3\\deg$ in the dipole tilt and $1.5\\%$ in Jupiter's dipole moment may have taken place between 1975 and 2000. The former could be accounted for by meridional flows with speeds in the deep interior of order $0.1\\cm\\s^{-1}$ \\citep{guillot04}. \\\\  \\citet{busse76,busse83,busse94} advocates deep flows. He applies the Taylor-Proudman theorem \\citep{taylor23} to deduce that zonal flows extend along cylinders oriented parallel the rotation axis in the molecular envelope, and then terminate at the outer boundary of the metallic core where he assumes that hydrogen undergoes a first order phase transition. But data from shock wave experiments show that hydrogen undergoes a continuous transition from a semi-conducting molecular state to a highly conducting metallic state as the pressure increases. This contradicts the assumption of a first order phase transition at the core-envelope boundary.  \\\\  Recently, a modified deep flow model for Jovian zonal flows has been proposed based on simulations of convection in a shell with a lower boundary near $0.9R_J$ \\citep{aurnou04,heimpel05}. The physical meaning of the lower boundary in the modified deep flow model is obscure. Hydrogen cannot undergo a phase change at that radius \\citep{guillot04}. So how might the Taylor-Proudman constraint be violated in order to reduce the zonal flow to a near zero value below that boundary? We address related issues in \\S \\ref{sec:truncate}. \\\\  In shallow flow models, the observed high-speed flow is confined to a thin, baroclinic layer near the cloud level; the interior flow is much slower. Even if the high velocity flow is confined to a shallow layer, its forcing may occur at depth. For example, if the flow were to arise from a process that conserved angular momentum per unit volume, $\\rho U$ would be approximately conserved, where $\\rho$ is the density and $U$ is the magnitude of the flow velocity. Since the density in the interior is several orders of magnitudes larger than that near the surface, the flow velocity could then be much greater near the surface. On the other hand, the observed zonal flow might be generated by shallow forcing due to the turbulence injected at the cloud level by moist convection, differential latitudinal solar heating, latent heat release from condensation of water, or other weather layer processes \\citep{vasavada05}. From the thermal wind equation, a latitudinal temperature gradient of about $5-10K$ across a few pressure scale heights below the cloud level would cause substantial vertical shear, which makes the flow velocity much greater near the surface than deeper down \\citep{ingersoll69,ingersoll84,vasavada05}. \\\\  The plan of our paper is as follows. Relevant details of the electrical conductivity of molecular hydrogen as measured in shock wave experiments are presented in \\S \\ref{odsec:conductivity}. Sections \\ref{sec:extrapolation} and \\ref{sec:cylindrical} are devoted to the calculation of Ohmic dissipation based on the assumption that the zonal wind penetrates the planet along cylinders. In the former, the poloidal magnetic field is determined by downward extrapolation of the external field. This procedure is appropriate in regions where the magnetic Reynolds based on the convective velocity, $R_m^c\\ll 1$. In the latter, we examine the consequences of assuming that the poloidal magnetic field is parallel to the rotation axis in regions where $R_m^c\\gg 1$. The requirement that the total Ohmic dissipation be bounded from above by the planet's net luminosity, $\\cal L$, limits the depth to which the observed zonal winds can penetrate. Section \\ref{sec:truncate} asks whether the zonal winds might be truncated within the convective envelope. A short summary of our main results is given in \\S \\ref{sec:discussion}. A few technical details are relegated to the Appendix. \\\\ ",
        "conclusions": "\\label{sec:discussion}   The condition that the total Ohmic dissipation not exceed the planet's net luminosity sets an upper bound on the depth to which the zonal flows observed in the atmospheres of Jupiter and Saturn could penetrate. At these depths, the magnetic Reynolds number, based on the observed zonal winds and the scale height of the magnetic diffusivity, is of order unity.  We consider it unlikely that the observed flows extend to the depth of maximum penetration because that would require the Taylor-Proudman constraint to be violated in the convective envelope. We have been unable to identify any plausible mechanism that could do this. Lack of a rigorous model for convection is a weakness. We base our analysis of the robustness of the Taylor-Proudman constraint on the mixing length model designed by \\citet{ingersoll82} to apply under conditions of strong differential rotation. However, we have confirmed that the conclusions we draw from it also follow from other convection models that consider only solid body rotation, or even those that neglect the effects of rotation entirely.  The boundaries of the cylindrical extensions of the equatorial jets essentially coincide with the maximum penetration depths. Thus these jets could maintain constant velocities along cylinders through the planets. Winds measured at $7.4^{\\circ} N$ by the Galileo probe as it descended into Jupiter are consistent with this possibility.  \\paragraph"
    },
    "0711/0711.3780_arXiv.txt": {
        "abstract": "Following our discovery of radio pulsations from the newly recognized Anomalous X-ray Pulsar (AXP) \\mag, we initiated X-ray monitoring with the \\swift\\ X-ray Telescope, and obtained a single target-of-opportunity observation with the {\\em Newton X-ray Multi-Mirror Mission} (\\xmm).  In comparison with its historic minimum flux of $3 \\times 10^{-13}$ ergs cm$^{-2}$ s$^{-1}$, the source was found to be in a record high state, $f_X(1$--$8\\,\\mbox{keV}) = 5 \\times 10^{-12}$ ergs cm$^{-2}$ s$^{-1}$, or $L_X = 1.7 \\times 10^{35}(d/9\\ {\\rm kpc})^2$ ergs~s$^{-1}$, and declining by 25\\% in 1 month. Extrapolating the decay, we bound the total energy in this outburst to $10^{42} < E < 10^{43}$ ergs. The spectra (fitted with a Comptonized blackbody) show that an increase in the temperature and area of a hot region, to 0.5~keV and $\\sim 16\\%$ of the surface area of the neutron star, respectively, are primarily responsible for its increase in luminosity. The energy, spectrum, and timescale of decay are consistent with a deep crustal heating event, similar to an interpretation of the X-ray turn-on of the transient AXP \\xte. Simultaneous with the 4.6 hour \\xmm\\ observation, we observed at 6.4~GHz with the Parkes telescope, measuring the phase relationship of the radio and X-ray pulse. The X-ray pulsed fraction of \\mag\\ is only $\\sim 7\\%$, while its radio pulse is relatively broad for such a slow pulsar, which may indicate a nearly aligned rotator. As also inferred from the transient behavior of \\xte, the only other AXP known to emit in the radio, the magnetic field rearrangement responsible for this X-ray outburst of \\mag\\ is probably the cause of its radio turn-on. ",
        "introduction": "\\label{sec:intro}  Anomalous X-ray pulsars (AXPs) and soft gamma-ray repeaters (SGRs) are young neutron stars (NSs) with rotation period of 2--12\\,s and inferred surface magnetic field strength $B \\approx 10^{14-15}$\\,G.  See \\citet{wt06} and \\citet{kas07} for recent reviews.  In the magnetar model \\citep[][1996]{dt92a,td95}\\nocite{td96a}, the rearrangement and decay of their extreme fields is responsible for their large and variable X-ray luminosity, which exceeds that available from rotational braking. Thirteen magnetars are confirmed\\footnote{Nine AXPs and four SGRs; there are two more candidates.  See catalog at http://www.physics.mcgill.ca/$\\sim$pulsar/magnetar/main.html.}, of which the AXP \\mag\\ is the most recent. Discovered with the \\einstein\\ X-ray satellite in 1980 \\citep{lam81}, \\mag\\ was identified as a magnetar candidate in the center of the small candidate SNR \\snr\\ by \\citet{gg07}.  It was subsequently observed to be emitting radio pulsations at a period of $P=2$~s \\citep{cam07b}, with spin-down properties of a magnetar ($\\dot P = 2.3 \\times 10^{-11}, B = 2.2 \\times 10^{14}$~G), and a distance estimate of 9~kpc from its dispersion measure (DM). \\mag\\ (\\psr) and the 5.5\\,s AXP \\xte\\ are the only magnetars known to emit in the radio \\citep{crh+06}.  Both are demonstrably transient radio sources, having not been detected in previous surveys of adequate sensitivity.  Here, we report on new X-ray observations of \\mag\\ made shortly after the radio discovery of its pulsations, and compare them with archival X-ray data in order to understand the nature of its apparent X-ray outburst from its spectral and pulse properties. Radio pulsar coverage was also obtained at the same time as one of the X-ray observations, and the phase relationship of the radio and X-ray pulse was determined. ",
        "conclusions": "Prompted by its discovery as a transient radio pulsar, we obtained new X-ray observations of \\psr\\ that reveal it in the highest state yet observed, with $L_X \\approx 1.7 \\times 10^{35}$ ergs~s$^{-1}$ and declining. The peak of the outburst was not observed in X-rays, but it could have been several times higher.  As this is the second AXP to be detected as a transient radio source, after \\xte, we infer that magnetic field rearrangement accompanying an X-ray outburst may be a necessary (although not sufficient) condition for radio turn-on of an AXP.  We also detected X-ray pulsations from \\psr\\ for the first time, although the pulsed fraction $\\sim 7\\%$ is among the smallest observed from AXPs. Together with the properties of the simultaneously observed radio pulse, this may indicate a nearly aligned geometry.  In both quiescence and outburst, the X-ray spectrum is dominated by a small blackbody hot spot of $kT_{\\rm BB} = 0.4-0.5$~keV, with a lesser Comptonized component, as is the case for most AXPs. The changing luminosity is due mostly to the increase in area of the blackbody, and partly to its increase in temperature.  This, plus the energy and time scale of the decay, suggests an origin of the brightening in a deep crustal heating event rather than a change in magnetospheric currents.  The historical X-ray luminosities of \\mag\\ span only a factor of 16, including some intermediate states, which is not as extreme a range as experienced by \\xte\\ and CXOU J164710.2$-$455216 (a factor of $\\sim 100$). It is not yet clear if we have sampled the total range of luminosity of \\mag, or what the duty cycle and decay time of its largest outbursts are."
    },
    "0711/0711.4675_arXiv.txt": {
        "abstract": "We report on a comprehensive analysis of the kilo-Hz ($\\ge 600$ Hz) quasi-periodic oscillations (kHz QPOs) detected from the neutron star X-ray transient Aquila X-1 (Aql X-1) with the Rossi X-ray Timing Explorer, between 1997 and 2007. Among kHz QPO sources, Aql X-1 is peculiar because so far only one kHz QPO has been reported, whereas in most sources, two kHz QPOs are usually detected (a lower and an upper kHz QPO). The identification of the QPOs reported so far has therefore been ambiguous, although it has been proposed that they were likely to be the lower QPO. Following up on previous work, we confirm the identification of the QPOs previously reported as lower QPOs, because of their high quality factors and the quality factor versus frequency dependency, which are similar to those observed in other sources. Combining all segments of data containing a lower QPO, we detect for the first time an upper kHz QPO. As in other sources for which the neutron star spin frequency is larger than 400 Hz (550.25 Hz in Aql X-1), the frequency difference between the two kHz QPOs is close to half the spin frequency. Based on this result, we re-examine the link between the neutron star spin and the frequency of the kHz QPOs, to show that a model in which the separation of the lower and upper QPOs relates to the neutron star spin frequency is still as good as any comparably simple model. ",
        "introduction": "Since its launch in 1995, the Rossi X-ray Timing Explorer (Bradt, Swank and Rothschild, 1993) has detected kHz QPOs in more than 25 accreting systems containing a weakly magnetized neutron star (see van der Klis 2006 for a review). In most sources, two kHz QPOs are usually detected; the lower of the two is seen as a relatively narrow peak in the Fourier Power Density Spectrum (PDS) with a quality factor ($Q=\\nu/\\Delta\\nu$) exceeding 200 (e.g., Barret, et al. 2005a), whereas the upper QPO is typically much broader ($Q\\sim 5-20$). This makes the lower QPO easier to detect, especially on short timescales. Both the lower and upper QPOs vary in frequency with time, but the frequency difference remains always close to the spin frequency of the neutron star (or half its value). Removing the contribution of the frequency drift to the measured QPO width, Barret, Olive, Miller (2005) have demonstrated that the lower and upper QPOs follow a different path in a quality factor versus frequency diagram. In particular, they found that the quality factor of the lower kHz QPOs increases smoothly with frequency, saturates at a maximum value, beyond which a sharp drop is observed (Barret, Olive, Miller 2005). The same behaviour has been observed in several different systems (Barret, Olive, Miller 2006), suggesting that the drop is related to a special location in spacetime, e.g., the innermost stable circular orbit (ISCO) (Barret, Olive, Miller, 2006, 2007).  Previous investigations of the kHz QPOs detected from the recurrent X-ray transient Aql X-1 (Zhang et al. 1998, Cui et al. 1998, Reig et al. 2000, M\\'endez et al. 2001, Reig et al. 2004) have reported only a single QPO, making its identification difficult. By comparison with the properties of lower QPOs seen in twin QPO sources (correlation of the frequency and spectral colors, QPO width, RMS-frequency dependency, energy spectrum of the QPOs), it was however proposed that the QPOs detected were likely to be lower QPOs (M\\'endez et al. 2001).  Aql X-1 contains a rapidly rotating neutron star, spinning at 550.25 Hz, as inferred from the discovery of an episode of coherent pulsation in its persistent emission (Casella et al. 2007).  This frequency is close to the previously detected frequency of X-ray burst oscillations (Zhang et al. 1998). By analogy with other sources spinning at a frequency above 400 Hz, one would therefore expect an upper QPO to be detected with a frequency separation close to 275 Hz (see van der Klis 2006 for a review). Despite extensive searches, no such QPO has yet been reported.  In this paper, we analyze in a homogenous way all archival RXTE data on Aql X-1, with the goal of studying the quality factor of the QPOs and searching for an upper QPO. The data used here have been presented in Zhang et al. (1998), Cui et al. (1998), Reig et al. (2000), M\\'endez et al. (2001) and Reig et al. (2004). In the next section, we present our analysis scheme, which builds upon previous previously described procedures in Barret, Olive \\& Miller (2005, 2006, 2007).  We then present our results, which reveal for the first time 1) that the quality factor of the lower QPOs follows the same pattern seen in other sources (e.g. 4U~1608--522 \\& 4U~1636--536, Barret et al. 2006), and 2) Aql X-1 displays an upper QPO that is detected close to half the spin frequency of the neutron star (275 Hz), when combining all the data available. Starting with this result, we then discuss on the link between kHz QPOs and neutron star spin.  \\begin{table*} \\centering \\caption{QPOs from Aql X-1. The observations are ObsID sorted. The name of the ObsID, the starting date of the observation, the cumulative integration time of all the 16 PDS shifted (T$_{obs}$, the fraction of PDS per ObsID, for which an instantaneous frequency could be estimated varies from about 40\\% up to 100 \\%), the total source count rate (Cts/s), the QPO frequency  ($\\bar{\\nu}$) to which all the PDS were shifted, the range of variation of the QPO frequency within the ObsID ($\\nu_{min}-\\nu_{max}$), the mean quality ($\\bar{Q}$) and mean amplitude (RMS) are listed. All errors are computed such that $\\Delta\\chi^2=1$. The RMS is computed as $\\sqrt{I_{lor}/S}$, where $I_{lor}$ is the fitted Lorentzian integral and $S$ is the source count rate. Because the error on the source count rate is negligible, the error on RMS is computed as $1/2 \\times RMS \\times \\Delta I_{lor}/I_{lor}$, where $ \\Delta I_{lor}$ is the error on $I_{lor}$, computed with $\\Delta\\chi^2=1$.  The significance of the QPO which is $I_{lor}/\\Delta I_{lor}$ can thus be retreived from the error quoted on the RMS: it exceeds typically $\\sim 7\\sigma$ and goes up to $\\sim 20\\sigma$. The high quality factor recovered suggests that the QPOs detected are all lower kHz QPOs.  Previous analysis of these data have been presented in Zhang et al. (1998), Cui et al. (1998), Reig et al. (2000), M\\'endez et al. (2001) and Reig et al. (2004).} \\begin{tabular}{@{}|cccccccc|@{}} \\hline ObsID & Date & T$_{obs}$ & Cts/s & $\\bar{\\nu}$ & $\\nu_{min}-\\nu_{max}$ &$\\bar{Q}$ & RMS (\\%)   \\\\ \\hline 20092-01-01-02 & 1997/08/13-11:09:35 &  656.0 & 1380.2 &  833.7 & $ 829.0- 837.6$ & $ 129.7\\pm  14.7$ & $ 7.2\\pm 0.3$ \\\\ 20092-01-02-01 & 1997/08/15-19:03:17 & 1232.0 & 1520.1 &  875.2 & $ 869.5- 883.0$ & $  96.8\\pm  15.3$ & $ 5.6\\pm 0.3$ \\\\ 20092-01-02-03 & 1997/08/17-06:06:08 &  592.0 & 1571.9 &  848.3 & $ 846.3- 850.6$ & $ 115.3\\pm  16.8$ & $ 6.6\\pm 0.4$ \\\\ 20092-01-05-01 & 1997/09/06-12:50:24 & 1360.0 & 1313.1 &  877.8 & $ 874.3- 885.5$ & $ 165.0\\pm  25.1$ & $ 5.6\\pm 0.3$ \\\\ 20092-01-05-01 & 1997/09/06-14:14:55 & 1504.0 & 1380.3 &  886.2 & $ 883.1- 892.2$ & $ 157.0\\pm  27.6$ & $ 5.2\\pm 0.3$ \\\\ 20092-01-05-01 & 1997/09/06-15:52:47 &  864.0 & 1360.4 &  890.9 & $ 885.5- 896.6$ & $ 104.4\\pm  22.7$ & $ 5.5\\pm 0.4$ \\\\ 20098-03-07-00 & 1997/02/27-05:54:24 & 1040.0 & 1098.4 &  773.4 & $ 761.9- 787.5$ & $ 150.6\\pm  15.8$ & $ 7.8\\pm 0.3$ \\\\ 20098-03-07-00 & 1997/02/27-06:54:23 & 3152.0 & 1111.8 &  801.9 & $ 766.7- 832.6$ & $ 129.3\\pm   7.8$ & $ 8.1\\pm 0.2$ \\\\ 20098-03-08-00 & 1997/03/01-21:34:47 & 2192.0 &  778.4 &  783.0 & $ 749.0- 802.1$ & $ 102.2\\pm  10.8$ & $ 8.9\\pm 0.4$ \\\\ 20098-03-08-00 & 1997/03/01-22:56:32 & 3088.0 &  765.0 &  797.9 & $ 769.5- 821.5$ & $ 111.2\\pm   9.9$ & $ 9.3\\pm 0.3$ \\\\ 30072-01-01-01 & 1998/03/03-14:01:46 & 1104.0 & 1603.2 &  678.2 & $ 673.0- 684.1$ & $ 108.1\\pm   9.0$ & $ 7.0\\pm 0.2$ \\\\ 30072-01-01-02 & 1998/03/04-14:00:40 & 1232.0 & 1994.6 &  856.3 & $ 849.1- 867.5$ & $ 115.5\\pm  12.4$ & $ 5.8\\pm 0.2$ \\\\ 30072-01-01-03 & 1998/03/05-12:24:38 & 1104.0 & 2212.4 &  838.3 & $ 833.6- 846.1$ & $ 138.2\\pm  11.4$ & $ 5.9\\pm 0.2$ \\\\ 40047-02-05-00 & 1999/05/31-15:30:26 & 2640.0 & 1112.4 &  852.7 & $ 843.9- 857.8$ & $ 131.9\\pm  13.1$ & $ 6.2\\pm 0.2$ \\\\ 40047-03-02-00 & 1999/06/03-15:33:09 & 2640.0 &  975.0 &  861.1 & $ 852.0- 874.7$ & $ 158.4\\pm  16.2$ & $ 6.8\\pm 0.2$ \\\\ 40047-03-03-00 & 1999/06/04-13:48:20 & 3152.0 & 1035.7 &  845.2 & $ 834.9- 859.9$ & $ 141.0\\pm  11.2$ & $ 7.4\\pm 0.2$ \\\\ 40047-03-03-00 & 1999/06/04-15:23:28 & 3024.0 & 1032.2 &  862.8 & $ 851.6- 873.6$ & $ 141.4\\pm  12.2$ & $ 7.3\\pm 0.2$ \\\\ 50049-02-13-00 & 2000/11/07-07:01:35 & 1952.0 & 1750.6 &  833.3 & $ 823.8- 839.5$ & $ 141.0\\pm  19.3$ & $ 4.4\\pm 0.2$ \\\\ 50049-02-15-03 & 2000/11/13-20:57:38 & 2064.0 & 1164.4 &  870.3 & $ 867.3- 874.9$ & $ 136.8\\pm  16.1$ & $ 6.2\\pm 0.3$ \\\\ 50049-02-15-03 & 2000/11/13-22:14:40 & 3024.0 & 1121.6 &  834.9 & $ 822.1- 848.5$ & $ 158.8\\pm  11.7$ & $ 7.1\\pm 0.2$ \\\\ 50049-02-15-04 & 2000/11/14-01:42:39 & 1744.0 &  836.5 &  811.7 & $ 808.3- 815.9$ & $ 137.3\\pm  17.4$ & $ 7.3\\pm 0.3$ \\\\ 50049-02-15-04 & 2000/11/14-03:02:23 & 2704.0 &  785.3 &  754.3 & $ 734.8- 767.1$ & $ 132.8\\pm  13.2$ & $ 7.4\\pm 0.3$ \\\\ 50049-02-15-05 & 2000/11/15-04:30:45 & 3152.0 &  841.3 &  631.3 & $ 613.7- 642.9$ & $  84.9\\pm   8.8$ & $ 7.6\\pm 0.3$ \\\\ 50049-02-15-05 & 2000/11/15-06:06:40 & 2784.0 &  825.6 &  643.8 & $ 627.7- 656.8$ & $  82.9\\pm   9.4$ & $ 7.4\\pm 0.3$ \\\\ 50049-02-15-06 & 2000/11/16-04:47:43 & 1808.0 &  745.0 &  717.6 & $ 712.0- 723.6$ & $ 117.0\\pm  11.8$ & $ 8.7\\pm 0.3$ \\\\ 50049-02-15-07 & 2000/11/16-06:33:43 & 1104.0 &  564.7 &  698.6 & $ 686.0- 709.1$ & $ 131.5\\pm  26.9$ & $ 8.1\\pm 0.6$ \\\\ 60054-02-03-03 & 2001/07/10-18:59:27 & 1552.0 &  452.2 &  770.4 & $ 755.0- 787.2$ & $ 129.9\\pm  19.0$ & $11.4\\pm 0.6$ \\\\ 60054-02-03-05 & 2001/07/12-17:15:12 & 1232.0 &  579.7 &  801.9 & $ 798.8- 813.3$ & $ 177.3\\pm  28.5$ & $ 8.9\\pm 0.5$ \\\\ 60429-01-05-00 & 2002/02/18-22:14:08 & 1824.0 &  558.0 &  787.9 & $ 769.7- 819.6$ & $ 111.3\\pm  16.1$ & $ 9.4\\pm 0.5$ \\\\ 60429-01-09-00 & 2002/02/27-21:42:07 & 1472.0 &  775.3 &  843.1 & $ 835.6- 850.2$ & $ 121.8\\pm  20.7$ & $ 6.9\\pm 0.4$ \\\\ 70069-03-01-01 & 2002/03/07-10:26:39 & 2192.0 &  945.8 &  807.0 & $ 786.1- 828.4$ & $ 183.5\\pm  16.0$ & $ 7.1\\pm 0.2$ \\\\ 70069-03-01-02 & 2002/03/07-15:17:35 & 1296.0 &  941.0 &  781.8 & $ 771.6- 797.2$ & $ 153.0\\pm  18.4$ & $ 7.2\\pm 0.3$ \\\\ 70069-03-02-00 & 2002/03/11-22:17:20 & 1648.0 &  815.4 &  627.7 & $ 607.4- 637.3$ & $  83.0\\pm  14.1$ & $ 7.1\\pm 0.4$ \\\\ 70069-03-02-01 & 2002/03/10-00:54:39 &  656.0 & 1593.4 &  768.5 & $ 763.4- 775.9$ & $ 169.8\\pm  17.3$ & $ 7.3\\pm 0.3$ \\\\ 70069-03-03-03 & 2002/03/17-20:48:59 & 1120.0 &  859.5 &  876.6 & $ 870.2- 886.3$ & $  95.6\\pm  18.5$ & $ 7.2\\pm 0.5$ \\\\ 70069-03-03-06 & 2002/03/18-12:45:43 &  656.0 &  819.6 &  821.1 & $ 813.4- 832.4$ & $ 103.1\\pm  19.3$ & $ 8.3\\pm 0.6$ \\\\ 70069-03-03-07 & 2002/03/18-20:34:39 & 2336.0 &  797.1 &  827.0 & $ 816.6- 844.3$ & $ 112.6\\pm  10.4$ & $ 8.5\\pm 0.3$ \\\\ 70069-03-03-09 & 2002/03/19-14:07:43 &  976.0 &  882.4 &  700.9 & $ 692.0- 711.7$ & $ 112.7\\pm  15.0$ & $ 8.8\\pm 0.4$ \\\\ 70069-03-03-14 & 2002/03/21-20:03:39 & 1952.0 &  600.5 &  764.7 & $ 756.6- 771.8$ & $ 119.9\\pm  15.2$ & $ 9.4\\pm 0.4$ \\\\ \\hline \\end{tabular} \\label{dbarret_tab1} \\end{table*} ",
        "conclusions": "We have shown that the properties of the QPOs detected from Aql X-1 so far are consistent with those seen from similar systems, in particular the quality factor of the lower QPOs and the frequency separation between the lower and upper peaks, which we have measured for the first time to be close to half the spin frequency of the neutron star. It would be worth following up the lower QPOs closer to the saturation frequency (at 900 Hz), with adequate sensitivity to estimate the quality factor of the QPOs, and determine whether it drops as in other sources. This may become possible with RXTE during the next outburst of this very active transient."
    },
    "0711/0711.1050_arXiv.txt": {
        "abstract": "{It is usual in helioseismology to remove unwanted instrumental low-frequency trends by applying high-pass filters to the time series. However, the choice of the filter is very important because it can keep the periodic signals throughout the spectrum. At the same time, these filters should not introduce any spurious effects on the remaining signal, which would modify the periodic signatures in the Fourier domain.} {One of the most used filters is the so-called backwards difference filter that can be applied when the time series are regularly sampled.  The problem of this filter is that the amplitudes and the phases of the periodic signals in the Fourier domain are modified and, therefore, their extracted amplitudes are biased when the frequency of the periodic signals decreases. The objective of this research note is to give a correction that could be applied to the power spectrum to over come this problem.} {We properly derive the transfer function of the backwards difference filter and we show how this can be applied to the power spectrum to correct the amplitudes. For the amplitude spectrum and the time series, we also derive a correction for the phase.} {The amplitudes of the periodic signals in the resultant power spectrum density are corrected from the cut-off frequency down to zero.  } {} ",
        "introduction": "One of the most common problems to be faced when analyzing oscillatory components of an observed signal is the presence of low-frequency drifts which can be of instrumental or natural origin. These low-frequency trends introduce a background on the Fourier domain which reduces the signal-to-noise ratio (S/N) of searched periodic signals. To avoid this, high-pass filters are commonly used. The ideal filter is one that reduces the effect of the trend while preserving the periodic signals even at low frequencies. Moreover, it is important to keep all the parameters of periodic signals untouched, including their amplitudes.  In this research note we study a very easy-to-use filter: the backwards difference filter which is commonly used in helioseismology \\citep[see for example][for a complete review on this research topic]{JCD2002} to reduce the low-frequency unwanted stochastic contribution of regularly sampled observations. We also deduce the transfer function of this filter. This transfer function allows us to correct the filtered power spectrum in order to retrieve the correct amplitudes of the periodic signals in the spectrum. Finally, we determine a correction for the resultant phase for when we are interested in the amplitude spectrum or in the time series. ",
        "conclusions": "In this research note we have shown the advantage of removing low-frequency trends by applying the backwards difference filter. The background in the PSD produced by this kind of trend is reduced and the S/N ratio of the periodic signals are improved. Moreover, the transfer function of this filter, $Q(\\omega)$, has been deduced in order to correct -- by dividing the spectrum by this quantity -- the power of the periodic signals in the PSD. The resulting amplitudes inferred from the PSD are the original ones.  Filtering with a smoothing is often used, because it is a powerful detrending tool. However, it must be carefully applied and the choice of the boxcar size needs a preliminary study, according to frequencies of searched signals. The appearance of the wriggles must be taken into account, ideally by correcting the spectrum with the transfer function of this filter.  Even though this presentation does not deal with effects of stochastic noise, it emphasizes the advantages of the BDf and explains the manner to use properly this filter in a manner which is very easy to implement."
    },
    "0711/0711.1861_arXiv.txt": {
        "abstract": "We report the results of our HST snapshot survey with the ACS HRC PR200L prism, designed to measure the near-UV emission in a sample of nearby M dwarfs\\footnote{Based on observations made with the NASA/ESA Hubble Space Telescope, obtained at the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS 5-26555. These observations are associated with program 10525.}. 33 stars were observed, spanning the mass range from 0.1 - 0.6 solar masses (T$_{eff}\\sim$ 2200K - 4000K) where the UV energy distributions vary widely between active and inactive stars. These observations provide much-needed constraints on models of the habitability zone and the atmospheres of possible terrestrial planets orbiting M dwarf hosts, and will be useful in refining the target selection for future space missions such as TPF. We compare our data with a new generation of M dwarf atmospheric models and discuss their implication for the chromospheric energy budget. These NUV data will also be valuable in conjunction with existing optical, FUV and X-ray data to explore unanswered questions regarding the dynamo generation and magnetic heating in low-mass stars. ",
        "introduction": "M dwarfs are by far the most numerous of nearby stars, comprising roughly 75\\% of the Galactic stellar population. The sheer number of M dwarfs in our local solar neighborhood makes them attractive targets for planet-finding missions, while their low masses make them a likely place for discovering Earth-mass worlds in radial velocity surveys. However, these cool neighbors also pose a unique set of problems for their resident planets, and so have been largely dismissed in years past in favor of more tractable cases.  Recently, however, there has been renewed interest in M dwarfs as hosts for habitable worlds. The case for M dwarf planets has been outlined in \\citet{t07}, bringing together studies by \\citet{jo97}, \\citet{jo03}, \\citet{seg05} and others to propose that perhaps the case for habitability is not so dire after all. As further timely motivation, \\citet{udry} recently announced the discovery of the lowest mass planet to date (M$_p \\sim$ 5M$_{\\earth}$), orbiting in the habitable zone of a M3.5 dwarf, Gl 581.  The near-UV emission of M dwarfs remains an outstanding unknown, affecting both our understanding of the basic atmospheric structure in these stars and the potential habitability of attendant planets. Due to their cool effective temperatures, M dwarfs exhibit almost no continuum emission in the near-UV. However, a large fraction of these stars possess surface magnetic fields, which give rise to a steady-state, hot outer atmosphere and transient events such as starspots and flares. Non-radiative heating processes transfer energy from the magnetic field into the stellar atmosphere to power these phenomena, producing significant UV output in the form of chromospheric (T$\\sim 10^4$K), transition region (T$\\sim 10^5$K) and coronal (T$\\sim 10^6$K) emission. Chromospheric near-UV emission can have critical consequences for planetary atmospheres, as it directly impacts the habitability of a neighboring planet through its role in the production of ozone, as well as its implications for long-term survival (and thus detectability) of biomarker gases (such as N$_2$O and CH$_4$). Preliminary studies of M dwarf habitability zones indicate that the typical chromospheric UV flux from an active star will produce observable ozone layers, comparable to that of Earth, on terrestrial planets around M dwarfs (see \\citet{seg05}, Figure 3). However, models of planetary atmospheres around relatively inactive M stars are presently based on stellar model photosphere predictions due to the lack of near-UV observational data.  The data described in this paper also provide crucial constraints for new model atmospheres, which may be used to address a number of outstanding problems in stellar astrophysics. For example, optical (H$\\alpha$) observations of nearby stars (\\citet{ja74}; \\citet{rhg95}; \\citet{hgr96}; \\citet{gi00};\\citet{we04}) show that the active fraction (as measured by H$\\alpha$ emission with equivalent width $>$ 1$\\mbox{\\AA}$) increases dramatically from types M0 to M6, peaks near spectral type M7-M8 and declines thereafter. The strength of activity, as measured by the ratio of the luminosity in H$\\alpha$ compared to the bolometric luminosity ($L_{H\\alpha}/L_{bol}$), also  declines at late types (Burgasser et al. 2002; Cruz \\& Reid 2002; West et al. 2004). In addition, it has been suggested that activity in ultracool dwarfs is confined primarily to large flares, with little or no ``quiescent'' emission (\\citet{ru00}; \\citet{li95}; \\citet{fgs00}). In contrast, STIS/MAMA observations of active M7-M9 dwarfs by \\citet{hjk03} showed that quiescent transition region (C IV) emission is present at levels comparable to those seen in earlier active M dwarfs.  The difficulty in interpreting these intriguing empirical data lies primarily in our lack of a theory of magnetic dynamo field production and subsequent atmospheric heating. At types later than about M3, stars  become fully convective and a solar-type shell dynamo can no longer operate. The theory of $\\alpha^2$ or turbulent dynamos is not yet developed enough to provide predictions or guidance regarding the efficacy  of magnetic field production.  Progress may be made, however, on the atmospheric heating problem. It has not been established whether the drop in active fraction and activity level seen at types later than M7 is a result of the particular diagnostic we are using (H$\\alpha$). If magnetic heating is still present, but the atmospheric structure has changed in these late type stars, emission may occur preferentially in other lines. In particular, the relatively unobserved 1700-3000$\\mbox{\\AA}$ wavelength region contains numerous strong emission lines important to the energy balance of the outer atmosphere, including the resonance lines of Mg II and a host of Fe II emission lines. These lines provide additional information on the atmospheric structure at chromospheric and transition region temperatures, and therefore to the underlying heating processes. ",
        "conclusions": "We obtained 31 ACS PR200L prism spectra of stars, covering spectral types K7 to M5, and including a variety of activity levels. The near-UV data presented here provide valuable constraints for both model stellar chromospheres and photochemical models of planetary atmospheres around M dwarfs.  We measured Mg II fluxes for our spectra and compared our results with complementary optical, near-UV and X-ray data. For those stars in our sample with IUE data, we used the IUE Mg II flux as a check of our ACS prism spectra measurements, and found them roughly consistent. We also examined the relationship between Mg II emission and ROSAT measurements of the X-ray flux. Our results agree with those of \\citet{md89}, showing a well-defined correlation for the most active dMe stars. Using existing H$\\alpha$ measurements for our sample in conjunction with the $\\chi$ factor of \\citet{wa04}, we computed H$\\alpha$ fluxes and compared these to our measured Mg II fluxes, finding a complex relationship between the optical and near-UV emission. Together, the Mg II, X-ray and H$\\alpha$ data discussed here present a complex picture of the chromospheric and coronal emission in these stars.  Using the `RH' NLTE radiative transfer code, we calculated Mg II and Fe II spectra for a suite of model atmospheres. The resulting model spectra were then degraded to the resolution of the ACS prism data using SLIM, the ACS prism spectra simulator. We compared our results with several of the observed spectra to infer potential chromospheric structure and active region coverage, finding that most stars could be matched by an active region area coverage of a few percent of the stellar surface. This result is consistent with our findings in \\citet{yzcmi}, where we determined that $\\sim$0.5\\% of the visible stellar surface was involved in the observed flares.  While in many cases the qualitative shape of the model spectra matched the observations, Fe II was overpredicted for the most active objects. The empirical results we find for Mg II, X-ray and H$\\alpha$ pose an interesting challenge for the models. The model atmospheres described in this paper will be used in future calculations of H$\\alpha$ and Ca II K, which will then be compared with our results for the near-UV emission.  In the case of Gl 876, a dM4 star that hosts three extrasolar planets, we scaled our observed spectra to match the expected flux at the distance of a planet in the habitable zone. We compared our results to a similar scaling of our AD Leo observation, and then discussed these results in the light of previous studies of M dwarf planet habitability. We find that our observation of what might be called an ``inactive'' star (on the basis of H$\\alpha$, the usual proxy for activity in late type stars) actually shows a substantial amount of near-UV flux. Although Gl 876 is a lower luminosity and lower activity star than AD Leo, the non-negligible near-UV flux and proximity of the habitable zone conspire to make it a similar host for planets. The spectra described in this paper will be further used as input stellar flux for the Virtual Planet Laboratory's photochemical planetary atmosphere models."
    },
    "0711/0711.0657_arXiv.txt": {
        "abstract": "{} {We study a sample composed of 28 of the brightest stars in the Arches cluster. Our aim is to constrain their stellar and wind properties and to establish their nature and evolutionary status.} {We analyze K-band spectra obtained with the integral field spectrograph SINFONI on the VLT. Atmosphere models computed with the code CMFGEN are used to derive the effective temperatures, luminosities, stellar abundances, mass loss rates and wind terminal velocities.} {We find that the stars in our sample are either H-rich WN7--9 stars (WN7--9h) or supergiants, two being classified as OIf$^{+}$. All stars are 2--4 Myr old. There is marginal evidence for a younger age among the most massive stars. The WN7--9h stars reach luminosities as large as $2 \\times 10^{6} L_{\\odot}$, consistent with initial masses of $\\sim$ 120 \\msun. They are still quite H-rich, but show both N enhancement and C depletion. They are thus identified as core H-burning objects showing products of the CNO equilibrium at their surface. Their progenitors are most likely supergiants of spectral types earlier than O4--6 and initial masses $>$ 60 \\msun. Their winds follow a well defined modified wind momentum -- luminosity relation (WLR): this is a strong indication that they are radiatively driven. Stellar abundances tend to favor a slightly super solar metallicity, at least for the lightest metals. We note however that the evolutionary models seem to under-predict the degree of N enrichment. } {} ",
        "introduction": "\\label{s_intro}  The center of our Galaxy is a unique environment to study massive stars. It harbors three massive clusters -- the Arches, Quintuplet and central cluster -- which together contain about 30\\% of the number of Wolf-Rayet stars known in the Galaxy \\citep{vdh06}. Interestingly, the three clusters have different ages, ranging from $\\sim$ 2 Myr for the Arches to $\\sim$ 6 Myr for the central cluster. Consequently, they host different populations of massive stars and sample the entire upper part of the HR diagram. Studying their stellar content gives us a unique opportunity to understand how massive stars evolve.  Although there is a global framework explaining the evolution of stars more massive than $\\gtrsim$ 20 \\msun, a quantitative description is still lacking. According to \\citet{paul95}, stars with masses in the range 25--60 \\msun\\ experience the sequence O $\\rightarrow$ Of $\\rightarrow$ LBV or RSG $\\rightarrow$ WN8 $\\rightarrow$ WNE $\\rightarrow$ WC, while for more massive stars, the sequence O $\\rightarrow$ Of $\\rightarrow$ WNL+abs $\\rightarrow$ WN7 ($\\rightarrow$ WNE) $\\rightarrow$ WC is preferred. \\citet{langer94} favor another scenario in which all stars have a H-rich WN phase prior to a LBV event: O $\\rightarrow$ H-rich WN $\\rightarrow$ LBV $\\rightarrow$ H-poor WN $\\rightarrow$ H-free WN $\\rightarrow$ WC.  We see that there are still some qualitative differences between the proposed scenarios. Further, the evolutionary sequences are not mature enough to allow a refinement of the classification of the stars in the different evolutionary states. For example, the spectral types of the O or WC stars entering the above scenarios are not specified. The question of whether or not all massive stars go through a LBV phase is also not answered. This is an important caveat, especially since this phase has recently been argued to be the one in which most of the mass removal happens \\citep{so06}.  In a previous study \\citep{gc07}, we analyzed 18 massive stars in the central cluster of the Galaxy. This cluster is especially intriguing since it hosts the supermassive black hole SgrA* \\citep{genzel96,ghez98}. In spite of the drastic tidal forces, several tens of massive stars formed recently \\citep{ahh90,krabbe95,pgm06}. Some of them are approaching the black hole at distances of only a few light hours \\citep{ghez03,frank05}. The presence of young stars in the Galactic Center together with the apparent implausibility of forming stars so close to the central supermassive black hole is a puzzle usually referred to as ``the paradox of youth''. Studying the dynamics of these young stars, \\citet{pgm06} \\citep[see also][]{lb03,genzel03} have shown that they orbit SgrA* in two counter-rotating disks. Together with other evidences (total mass and structure of the disks), this points to a local, ``in-situ'' star formation event. The detailed analysis of the post-main sequence massive stars has revealed that, surprisingly, their evolution follows almost perfectly the predictions of evolutionary models \\citep{gc07}. This implies that whatever the exact formation mechanism is, the subsequent evolution is not different from that predicted for normal stars. We found that all stars seem to have progenitors in the mass range 25--60 \\msun\\ and that they follow relatively well the evolutionary scenario proposed by \\citet{paul95} for this mass range. We have been able to refine this scenario, pinpointing the relation between different spectral types: O $\\rightarrow$ (Ofpe/WN9 $\\leftrightarrow$ LBV) $\\rightarrow$ WN8 $\\rightarrow$ WN8/WC9 $\\rightarrow$ WC9. This was made possible by the detailed study of stellar abundances in various Wolf-Rayet stars and related objects. Abundance analysis is a powerful tool to constrain stellar evolution since it gives direct access to the evolutionary state of a star.  The above study focused on stars in the mass range 25--60 \\msun\\ due to the age of the central cluster (more massive stars do not exist any more). In order to constrain stellar evolution at very high mass, we need to study younger clusters. The Arches cluster in the Galactic Center, only 30 pc away from the central cluster, is the ideal target. Not only is it believed to be quite young (2 to 4.5 Myr, see Figer et al.\\ 1999, Blum et al. 2001 and the present study), but it also shares the same environment as the central cluster, and hence have the same metallicity. Hence, its study ensures to obtain a homogeneous view of stellar evolution among all types of massive stars in the Galactic Center. The Arches cluster, first discovered by \\citet{nagata95} and \\citet{cotera96}, is also one the most massive and densest cluster of the Galaxy. \\citet{figer99} first showed that the mass function (MF) of its central regions might be shallower than the typical Salpeter IMF. This result was confirmed by \\citet{stolte02}. Although there are indication that high mass star formation might be favored in the Galactic Center \\citep{ms96,klessen07}, recent simulations of the dynamical evolution of the cluster by \\citet{kim06} indicate that we might in fact witness the effects of mass segregation in the cluster core rather than an actual top-heavy initial mass function. Whatever the physical reason, there are nearly 100 massive stars in the Arches cluster. From K-band spectroscopy of the brightest members, \\citet{blum01} and \\citet{figer02} identified several late WN stars and early O supergiants. The analysis of five of these stars by \\citet{paco04} revealed, by an indirect method, that their metallicity was close to solar.  Here, we analyze a much larger sample (28 stars in total) in order to better constrain their stellar and wind parameters. We rely on high quality data obtained with the integral field spectrograph SINFONI on the VLT. The K-band spectra extracted from this data set are analyzed with atmosphere models computed with the code CMFGEN \\citet{hm98}. In Sect.\\ \\ref{s_spec} we describe the observations, our sample and the spectroscopic classification; in Sect.\\ \\ref{s_analysis} our method to analyze the stars are presented; the results are summarized in Sect.\\ \\ref{s_res} and discussed in Sects.\\ \\ref{s_hr}, \\ref{s_nature_wn7} and \\ref{s_Z}. We give our conclusions in Sect.\\ \\ref{s_conc}. ",
        "conclusions": "\\label{s_conc}  We have presented a quantitative study of the most massive stars in the Arches cluster. K--band spectra have been obtained with SINFONI on the VLT. A detailed spectral classification has revealed the presence of WN7--9h stars as well as O supergiants, including two extreme OIf$^{+}$ stars. We have quantified the main stellar and wind parameters of 28 stars using the atmosphere code CMFGEN. The main results of our study are:  \\begin{itemize} \\item the massive star population of the Arches cluster is 2--4 Myr old. Although marginal, there seems to be the trend that the most massive stars are also the youngest: WNLh stars are 2--3 Myr old, while less massive O supergiants might have an age of up to 4 Myr. Initial masses as large as 120 \\msun\\ are derived for the WNLh stars from the HR diagram. \\item The WN7--9h stars are identified as core H burning stars which show chemical enrichment typical of the CNO equilibrium: they still contain a significant amount of hydrogen and show both N enhancement and C depletion. They have supergiants of spectral types earlier than O4-6 as progenitors. \\item The level of N enrichment suggests either a super-solar initial content and/or a too low efficiency of N enrichment in the evolutionary models. The indication that Si might be overabundant by a factor 2 compared to the solar abundance argues in favor of a super-solar metallicity at least for the lightest metals. \\item The properties of the Arches massive stars argue in favor of the evolutionary scenario of \\citet{paul95} for the most massive stars: O $\\rightarrow$ Of $\\rightarrow$ WNL+abs $\\rightarrow$ WN7 \\item The winds of the WN7--9h stars follow a well defined modified wind momentum -- luminosity relation. This is a strong indication that they are radiatively driven. It also seems that the winds are less clumped in the radio continuum emitting region than in the near-IR line emitting region, in agreement with \\citet{puls06}. \\end{itemize}  In order to test the indication that the most massive stars are the last to form, intermediate mass stars must be analyzed. The question of the metallicity in the Galactic Center, and the Arches cluster in particular, is clearly not answered yet. New observations/analysis of cooler stars showing a larger number of metallic lines are needed. We will present such a study in a forthcoming paper."
    },
    "0711/0711.2839_arXiv.txt": {
        "abstract": "We use the ROSAT All-Sky Survey to study the X-ray properties of a sample of 625 groups and clusters of galaxies selected from the Sloan Digital Sky Survey. We stack clusters with similar velocity dispersions and investigate whether their average X-ray luminosities and surface brightness profiles vary with the radio activity level of their central galaxies. We find that at a given value of $\\sigma$, clusters with a central radio AGN have more concentrated X-ray surface brightness profiles, larger central galaxy masses, and higher X-ray luminosities than clusters with radio-quiet central galaxies. The enhancement in X-ray luminosity is more than a factor of two, is detected with better than 6$\\sigma$ significance, and cannot be explained by X-ray emission from the radio AGN itself. This difference is largely due to a subpopulation of radio-quiet, high velocity dispersion clusters with low mass central galaxies. These clusters are underluminous at X-ray wavelengths when compared to otherwise similar clusters where the central galaxy is radio-loud, more massive, or both. ",
        "introduction": "There is increasing evidence that the majority of radio-loud AGN at low redshift may be triggered by the accretion of hot gas. Using a sample of 625 nearby groups and clusters, \\cite{Best07} showed that the galaxies located closest to the centres of the clusters are more likely to host a radio-loud AGN than other galaxies of similar stellar mass.  \\cite{Allen06} analyzed \\emph{Chandra} X-ray images of nine nearby X-ray luminous elliptical galaxies and showed that the jet power is tightly correlated with the Bondi accretion rate onto the central black hole as estimated from the observed gas temperature and density profiles.  It is also now recognized that the `over-cooling problem' at the centre of many galaxy clusters could be solved if radio-loud AGN can heat surrounding gas. Direct evidence of AGN heating came with the discovery of X-ray cavities in the hot intracluster medium of a number of clusters and groups \\citep{Boehringer93,Churazov00,Fabian00,McNamara00,Blanton03,Gitti07}. Some of these cavities coincide with extended lobes of radio emission produced by an AGN in the central cluster galaxy\\footnote{Some cavities do not have associated radio emission above current detection limits [e.g. \\cite{Ettori02}, \\cite{Clarke05} and \\cite{Jetha08}] but these hot bubbles may have been inflated by a previous cycle of nuclear activity.}. The observations suggest that relativistic radio plasma displaces the X-ray emitting gas, generating turbulence and wave activity which heat the intracluster medium (ICM).  AGN feedback may also explain why low-temperature clusters are less X-ray luminous than predicted by a homologous scaling of the properties of hotter and more massive systems \\citep{Nath02,Best07}. The energy available from the central AGN depends primarily on black hole accretion rate. In low-mass groups, it can be comparable to the total gravitational binding energy of the X-ray emitting gas, whereas high-mass clusters will be less strongly affected. As a result, relations such as $L_X-T_X$ (X-ray luminosity vs temperature) and $L_X-\\sigma$ (X-ray luminosity vs velocity dispersion) may be modified by radio-source heating. At smaller scales, AGN feedback may play an important role in the formation and evolution of galaxies. Recent work has demonstrated that it can plausibly explain the exponential cutoff at the bright end of the galaxy luminosity function, as well as the `down-sizing' of galactic star-formation activity at recent cosmic epochs \\citep{Croton06, Bower06, Kang06}.  It may be that hot gas accretion, AGN triggering, reheating of ambient gas, and suppression of AGN activity occur in a cycle. Accretion onto the central black hole may cause an outburst which removes the fuel supply for further radio activity.  Without AGN feedback, the ICM reverts to the state that triggered the outburst \\citep{Churazov05}. Although this picture is attractive, it remains to be verified in detail. We do not know the necessary conditions to trigger or to quench radio activity. We also do not understand the extent to which the global X-ray properties of the gas are modified by AGN feedback\\citep[e.g.][]{Rizza00,Omma04,Nusser08}. Only a few nearby clusters have deep enough X-ray images to reveal low-density cavities and permit the energetics of the gas to be studied in a spatially resolved fashion. If we wish to study how the global state of the X-ray gas is linked to AGN activity in the central galaxy, we are forced to adopt a more statistical approach.  It has long been known that the radio activity of cluster central galaxies is linked with the cooling flow or ``cool core'' phenomenon \\citep{Burns90,Fabian94a}. This is the fact that many but not all massive clusters show a strong central peak in X-ray surface brightness, almost always coincident with the brightest cluster galaxy (BCG), within which the directly inferred cooling time is less than the Hubble time. The fraction of cool-core clusters with a radio-loud BCG is much higher than the fraction in the rest of the population. Furthermore, because the cores add significantly to cluster luminosity and reduce the emission-weighted temperature, cool-core clusters fall systematically to one side of many of the X-ray scaling relations for clusters, e.g.  the $L_X-T_X$ and $L_X-\\sigma$ relations\\citep{Fabian94b, OHara06, Chen07}. Thus, in clusters of given velocity dispersion (or mass), cool cores tend to be associated with enhanced X-ray luminosity, with massive central BCGs, and with radio-loud BCGs. Given that the probability of radio activity increases with BCG mass\\citep{Best05}, that both BCG mass and cluster X-ray luminosity increase with cluster mass\\citep{Edge91}, and that central black hole mass correlates tightly with galaxy mass\\citep{Tremaine02}, it seems clear that radio activity is tightly related to both black hole mass and the presence of a cooling hot atmosphere which can provide fuel. The exact causal relation between these phenomena remains unclear, however, and recent observational studies of lower mass systems appear, at least superficially, in conflict with the trends established for relatively massive clusters [e.g. compare \\cite{Croston05} and \\cite{Jetha07} with the above references].  In this paper, we provide improved statistics for the relationship between cluster properties, both X-ray and optical, and the radio activity of their BCGs. We use a low redshift ($z<0.1$) sample of 625 groups and clusters with carefully controlled BCG identifications. These were selected by \\cite{Linden07} from the Sloan Digital Sky Survey\\citep[SDSS,][]{York00}. Radio properties of these BCGs were determined following \\cite{Best05}, who cross-matched the galaxies with the National Radio Astronomy Observatory Very Large Array Sky Survey \\citep[NVSS,][]{Condon98} and the Faint Images of the Radio Sky at Twenty Centimeters Survey \\citep[FIRST,][]{Becker95}. By combining these two radio surveys, the identification of radio galaxies is both reasonably complete ($\\sim95$ per cent) and highly reliable ($\\sim99$ per cent). Radio-loud AGN are then separated from star-forming galaxies with detectable radio emission on the basis of 4000\\AA\\ break strength\\citep{Best05}.  The sky coverage of this cluster sample is $\\sim4000$ square degrees.  The only X-ray survey with enough sky coverage to provide a reasonable match is the ROSAT All-Sky survey(RASS). Because the RASS is quite shallow, relatively few individual groups and clusters have sufficient X-ray flux for unambiguous detection, particularly at low velocity dispersion. The mean flux of such objects {\\it can} be detected, however, by stacking their X-ray images \\citep[e.g.][]{Bartelmann03, Shen06, Dai07, Rykoff08}. Such stacking avoids selection biases that can occur if one analyses samples in which only the most X-ray luminous systems are detected. The limited resolution of RASS maps and our stacking strategy, make it difficult to study the structure of the ICM in detail, but we will see that concentration differences, reflecting the presence of cool cores, are nevertheless detectable.  Our paper is organized as follows. In Section 2, we introduce our sample of groups and clusters and describe their radio properties. In Section 3, we describe our X-ray detection techniques, both our method for detecting individual groups and clusters and our stacking technique. In section 4, we study and compare the X-ray properties of clusters, emphasising how the $L_X-\\sigma$ relation and the surface brightness profiles of clusters vary with the radio properties of their BCGs. We discuss the contribution of the radio AGN to the total X-ray emission in Section 5.1, while in Section 5.2, we discuss how the $L_X-\\sigma$ relation of clusters depends on the stellar properties of their BCGs. We present our conclusions in Section 6. ",
        "conclusions": "We have used the ROSAT All Sky Survey to study the X-ray properties of a sample of 625 groups and clusters of galaxies selected from the Sloan Digital Sky Survey. We focus on the $L_X-\\sigma$ relation of these groups and clusters, and we study whether this relation depends on the radio properties of the central galaxy (BCG). A cluster is termed `radio-loud' if its central BCG is a radio-loud AGN, and `radio-quiet' if the BCG is not detected at radio wavelengths. We find that the fraction of clusters with individual X-ray detections depends strongly on whether the BCG is radio-loud. Radio-loud clusters are detected more frequently than radio-quiet clusters of the same velocity dispersion and redshift.  The $L_X-\\sigma$ relations for individually detected radio-loud and radio-quiet clusters are very similar, but this is purely a selection bias. The majority of detected clusters are just above the X-ray detection limit in both samples. Since the redshift distribution at each velocity dispersion is similar for the two types of cluster, the mean relations for {\\it detected} objects are forced to be similar.  By stacking the X-ray images of clusters with similar velocity dispersion, we studied the {\\em average} X-ray luminosities and surface brightness profiles of our clusters as function of velocity dispersion. The average X-ray luminosities of radio-loud clusters are systematically higher and their luminosity profiles are more concentrated than those of radio-quiet systems. X-ray emission from the radio AGN itself is by far insufficient to explain this doubling of the X-ray luminosity.  Our results demonstrate convincingly and quantitatively that the X-ray properties of the intracluster gas correlate with the presence of a central radio source in the way first suggested by Burns (1990).  The stellar masses of the BCGs also correlate with their radio properties; radio-loud clusters tend to mave more massive BCGs than radio-quiet clusters of the same velocity dispersion. Those clusters of given velocity dispersion which have unusually low mass BCGs tend to be both underluminous in X-rays and radio-quiet. This effect is particularly pronounced for large velocity dispersion clusters.  Among radio-loud clusters we find no dependence of X-ray luminosity on BCG stellar mass. These results can be summarised by saying that the only clear dependence of X-ray properties on BCG properties is that high velocity dispersion clusters in which the BCG is both low mass and radio-quiet tend to be several times less X-ray luminous than otherwise similar clusters in which the BCG is massive and/or radio-loud.  Clearly a high central hot gas density is needed for effective fuelling of the radio source. If the radio activity is in turn able to heat the surrounding gas and cause its re-expansion, the feedback cycle needed to control the growth of the central galaxy could be established. The conditions which lead to effective AGN fuelling are clearly related to the presence of a massive central galaxy, although the exact causal relationship between BCG growth, AGN activity and ICM structure remains to be clarified."
    },
    "0711/0711.3524_arXiv.txt": {
        "abstract": "Recent multi-wavelength observations of \\object{3C454.3}, in particular during its giant outburst in 2005, put severe constraints on the location of the 'blazar zone', its dissipative nature, and high energy radiation mechanisms.  As the optical, X-ray, and millimeter light-curves indicate, significant fraction of the jet energy must be released in the vicinity of the millimeter-photosphere, i.e. at distances where, due to the lateral expansion, the jet becomes transparent at millimeter wavelengths.  We conclude that this region is located at $\\sim 10$ parsecs, the distance coinciding with the location of the hot dust region.  This location is consistent with the high amplitude variations observed on $\\sim 10$ day time scale, provided the Lorentz factor of a jet is $\\Gamma_j \\sim 20$. We argue that dissipation is driven by reconfinement shock and demonstrate that X-rays and $\\gamma$-rays are likely to be produced via inverse Compton scattering of near/mid IR photons emitted by the hot dust.  We also infer that the largest gamma-to-synchrotron luminosity ratio ever recorded in this object -- having taken place during its lowest luminosity states -- can be simply due to weaker magnetic fields carried by a less powerful jet. ",
        "introduction": "Multi-wavelength coverage of recent activity of quasar \\object{3C454.3} provided exceptional data to investigate the structure and physics of its blazar zone. Prior to year 2000, this object spent most of its time in the low, relatively quiescent state. Starting in 2000, \\object{3C454.3} entered a highly active state, changing optical flux by a factor tens on time scales of a few months \\citep{Fuhrmann06,Villata06}.  The most powerful event took place in the middle of 2005.  This event was monitored also in the X-ray bands (Swift/XRT/BAT: \\citealt{Giommi06}; INTEGRAL: \\citealt{Pian06}; Chandra: \\citealt{Villata06}), and at millimeter wavelengths \\citep{Krichbaum07}.  These data allow a construction of quasi-simultaneous broadband spectrum around the outburst peak.  As is the case for other blazars, the spectrum is composed of two humps, the lower energy one produced via synchrotron mechanism and peaking in the far-infrared band, and the higher energy one most likely generated by inverse-Compton process and peaking in the $\\gamma$-ray band.  The lack of coverage of the event by $\\gamma$-ray observatories does not allow us to determine the luminosity of the high energy component.  Nevertheless, X-ray data suggest that luminosity ratio of the high- to the low-energy components was much smaller during the outburst than during low states monitored in $\\gamma$-rays by CGRO \\citep{Mukherjee97,Hartman99,Zhang05}.  This difference was theoretically investigated by \\citet{Pian06} and by \\citet{Katarzynski07}.  \\citet{Pian06} suggested that during the low states the blazar zone is located inside the broad line region (BLR) and that high energy spectra are produced by the External Radiation Compton (ERC) process involving scattering of broad line photons \\citep[via scenario described in][]{Sikora94}, while during the 2005 outburst the dissipation zone moved outside the BLR where the ERC becomes inefficient.  In such a model, production of the optical outburst doesn't require increase of a jet power.  Similarly, in the scenario proposed by \\citet{Katarzynski07} the jet power is constant, but the drop of luminosity of the high energy component is explained by decrease of the Lorentz factor.  The idea of the constant jet power might be challenged by the most recent optical outburst which in July 2007 was also detected in $\\gamma$-rays by AGILE \\citep{Vercellone07}.  Bolometric luminosity of this outburst was 4-5 times larger than bolometric luminosity during the low optical states and the radiative output was strongly dominated by the $\\gamma$-ray flux.  The currently available millimeter-band light curves \\citep{Krichbaum07} do not indicate any significant delay of the millimeter flux after the bolometric flux as inferred from the infrared and optical data presented in \\citet{Bach07}.  All of the above motivated us to investigate a different scenario, with the origin of the high energy peak involving ERC with IR seed photons and operating in the vicinity of the millimeter photosphere of the source. Basic assumptions of the scenario are described in \\S\\ref{sec:assum}; results of modeling of the broadband spectrum of the 2005 outburst are presented in \\S\\ref{sec:outburst}; explanation of the large $\\gamma$-ray dominance in the low optical states is provided in \\S\\ref{sec:model}; and the main results are summarized in \\S\\ref{sec:discussion}. ",
        "conclusions": "} We demonstrated in this paper that broadband spectra of \\object{3C454.3} can be reconstructed assuming that they are produced at distances $r \\sim 3-9$ parsecs.  By the end of this distance range the jet becomes transparent at millimeter wavelengths.  Blazar activity historically has been defined via observations in the IR/optical bands, while ``blazar-zone'' is often considered to be located deeply within the millimeter photosphere.  However, the optical and millimeter light-curves seem to indicate a significant overlap of the blazar-zone with a region where the jet becomes transparent at millimeter wavelengths (see \\S\\ref{sec:location}). This is further supported by very large millimeter luminosities which require high, {\\it in situ}, dissipation rate of energy, and is consistent with time scales of the fastest high amplitude variations, of the order of $10$ days in both spectral bands. Furthermore, at such distances the co-spatial model self-consistently incorporates production of X- and $\\gamma$-rays, via scatterings of near/mid IR photons emitted by hot dust.  It should be emphasized here that the input-parameter set for ERC models is not unique and that high energy spectra can be reproduced also by scattering of broad emission photons if taking place in the sub-parsec region.  However, then the high energy non-thermal radiation should be accompanied by bulk-Compton features \\citep{Sikora00,Moderski04,Celotti07}, which so far have not been observationally confirmed.  Their lack or weakness can be explained by assuming that in the sub-parsec region jet is still in acceleration phase and the blazar zone is located at larger distances \\citep{Kataoka07}.  We identify the ``blazar zone'' with a reconfinement shock.  That, together with optical polarization data imply domination of the toroidal magnetic field over chaotic/turbulent magnetic fields. However it should be noted that domination of the toroidal component doesn't necessary indicate the domination of the Poynting flux over the matter energy flux. It is very likely that the conversion of the Poynting flux dominated jet into matter dominated jet -- and hence the jet acceleration process -- are accomplished on sub-parsec scales \\citep{Sikora05,Komissarov07}.  Similar conclusions are reached by \\citet{Jorstad07}, following multi-waveband polarimetric observations of 15 AGN.  During its 2005 outburst, \\object{3C454.3} was the most luminous object ever recorded in the optical band.  To explain such an outburst, the jet power larger than $7 \\times 10^{47}\\,$erg s$^{-1}$ is required (see Table~\\ref{tab:param}).  Is it feasible?  Noting that the estimates of the black hole mass in this object give $\\sim 4 \\times 10^9 M_{\\odot}$ \\citep{Gu01}, we infer that the jet power is on the order of the Eddington luminosity.  This, however, is at least by a factor of few larger than the accretion luminosity, which in turn, as determined from the optical luminosity of the thermal component detected during the low state \\citep{Smith88}, and after application of the bolometric correction, is likely to be of the order $10^{47}\\,$erg s$^{-1}$.  \\object{3C454.3} is in this respect not exceptional among most powerful radio-loud quasars: powers of jets larger than $10^{47}\\,$erg s$^{-1}$ have been inferred for several other quasars from analysis of the lobe energetics \\citep{Rawlings91}, as well as from Chandra and HST observations of gamma-ray blazars \\citep{Tavecchio07}."
    },
    "0711/0711.4903_arXiv.txt": {
        "abstract": "In this paper, we first summarize the results of a large-scale double-blind tests campaign carried out for the realistic estimation of the Gaia potential in detecting and measuring planetary systems. Then, we put the identified capabilities in context by highlighting the unique contribution that the Gaia exoplanet discoveries will be able to bring to the science of extrasolar planets during the next decade. ",
        "introduction": "Despite a few important successes (e.g., Bean et al. 2007, and references therein), astrometric measurements with mas precision have so far proved of limited utility when employed as either a follow-up tool or to independently search for planetary mass companions orbiting nearby stars (see for example Sozzetti 2005, and references therein).  In several past exploratory works (Casertano et al. 1996; Lattanzi et al. 1997, 2000; Sozzetti et al 2001, 2003), we have shown in some detail what space-borne astrometric observatories with $\\mu$as-level precision, such as Gaia (Perryman et al. 2001), can achieve in terms of search, detection and measurement of extrasolar planets of mass ranging from Jupiter-like to Earth-like. In those studies we adopted a qualitatively correct description of the measurements that each mission will carry out, and we estimated detection probabilities and orbital parameters using realistic, non-linear least squares fits to those measurements.  Those exploratory studies, however, need updating and improvements. In the specific case of planet detection and measurement with Gaia, we have thus far largely neglected the difficult problem of selecting adequate starting values for the non-linear fits, using perturbed starting values instead.  The study of multiple-planet systems, and in particular the determination of whether the planets are coplanar---within suitable tolerances---is incomplete.  The characteristics of Gaia have changed, in some ways substantially, since our last work on the subject (Sozzetti et al 2003). Last but not least, in order to render the analysis truly independent from the simulations, these studies should be carried out in double-blind mode.  We present here a substantial program of double-blind tests for planet detection with Gaia (preliminary findings were recently presented by Lattanzi et al. (2005)), with the three-fold goal of obtaining: a) an improved, more realistic assessment of the detectability and measurability of single and multiple planets under a variety of conditions, parametrized by the sensitivity of Gaia; b) an assessment of the impact of Gaia in critical areas of planet research, in dependence on its expected capabilities; and c) the establishment of several Centers with a high level of readiness for the analysis of Gaia observations relevant to the study of exoplanets. ",
        "conclusions": ""
    },
    "0711/0711.1521_arXiv.txt": {
        "abstract": "We present new results on the nature of the Herbig--Haro-like object 73 (HHL~73, also known as [G84b]~11) based on narrow-band CCD H$\\alpha$ and \\sii\\ images of the HHL~73 field, and Integral Field Spectroscopy and radio continuum  observations at 3.6 cm covering the emission of the  HHL~73 object. The CCD images allow us to resolve the HHL~73 comet-shaped  morphology into two components and a collimated emission feature of $\\sim4$-arcsec long, reminiscent of a microjet. The IFS spectra of HHL~73 showed emission lines characteristic of the spectra of Herbig--Haro objects. The kinematics derived for HHL~73 are complex. The profiles of the \\sii\\ $\\lambda\\lambda$6717, 6731 \\AA\\ lines were well fitted with a model of three Gaussian velocity components peaking at $\\mbox{\\vlsr}\\simeq-100$, $-20$ and +35 \\kms. We found differences among the spatial distribution of the kinematic components that are compatible with the emission from a bipolar outflow with two blueshifted (low- and high-velocity) components. Extended radio continuum emission at 3.6 cm was detected showing a distribution in close agreement with the HHL~73 redshifted gas. From the results discussed here, we propose HHL~73 to be a true HH object. \\iras~21432+4719, offset 30 arcsec northeast from the HHL~73 apex, is the most plausible candidate to be driving HHL~73, although the evidence is not conclusive. ",
        "introduction": "Herbig--Haro-like objects (HHLs) appear as bright nebulosities in the red optical images. Spectroscopy of HHLs indicates that a fraction of them are reflection nebulae originated by the scattered light from a stellar source associated with the nebulosity, while only a fraction of HHLs are Herbig--Haro (HH) objects arising from shock-excited gas. HHL~73, also known as [G84b]~11, was first reported by \\citet{Gyu84} as object 11 found in his survey of \\textit{DSS} red plates. In these plates, the object shows a comet-shaped morphology, being located close to the western border of a region with high visual extinction. In fact, HHL~73 is found associated with a  dark cloud complex in Cygnus, close to the Lynds Dark Cloud 1035A (\\textit{LDN} 1035A), and lying in a region with clear signs of recent star formation activity. The estimated distance to the region is $900\\pm100$~pc \\citep{Eli78}.  The source \\iras~21432+4719 is located  close to HHL~73. Both sources are engulfed within a clump of high-density molecular gas, traced by the emission of the NH$_3$ and CS molecules, whose emission peaks around the HHL~73 and \\iras~21432+4719 positions (see \\citealt{Ang97} and references therein for a wider description on the high-density molecular gas emissions in the HHL~73 field). Emission from high-velocity molecular gas in this field is also reported  by \\citet{Dob93}, who detect a highly asymmetric CO bipolar outflow. The emission of the redshifted outflow lobe is spatially more extended than the blueshifted lobe and encompasses the region around \\iras\\ and HHL~73. In addition, the detection of a water maser near HHL~73 was reported by \\citet{Gyu87}.  \\citet{Dev97} imaged the HHL~73 field through narrow-band H$\\alpha$ and \\sii\\ $\\lambda\\lambda$6717, 6731 \\AA\\ filters. They find two stellar-like condensations $\\sim$ 2.5 arcmin west of HHL~73 with similar H$\\alpha$ and \\sii\\ surface brightness, surrounded by diffuse emission. They designate their finding as the new Herbig--Haro object HH 379. In addition, they suggest that the exciting source of HH~379 may be embedded within the cometary nebula which appears in their images at a position  close to \\iras~21432+4719 (thus, at the position given by \\citet{Gyu84} for HHL~73). These authors also discuss a plausible association between the \\iras\\ source  and the undetected  exciting source of HH~379. It should be noted that  \\citet{Dev97} never refer to the optical cometary nebula as the HHL~73 object reported by \\citet{Gyu84}. Later on, \\citet{Dav01, Dav03} performed \\nir\\ echelle spectroscopy at the HHL~73 position. They detect emission  from H$_2$ $\\lambda$2.122 $\\mu$m and \\fe\\ $\\lambda$1.644 $\\mu$m lines, while no continuum  emission is detected at the HHL~73 position. Both lines show a rather complex profile, the width of the \\fe\\ line being roughly twice that of the H$_2$ line. \\citet{Dav03} suggest that the \\fe\\ emission may trace the higher-velocity component of a jet. In addition, \\citet{Dav01, Dav03} also suggest that the optical nebula  found in the \\citet{Dev97} images harbours the true exciting source of HH~379, thus discarding the possibiility that \\iras~21432+4719 was driving HH~379. This yet undetected exciting source of HH~379 is named as HH 379-IRS by these authors. In order to help with the identification of the sources mentioned here, we show in Fig.\\ \\ref{fig1} a CCD H$\\alpha$ image of the HHL~73 field where all these sources have been marked.  \\begin{figure} \\centering \\includegraphics[width=\\hsize]{Fig1_HHL73_int_g.eps} \\caption{ CCD H$\\alpha$-band image of the HHL~73 field obtained at the Isaac Newton Telescope (INT, ORM). Arrows point to the HHL~73 and HH~379 objects. The position of \\iras~21432+4719 and the error ellipse are marked. The tilted cross marks the position given by \\citet{Gyu87} for the H$_2$O maser. The image FOV is $\\sim8.5\\times6$ arcmin$^2$. North is up and east is to the left. \\label{fig1}} \\end{figure}  With the aim of unravelling the nature of HHL~73, we included this target within a programme of Integral Field Spectroscopy (IFS) of Herbig--Haro objects using the  Potsdam Multi-Aperture Spectrophotometer (PMAS) in the wide-field IFU mode PPAK. Additional CCD narrow-band images and radio continuum 3.6 cm data have been analysed to complement the study. We present in this paper the results on the kinematics and physical conditions of HHL~73 obtained from this IFS pilot  programme.   The paper is organized as follows. The observations and data reduction of CCD imaging, radio continuum and IFS observations  are described in \\S\\ 2.1, 2.2 and 2.3 respectively. Results are given in \\S\\ 3: CCD imaging in \\S\\ 3.1, radio continuum in \\S\\ 3.2 and IFS in \\S\\ 3.3, including results from integrated emission (\\S\\ 3.4) and a line-profile decomposition model for \\sii\\  profiles (\\S\\ 3.5). A global discussion is made in \\S\\ 4. In \\S\\ 5 the main conclusions are summarized. ",
        "conclusions": "From the analysis carried out of the CCD and IFS HHL~73 observations, complemented with those available from the literature we found: \\begin{description}  \\item The bulk of the HHL~73 optical red emission is mainly line emission from shock-excited gas, such as is usually found in HH objects.  Continuum emission from dust scattered starlight, if present, seems to be weak.  \\item High-spatial resolution, narrow-band images of HHL~73 resolved its bow-shaped morphology into two components and a collimated emission $\\sim4$ arcsec long emerging towards the west, which is reminiscent of a microjet.  \\item The kinematics of HHL~73 derived from the IFS data show that the emission is supersonic, with both blueshifted and redshifted velocities. From a line-profile multicomponent model, we were able to resolve three velocity components for the \\sii\\ $\\lambda\\lambda$6717, 6731 \\AA\\ emission lines. These kinematic components correlate with the morphological components discovered from the high-resolution CCD images. The redshifted component, being the most extended one, spatially correlates with the bright, extended HHL~73 emission, elongated  in the north--south direction. Its spatial distribution is also in good agreement with that of the radio continuum emission. The two blueshifted (low- and high-velocity) components are morphologically well correlated with the northwestern HHL~73 emission and the microjet. In order to mimic the complex kinematic pattern derived from the IFS spectra, we tentatively propose a scenario involving both low- and highly-collimated supersonic gas ejected from a YSO.  \\item In the light of the data here discussed, we propose that the exciting source of HHL~73 (and possibly of HH 379) is deeply embedded in the dark cloud causing the high extinction, lying to the northeast of HHL~73, close to its apex. In spite of its angular separation, \\iras~21432+4719 is the most plausible candidate to be tracing the low-/intermediate YSO that is driving the centimetre, \\textit{IR} and optical emission associated with HHL~73, although high sensitivity observations with high angular resolution in the millimetre wavelength range, not affected by extinction, are needed to confirm this issue and to establish clearly the location and properties of the HHL~73 driving source.  \\item Further observations, both in the centimetre wavelength range and IFS data covering the wavelength range blueward of H$\\alpha$, to include high-excitation HH emission lines, are needed to settle the nature of the HHL~73 object completely. \\end{description}"
    },
    "0711/0711.4845_arXiv.txt": {
        "abstract": "Using a set of deep imaging data obtained by the Advanced Camera for Surveys (ACS) on the {\\it Hubble Space Telescope} (\\hst) shortly after its deployment, Yan, Windhorst \\& Cohen (2003) found a large number of F775W-band dropouts ($i$-dropouts), which are consistent with being galaxies at $z\\approx 6$. The surface density of $i$-dropouts thus derived, however, is an order of magnitude higher than those subsequent studies found in other deep ACS fields, including the {\\it Hubble} Ultra-Deep Field (HUDF). Here we revisit this problem, using both the existing data and the new data obtained for this purpose. We confirm that the large overdensity of $i$-dropouts does exist in this field, and that their optical-to-IR colors are similar to those in the HUDF. However, we have discovered that the $i$-dropout overdensity is accompanied with an even larger excess of faint field objects in this region and its vicinity. This large excess of field objects is most likely caused by the fact that we have resolved the faint diffuse light extending from an interacting galaxy pair in the Virgo Cluster, M60/NGC4647, which lies several arcminutes away from the region where the excess is found. The integrated light from the excess is a few percent of the luminosity of the galaxy pair. We argue that this excess population is most likely tidal ``debris'' and/or halo stars related to the galaxy pair rather than to the Virgo Cluster in general. The $i$-dropouts in this field are within the magnitude range where this excess of field objects occurs. The spatial distribution of the $i$-dropouts seems to follow the same gradient as the excess field population. This excess population is also red in color, and the red wing of its color distribution continuously extends to the regime where the $i$-dropouts reside. While we still cannot completely rule out the possibility that the overdensity of $i$-dropouts might be a genuine large-scale structure of galaxies at $z\\approx 6$, we prefer the interpretation that most of them are part of the excess stellar population related to M60/NGC4647. Future spectroscopic work will be needed to unambiguously identify the nature of this $i$-dropout overdensity. ",
        "introduction": "The so-called ``Lyman-break'' technique (Steidel \\& Hamilton 1992) is now widely used to select galaxies at high redshifts. This method relies on multi-band imaging to identify the characteristic discontinuity --- the {\\it Lyman-break} --- in the spectral energy distributions (SED) of high-redshift galaxies, which is largely caused by the Lyman-limit and \\Lya\\, absorptions of intervening neutral hydrogen along the sightlines to such galaxies (e.g., Madau 1995). The technique requires imaging in at least two passbands, one to the blue side of the break and the other to the red side. Lyman-break makes high-redshift galaxies much fainter in the blue band than in the red one, or in other words, it makes them seem to ``drop-out'' from the blue band. For this reason, this method is also known as the ``drop-out'' selection, and the candidates found in this way are usually referred to as ``dropouts''.  As its first application in the $z\\approx 6$ regime, Yan, Windhorst \\& Cohen (2003; hereafter YWC03) used this technique in a deep field observed by the Advanced Camera for Surveys (ACS; Ford et al. 2003) in the default pure parallel mode (Sparks et al. 2001) soon after its deployment on-board the Hubble Space Telescope (\\hst). In this redshift range, Lyman-break occurs at around 8512\\AA\\, in observer's frame (restframe 1216\\AA), which is well targeted by the F775W and F850LP filters through which this deep parallel field was imaged. YWC03 found a large number of F775W-band dropouts, and argued that the vast majority of them were very likely galaxies at $z\\approx 6$.  The number density of $z\\approx 6$ galaxies inferred from YWC03, however, seem to be much higher than those in other deep ACS fields (e.g., Bouwens et al. 2003; Stanway, Bunker \\& McMahon 2003; Dickinson et al. 2004; Bouwens et al. 2004; Bunker et al. 2004; Yan \\& Windhorst 2004, hereafter YW04). While the cosmic variance could result in significant fluctuation in the number density (e.g., Somerville et al. 2004; Bouwens et al. 2004), it cannot account for the large difference between the result of YWC03 and those of others. As the true number density of galaxies at $z\\approx 6$ is closely related to a series of important cosmological questions (e.g., the formation of early galaxies; the reionization history of the universe; etc.), it is prudent to closely examine the cause of the observed difference.  In this paper, we present new results that shed light to the nature of this overdensity. We first describe the existing and new \\hst\\, observations in \\S 2 and the data reduction in \\S 3. The overdensity of F775W-band dropouts in YWC03 is scrutinized in \\S 4, and a new interpretation is given in \\S 5. We conclude with a summary in \\S 6. For simplicity, we refer to the \\hst\\, F475W, F775W, F850LP, F110W, and F160W passbands as \\acsg, \\acsi, \\acsz, \\nicJ, and \\nicH, respectively. The F775W-band dropouts are then referred to as $i$-dropouts. The magnitudes are in AB system. All coordinates quoted are in J2000. The following cosmological parameters from the first-year Wilkinson Microwave Anisotropy Probe (WMAP) results in Spergal et al. (2003) are adopted: $(\\Omega_M, \\Omega_\\Lambda, H_0)=(0.27,0.73,71)$. Using the three-year WMAP results (Spergal et al. 2007) would not change any of our results. ",
        "conclusions": "In this paper, we use new and existing \\hst\\, data to closely examine the large overdensity of $i$-dropouts that YWC03 found in a deep ACS parallel field. The new data include NIC3 imaging of a subset of the $i$-dropouts, and a set of new ACS parallel data in a nearby region obtained during the NIC3 observations. We confirm that a factor of $\\sim 20$ overdensity of $i$-dropouts does exist in the parallel field of YWC03, and that the optical-to-IR colors of these objects are similar to those selected in other deep ACS fields such as the HUDF.  However, these $i$-dropouts might {\\it not} be galaxies at $z\\approx 6$ as one would expect. During the course of our investigation, we have found a large excess of discrete sources in both the field of YWC03 and the new ACS parallel field. This excess is most likely due to a stellar population that is related to M60/NGC4647, an interacting galaxy pair in the Virgo Cluster, which lies several arcminutes to the south of the ACS parallel fields. At the distance of the Virgo Cluster, the ACS parallel fields are several tens of kpc away from the galaxy pair, and hence the stellar objects are likely tidal ``debris'' expelled from the pair and/or distant halo stars belong to these two galaxies. Both the spatial distribution and the amplitude of the excess indicate that it is a population related more closely to the galaxy pair rather than to the Virgo's intracluster populations.  The $i$-dropouts in the field of YWC03 are within the magnitude range where the field object excess occurs. Their distribution also shows similar spatial gradient of the field object excess. Most importantly, the field object excess is red in color, and the red wing of its color distribution extends well into the regime where the $i$-dropouts were selected. For these reasons, we believe these $i$-dropouts are part of the stellar objects that constitute the field object excess rather than galaxies at $z\\approx 6$.  Nevertheless, we should caution that we still cannot completely rule out the possibility that these $i$-dropouts might indeed be $z\\approx 6$ galaxies. On one hand, we do not have any definite evidence against this interpretation. On the other hand, we have not yet found a stellar population that can reasonably explain the colors of these $i$-dropouts. To unambiguously identify their nature, spectroscopic observations of these objects are in demand. As the brightest sources in our sample are brighter than 26.0 mag in \\acsz, this might be feasible with the existing spectrographs at our current 8--10m class telescopes."
    },
    "0711/0711.3148_arXiv.txt": {
        "abstract": "We review the consequences of the growth and evolution of Black Holes on the distribution of stars and Dark Matter (DM) around them. We focus in particular on Supermassive and Intermediate-Mass Black Holes, and discuss under what circumstances they can lead to significant overdensities in the surrounding distribution of DM, thus effectively acting as DM Annihilation Boosters. ",
        "introduction": "\\label{sec:one}  To identify the nature of Dark Matter (DM) particles, three different experimental strategies have been devised. First, high-energy colliders like the Large Hadron Collider at CERN \\cite{LHC} may soon be able to produce DM particles, and possibly identify them (see e.g. Refs. \\cite{Baltz_Battaglia_Peskin_Wizanski} and \\cite{Datta_Kong_Matchev} and references therein). Another possibility is {\\it Direct Detection}, based on the search for rare events in which a DM particle scatters off the nuclei of large detectors (see e.g. Ref. \\cite{Munoz} for a review). Finally, {\\it Indirect Detection} is based on the search for the products of annihilation or decay of DM particles (such as photons, neutrinos and anti-protons). Many candidates have been proposed for DM, but the most studied ones belong to the class of Weakly Interacting Massive Particles (WIMPs) \\cite{Bertone_Hooper_Silk, Bergstrom}, that includes, among others, the lightest neutralino in Minimal or Non-Minimal Supersymmetric extensions of the Standard Model \\cite{Jungmann_Kamionkowski_Griest}, and the lightest Kaluka-Klein state in models with Unified Extra Dimensions \\cite{Appelquist_Cheng_Dobrescu,Servant_Tait,Cheng_Matchev_Schmaltz}.  Here we focus on Indirect DM searches. Ideal targets for Indirect DM searches are nearby, high DM density regions, such as local celestial bodies like the Earth and the Sun, where DM can be captured, or centers of galactic halos (see e.g. Refs. \\cite{Aharonian_et_al,Zaharijas_Hooper,Profumo,Mambrini_Munoz_Nezri_Prada,Bertone_Merritt_two,Cesarini_Fucito_Lionetto_Morselli_Ullio,Bouquet_Salati_Silk}). Two classes of astrophysical objects of particular interest are SuperMassive Black Holes (SMBHs) \\cite{Ferrarese_Ford,Kormendy_Ho}, with masses from $ 10^6 $ to $ 10^9 M_{\\odot} $ and the more speculative Intermediate Mass Black Holes (IMBHs), with a mass from $ 20 M_{\\odot} $ to $ 10^6 M_{\\odot} $ (see e.g. Refs. \\cite{Miller_Colbert,Koushiappas_Bullock_Dekel,Bertone_Zentner_Silk} and references therein). Both these classes of compact objects can influence the distribution of DM in which they are embedded, leading to strong overdensities: we review here the impact of the formation and growth of BHs on the surrounding distribution of matter, and the consequences for Indirect DM searches. Being the annihilation flux proportional to the integral of the DM density squared, scenarios where the density is {\\it boosted} by the presence, or the growth, of a central BH, are very promising for Indirect searches. BHs can thus be considered as {\\it DM Annihilation Boosters}.  The paper is organized as follows: in Section \\ref{sec:two} we discuss our current understanding of DM profiles in halos without central BHs. In Section \\ref{sec:three} we address the case of a DM distributon evolving around an already-formed BH. In Section \\ref{sec:four}, we discuss the case of adiabatic growth and the conditions under which DM ``spikes'' can form. In Section \\ref{sec:five} we will focus on IMBHs, and show that some of the problems associated to the formation of spikes around SMBHs can be evaded, with important consequences for Indirect DM searches. Conclusions will be drawn in Section VI. ",
        "conclusions": "\\label{sec:six} We have reviewed the consequences of the growth and evolution of BHs on the distribution of stars and DM around them. Initially, the BH mass was considered to be fixed in time, with particles evolving around it under the influence of its gravitational potential and their gravitational collisions. A steady-state solution is predicted for stars only for collisional nuclei: it is a zero-flux solution characterized by a central cusp with a slope of $ -7/4 $ \\cite{Bahcall_Wolf} within $ 0.2 \\mbox{ } r_h $. DM particles, being collisionless, are not affected by self gravitational interactions but by encounters with stars \\cite{Merritt_Harfst_Bertone}: in the same time that the $ -7/4 $ cusp forms for stars, the DM distribution is rearranged in a profile with a central slope of $ -3/2 $, that will decay due to heating. Other effects can significantly reduce collisional overdensities, e.g. both stellar and DM cusps react to the presence of a BH binary \\cite{Merritt_binaries} and the stellar number density can be influenced by loss-cone dynamics \\cite{Merritt}.  In the case of adiabatic growth of a BH, instead, the final DM and stellar overdensities are much steeper, reaching slope as large as $ -2.25 $ \\cite{Quinlan_Hernquist_Sigurdsson,Gondolo_Silk}. But again, the evolution of the host galaxy drives a suppression of the spike, due to cumulative mergers \\cite{Merritt_Milosavljevic_Verde_Jimenez} and possible off-center BH formation \\cite{Ullio_Zhao_Kamionkowski}. Even if DM overdensities around BHs may seem to be very promising for Indirect detection, the presence of spikes is probably not realistic \\cite{Ullio_Zhao_Kamionkowski}. Alternative strategies include the study of DM contribution to the EGB and the search for IMBHs. In the first case, the gamma-ray background receives contribution from BHs at high redshift, when spikes and mini-spikes were already present, but damping mechanisms were not effective yet. DM annihilations from spikes around SMBHs at high redshift contribute significantly to the EGB \\cite{Ahn_Bertone_Merritt_Zhang}. A population of IMBHs is then able to reproduce the value detected by EGRET, even for standard assumptions for the initial DM profile \\cite{Horiuchi_Ando}.  Mini-spikes around IMBHs can effectively be searched for with GLAST and can be identified with a careful spectral analysis of the signal \\cite{Bertone_Zentner_Silk}. They can also been searched for in the Andromeda galaxy, distributed in a quite characteristic fashion that leaves no room for alternative astrophysical interpretations \\cite{Fornasa_Taoso_Bertone}.  We show in Fig. 11 the constraining power of the different detection strategies discussed above: \\begin{itemize} \\item the dotted and solid lines are relative to the population of IMBHs in our Galaxy \\cite{Bertone_Zentner_Silk}. For each value of the mass of the DM candidate, we estimate the value of $ \\sigma v $ for which all the aformentioned realizations of the Milky Way have {\\it at least one} IMBH that should have been detected by EGRET at $ 5 \\sigma $, as discussed in Ref. \\cite{Bertone_Zentner_Silk}. In other words, if none of the EGRET unidentified sources can be interpreted as IMBH, a model with a combination $ (m_{\\chi},\\sigma v) $ above the dotted line is compatible with the data with a probability less than 0.5\\%. The solid line shows the prospectesd for detection with GLAST \\item the thin solid line is relative to the population of IMBHs in M31 \\cite{Fornasa_Taoso_Bertone}. For each value of $ m_{\\chi} $ we determined the value of $ \\sigma v $ for which all the 200 realizations of M31 have at least one IMBH with an annihilation flux larger than the $ 5 \\sigma $ sensitivity of GLAST (parametrized as in Ref. \\cite{Fornasa_Taoso_Bertone}), for an exposure of 2 months, and an energy threshold of 4 GeV. Null searches of IMBHs in Andromeda will exclude the region above the thin line \\item the dashed line is relative to the contribution of IMBHs to the EGB \\cite{Horiuchi_Ando}. Each point above the line corresponds to a DM candidate for which annihilations in IMBHs mini-spikes at high redshift exceeds the EGB values as measured by EGRET in Ref. \\cite{Sreekumar_et_al}. \\item the empty dots represent Supersymmetric Models (generated with DarkSusy \\cite{DarkSusy}) for which the neutralino can account for all the DM, being within $ 3\\sigma $ from the WMAP determination of the DM relic density \\cite{WMAP} \\item the full dots represent SUSY Models within $ 5 \\sigma $ from WMAP \\end{itemize}    GLAST can thus probe regions of the parameter space at very low cross section, so that IMBHs can be soon discovered or ruled out.  If none of these sources is detected by GLAST, it is very unlikely than such scenario will survive, at least for a WIMP DM candidate.  The gamma-ray flux from DM annihilations in spikes and mini-spikes has been widely discussed in literature (Refs. \\cite{Ullio_Zhao_Kamionkowski,Gondolo_Silk,Merritt_Milosavljevic_Verde_Jimenez,Merritt_Harfst_Bertone,Bertone_Merritt,Zaharijas_Hooper,Pieri_Branchini}, just to name a few) and many experimental collaborations are actively working to improve their discovery potential. But so far, the gravitational effects of BHs have only been observed through their  impact on stellar populations, since there is not a detected signal currently and clearly interpretated as DM annihilation. Photons from DM annihilations could have been detected in the past as subdominant contributions from gamma-ray sources \\cite{Aharonian_et_al} but the main challenge remains to disentagle such a contribution from other photon production mechanisms.  Instead, the effect of a central BH on the stellar population has been detected for the MW \\cite{Genzel_et_al,Schodel_et_al} and such interpretation is proposed also for other galaxies of the Local Group \\cite{Lauer,Young_et_al}, encouraging further attempts of detecting BH effects on DM and, eventually, of detecting DM indirectly through its annihilations products."
    },
    "0711/0711.3462_arXiv.txt": {
        "abstract": "We describe a map-making method which we have developed for the Balloon-borne Large Aperture Submillimeter Telescope (BLAST) experiment, but which should have general application to data from other submillimeter arrays.  Our method uses a Maximum Likelihood based approach, with several approximations, which allows images to be constructed using large amounts of data with fairly modest computer memory and processing requirements.  This new approach, Signal And Noise Estimation Procedure Including Correlations (SANEPIC), builds upon several previous methods, but focuses specifically on the regime where there is a large number of detectors sampling the same map of the sky, and explicitly allowing for the the possibility of strong correlations between the detector timestreams. We provide real and simulated examples of how well this method performs compared with more simplistic map-makers based on filtering.  We discuss two separate implementations of SANEPIC: a brute-force approach, in which the inverse pixel-pixel covariance matrix is computed; and an iterative approach, which is much more efficient for large maps. SANEPIC has been successfully used to produce maps using data from the 2005 BLAST flight. ",
        "introduction": "The problem of optimal ``map-making'' or ``image reconstruction'' is complex and multi-faceted, with the basic procedures and even the terminology differing dramatically between different sub-fields of astronomy.  The method adopted depends on the form in which the data are gathered, and on the dominant source of systematic effects.  From the optical to the near-IR one talks about combining ``frames'', along with measurements of ``darks'' and ``flat-fields''.  For the reduction of Cosmic Microwave Background (CMB) data the now conventional method is to start from the principle that there is a linear algebra approach to solving the Maximum Likelihood problem.  However, this has only been feasible up until now because of the limited number of detectors in the typical CMB experiment, and the fact that correlated signals among the detectors can be effectively ignored.  Because of the rapid development of large bolometer arrays, the question which arises is: how does one adapt the CMB approach to dealing with substantial numbers of detectors, and where there are significant cross-correlations of noise between the detector timestreams.  Data from the Balloon-borne Large Aperture Submillimeter Telescope (BLAST, Devlin et al.~2004) represent a new challenge for bolometric timestream-to-map algorithms.  Recent CMB experiments which use detectors similar to those used on BLAST, such as BOOMERANG \\citep{crill} and Archeops \\citep{benoit}, only use a handful of separate bolometers.  Furthermore, these experiments' off-axis designs lead to small correlations between detectors.  Consequently, the correlations could be ignored at the map-making stage, and each detector timestream could be treated as an independent sub-set of the data.  This has changed for BLAST, which has up to 139 detectors per band, with significant correlations induced by the on-axis design, as well as the higher frequencies of the observations.  Just by itself, the large number of channels increases the impact of even small time stream correlations, the contribution from which does not integrate down with increasing number of detectors, unlike the uncorrelated noise.  The high level of correlation (largely induced by temperature drifts in the obscuring secondary supports in BLAST) make it important that the correlations be handled carefully in the map-making process.  This paper describes a Signal And Noise Estimation Procedure Including Correlations  (SANEPIC) which has been developed for the analysis of BLAST.  This algorithm will also have application to many next generation experiments which will involve both noise correlations between channels (including correlations from the atmosphere) and very large numbers of detectors.  This includes the next generation of larger-format arrays for use in ground- and balloon-based instruments at microwave and millimeter wavelengths, which will have typically thousands of detectors.  This paper is arranged as follows.  We next describe the pertinent aspects of BLAST.  In Section~3, the longest section, we set out our basic map-making method.  Simulations we use for testing the method are described in Section~4, with results presented in Section~5, demonstrating the benefit of accounting for the correlated noise in BLAST-like observations.  Finally some of the maps obtained from the June 2005 BLAST flight are presented in Section~6. ",
        "conclusions": "Large format detector arrays operating at far-IR and submillimeter wavelengths are becoming the norm, rather than the exception. Ground-based instruments are plagued with common-mode emission arising from the Earth's atmosphere.  And as we have found with BLAST, the same applies to high frequency balloon-borne instruments, where we see correlated noise from thermal as well as atmospheric effects.  There is an expectation that even upcoming satellites might be faced with similar issues, because of thermal variations in the spacecraft, for example.  In general, we expect that correlated noise between detectors will be a major issue which all such experiments have to deal with, and we expect that the SANEPIC approach, which we have described here, will be widely applicable.  Indeed, there is evidence from existing arrays (e.g.,~SHARC-II and AzTEC) that once there are many detectors, there are {\\it multiple\\/} correlations between sub-sets of the detectors, as well as an overall common-mode term.  Consequently, one sees correlations between contiguous blocks of detectors on the array, or sets of detectors which share amplifiers or are otherwise coupled through the electronics.  Provided that these correlations can be investigated and their behavior modelled, it is straightforward to extend the SANEPIC approach to deal with several distinct sources of correlated noise.  Hence we expect the SANEPIC approach to be applicable to future instruments such as SCUBA-2, SPIRE, ACT, {\\sl Planck\\/} HFI and others.  There are also many experiments being planned which use large detector arrays to perform sensitive polarization measurements, and we see no reason why the SANEPIC approach could not also be extended to polarimetry."
    },
    "0711/0711.3181_arXiv.txt": {
        "abstract": "Dynamical interactions that occur between objects in dense stellar systems are particularly important for the question of formation of X-ray binaries. We present results of numerical simulations of 70 globular clusters with different dynamical properties and a total stellar mass of $2\\times 10^7 M_\\odot$.   We find that in order to retain enough neutron stars to match observations we must assume that NSs can be formed via electron-capture supernovae. Our simulations explain the observed dependence of the number of LMXBs on ``collision number'' as well as the large scatter observed between different globular clusters.  For millisecond pulsars, we obtain good agreement between our models and the numbers and characteristics of observed pulsars in the clusters Terzan 5 and 47 Tuc. ",
        "introduction": "In globular clusters (GCs), neutron stars (NSs) are seen as low-mass X-ray binaries (LMXBs), bright or  quiescent, and as binary or single millisecond pulsars (MSPs). As the numbers per unit mass of LMXBs and MSPs in clusters greatly exceed their numbers in the Galaxy, their origin has been linked to stellar encounters that should occur frequently in an environment of high stellar density (Clark 1975). Galactic GCs are known to contain thirteen bright X-ray sources, and as many as $\\sim$100-200 quiescent LMXBs (qLMXBs) are thought to exist in the Galactic GC system (Heinke et al. 2003). Millisecond pulsars, likely descendants of LMXBs (for a review see, e.g., Bhattacharya and van den Heuvel 1991), also are present in GCs in great numbers: about 140 GC millisecond radio pulsars have been detected by now\\footnote{See http://www.naic.edu/~pfreire/GCpsr.html for an updated list.}, with more than a dozen in several GCs -- in 47~Tuc (Camilo et al. 2000), M28 (Stairs et al. 2006),  and Terzan~5 (Ransom et al. 2005). As only a few X-ray binaries or few dozen MSPs are present per fairly massive (more than a million stars) and dense cluster, computationally this problem is  very challenging. The target time for direct $N$-body methods to address the million-body problem is around 2020 (Hut 2006). In our studies, we use a modified encounter rate technique method, described in detail in Ivanova et al. (2005), and with the updates described in Ivanova et al. (2007). ",
        "conclusions": "We studied the formation and evolution of NSs in globular clusters. We find that NS formation via different channels of ECS is very effective in GCs and provides most of the retained NSs. Having as many as a few hundred retained NSs per typical GC, or about 1100 per cluster like 47 Tuc, produces LMXBs in numbers comparable to observations. We note that if AIC does not lead to the formation of NSs, then the number of formed NSs is reduced only by $\\sim 20\\%$, but the number of appearing LMXBs is decreased by 2-3 times (per Gyr, at the cluster age of 11 Gyr), although it may still be consistent with the observations, given the large scatter in the simulations. We find that up to half of NSs could gain mass after their formation through mass transfer, hyper-accretion during a common envelope, or physical collision. It is likely that most of these mass-gaining events do not lead to  NS spin-up, and that only a few per cent of all NSs appear eventually as MSPs, implying that there is a large underlying population of unseen NSs in GCs."
    },
    "0711/0711.4242_arXiv.txt": {
        "abstract": "We exploit the gauge--invariant formalism to analyse the perturbative behaviour of two cosmological models based on the generalized Chaplygin gas describing both dark matter and dark energy in the present Universe. In the first model we consider the generalized Chaplygin gas alone, while in the second one we add a baryon component to it. We extend our analysis also into the parameter range $\\alpha > 1$, where the generalized Chaplygin gas sound velocity can be larger than that of light.  In the first model we find that the matter power spectrum is compatible with the observed one only for $\\alpha < 10^{-5}$, which makes the generalized Chaplygin gas practically indistinguishable from $\\Lambda$CDM.  In the second model we study the evolution of inhomogeneities of the baryon component. The theoretical power spectrum is in good agreement with the observed one for almost all values of $\\alpha$. However, the growth of inhomogeneities seems to be particularly favoured either for sufficiently small values of $\\alpha$ or for $\\alpha \\gtrsim 3$. Thus, it appears that the viability of the generalized Chaplygin gas as a cosmological model is stronger when its sound velocity is superluminal. We show that in this case the generalized Chaplygin gas equation of state can be changed in an unobservable region in such a way that its equivalent $k$-essence microscopical model has no problems with causality. ",
        "introduction": "Together with quintessence (a scalar field with some potential minimally coupled to gravity) \\cite{PR88,W88,Cald,S98} and other physical and geometrical (modified gravity) models of dark energy (see, e.g., the reviews \\cite{SS00,SS06} for exact definitions), the generalized Chaplygin gas \\cite{Kam2} (hereafter gCg) is one widely studied model among those proposed to describe the observed accelerated expansion of the universe \\cite{Riess,Perlmutter}. In its generalized form, the Chaplygin gas is a barotropic fluid with the following equation of state: \\begin{equation}\\label{gcgeos} \\pp  = - A\\rr^{-\\alpha}, \\end{equation} where $A$ and $\\alpha$ are positive constants (for the original Chaplygin gas $\\alpha = 1$). The generalized equation of state (\\ref{gcgeos}) has been introduced in \\cite{Kam2} and analysed in \\cite{Bilic,bentobertosen,Gor} and in many other subsequent papers. In contrast to many models describing dark energy alone, the gCg gives a unified description of dark matter and dark energy, enrolling itself in the class of so-called quartessence or unified dark matter (UDM) cosmological models. It allows to interpolate between a dust--dominated phase of the evolution of the Universe in the past and an accelerated one at recent time, see Eq. (\\ref{gcgrhoevo}) below. That is why the gCg has attracted much attention in cosmology.  The evolution of the energy density as a function of the cosmic scale factor $a(t)$ of the Friedmann--Lema\\^{i}tre--Robertson--Walker (FLRW) cosmological model is easily obtained from (\\ref{gcgeos}) and from the energy conservation equation: \\begin{equation}\\label{gcgrhoevo} \\rr = \\left(A + \\frac{B}{a^{3(1 + \\alpha)}}\\right)^{\\frac{1}{1 + \\alpha}}, \\end{equation} where $B$ is an integration constant usually chosen to be positive (if $B < 0$, the weak energy condition is violated and phantom cosmology takes place, see for example \\cite{Zhang,Bou}).  From (\\ref{gcgeos}) and (\\ref{gcgrhoevo}), the parameter $\\ww \\equiv p_{\\rm Ch}/\\rho_{\\rm Ch}$ and the square sound velocity $c_{\\rm sCh}^{2}$ have the following expressions: \\numparts \\begin{eqnarray} w_{\\rm Ch} &=& -\\left[1 + \\frac{B}{A}\\frac{1}{a^{3\\left(\\alpha + 1\\right)}} \\right]^{-1}, \\label{gcgw}\\\\ c_{\\rm sCh}^{2} &\\equiv& {\\frac {{\\rm d} \\pp}{{\\rm d} \\rr}} = -\\alpha w_{\\rm Ch}. \\label{gcgcs} \\end{eqnarray} \\endnumparts Note that for $\\alpha > 1$, $c_{\\rm sCh}^{2}$ may exceed the velocity of light (set to unity in our notations) in the course of the recent or future evolution of the Universe, while it was non-relativistic at large redshifts, during the matter--dominated stage, for all values of $\\alpha$.  Cosmological models based either on one or on more than one fluid are essentially different. The reason for this is that the various fluid components interact indirectly via geometry and may also interact by direct energy exchange. Direct non-gravitational interactions are typically neglected but the indirect ones are always present and can give rise to totally different evolutions, especially at the perturbative level. We will see how this works for the gCg together with baryons.  The pure gCg, namely the cosmological model based on the gCg alone, has passed many tests of standard cosmology. Quite recently, updated constraints for the gCg parameters have been published \\cite{Puxun,WY07,Davis}. However, the behaviour of the gCg under perturbations is still problematic \\cite{Bean, Sandvik}. In \\cite{Sandvik} the study of the power spectrum of large scale structures seems to indicate that the best fit value of $\\alpha$ is very close to zero, rendering the gCg indistinguishable from $\\Lambda$CDM [the latter is indeed the limit of gCg for $\\alpha \\to 0$, as one can infer from (\\ref{gcgH}) below]. This feature has to be attributed to the gCg sound velocity which, during the cosmic evolution, grows from 0 to $\\sqrt{\\alpha}$ driving inhomogeneities to oscillate (if $\\alpha > 0$) or to blow--up (if $\\alpha < 0$). This characteristic seems to be common to all UDM models \\cite{Sandvik} which, as a consequence, might appear to be ruled out.  However, as shown in \\cite{Beca}, it appears that adding a baryon component to the gCg the problem disappears and the model is in agreement with observations. But the debate still goes on and the authors of \\cite{Sandvik} have corroborated their claim about ruling out gCg and, in general, all quartessence models through gravitational lensing measurements on the basis of the current value of the cosmological parameter $\\sigma_{\\rm 8}$ (the rms mass dispersion on a sphere of radius $8h^{-1}$ $Mpc$).  Still, the parameter $\\sigma_{\\rm 8}$ need not necessarily be a confident discriminant among different cosmological models since it is clearly a quantity which is strongly affected by non--linear effects, while all the calculations performed in \\cite{Sandvik} are carried out in the linear regime. This point was also raised much earlier in \\cite{peacock} in a more general context.  In the present paper we investigate the issue of cosmological perturbations in gCg models using the gauge--invariant formalism \\cite{Bardeen}. We consider both the gCg alone and in the presence of baryons, confirming the results of \\cite{Sandvik} and \\cite{Beca}. In addition, the gCg seems to favour structure formation when its sound velocity is superluminal, in particular when $\\alpha \\gtrsim 3$.  The structure of the paper is the following: \\begin{enumerate} \\item In section \\ref{Giform} we briefly outline Bardeen gauge--invariant formalism and derive the equations that we will numerically solve in the rest of the paper. \\item In section \\ref{gcgruled} we investigate structure formation in a gCg--dominated universe and display the role of the sound velocity. \\item In section \\ref{gcgpowers} we calculate the theoretical power spectrum in a gCg--dominated universe and compare it with the observed one. \\item In section \\ref{gcgandbar} we consider a two--fluid model based on gCg in presence of baryons. For this variant of the model we perform the same analysis of the previous two sections and compare the results. \\item In section \\ref{concl} we present our conclusions and discuss in more detail why a possible superluminal sound velocity of the gCg is not prohibited by causality arguments. \\item For completeness, we add an appendix in which, using the technique developed in \\cite{SolovStar} we write down the exact solution for perturbations of dust-like matter and generalized Chaplygin gas in the Newtonian approximation at the matter--dominated stage. \\end{enumerate} ",
        "conclusions": "From our analysis it emerges that: \\begin{itemize} \\item The generalized Chaplygin gas cosmological model, with no additional fluid components, is compatible with structure formation and large scale structure clustering properties (encoded in the power spectrum) only for $\\alpha$ sufficiently small ($\\alpha < 10^{-5}$), in which case it is practically indistinguishable from the $\\Lambda$CDM model. \\item Adding to the generalized Chaplygin gas model a baryon component we find that the growth of the density contrast and the large scale structure clustering properties are compatible with observations for all values of $\\alpha$. However very small values of $\\alpha$ and the case $\\alpha \\gtrsim 3$ are favoured. \\end{itemize}  Note, however, that the integrated Sachs-Wolfe (ISW) effect (not considered in this paper) leads to rather large enhancement of low CMB multipoles $C_l$ in UDM models, in contrast to a modest increase of $C_l$ in the $\\Lambda$CDM model \\cite{KS85}, that results in the exclusion of the gCg with $0.01< \\alpha \\le 1$ \\cite{Bean,CF03,AFBC03} (see also the recent paper \\cite{BB07} for an analytical treatment of the ISW effect for UDM models). It may be interesting to investigate if it remains so large for all $\\alpha > 1$.  Now, since most of the new results of this paper refer to the superluminal case $\\alpha > 1$, it is natural to consider the question if this range is physically admissible in more detail. The fact that the group velocity $c_{\\rm sCh}$ in media may exceed the light velocity in vacuum (unity in our notations) is not, in itself, unphysical. Not only it is not prohibited from the theoretical point of view, but this phenomenon has been actually observed in laboratories in case of light propagation through dispersive media. As was proved long ago (see e.g. \\cite{Brill}), what is really necessary in order not to violate causality is that the signal (or, the wavefront) velocity should not exceed $1$.  Then, what is the signal velocity for the gCg? It does not appear possible to answer this question at the hydrodynamic level since the gCg phenomenological equation of state in the normal (non-phantom) case requires $\\rho_{\\rm Ch}\\ge A^{1/(1+\\alpha)}$, so that the limit $\\rho_{\\rm Ch}\\to 0$ corresponding to a wavefront expanding into vacuum cannot be studied. Therefore, to determine the value of the signal velocity for the gCg, one has to use some underlying microscopic field--theoretical model from which the equation of state (\\ref{gcgeos}) arises in macroscopic hydrodynamics.  We will use the so-called tachyon representation of the gCg for this purpose where the gCg is described by a scalar field $\\phi$ ($k$-essence) with the Lagrangian density \\cite{bentobertosen} \\begin{eqnarray}\\label{tach} {\\cal L}\\equiv p_{\\rm Ch}=-V_0\\left[ 1-\\left(\\frac {X}{V_0}\\right)^ {\\frac {1+\\alpha}{2\\alpha}}\\right]^{\\frac {\\alpha}{1+\\alpha}}~, ~~ X= \\frac {1}{2} \\phi_{,\\mu}\\phi^{,\\mu} < V_0~, \\\\ \\rho_{\\rm Ch}=2X\\frac {d{\\cal L}}{dX} - {\\cal L} = V_0 \\left[ 1-\\left(\\frac {X}{V_0}\\right)^{\\frac {1+\\alpha}{2\\alpha}} \\right]^{-\\frac {1}{1+\\alpha}}~, ~~~A=V_0^{1+\\alpha}~. \\end{eqnarray} Though other microscopic models are also possible (see e.g. discussion in \\cite{GKMPS05} in the case $\\alpha = 1$), the present one has the following advantages: a) its solutions are in one--to--one correspondence to those of the hydrodynamic gCg if $X>0$, and b) it has the same equations for perturbations, too.  Now we are able to consider the limit $X\\to 0$ corresponding to a wavefront expanding into vacuum with no $k$-essence in front of it. To prove that the signal velocity is equal to unity in this case, it is sufficient to show that the wave equation corresponding to (\\ref{tach}), namely \\begin{equation} \\label{tach-we} {\\cal L}_X\\phi_{;\\mu}^{;\\mu} + {\\cal L}_{XX}\\phi^{,\\mu}\\phi^{,\\nu} \\phi_{;\\mu;\\nu} = 0, \\end{equation} has plane-wave solutions of the type $f(x-t)$ in flat space-time (since the signal velocity is also the infinite momentum limit of the phase velocity of a wave). It is straightforward to check that the sufficient condition for this to hold is that both ${\\cal L}_X$ and ${\\cal L}_{XX}$ are finite at $X=0$ (a similar remark is made in the recent paper \\cite{Vik}, see also \\cite{KVW07}). In addition, ${\\cal L}_X(X=0)$ should be non-negative in order for ghosts not to be present.  In the superluminal case $\\alpha > 1$, the Lagrangian density (\\ref{tach}) does not satisfy this condition since ${\\cal L}_X$ and ${\\cal L}_{XX}$ diverge at $X\\to 0$. However, without going into a more detailed analysis of the singularity structure at a null hyper--surface in this case, we propose a simple {\\em sufficient} way to cure the causality problem for the gCg. Namely, let us assume that the expression (\\ref{tach}) changes to \\begin{equation} \\label{change} {\\cal L} = -V_0 +  X\\left(\\frac {V_0}{V_1}\\right)^{\\frac {\\alpha -1}{2\\alpha}} \\end{equation} for $X\\ll V_1$ where $V_1\\ll V_0$ and that it smoothly matches (\\ref{tach}) at $X\\sim V_1$. Then it is clear that the signal velocity is exactly unity. The Lagrangian density (\\ref{change}) describes a massless scalar field minimally coupled to gravity and a cosmological constant.  The correspondingly corrected macroscopic gCg remains the same in the observable region $X\\gtrsim V_0$, or $\\rho_{\\rm Ch}+p_{\\rm Ch}\\gtrsim V_0$ (note that $V_0$ is of the order of the present day critical energy density in the Universe) but will become a mixture of a cosmological constant and the extremely stiff ideal fluid ($p=\\rho$) in the far future when $X$ will drop to $V_1$. With the same accuracy, one can say that the corrected gCg becomes the usual Chaplygin gas when its density is very close to $V_0$, much closer than the one which occurs presently. Thus, by slightly changing the gCg equation of state in the unobservable region, the causality problem is cured. As a result, the present gCg superluminal sound velocity is a {\\em transient} phenomenon which disappears at $t\\to\\infty$.  Finally, note that in the limit $\\alpha\\to\\infty$, the gGg realizes the Parker-Raval scenario \\cite{PR99} of a sudden transition from the matter--dominated stage with $a(t)\\propto t^{2/3}$ to the de Sitter expansion with a constant curvature (modulo a small term $\\propto \\Omega_b$). The original idea of Parker and Raval to produce such a scenario using non-perturbative vacuum polarization of a light minimally coupled scalar field is questionable, both regarding derivation of the effective action and stability of the resulting specific kind of the $f(R)$ theory of gravity (here $R$ is the Ricci scalar). But the same scenario based on the limiting case of the gCg does not have these problems. This shows once more that the superluminal gCg deserves further, more detailed study.  \\ack The work of AYK and AAS was partially supported by the Research Programme \"Astronomy\" of the Russian Academy of Sciences, by the RFBR grant 05-02-17450 and by the grant LSS-1157.2006.2. AAS also thanks Prof. Misao Sasaki and the Yukawa Institute for Theoretical Physics, Kyoto University, for hospitality during the middle part of this project.  \\appendix"
    },
    "0711/0711.2244_arXiv.txt": {
        "abstract": "{}{}{}{}{}  \\abstract { \\object{PG~1159-035}, a pre-white dwarf with $T_{\\mathrm{eff}}\\simeq 140\\,000$~K, is the prototype of both two classes: the PG~1159 spectroscopic class and the DOV pulsating class. Previous studies of \\object{PG~1159-035} photometric data obtained with the Whole Earth Telescope (WET) showed a rich frequency spectrum allowing the identification of 122 pulsation modes. Analyzing the periods of pulsation, it is possible to measure the stellar mass, the rotational period and the inclination of the rotation axis, to estimate an upper limit for the magnetic field, and even to obtain information about the inner stratification of the star. } { We have three principal aims: to increase the number of detected and identified pulsation modes in \\object{PG~1159-035}, study trapping of the star's pulsation modes, and to improve or constrain the determination of stellar parameters. } { We used all available WET photometric data from 1983, 1985, 1989, 1993 and 2002 to identify the pulsation periods. } { We identified 76 additional pulsation modes, increasing to 198 the number of known pulsation modes in  \\object{PG~1159-035}, the largest number of modes detected in any star besides the Sun. From the period spacing we estimated a mass $M/M_\\odot = 0.59\\pm 0.02$ for \\object{PG~1159-035}, with the uncertainty dominated by the models, not  the observation. Deviations in the regular period spacing suggest that some of the pulsation modes are trapped, even though the star is a pre-white dwarf and the gravitational settling is ongoing. The position of the transition zone that causes the mode trapping was calculated at $r_c/R_\\star = 0.83\\pm 0.05$. From the multiplet splitting, we calculated the rotational period $P_{\\rm rot} = 1.3920\\pm 0.0008$ days and an upper limit for the magnetic field, $B< 2000$ G. The total power of the pulsation modes at the stellar surface changed less than 30\\% for $\\ell=1$ modes and less than 50\\% for $\\ell=2$ modes. We find no evidence of linear combinations between the 198 pulsation mode frequencies. \\object{PG~1159-035} models have not  significative convection zones, supporting  the hypothesis that nonlinearity arises in the convection zones in cooler pulsating white dwarf stars. } {} ",
        "introduction": "%  The star \\object{PG~1159-035} was identified by R. F. Green in 1977 in a survey for objects with ultraviolet excess, known as the Palomar-Green Survey (Green et al.  1986). The presence of lines of He II in the \\object{PG~1159-035} spectrum suggested a high superficial temperate (McGraw et al.  1979). The analysis of the far ultraviolet flux distribution --- from $\\sim 1200$~\\AA\\ to the Lyman limit at $912$~\\AA\\ --- obtained with the {\\it Voyager 2} ultraviolet spectrophotometer indicated an effective temperature above $100\\,000$ K (Wegner et al.  1982). Later analysis with the {\\it IUE} and {\\it EXOSAT} show that \\object{PG~1159-035}  is one of the hottest stars known (Sion et al.  1985, Barstow et al.  1986); the current estimated temperature for \\object{PG~1159-035} is $140\\,000\\pm 5\\,000$ K (Werner et al.  1991, Dreizler et al.  1998 and Jahn et al.  2007) and $\\log g(cgs) = 7.0\\pm 0.5$ (Werner et al. 1991), placing it in the class of the pre-white dwarf stars.  McGraw et al.  (1979) discovered that \\object{PG~1159-035} is a variable star and identified at least two pulsation periods. The Fourier transform of more extensive light curves obtained in the following years, between 1979-1985, allowed the detection of eight pulsation modes (Winget et al. 1985); the highest amplitude mode has a period of $516\\,s$.  A long light curve is necessary to resolve two nearby frequencies in the  Fourier transform (FT) of a multiperiodic pulsating star.  The Fourier transform resolution is roughly proportional to the inverse of the light curve length. For instance, if the difference between two frequencies is equal to 100 $\\mu$Hz, just three hours of photometric data are needed to resolve them, but, more than 10 days are needed  if the difference between them is of 1 $\\mu$Hz. On other hand, the presence of gaps in the light curve  introduces in the FT an intricate structure of side-lobes, which may hinder the detection and identification of real pulsation frequencies.  With the establishment of the WET ({\\it Whole Earth Telescope}) in 1988 (Nather et al.  1990), \\object{PG~1159-035} was observed for about 12 days with an effective coverage around $60\\%$, resulting in a quasi-continuous 228 hours of photometric data. The high-resolution Fourier transform of the light curve allowed the detection and identification of 122 peaks (Winget et al.  1991).  In white dwarf and pre-white dwarf stars, gravity plays the role of the restoring force in the oscillations. Any radial displacement of mass suffers the action of the gravitational force causing the displaced portion of mass to be scattered inwards and sideways. This type of nonradial pulsation modes are called {\\it g-modes}.  General nonradial pulsations are characterized by three integer numbers: $k$, $\\ell$, $m$. The number $k$ is called {\\it radial index} and is related with the number of {\\it nodes} in the radial direction of the star. The number $\\ell$ is called {\\it index of the spherical harmonic} (or {\\it degree of the pulsation mode}). For nonradial modes, $\\ell > 0$, while a radial pulsation has $\\ell = 0$. In white dwarf stars the pulsations are dominated by temperature variations (Robinson, Kepler, and Nather 1982). The index $\\ell$ is related with the total number of hotter and colder zones relative to the mean effective temperature on the stellar surface. Finally, the number $m$ is a number between $-\\ell$ and $+\\ell$ and is called the {\\it azimuthal index}. The degeneracy of modes with different $m$ is broken when the spherical symmetry is broken, for example, by rotation of the star, or the presence of magnetic fields. The absolute value of the azimuthal index, $|m|$, is related with the way  the cold zones and hot zones are arranged on the stellar surface. The sign of the index $m$ indicates the direction of the temporal pulsation propagation. We adopted the convention used by                   % Winget et al.  (1991): $m$ is positive if the pulsation and the rotation have the same direction and negative if they have opposite directions.  The \\object{PG~1159-035} Fourier transform published in 1991 also revealed the presence of triplets and multiplets, caused by {\\it rotational splitting}, allowing the determination of the rotational period of the star ($P_{rot}=1.38$ days). Stellar rotation causes the g-modes with $m\\ne 0$ to appears in the FTs as frequencies slightly higher ($m>0$) or lower ($m<0$) than the frequency of the $m=0$ mode, depending on whether the pulsation is travelling in the same direction of the rotation of the star (higher frequency) or in opposite direction (lower frequency). Slow rotation splits a mode in a multiplet of $2(\\ell+1)$ peaks. For  $\\ell=1$ modes, the multiplets have three peaks (triplets) and for $\\ell=2$ modes they have five peaks (quintuplets). But not necessarily all components are seen in the FTs, because some of them might be excited with amplitudes bellow the detection limit. Besides stellar rotation, a weak magnetic field can also break the degeneracy and cause an observable splitting of the pulsation modes into $(\\ell+1)$ components in first order. However, no notable magnetic splitting has been observed in \\object{PG~1159-035} (Winget et al. 1991).  The immediate goal of this work was to detect and identify a larger number of pulsating modes in \\object{PG~1159-035} from the analysis and comparison of the FTs of photometric data obtained in different years. A consequence is the improvement in the determination of the spacing between the periods of the pulsation modes used in the calculation of the stellar mass and in the determination of the inner stratification of the star. The analysis of the splitting in frequency in the multiplets of the combined data allows the calculation of the rotation period with higher accuracy and a better estimate of a upper limit for the strength of the star's magnetic field. We are also interested in the search for possible linear combination of frequencies, as an indication of nonlinear behavior.  This paper is organized as follow: in next Section we present some basic background in pulsation theory. In \\S3 and \\S4 we discuss the observational data and the data reduction process used in this work. In \\S5 we discuss the detection of pulsation modes from the \\object{PG~1159-035} FTs. Then, in \\S6, we present the calculation of the period spacing for the detected pulsation modes. The mode identification, i.e., the determination of the numbers $k$, $\\ell$ and $m$ of the detected pulsation modes is discussed in \\S7 and  in \\S8 \\S9 we calculate the rotational and the magnetic splittings and the rotation period of the star. An estimate of the inclination angle of the rotational axis of the star is done in \\S10 and in \\S11 we use the  splitting to obtain an upper limit for the \\object{PG~1159-035} magnetic field. In \\S12 we present  the estimate of the mass of \\object{PG~1159-035} from the period spacing in comparison with the masses calculated from spectroscopic models. The analysis of a possible trapping of pulsation modes in \\object{PG~1159-035} is presented in \\S13 and in \\S14 we use the results to calculate the position of a possible trapping zone inside the star. In \\S15 we comment on the absence of linear combination of frequencies in \\object{PG~1159-035} and in \\S16 on the energy conservation of the pulsation modes in the star surface. Finally, in \\S17 we summarize our main results. ",
        "conclusions": "%  Winget et al. (1991), analyzing the WET 1989 data set of \\object{PG~1159-035}, found 122 pulsation modes, with frequencies between 1000 and 3000 $\\mu$Hz in the star's FT, with a spacing in period of $\\Delta P_1=21.50 \\pm 0.03\\,$s and $\\Delta P_2 = 12.67 \\pm 0.03\\,$s. The seismological mass calculated from the spacing in period and using a theoretical model for a PG~1159 star was $M/M_\\odot = 0.586\\pm 0.003$. The analysis of the fine structure of the multiplets shown a spacing in frequency $\\delta\\nu_{\\rm rot,1}=4.22\\pm 0.04\\mu$Hz and   $\\delta\\nu_{\\rm rot,2}=6.92\\pm 0.07\\mu$Hz for $\\ell=1$ and $\\ell=2$ modes, respectively, which allowed the calculation of the star's rotation period, $P_{\\rm rot}=1.38\\pm 0.01$ days; of the rotation axis inclination, $i\\sim 60^o$; and to estimate an upper limit for the magnetic strength, $B<6000\\,$G.  In this work, we followed the same steps of Winget, but using a larger number of data sets and improved  data reduction and  data analysis techniques. The combination of the Fourier transforms of the data sets from different years (1983, 1985, 1989, 1993 and 2002) allows us to refine the determination or put new constraints over several stellar parameters of \\object{PG~1159-035}. Our main results are:   \\begin{enumerate} \\item We identified 76 new pulsation modes, increasing to 198 the total number of known pulsation modes for \\object{PG~1159-035}. Only 14 of them (all with $\\ell=1$) are present in the FTs of all the years, but with different amplitudes. The comparison of the annual FTs shows that the amplitudes of the pulsation modes are changing in time, and can reach amplitudes below our detection limits. No evidence of $\\ell>2$ modes was found in the combined FTs of all the years.  \\item The period spacings are $\\Delta P_1=21.43\\pm 0.03\\,s$ for $\\ell=1$ modes and $\\Delta P_2 = 12.38\\pm 0.01\\,s$ for $\\ell=2$ modes. The period constant, $P_o=\\sqrt{\\ell (\\ell+1)}\\, \\Delta P_\\ell$, is $P_o=30.32\\pm 0.03$ sec.  \\item We found a mass $M/M_\\odot = 0.59\\pm 0.02$ from the KB94 parameterization. The apparent lower accuracy in the mass determination relative to the Winget determination, although we have substantially decreased the uncertainty in $P_o$, is because we took into account the dominant uncertainty in the theoretical models. As pointed out by Winget, the stated error in their determination of the \\object{PG~1159-035} mass reflects only the uncertainty in the measured period spacing, and not the systematic errors associated to the models. In this sense, even though our result is formally less accurate, it is more realistic.  \\item Analyzing the spacing in frequency inside the multiplets we found that the splitting due to the stellar rotation is $\\delta\\nu_{\\rm rot,1}=4.154\\pm 0.002\\mu$H\\ for $\\ell=1$ modes and $\\delta\\nu_{\\rm rot,2}=6.90\\pm 0.01\\mu$Hz for $\\ell=2$ modes.  \\item We also estimated that the splitting in frequency due to the magnetic field for $\\ell=1$ modes is $\\delta\\nu_{\\rm mag,1}=0.007\\pm 0.002\\mu$Hz, contributing with less than 1\\% of the total splitting. Unfortunately, it was not possible to calculate the magnetic spacing in frequency for the $\\ell=2$ modes, due to the absence of several components in the multiplets.  \\item From the rotational  frequency splitting we calculated the rotational period $P_{\\rm rot} = 1.3920\\pm 0.0008$ days.  \\item The magnetic splitting in frequency suggests a upper limit for the magnetic strength lower than the previous estimates: $B\\lesssim 2000\\,$ G.  \\item The analysis of the fine structure of the combined data multiplets ($\\ell=1$ {\\it and} $\\ell=2$) suggests that the inclination angle of the rotational axis is $i\\sim 70^o\\pm 6^o$.  \\item The $\\Delta P$ diagram of \\object{PG~1159-035} for $\\ell=1$ modes suggests that \\object{PG~1159-035} is already a stratified star. The $\\Delta P$ diagram presents five minima that can be interpreted as  the indication of trapped modes with periods of 452.43s, 538.12s, 668.1s, 773.7s and 883.6s.  \\item For this sequence of trapped modes, we calculated the position of the transition zone that causes the trapping mode at $r_c/R_\\star = 0.83\\pm 0.05$ for a star with $M/M_\\odot = 0.59$ and $R_\\star/R_\\odot = 0.025$.  \\item There is no evidence of linear combinations of frequencies in \\object{PG~1159-035}. As models of PG 1159-035 do not have any significant convective layer, this provides indirect support to the hypothesis that nonlinearity arises in the convection zone.  \\item Comparing the total power of the pulsation modes of the different years between 1983 and 2002, we observe that the differences relative to the mean value are less than $30\\%$ for the $\\ell=1$ modes and less than $50\\%$ for the $\\ell=2$ modes, indicating that the total surface kinetic energies are not conserved.  \\end{enumerate}   In continuation to the present work, we measured the temporal changing of the pulsation period ($\\dot P$) of several pulsation modes in \\object{PG~1159-035}. As PG~1159-035 is a very hot star, it is quickly evolving and its pulsation periods are changing in time. The period changes are large enough to be directly measured and for some we can  derive the second order temporal variation ($\\ddot P$). The results are presented in a separated paper (Costa et al.  2007).  We are currently performing a detailed asteroseismological study on \\object{PG~1159-035} based on an enlarged set of full PG1159 evolutionary models. Preliminary results, obtained from the comparison of the average period spacings of the models and of the observed periods, suggest that the mass of PG~1159-035 is $\\approx 0.56 - 0.57\\,M_\\odot$ (see Sect.\\ref{sect:12_2}). The next step is to fit the models to the individual observed periods in PG~1159-035. This study will be published in a separated paper (C\\'orsico et al.  2007 in preparation)."
    },
    "0711/0711.0594_arXiv.txt": {
        "abstract": "{The main goal of this work is to calculate the contributions to the cosmological recombination spectrum due to bound-bound transitions of helium. We show that due to the presence of helium in the early Universe unique features appear in the total cosmological recombination spectrum. These may provide a unique observational possibility to determine the relative abundance of primordial helium, well before the formation of first stars. We include the effect of the tiny fraction of neutral hydrogen atoms on the dynamics of $\\ion{He}{ii}\\rightarrow \\ion{He}{i}$ recombination at redshifts $z\\sim 2500$. As discussed recently, this process significantly accelerates $\\ion{He}{ii}\\rightarrow \\ion{He}{i}$ recombination, resulting in rather narrow and distinct features in the associated recombination spectrum. In addition this process induces some emission within the hydrogen Lyman-$\\alpha$ line, before the actual epoch of hydrogen recombination round $z\\sim 1100-1500$. We also show that some of the fine structure transitions of neutral helium appear in absorption, again leaving unique traces in the Cosmic Microwave Background blackbody spectrum, which may allow to confirm our understanding of the early Universe and detailed atomic physics. } ",
        "introduction": "The recombination of helium practically does not influence the Cosmic Microwave Background (CMB) angular fluctuations, as measured with great success by {\\sc Wmap} \\citep{WMAP_params}, since it occurred well before the Thomson visibility function defined by hydrogen recombination \\citep{Sunyaev1970} reaches its maximum. However, similar to the release of photons during the epoch of cosmological hydrogen recombination \\citep{Jose2006, Chluba2007, Chluba2006b}, one does expect some emission of photons by helium, and the main goal of this paper is to calculate the contributions to the cosmological recombination spectrum due to bound-bound transitions of helium.  In our recent papers we computed the detailed cosmological recombination spectrum of hydrogen resulting from bound-bound \\citep{Jose2006, Chluba2007} and bound-free \\citep{Chluba2006b} transitions between atomic levels, including up to 100 shells, also taking into account the evolution of individual energetically degenerate angular momentum sub-states. We followed the ideas and suggestions of earlier investigations \\citep{Zeldovich68, Peebles68, Dubrovich1975, Bernshtein1977, Beigman1978, RybickiDell1993, Dubrovich95, Burgin2003, Dubrovich2004, Kholu2005, Wong2006}. Observations of these recombinational lines might provide an additional unbiased way to {\\it directly} determine the baryon density of the Universe (e.g.  see \\citet{Dubrovich1975} and \\citet{Bernshtein1977}, or more recently \\citet{Kholu2005} and \\citet{Chluba2007d}) and to obtain some additional information about the other key cosmological parameters, facing different degeneracies and observational challenges.  Obviously, direct evidence for the emission of extra $\\sim 5$ photons per recombining hydrogen atom \\citep{Chluba2006b} will be an {\\it unique proof} for the completeness of our understanding of the processes occurring at redshifts $z\\sim 1400$, i.e. before the CMB angular fluctuations were actually formed. From this point of view an observation of lines emitted during $\\ion{He}{iii}\\rightarrow \\ion{He}{ii}$ close to $z\\sim 6000$, and $\\ion{He}{ii}\\rightarrow \\ion{He}{i}$ around $z\\sim 2500$ will be an even more impressive confirmation of the predictions within the standard hot big bang model of the Universe, realising that nowadays exact computations using the full strength of atomic physics, kinetics and radiative transfer in principle should allow a prediction of the cosmological recombination spectrum from both epochs with very high precision.  The first attempt to estimate the emission arising from helium recombination was made by \\citet{DubroStol1997}. However, only now detailed numerical computations are becoming feasible, also due to the fact that atomic physicists began to publish {\\it accurate} and {\\it user-friendly} transition rates \\citep{DrakeMorton2007, PCBeigman2007} for neutral helium, including singlet-triplet transitions, which very strongly influence the recombination of helium.  According to the computations of nuclear reactions in the early Universe \\citep{Olive1995, Cyburt2004}, the abundance of helium is close to $8\\%$ percent of the number of hydrogen atoms, so naively only small additional distortions of the CMB blackbody spectrum due to helium recombination are expected. However, for helium there are {\\it two} epochs of recombination, a fact that at least doubles the possible amount of additional photons. Furthermore, $\\ion{He}{iii}\\rightarrow\\ion{He}{ii}$ recombination is very fast, in particular because there is a large quasi-constant amount of free electrons belonging to hydrogen \\citep{DubroStol1997}. This implies that photons are emitted in a much shorter period, so that more narrow features are produced\\footnote{As we will demonstrate here, even the scattering of photons by free electrons cannot change this conclusion (see Sect.~\\ref{sec:spec.e-scatt}).}. It is also very impressive that the $\\ion{He}{iii}\\rightarrow\\ion{He}{ii}$ recombination lines practically coincide and therefore amplify the corresponding hydrogen line (see Fig.~\\ref{fig:final}). This is because the difference in the redshifts of the two recombinations is close $\\sim 4.3$, when on the other hand the energy of similar transitions scales as $Z^2=4$ for $\\ion{He}{ii}$, such that the two effects practically compensate eachother.  The spectral distortion due to $\\ion{He}{ii}\\rightarrow\\ion{He}{i}$ recombination should have a completely different character. First, for small $n$ neutral helium has a much more complicated spectrum than hydrogenic atoms (e.g.  highly probable fine-structure transitions). In addition, the ratio of the energies for the second and first shell is $\\sim 2.1$ times higher than for hydrogenic atoms, while the energies of the highly excited levels are very close to hydrogenic. Since the transitions from the second to the first shell are controlling helium recombination, this leads to the situation that even for transitions among highly excited levels the corresponding $\\Delta n=1$-lines do not coincide with those emitted during hydrogen or $\\ion{He}{iii}\\rightarrow\\ion{He}{ii}$ recombination.  Also it will be shown below (Sect.~\\ref{sec:He.spectra}), that in the recombinational spectrum some fine-structure lines become very bright and that two of them are actually appearing in {\\it absorption}. These features lead to additional non-uniformities in the spectral variability structure of the {\\it total} CMB spectral distortion from recombination, where some of the maxima are amplified and others are diminished. This may open an unique possibility to separate the contributions of helium and hydrogen, thereby {\\it allowing to measure the pre-stellar abundance of helium in the Universe}. Until now not even one {\\it direct} method for such a measurement is known.  For the computations of the recombinational helium spectrum we are crucially dependent in the recombination history of helium and additional processes that affect the standard picture strongly. In this context, probably the most important physical mechanism is connected with the continuum absorption of the permitted 584$\\,\\AA$ and intercombinational 591$\\,\\AA$ line by a very small amount of neutral hydrogen present in ionizational equilibrium during the time of $\\ion{He}{ii}\\rightarrow\\ion{He}{i}$ recombination \\citep{Hu1995, Switzer2007I, Kholupenko2007}. \\citet{Switzer2007I} and \\citet{Kholupenko2007} recently made the first detailed analysis of this problem, and included it for the computations of the $\\ion{He}{ii}\\rightarrow\\ion{He}{i}$ recombination history, showing that the recombination of neutral helium is significantly faster. Here we reanalyse this process, and discuss in detail some physical aspects of the escape problem in the aforementioned lines.  We first consider two ``extreme'' cases for the escape problem, which can be treated analytically: (i) where line scattering leads to {\\it complete redistribution} of photons over the line profile, and (ii) where there is {\\it no redistribution}\\footnote{In this case line scattering is totally coherent in the lab frame. During recombination this is a very good approximation in the very distant wings of the line.}. Moreover, we develop an useful 1D integral approximation for the escape probability which permit us to treat any of these two cases without increasing the computation time significantly. Our final results for the escape probability in the more realistic case of {\\it partial redistribution} (or equivalently for coherent scattering in the rest frame of the atom) are based on detailed numerical computations which will be presented in a separate paper (Chluba et al. 2007, in preparation). One can then obtain a sufficiently accurate description of the real dynamics by {\\it fudging} the escape probability using the ``no redistribution'' case mentioned above with a certain function that can be obtained by comparison with the full numerical computations. Our final results for the escape probability are in very good agreement with those obtained by \\citet{Switzer2007I} (see discussion in Sect.~\\ref{sec:HIopacity}).   \\begin{figure*} \\centering \\resizebox{0.95\\hsize}{!}{\\includegraphics[angle=90]{eps/final_spectrum.ps}} \\caption{Full helium and hydrogen (bound-bound) recombination spectra. The following cases are shown: (a) the $\\ion{He}{ii}\\rightarrow \\ion{He}{i}$ recombination spectrum (black solid line), which has been obtained including up to $\\nmax=30$ shells, and considering all the J-resolved transitions up to $n=10$. In this case, there are two negative features, which are shown (in absolute value) as dotted lines; (b) the $\\ion{He}{iii}\\rightarrow\\ion{He}{ii}$ recombination spectrum (red solid line), where we include $\\nmax=100$ shells, resolving all the angular momentum sub-levels and including the effect of Doppler broadening due to scattering off free electrons; (c) the \\ion{H}{i} recombination spectrum, where we plot the result from \\citet{Chluba2007} up to $\\nmax=100$. The \\ion{H}{i} Lyman-$\\alpha$ line arising in the epoch of \\ion{He}{i} recombination is also added to the hydrogen spectrum (see the feature around $\\nu=1300$~GHz). In all three cases, the two-photon decay continuum of the $n=2$ shell was also incorporated. Feedback processes for the \\ion{He}{ii} and \\ion{He}{i} recombinations are not taken into account. Blue line shows the total recombination spectrum.  } \\label{fig:final}% \\end{figure*} Interestingly, the hydrogen continuum process leaves additional distinct trace in the cosmological recombination spectrum, because the continuum absorption of the \\HeI photons leads to significant early emission in the \\ion{H}{i} Ly-$\\alpha$ transition at $\\nu\\sim 1300\\,$GHz, well separated from the Ly-$\\alpha$ line originating during hydrogen recombination at $\\nu\\gtrsim 1500\\,$GHz, and containing about $7\\%$ of all photons that were released in the hydrogen 2p-1s transition. The amplitude and width of this feature is completely determined by the conditions under which the above process occurs.  The main result of this paper are the bound-bound spectra of \\ion{He}{ii} and \\ion{He}{i} from the epoch of cosmological recombination (see Fig.~\\ref{fig:final}). The strongest additions to the cosmological hydrogen recombination spectrum due to the presence of helium lines reach values up to 30-40\\% in several frequency bands. This strongly exceeds (roughly by a factor of four) the relative abundance ratio of helium to hydrogen, raising hopes that these distortions will be found once the recombinational lines will become observable.  It is important to note that for the computations in this paper (see Fig.~\\ref{fig:final}), we do not include the impact of feedback processes on the computed recombinational lines. Among all the possible feedback mechanisms, the most relevant for the recombinational spectrum is the pure continuum absorption (far away from the resonances) of the remaining \\HeIInt-line and \\HeILya photons. Due to this process, these photons will be finally absorbed, and the corresponding features on the final spectrum will disappear, producing additional photons that will emerge mainly through the Ly$\\alpha$ line in the hydrogen spectrum. ",
        "conclusions": "\\label{sec:con} We have presented detailed computations of the contributions to the cosmological recombination spectrum due to bound-bound transitions in primordial helium. The re-processing of \\HeILya and \\HeIInt-photons by neutral hydrogen has been taken into account, yielding a significant acceleration of $\\ion{He}{ii}\\rightarrow\\ion{He}{i}$ and hence much more narrow features than without the inclusion of this process. In addition, some hydrogen photons are released prior to the actual epoch of hydrogen recombination around $z\\sim 1100- 1500$, with distinct traces due to the hydrogen Ly-$\\alpha$ transition (see Fig.~\\ref{fig:final}).  Probably the most interesting result is the presence of two negative features in the $\\ion{He}{ii}\\rightarrow\\ion{He}{i}$ recombinational spectrum.  This is qualitatively different from any of the other spectra discussed so far ($\\ion{H}{i}$ and $\\ion{He}{ii}$). One of those negative features is associated to fine-structure transitions in neutral helium.  As illustrated in Fig.~\\ref{fig:final}, the {\\it total} cosmological recombination spectrum contains non-trivial signatures of all recombination epochs. We emphasize this fact in Figure~\\ref{fig:details}, were we present a detailed view, using linear intensity scale, of three regions in the recombination spectrum covering the low, intermediate and high frequency domain. Although the relative number of helium to hydrogen nuclei is rather small ($\\sim 8\\%$), constructive and destructive interference of the oscillatory emission patterns at low frequencies, and strong non-overlapping lines at high frequencies may provide a unique opportunity to determine some of the key cosmological parameters, and to confront our current picture of recombination with experimental evidence. Interestingly the signatures due to helium may allow a direct determination of its relative abundance, much before the first appearance of stars, and as pointed out in \\citet{RichnessBeauty}, these measurements do not suffer from limitations set by cosmic variance.  As we outlined in Sect.~\\ref{sec:dis}, several neglected processes have to be studied in connection with helium recombination, in order to obtain definite predictions, possibly with additional revisions. Nevertheless, all the results presented here strongly depend on our understanding of atomic physics {\\it and} the processes in the early Universe. Currently in particular the data for neutral helium may still not be sufficient. Here help from atomic physicist is required in order to increase the availability of more complete {\\it accurate} and {\\it user-friendly} atomic data, in particular for the photoionization cross-sections and transition rates.  All the numerical predictions for the recombinational lines obtained in this paper, which were used to produce all the figures in this paper, can be downloaded from \\verb!http://www.iac.es/galeria/jalberto/recomb!.   \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{eps/detail1.eps} \\\\ \\includegraphics[width=\\columnwidth]{eps/detail2.eps} \\includegraphics[width=\\columnwidth]{eps/detail3.eps} \\caption{Relative contribution of the \\HI, \\HeI and \\HeII bound-bound recombination spectra to the total spectrum at high- (top), intermediate- (middle) and low- frequencies (low). Helium recombination spectra (both \\HeI and \\HeII) modify the shapes of the existing hydrogen features, shift the peaks positions and introduce new features which represent changes of 30-40\\% respect to the \\HI recombination spectrum alone. } \\label{fig:details}% \\end{figure}     \\acknowledgement{ The authors thank I.~L.~Beigman and L.~A.~Vainshtein for many useful discussions on the physics of neutral helium, and R. Porter for useful discussions about the details of the paper by \\citet{Bauman2005}. We also acknowledge use of the {\\sc Cuba}-Library \\citep{Hahn2004}. }   \\begin{appendix}"
    },
    "0711/0711.2302_arXiv.txt": {
        "abstract": "Galaxy formation requires a process that continually heats gas and quenches star formation in order to reproduce the observed shape of the luminosity function of bright galaxies. To accomplish this, current models invoke heating from supernovae, and energy injection from active galactic nuclei. However, observations of radio-loud active galactic nuclei suggest that their feedback is likely to not be as efficient as required, signaling the need for additional heating processes. We propose the self-annihilation of weakly interacting massive particles that constitute dark matter as a steady source of heating. In this paper, we explore the circumstances under which this process may provide the required energy input. To do so, dark matter annihilations are incorporated into a galaxy formation model within the Millennium cosmological simulation. Energy input from self-annihilation can compensate for {\\bf all} the required gas cooling and reproduce the observed galaxy luminosity function only for what appear to be extreme values of the relevant key parameters. The key parameters are: the slope of the inner density profile of dark matter haloes and the outer spike radius. The inner density profile needs to be steepened to slopes of $-1.5$ or more and the outer spike radius needs to extend to a few tens of parsecs on galaxy scales and a kpc or so on cluster scales. If neutralinos or any thermal relic WIMP with s-wave annihilation constitute dark matter, their self-annihilation is inevitable and could provide enough power to modulate galaxy formation. Energy from self-annihilating neutralinos could be yet another piece of the feedback puzzle along with supernovae and active galactic nuclei. ",
        "introduction": "\\label{section1}  A major challenge to our current understanding of structure and galaxy formation in the Universe is the discrepancy between the theoretically predicted mass function of dark matter haloes and the observed shape of the luminosity function of galaxies (Kauffmann \\& White 1993; Cole et al. 1994; Somerville, Primack \\& Faber 2001; Benson et al. 2003). Simply put, the number density of dark matter haloes falls off as a power-law at high masses, whereas the luminosity function of galaxies occupying such haloes terminates exponentially above a characteristic luminosity (e.g. Cole et al. 2001; Huang et al. 2003). This implies that the supply of gas to a galaxy and the conversion of this gas into stars becomes preferentially inefficient in more massive systems (White \\& Rees 1978). A key issue for galaxy formation theory is thus to illuminate the physical processes that heat and cool gas in massive galaxies as this cycle regulates the formation of new stars. In addition, one of the challenges for galaxy formation is to unravel the inter-play of baryons with the ubiquitous dark matter in galaxies.  To this end, feedback processes operating in galaxies at both the low and high-mass ends of the halo mass function are required to explain the faint and bright-end slopes of the observed galaxy luminosity function. Energy input from supernovae is thought to play a significant role in the regulation of star formation in low mass galaxies (Dekel \\& Silk 1986; Mac Low \\& Ferrara 1999), however the supernova energy injected in high mass galaxies is too small to suppress gas cooling effectively. For such objects, to reconcile theory with observations a more energetic process that continually heats the gas and operates independently of star formation appears to be required (Croton et al. 2006; Bower et al. 2006).  Recent observations in nearby galaxies reveal a correlation between the masses of supermassive black holes and the velocity dispersion of the stellar component.  This suggests that black holes might play a role in regulating star formation (Magorrian et al. 1998; Tremaine et al. 2002; Ferrarese \\& Merritt 2002). Energy input from nuclear outflows driven by accreting black holes are currently favoured as the principal source of feedback driving the truncation of star formation in massive galaxies (Di Matteo et al. 2005; Croton et al. 2006; Bower et al. 2006). Unlike supernovae feedback, the observed energy output from Active Galactic Nuclei (AGN) can far exceed that liberated from cooling gas in the hot halo. AGN heating in two distinct forms has been proposed.  Episodic feedback from the AGN via outflows generated during the merger process have been proposed by Di Matteo et al. (2005); Sijacki \\& Springel (2006), the so-called `quasar-mode'. However, it is known that AGN spend most of their lifetimes in a low accretion rate state, and therefore Croton et al. (2006) and Bower et al. (2006) argue for a more steady, so called `radio-mode' feedback that is long-lived.  The details of both these processes are complex and the micro-physics is currently not well understood.  Here, we focus on an alternative heating mechanism, the energy steadily generated by the self-annihilation of dark matter particles in the inner regions of haloes.  Although the standard cosmological paradigm is predicated on the existence of non-baryonic dark matter (DM) particles, the precise nature of these particles and their interactions remain a puzzle. The neutralino is the current leading dark matter candidate. In this paper we explore if the energy supplied from the self-annihilation of dark matter in the centres of galaxies and cluster haloes could possibly play a significant role in the baryonic cooling/heating cycle (Ascasibar 2007; Totani 2004; 2005; Colafrancesco et al. 2006). Detailed astrophysical implications of neutralino dark matter annihilations in galaxy clusters, with a specific application to the Coma cluster have been calculated by Colafrancesco, Profumo \\& Ullio (2006).  They performed a thorough analysis of the transport and diffusion properties of neutralino annihilation products, and investigated the resulting multi-frequency signals, from radio to gamma-ray frequencies. They also study other relevant astrophysical effects of neutralino annihilations, like the DM-induced Sunyaev-Zel'dovich effect and the intracluster gas heating.  Annihilation luminosity is expected to be produced as a result of an enhancement in the density of the dark matter distribution in the inner-most regions of galaxies, due to the response of the dark matter to the adiabatic growth of a central black hole (Gondolo \\& Silk 1999) and/or to the adiabatic compression suffered by the dark matter due to the collapse of baryons into the dark matter potential well (Blumenthal et al. 1986; Ryden \\& Gunn 1987; Gnedin et al. 2004). By including this steady heating source of annihilating neutralinos to suppress cooling flow gas in a sophisticated model of galaxy formation, we explore the effect of various key parameters in reproducing the bright-end of the observed galaxy luminosity function. The question we address is whether self-annihilation can provide an alternative mechanism to AGN driven feedback within the context of the current paradigm.  The outline of this paper is as follows. In Section~\\ref{section2} we discuss the physics of neutralino annihilation; a scenario for the coupling of this energy source to baryons in galactic nuclei is explored in Section~\\ref{section3}. With this phenomenology developed, we incorporate this feedback scheme into a galaxy formation model. The details of the N-body simulation and galaxy formation model used in this implementation are discussed in Section~\\ref{section4}. Our results are presented in detail in Section~\\ref{section5} and summarized in Section~6. We conclude with a discussion of the implications of our results and comparison of annihilation heating with other modes of feedback. Unless otherwise stated we assume Hubble constant parametrised as $H_0=h\\,100\\,{\\rm km}\\,{\\rm s}^{-1}\\,{\\rm Mpc}^{-1}$ with $h=0.73$. ",
        "conclusions": "\\label{section6}  Using the Millennium Run N-body simulation coupled with a sophisticated model of galaxy formation that includes the heating of cooling flow gas through neutralino annihilation, we have shown that:  \\begin{itemize} \\item Density spikes that support the annihilation flux at the levels required to offset cooling flows are stable enough over long enough time-scales to maintain a reasonably constant heating source (Fig.~\\ref{fig:merger_rate}).  \\item Models that appear to be extreme at the present time (given our current understanding of DM density profiles) are required to produce enough heating flux to offset the predicted cooling rates. To obtain the required heating rates we either need to steepen the inner DM density slope to values $\\gamma\\,>\\,1.5$ or increase the outer spike radius. For galaxy sized haloes the outer spike radius is required to be of the order of $\\sim\\,25\\,r_{\\rm bh}$, and for cluster haloes $\\sim\\,60\\,r_{\\rm bh}$ (Fig.~\\ref{fig:heating} and \\ref{fig:LFs1}).  \\item The efficiency of heating in the DM annihilation model scales with halo mass (and circular velocity), therefore this mechanism does provide preferential suppression of star formation in more massive haloes as required to explain current observations of the luminosity function (Fig.~\\ref{fig:heating}). \\end{itemize}  In this treatment, we have assumed that the mass of central black holes grows significantly at each merger (as the solution for adiabatic growth in valid only in the limit $M_{\\rm bh, initial}\\,<<\\,M_{\\rm bh, final}$. We note here that an alternative source of annihilations may be provided by mini-spikes around inter-mediate mass black holes as suggested by Bertone, Zentner \\& Silk (2005).  It is clear that feedback and energy injected into the inter-stellar medium of galaxies and the intra-cluster medium is a complex process, and that a combination of astrophysical processes, including the one explored here, are likely at play. One of the key uncertainties in the model explored in this paper arises from the fact that we lack an understanding of the physics through which the annihilation flux is expected to couple with the cooling hot halo gas. While we have discussed some possibilities, like coulomb collisions, brehmstrahlung and synchrotron radiation.  Colafrancesco et al. (2007) have explored these in more detail for the case of the heating of gas in the Coma cluster due to DM annihilations. Experimental confirmation of super-symmetery from the LHC at CERN might throw new light on the viability and likely couplings for the neutralino. Additionally while following the cumulative heating history is very challenging to do it is needed to really understand the detailed energetics of the gas. There is incontrovertible evidence for the presence of copious amounts of DM on all scales in the Universe, so DM annihilation is an inescapable phenomenon. However, how much energy is released in the process and how efficiently it couples to the baryonic component are unclear."
    },
    "0711/0711.1137_arXiv.txt": {
        "abstract": "Understanding the thermal equilibrium (stability) curve may offer insights into the nature of the warm absorbers often found in active galactic nuclei. Its shape is determined by factors like the spectrum of the ionizing continuum and the chemical composition of the gas. We find that the stability curves obtained under the same set of the above mentioned physical factors, but using recently derived dielectronic recombination rates, give significantly different results, especially in the regions corresponding to warm absorbers, leading to different physical predictions. Using the current rates we find a larger probability of having thermally stable warm absorber at $10^5 \\kel$ than previous predictions and also a greater possibility for its multiphase nature. the results obtained with the current dielectronic recombination rate coefficients are more reliable because the warm absorber models along the stability curve have computed coefficient values, whereas previous calculations relied on guessed averages for the same due to lack of available data. ",
        "introduction": "\\labsecn{introduction}  Warm Absorbers are highly photoionised gas found along our line of sight to the continuum source of active galactic nuclei (AGN). Their signatures are a wealth of absorption lines and edges from highly ionized species, notably OVII, OVIII, FeXVII, NeX, CV and CVI in the soft X-ray (0.3-1.5 keV) spectra. The typical column density observed for the gas is $\\nh \\sim 10^{22\\pm1}\\cmsqi$ and the temperature $T$ is estimated to be a few times $10^5\\kel$. For many objects the warm absorber exists as a multiphase absorbing medium with all the phases in near pressure equilibrium (Morales, Fabian \\& Reynolds, 2000; Collinge \\etal 2001; Kaastra \\etal 2002; Krongold \\etal 2003; Netzer \\etal 2003; Krongold \\etal 2005; Ashton \\etal 2006).  Any stable photoionised gas will lie on the thermal equilibrium curve or `stability' curve, in the temperature - pressure phase space, where heating balances cooling; this curve is often used to study the multiphase nature of the photoionised gas. The equilibrium depends on the shape of the ionizing continuum and the chemical abundance of the gas (Krolik, McKee \\& Tarter, 1981; Krolik \\& Kriss, 2001; Reynolds \\& Fabian, 1995 and Komossa \\& Mathur, 2001). We are investigating this dependences in details and will report elsewhere (Chakravorty \\etal, to be submitted).  Over the past two decades the estimates of dielectronic recombination rate coefficients have improved. In this letter we show that these have affected the stability curves significantly which may lead to quite different physical models for the warm absorber gas. We conclude with a caution on the reliability of earlier results. ",
        "conclusions": "\\labsecn{conclusion}  We have shown that stability curves for warm absorbers in AGN generated by two versions of the ionization code CLOUDY, C84 and C07, for the same physical conditions, are substantially different in shape, leading to different conclusions regarding the nature of the warm absorber. The differences in the results of the photoionisation calculations arise due to major changes in the dielectronic recombination rate coefficient (DRRC) data bases which have taken place over the last decade. The modern version C07 includes reliable computed DRRC values for many more ions than C84 for the entire part of the stability curve relevant for warm absorbers and does not rely on guessed average values. Thus, the physical nature of the warm absorber predicted by modern calculations are more reliable. We suggest that past calculations for photoionised plasma in thermal equilibrium should be reconsidered, taking into account these changes. Since, any other atomic physics code (eg. XSTAR), in popular use for warm absorber studies, is also likely to be affected by the changes in the DRRC data base, caution should be exercised while using or comparing results from them as well."
    },
    "0711/0711.3903_arXiv.txt": {
        "abstract": "Coherent properties of the baryon-photon fluid decoupling are considered in the terms of an effective nonlinear Schr\\\"{o}dinger equation for a macroscopic wave function that specifies the index of the coherent state. Generation of a transitional acoustic turbulence preceding formation of large-scale condensate in the plasma and its influence on the CMB power spectrum has been studied. A scaling $k^{-1}$ law is derived for the CMB Doppler spectrum $E(k)$ (angle-averaged) in the {\\it wavenumber} space, for sufficiently large wavenumber $k$ and for the weak nonlinear and completely disordered initial conditions. Using the recent WMAP data it is shown that the so-called first acoustic peak represents (in a compensated spectral form) a pre-condensate fraction of the spectrum $E(k)$ at a rather advance stage of the condensate formation process. ",
        "introduction": "The recent high resolution cosmic microwave background (CMB) radiation measurements (see, for instance \\cite{wmap}) provide a possibility for quantitative investigation of the nontrivial processes taking place at the {\\it decoupling} of the baryon-photon plasma (fluid) in the early universe. Before the decoupling the CMB photons were tightly coupled to the baryons by photon scattering (mainly from the free electrons) forming so-called baryon-photon fluid. At a certain stage of the early universe expansion the energy of the CMB photons becomes insufficient to keep hydrogen ionized and a recombination process becomes dominating in the baryon-photon fluid. The recombination is a gradual process, but it serves as a trigger and a catalyst of the decoupling of the baryon-photon fluid. Before the decoupling the baryons and photons were tightly coupled whereas after the decoupling the photons are effectively free. Physical processes at the decoupling itself can be described on different levels. In present paper we will concentrate on the {\\it coherent} properties of the decoupling. It is well known that the main energy of the plasma fluctuations at a certain (advance) stage of the decoupling process is concentrated in the coherent states (the acoustic coherent oscillations - stochastic standing waves). Originally the acoustic oscillations become phase-coherent when the causal horizon overtakes their wavelength. Around the last scattering surface the radiation density and ionization are high enough that the photon drag on the baryons can be still significant. Therefore, the baryons are involved in the coherent motion by the photon drag, which differentially accelerates the ions and electrons. The resulting significant electron-ion drift generates large-scale {\\it coherent} electric currents and fields, which rather quickly become comparable to the photon drag in their effect on the electrons (see, for instance, \\cite{bd},\\cite{ho}). While we can readily introduce the photon (Thomson) drag in the dynamical plasma equations it is rather nontrivial task to describe the coherent (collective) phenomena at the decoupling, which are crucial for the entire process. To tackle this problem we suggest to use an effective nonlinear Schr\\\"{o}dinger equation for macroscopic wave function, which is a complex-valued classical field that specifies the index of the coherent state. To the leading order each coherent state evolves along its 'classical' trajectory which is given by the corresponding nonlinear Schr\\\"{o}dinger equation. The idea is based on the Ehrenfest's theorem stating that the the expectation values of displacement and momentum given by the Schr\\\"{o}dinger equation obey time evolution equations which are analogous to 'classical' ones in the the Madelung representation. In this representation the macroscopic wave function of the free (non recombined) electrons can be defined as $\\psi = \\sqrt{n_e} \\exp[{\\rm i}S({\\bf x},t)/h_c]$, where $n_e$ is the electron number density, $\\nabla S({\\bf x},t)=m {\\bf v} ({\\bf x},t)$ with $m$ and ${\\bf v} ({\\bf x},t)$ as the mass and velocity of the electrons, $h_c$ here is not the Plank constant but a parameter characterizing coherence in the system (real quantum effects are negligible in the considered situation, cf recent papers \\cite{gri},\\cite{hab} and references therein). Then, momentum equation for the electron can be considered as a nonlinear Schr\\\"{o}dinger equation for the macroscopic wave function $\\psi ({\\bf x},t)$: $$ {\\rm i}h_c \\pd \\psi t = -\\left[ {\\frac{h_c^2}{2m}} \\nabla^2  +\\eta ({\\bf x}) - U_{eff}(|\\psi|^2) \\right] \\psi,   \\eqno{(1)} $$ where $\\eta ({\\bf x})$ is a real-valued potential, $U_{eff}(|\\psi|^2)$ is an effective potential represented by a real-valued function of the free electrons density $n_e=|\\psi|^2$. Analogous approach at rather different circumstances was used for a mean field description of dense Fermi plasmas (see, for instance, Refs. \\cite{m},\\cite{se},\\cite{ss}). Unlike the quantum plasma considered in Refs. \\cite{m},\\cite{se},\\cite{ss} the original (input) coherence at the decoupling comes from an external source (cf Section VI). At this description the addition of the effective potential $U_{eff}(|\\psi|^2)$ to the equation (1) is possible due to a specific gauge invariance of the equation. Namely, it can be readily shown that addition of the potential $U_{eff}(|\\psi|^2)$ (satisfying condition $\\lim_{n \\rightarrow 0} \\int U_{eff} (n) dn=0$) does not affect the 'classic' limit (i.e. the Ehrenfest's theorem) $$ m d\\langle {\\bf r}\\rangle/dt= \\langle {\\bf p} \\rangle,~~~d\\langle {\\bf p} \\rangle/ dt = - \\langle \\nabla \\eta \\rangle  \\eqno{(2)} $$ where ${\\bf p}= m {\\bf v}$ is momentum. The effective potential can be also directly depending on ${\\bf x}$, but the potential $U_{eff}({\\bf x},|\\psi|^2)$ should be slowly varying on ${\\bf x}$ (in the adiabatic terms of the Ehrenfest's theorem). About a possibility of nonlocal dependence of $U_{eff}(|\\psi|^2)$ on $|\\psi|^2$ see Section VI. It will be shown below that such type of invariance of the Schr\\\"{o}dinger equation is crucial for the condensation process.  We deliberately have the dissipation term not included into the equation because we will concentrate on the large- and intermediate-scale processes, while the dissipation term is usually related to the small scales (we will return to this problem below). The first term in the right-hand side of Eq. (1) provides, in particular, a possibility of tunneling at the decoupling. On the other hand, it is believed that the nonlinear Schr\\\"{o}dinger equation can describe development of the coherent regimes \\cite{ks}-\\cite{bs} at a certain stage of evolution. Theoretical considerations (see for instance, \\cite{ks},\\cite{kss}) show that this regime can appear in a long-wavelength region of wavenumber space after the breakdown of the regime of weak turbulence: the Kagan-Svistunov (KS-) scenario. This is supported by recent numerical simulations for the weakly interacting Bose gas \\cite{bs} and for the dense Fermi plasmas \\cite{ss}. It is expected that there exist three different stages in the time development for the KS-scenario: weak turbulence, transitional turbulence and condensate.  A pre-condensate begins its formation in a long-wavelength region at the transitional turbulent regime. This transitional turbulent regime is the most difficult for theoretical description \\cite{bs}. There is a problem to define a continuous vorticity filed in the quantum-like turbulence due to non-rotational nature of the velocity field defined on the wave function given by the Schr\\\"{o}dinger equation. Velocity field defined on the wave functions is a potential field. Vorticity of such field vanishes everywhere in a single-connected region. It is believed that in the condensate state itself all rotational flow is carried by quantized vortices (the circulation of the velocity around the core of such vortex is quantized). This idea turned out to be very fruitful, and recent experiments give direct support for existence of such vortices in the condensate \\cite{Donnelly},\\cite{vn},\\cite{sreeni}. However, for the transitional turbulence apparent absence of a continuous field associated with vorticity is a difficult problem. Moreover, it is clear that the only flux of the particles from the region of lager energies toward the future condensate in not sufficient for transformation of the completely disordered initial weak turbulence into the condensate with its local superfluid order and with the tangles of quantized vortices. A certain vorticity-like quantity have to be involved in this process just before appearance of the tangles of quantized vortices. In present paper we study a generalized vorticity defined on the {\\it weighted} velocity field. The generalized vorticity is not vanishing in the bulk of the flow. Then, we show that adiabatic invariance of an enstrophy (mean squared generalized vorticity) controls the transitional turbulence just before formation of the condensate with its tangles of quantized vortices. Since in the above described scenario the transitional turbulence already corresponds to a coherent regime (in this regime the phases of the complex amplitudes of the field $\\psi$ become strongly correlated) we will speculate that an advance stage of the decoupling can be associated with the transitional turbulence regime of nonlinear Schr\\\"{o}dinger equation (1) . Them we will consider significance of the generalized vorticity for the CMB photons visibility function, which is the main CMB characteristic of the decoupling stage. This scenario provokes an additional speculation: to associate the condensate regime itself (with its tangles of quantized vortices - the topological defects) with the universe state just after the decoupling. The recombination process considerably reduces the number of the free baryons before the last stage of the decoupling but a sufficient number of the free baryons can still exist at this stage in order to form a condensate due to the {\\it fast} character of the condensation process (Section VII and Fig. 5). The far reaching consequences of such additional speculation are quiet obvious, cf for instance \\cite{gam},\\cite{pietr},\\cite{bp}. ",
        "conclusions": "Different sources of turbulence in the early universe in the times preceding the decoupling were discussed in the literature, from the very early times \\cite{g} up to the recombination time \\cite{ber2}-\\cite{beo}. In present paper we have discussed a possibility that the decoupling process itself can generate a turbulence related to a condensation process in the primordial plasmas. Just this 'later' turbulence has a preferable chance to imprint itself both on the CMB and on the further dynamics of the luminous matter. Especially significant is the question: whether the condensate itself with its tangles of quantized vortices appeared at the end of decoupling or the decoupling was finished at the transitional stage of the condensation process (see the end of the Section I) . The apparent flatness of the pre-condensate fraction of spectrum $E(k)$ in Fig. 5 (in the range: $k_cR > kR > 30$) indicates a rather advance stage of the condensation process in this case. One can even speculate that the end of the transition from the pre-condensate to the condensate is the effective end of the decoupling (and just this is imprinted in the CMB spectrum shown in Fig. 5). But the question is too serious to make definite conclusion from this observation only.   On the other hand, the recombination and dissipation processes could effect the decoupling turbulence. However, the coherent nature of this comparatively large-scale turbulence makes it stable to these effects. First of all, the recombination obviously is much slower process then the particles (enstrophy) exchange between the pre- and above-condensate fractions. Therefore, the total number of the particles $N$ can be considered as an adiabatic invariant at the decoupling and the consideration of the section III can be preserved when one takes into account the recombination. The linear dissipation can be treated by the same way (considering $N$ as an adiabatic invariant). \\\\  The author acknowledges the use of the Legacy Archive for Microwave Background Data Analysis (LAMBDA) \\cite{wmap}."
    },
    "0711/0711.3418_arXiv.txt": {
        "abstract": "{} {We present a newly developed code that allows simulations of optical observations of galaxy fields with a variety of instruments. The code incorporates gravitational lensing effects and is targetted at simulating lensing by galaxy clusters. Our goal is to create the tools required for comparing theoretical expectations with observations to obtain a better understanding of how observational noise affects lensing applications such as mass estimates, studies on the internal properties of galaxy clusters and arc statistics.}  {Starting from a set of input parameters, characterizing both the instruments and the observational conditions, the simulator provides a virtual observation of a patch of the sky. It includes several sources of noise such as photon-noise, sky background, seeing, and instrumental noise. Ray-tracing through simulated mass distributions accounts for gravitational lensing. Source morphologies are realistically simulated based on shapelet decompositions of galaxy images retrieved from the GOODS-ACS archive. According to their morphological class, spectral-energy-distributions are assigned to the source galaxies in order to reproduce observations of each galaxy in arbitrary photometric bands.} {We illustrate our techniques showing virtual observations of a galaxy-cluster core as it would be observed with the space telescope DUNE, which was recently proposed to ESA within its \"Cosmic vision\" programme. We analyze the simulated images using methods applicable to real observations and measure the properties of gravitational arcs. In particular, we focus on the determination of their length, width and curvature radius.} {We find that arc properties strongly depend on several properties of the sources. In particular, our results show that compact, faint or low surface brightness galaxies that are barely detectable are more easily distorted as arcs with large length-to-width ratios. We conclude that realistic lensing simulations can be obtained with the method proposed here. They will be essential for evaluating and improving analysis techniques currently used for cosmological interpretations of cluster lensing.} ",
        "introduction": "Gravitational lensing of distant galaxies is one of the most powerful tools for investigating the content of the universe and for understanding how the matter is distributed within the cosmic structures.   Being the most massive objects in the universe, galaxy clusters are the most powerful existing gravitational lenses. This makes them suitable for a large number of purposes.  Clusters are able to magnify distant galaxies. Thus, they are natural gravitational telescopes, that can be used to investigate the first sources of light in the universe \\citep{EL01.1,SCHA03.1,KN04.1,LI07.1}.  Strong lensing features, like gravitational arcs and multiple images are observed in the central, most dense regions of the lenses \\citep{FO88.1,LY89.1}. Hence, they can be used to constrain the inner structure of clusters \\cite[e.g][]{LI07.1,CO05.1,OG03.1,SA03.1,ME07.1,ME07.2}.  Gravitational arcs are also used in a variety of cosmological applications.  Their frequency on the sky reflects the abundance, the concentration and the redshift distribution of massive lenses in the universe. Moreover, it has been shown that the probability of forming arcs is enhanced during cluster mergers \\citep{TO04.1,FE05.1}. Thus, arc statistics are a powerful tool for tracing structure formation and for constraining the cosmological parameters \\citep{BA98.2,LI05.1,ME05.2}.  Moreover, since they form along critical lines, whose size scales with the source and lens redshifts, lensed images arising from multiple sources at different redshifts can be used to constrain the geometry of the universe \\citep{GO02.2,ME07.2}, and thus also cosmological parameters.  In the outer regions of clusters, the distortions induced by the cluster potentials are weaker. Nevertheless, using suitable algorithms \\cite[e.g.][and references therein]{BA01.1}, the distortion field can be inverted to obtain two dimensional surface density maps of the the lenses \\cite[see][for some recent examples]{CL04.1,ST06.1,PH07.1}. Strong and weak lensing can also be used jointly to improve the mass reconstrution in a non-parametric way \\citep{BR04.1,CA06.1,DI07.1}.  Although the potential power of gravitational lensing for the above mentioned applications has been already demonstrated, different implementations of methods adopted for extracting the desired information from the lensing data often give contradictory results. An example is given by the so called \"arc statistics problem\", i.e. the claimed inconsistency between the observed and the predicted number of highly distorted gravitational arcs expected to be seen on the sky in the standard $\\Lambda$CDM cosmology \\citep{BA98.2,DA04.1,WA03.1,LU99.1,GL03.1,ZA03.1,LI05.1}. Moreover, testing a particular analysis technique and/or comparing the results to the theoretical expectations is difficult because observational noise introduces uncertainties that need to be properly modeled \\citep[see e.g.][]{RH07.1}.   Following an approach that turned out to be extremely successful in cosmic shear studies \\cite[see e.g][]{HEY06.1,MAS07.2}, we aim at overcoming these difficulties and facilitating the comparison between theory and observations by performing realistic simulations, including the relevant observational effects. Using suitable ray-tracing codes, the gravitational lensing produced by realistic mass distributions obtained from numerical N-body and hydrodynamical simulations can be included. Simulated images can finally be analyzed like real observations.  In this paper, we describe a simulation pipeline that we will use to simulate observations of lensing events in galaxy clusters, but which can easily be used also for simulating lensing by large scale structures in wide fields. These simulations will be used in forthcoming papers to evaluate the accuracy of several lensing techniques mentioned above.  The plan of the paper is as follows. In Sect.~\\ref{sect:simulator} we present the architecture of the simulator and describe the methods used to include lensing effects and observational noise. In Sect.~\\ref{sect:numtest}, we use a galaxy cluster obtained from N-body simulations and use our code to create mock optical observations of its central regions. We discuss how several strong lensing features can be extracted from the simulated images and how the properties of gravitational arcs can be measured. We also quantify how these properties depend on the assumed source models. Finally, in Sect~\\ref{sect:conclu} we report our conclusions. ",
        "conclusions": "\\label{sect:conclu}  We presented a new code for mimicking optical observations of distant galaxies. It includes all relevant sources of noise affecting real observations, such as sky background, photon noise and effects of the atmosphere and of the instrumental PSF. The code also incorporates ray-tracing routines that allow us to include lensing effects by matter distributed between the observer and the sources.  In this work we have detailed the thorough method that we have developed in order to model the source galaxies realistically. The source morphologies are modeled via the shapelet decomposition of real galaxy images observed with HST. Four spectral templates are used to simulate the spectral energy distributions of the galaxies. Their redshift and luminosity distributions as a function of their spectral classification are derived from observed galaxy luminosity functions in the VVDS.  Being particularly suitable for simulating lensing by galaxy clusters, the code also simulates the emission from the cluster galaxy population, whose properties are obtained from semi-analytic models of galaxy formation and evolution.  Many features of the simulator have been discussed in detail. In particular, we have shown that the code can be easily used to mimic observations with a variety of existing instruments, both from the space (HST) and from the ground (LBT).  In order to illustrate the code and to introduce some potential applications, we have simulated observations of a galaxy cluster obtained from N-body simulations. These virtual observations were made using a telescope with the characteristics of the proposed space telescope DUNE.  The simulated images were analyzed using standard observational techniques. After subtracting the light of the foreground galaxies, multiple images of several distant sources that were strongly lensed by the foreground cluster can be identified. Each detected arc-like image has been classified, measuring its length, width and curvature radius.  We discussed how the properties of arcs can be measured consistently in simulations and observations, in order to facilitate the comparison with theoretical predictions.  In particular, we focused on the determination of the arc widths and on how they can be corrected for the PSF broadening. For this purpose, we fit the radial profile of the arcs with multiple Gaussians. By comparing simulations with and without PSF convolution, we verified that the de-convolution technique allows one to determine the true width of the arcs with a typical error of one pixel.  By varying the characteristics of the sources, we studied the detectability limits of arcs in DUNE observations. With an exposure time of $1500$ seconds, we expect that DUNE will be able to observe arcs arising from sources whose magnitudes are fainter than $\\sim 27$ in the $riz$ band. The shape of the arcs are found to be very sensitive to the properties of the sources. In particular, we found that arcs tend to acquire larger length-to-width ratios as their sources become fainter or more compact.  These results are particularly intriguing for arc statistics. Indeed, they show that observational effects may have a large impact on the abundance of arcs with large length-to width ratios, and they deserve careful investigation. We will dedicate a forthcoming paper to this subject (Meneghetti et al., in prep.). The source morphologies, and in particular sub-structures of star-forming regions, are clearly more relevant for arc statistics than is usually assumed.  This is as a result of the small angular size of the bright star forming regions which are more easily distorted. Moreover, the number of very elongated images is expected to grow rapidly as a function of the depth of the observations. Thus, space missions like DUNE, capable of making deep and wide surveys thanks to their large sensitivity and field of view, are likely to discover a significant number of new gravitational arcs. The usage of efficient software for automatic arc identification \\citep[see e.g.][]{LE05.1,SE07.1,AL06.1,HO05.1,CA07.1} will facilitate the detection of strong lensing events in these large data sets."
    },
    "0711/0711.1889_arXiv.txt": {
        "abstract": "% Our presentation is aimed at data centers providing access to complex observational and theoretical data and to the users of these resources. We show how to visualize complex datasets stored in the VO enabled data archives using existing VO client software and PLASTIC, a prototype of an application messaging protocol, for interaction between archive query results and tools. We demonstrate how to display and explore observable IFU datasets, provided within the ASPID-SR archive, using CDS Aladin, ESA VOSpec, and VO-Paris Euro3D Client. In the second part of the paper we show how to use TOPCAT for displaying results of N-body simulations of galaxy mergers available in the HORIZON GalMer database. ",
        "introduction": "At present the International Virtual Observatory has become a rapidly growing initiative. Recently, several VO resources providing access to complex observational and theoretical datasets have appeared. Providing the transparent and efficient data access and visualization mechanisms are the crucial points for data sources to be used by the scientific community.  In this paper we demonstrate how to visualize complex observable and theoretical datasets stored in the VO-enabled data archives using a WEB-browser, existing VO client software and PLASTIC (PLatform for Astronomical Tools Inter-Connection), a prototype of an application messaging protocol, for interaction between archive query results and tools. The technical details of the middle layer software implementation are given in a paper ``Middleware for data visualization in VO-enabled data archives'' by Zolotukhin \\& Chilingarian (this volume). ",
        "conclusions": "The two implementations described above demonstrate that the observable and theoretical datasets having complex structure can be discovered and accessed at the present stage of the Virtual Observatory development, when not all the interoperability standards are yet established. A WEB-browser and existing client applications interacting via simple application messaging protocol such as PLASTIC provide an infrastructure powerful enough for scientific usage of the data sources in a VO framework."
    },
    "0711/0711.4181_arXiv.txt": {
        "abstract": "The cosmological principle, promoting the view that the universe is homogeneous and isotropic, is embodied within the mathematical structure of the Robertson-Walker (RW) metric. The equations derived from an application of this metric to the Einstein Field Equations describe the expansion of the universe in terms of comoving coordinates, from which physical distances may be derived using a time-dependent expansion factor. These coordinates, however, do not explicitly reveal properties of the cosmic spacetime manifested in Birkhoff's theorem and its corollary. In this paper, we compare two forms of the metric---written in (the traditional) comoving coordinates, and a set of observer-dependent coordinates---first for the well-known de Sitter universe containing only dark energy, and then for a newly derived form of the RW metric, for a universe with dark energy and matter. We show that Rindler's event horizon---evident in the co-moving system---coincides with what one might call the ``curvature horizon\" appearing in the observer-dependent frame. The advantage of this dual prescription of the cosmic spacetime is that with the latest WMAP results, we now have a much better determination of the universe's mass-energy content, which permits us to calculate this curvature with unprecedented accuracy. We use it here to demonstrate that our observations have probed the limit beyond which the cosmic curvature prevents any signal from having ever reached us. In the case of de Sitter, where the mass-energy density is a constant, this limit is fixed for all time. For a universe with a changing density, this horizon expands until de Sitter is reached asymptotically, and then it too ceases to change. ",
        "introduction": "The smoothness of the cosmic microwave background (CMB) provides direct evidence that the universe is isotropic about us. Invoking the Copernican principle---that we do not live in a special place---allows one to argue further that the universe should be isotropic around every point, i.e., that it is also (on average) homogeneous. Standard cosmology is therefore based on the Robertson-Walker (RW) metric for a spatially homogeneous and isotropic three-dimensional space, expanding or contracting as a function of time: $$ ds^2=c^2\\,dt^2-a^2(t)[dr^2(1-kr^2)^{-1}+ $$ $$ r^2(d\\theta^2+\\sin^2\\theta\\,d\\phi^2)]\\;. \\eqno(1) $$ \\vskip 0.1in\\noindent The coordinates for this metric have been chosen so that $t$ represents the proper time as measured by a comoving observer, $a(t)$ is the expansion factor, and $r$ is an appropriately scaled radial coordinate in the comoving frame. The geometric factor $k$ is $+1$ for a closed universe, $0$ for a flat, open universe, or $-1$ for an open universe.  Assuming that general relativity (GR) provides the correct framework for interpreting cosmological dynamics, one may then apply the RW metric to Einstein's Field Equations (EFE) to obtain the (FRW) differential equations of motion. These are the Friedman equation, written as $$ H^2\\equiv\\left({\\dot a\\over a}\\right)^2={8\\pi G\\over 3c^2}\\rho-{kc^2\\over a^2}\\;, \\eqno(2) $$ and the ``acceleration\" equation, $$ {\\ddot a\\over a}=-{4\\pi G\\over 3c^2}(\\rho+3p)\\;.\\eqno(3) $$ An overdot denotes a derivative with respect to cosmic time $t$, and $\\rho$ and $p$ represent the total energy density and total pressure, respectively. A further application of the RW metric to the energy conservation equation in GR yields the final equation, $$ \\dot\\rho=-3H(\\rho+p)\\eqno(4) $$ which, however, is not independent of equations (2) and (3).  But written in this form, equations (2), (3), and (4) do not explicitly reveal the spacetime curvature most elegantly inferred from the \\emph{corollary} to Birkhoff's theorem (Birkhoff 1923). This theorem states that in a spherically symmetric spacetime, the only solution to the Einstein equations in vacuum is the Schwarzschild exterior solution. Furthermore, a spherically symmetric vacuum solution in the exterior is necessarily static. Birkhoff's purpose was to prove that for general relativity, as with the Newtonian theory, the exterior gravitational field of a spherically symmetric distribution of matter is independent of interior radial pulsations.  However, what is relevant to our discussion here is not so much the Birkhoff theorem per se, but rather its very important corollary. The latter is a generalization of a well-known result of Newtonian theory (which also finds a direct application in electrodynamics), in which the gravitational field of a spherical shell vanishes inside the shell. The corollary to Birkhoff's theorem states that the metric inside an empty spherical cavity, at the center of a spherically symmetric system, must be equivalent to the flat-space Minkowski metric. Space must be flat in a spherical cavity even if the system is infinite. It matters not what the constituents of the medium outside the cavity are, as long as the medium is spherically symmetric.  If one then imagines placing a spherically symmetric mass at the center of this cavity, according to Birkhoff's theorem and its corollary, the metric between this mass and the edge of the cavity is necessarily of the Schwarzschild type. Thus the worldlines linked to an observer in this region are curved relative to the center of the cavity in a manner determined solely by the mass we have placed there. This situation may appear to contradict our assumption of isotropy, which one might naively take to mean that the spacetime curvature within the medium should cancel since the observer sees mass-energy equally distributed in all directions. In fact, the observer's worldlines are curved in every direction because, according to the corollary to Birkhoff's theorem, only the mass-energy between any given pair of points in this medium affects the path linking those points. This is why, of course, the universe does not expand at a constant rate, due to the spacetime curvature induced (in every direction) by its internal constituents.  This consequence of the corollary to Birkhoff's theorem is so important---and critical to the discussion in this paper---that it merits re-statement: the spacetime curvature of a wordline linking any point in the universe to an observer a distance $R$ away may be determined by calculating the mass-energy enclosed within a sphere of radius $R$ centered at the origin (i.e., at the location of the observer). The mass-energy outside of this volume has a net zero effect on observations made within the sphere.  In this paper, we will introduce a new set of coordinates (distinct from the co-moving coordinates $[r,t]$), that elicit this effect directly from the transformed metric, originally appearing in equation (1). We will show that, with the recent WMAP results (Spergel et al. 2003), our observational limit clearly corresponds to the distance beyond which the spacetime curvature prevents any signal from ever reaching us. ",
        "conclusions": "The cosmic horizon we have discussed in this paper coincides with the event horizon derived earlier by Rindler (1956) because they represent aspects of the same spacetime, though viewed in two different coordinates systems. Rindler formalized the definition of two horizons in cosmology, both functioning as hypersufaces in spacetime that divide ``things\" into two separate non-null classes. For a fundamental observer $A$, the {\\it event} horizon is defined to be a hypersurface that divides all events into the class in which these have been, are, or will be, observable by $A$, and the complementary class in which the events are forever outside of $A$'s power of observation. These event horizons are created by the universe's expansion, which produces distant regions receding from $A$ at such speeds that no causal contact can ever occur between them and the observer. Since the comoving coordinates represent the universe's ``freely falling\" frame, it is clear that this critical condition must therefore correspond to the radius $R_0$ at which the curvature in the cosmic spacetime causes $T$ to diverge.  Rindler's second type of horizon is called a {\\it particle} horizon. For the fundamental observer $A$ and cosmic instant $t_0$, the particle horizon is a surface in instantaneous 3-space at time $t=t_0$ that divides all fundamental particles into two classes: those that have already been observed by $A$ at time $t_0$, and those that have not. These horizons are therefore created by the finite propagation speed of signals that couple particles at distance $d$ to the observer $A$ for times $t>d/c$. The cosmic horizon $R_0$ is not related to the particle horizon.  An interesting property of $R_0$ emerges upon closer scrutiny of its value in the case of de Sitter and the alternate cosmology containing both matter and dark energy. First of all, the density $\\rho$ is constant in the former, so $R_0$ does not change with time. Thus, the horizon in de Sitter is fixed forever, and no events occurring beyond it can ever be in causal contact with the observer at the origin. This is presumably the situation that emerges as the universe approaches de Sitter asymptotically, in which limit $R_0$ must be calculated for $\\rho$ due to dark energy alone. Second, the density $\\rho$ in the alternate metric decreases with time and therefore $R_0$ correspondingly increases---i.e., $R_0(t)$ is a function of $t=T(0)$. This means, of course, that some portions of the universe that produced effects at $R<R_0(t)$ we measure at time $t$ have by now moved beyond the limiting radius. However, the effects of gravity travel at the speed of light, so what matters in setting the structure of the universe within the horizon at time $t$ is the mass-energy content within $R_0$. The influence of these distant regions of the universe ended once their radius from us exceeded $R_0$.  The light-travel time distance ($\\sim$13.7 billion light-years) could only have been identified with the cosmic horizon once the data (Spergel et al. 2003) confirmed that the universe is flat, and revealed that the density $\\rho$ is close to, or at, its ``critical\" value. It has never been quite clear why this critical condition exists, except for the possibility that inflation in the early universe could have led to this. But based on our analysis of the various metrics in this paper, an alternative interpretation---or perhaps simply an additional reason---for this criticality is that the universe has now expanded long enough for the condition $R_0\\approx ct$ to have been attained.  One can easily show from the FRW equations that once $R_0$ has been reached, the expansion thereafter proceeds at a constant rate. But this does not necessarily mean that this dynamical state must have been present since the Big Bang, so it is by no means trivial to see how such a scenario would affect Big Bang nucleosynthesis (Burles et al. 2001), and structure formation in the early universe (Springel et al. 2006).  But for the current universe, the near-equality $R_0\\approx ct$ would produce an observable signature because it is equivalent to the condition $\\ddot a\\ge 0$. At least out to a redshift of $\\sim 1.8$, the universe appears to be expanding with a slight positive acceleration (Riess et al. 2004). With ongoing observations of Type Ia supernovae, the value of $\\ddot a$ will continue to be refined and, with it, so too the value of $R_0-ct$.  Clearly, more work needs to be done. There is little doubt that a cosmic horizon exists. It is required by the application of the corollary to Birkhoff's theorem to an infinite, homogeneous medium, and there is some evidence that we have already observed phenomena close to it. However, it may be that observational cosmology is not entirely consistent with the condition $R_0\\approx ct$ in the current epoc. If not, there must be some other reason for this apparent coincidence. Perhaps the assumption of an infinite, homogeneous universe is incorrect. Whatever the case may be, the answer could be even more interesting than the one we have explored here."
    },
    "0711/0711.4979_arXiv.txt": {
        "abstract": "{The future new-generation radio telescope SKA (Square Kilometre Array) and its precursors will provide a rapidly growing number of polarized radio sources. Hundred and thousands polarized background sources can be measured towards nearby galaxies thus allowing their detailed magnetic field mapping by means of Faraday rotation measures (RM). We aim to estimate the required density of the background polarized sources detected with the SKA for reliable \\emph{recognition} and \\emph{reconstruction} of the magnetic field structure in nearby spiral galaxies. We construct a galaxy model which includes the ionized gas and magnetic field patterns of different azimuthal symmetry (axisymmetric (ASS), bisymmetric (BSS) and quadrisymmetric spiral (QSS), and superpositions) plus a halo magnetic field. RM fluctuations with a Kolmogorov spectrum due to turbulent fields and/or fluctuations in ionized gas density are superimposed. \\emph{Recognition} of magnetic structures is possible from RM towards background sources behind galaxies or a continuous RM map obtained from the diffuse polarized emission from the galaxy itself. Under favourite conditions, about a few dozens of polarized sources are sufficient for a reliable recognition. \\emph{Reconstruction} of the field structure without precognition becomes possible for a large number of background sources. A reliable reconstruction of the field structure needs at least 20 RM values on a cut along the projected minor axis which translates to $\\approx1200$ sources towards the galaxy. Radio telescopes operating at low frequencies (LOFAR, ASKAP and the low-frequency SKA array) may also be useful instruments for field recognition or reconstruction with the help of RM, if background sources are still significantly polarized at low frequencies. This work is a part of the science simulations of cosmic magnetism in the frame of the SKADS (SKA Design Studies).}   \\FullConference{From Planets to Dark Energy: the Modern Radio Universe\\\\ October 1-5 2007\\\\ The University of Manchester, UK}  \\begin{document} ",
        "introduction": "Faraday rotation measures (RM) in galaxies are generated by regular fields of the galaxy plus its ionized gas, both of which extend to large galactic radii (Fig.~\\ref{fig:m31}). RM towards polarized background sources can trace regular magnetic fields in these galaxies out to even larger distances, however, with the sensitivity of present-day radio telescopes, the number density of polarized background sources is only a few sources per solid angle of a square degree, so that only M~31 and the LMC could be investigated so far \\cite{han98,gaensler05}.  Future high-sensitivity radio facilities will observe polarized intensity and RM for a huge number of faint radio sources, thus providing the high density background of polarized point sources. This opens the possibility to study in detail the large-scale patterns of magnetic fields and their superpositions thus allowing the dynamo theory for field amplification and its ordering to be tested. A major step towards a better understanding of galactic magnetism will be achieved by the Square Kilometre Array (SKA, www.skatelescope.org) and its pathfinders. ",
        "conclusions": ""
    },
    "0711/0711.1779_arXiv.txt": {
        "abstract": "Using ISAAC/VLT, we have obtained individual spectra of all NIR-bright stars in the central 2'x2' of the cluster Westerlund 1 (Wd 1) with a resolution of R$\\approx$9000 at a central wavelength of 2.30$\\mu$m. This allowed us to determine radial velocities of ten post-main-sequence stars, and from these values a velocity dispersion. Assuming virial equilibrium, the dispersion of $\\sigma$=8.4 km/s leads to a total dynamical cluster mass of 1.25$\\times 10^5$M$_\\odot$, comparable to the photometric mass of the cluster. There is no extra-virial motion which would have to be interpreted as a signature of cluster expansion or dissolution. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.1586_arXiv.txt": {
        "abstract": "{The MAGIC telescope has observed very-high energy gamma-ray emission from the BL~Lac object PG 1553+113 in 2005 and 2006 at an overall significance is $8.8 \\sigma$. The light curve shows no significant flux variations on a daily timescale. The flux level during 2005 was, however, significantly higher as compared to 2006. The differential energy spectrum between $\\sim$90 GeV and 500 GeV is well described by a power law with a spectral index of $-4.2\\pm0.3$. The photon energy spectrum and spectral modeling allow to pose upper limits of $z=0.74$ and $z=0.56$, respectively, on the yet undetermined redshift of PG 1553+113. Recent VLT observations of this blazar show featureless spectra in the near-IR, thus no direct redshift could be determined from these measurements.}  \\begin{document} ",
        "introduction": "The active galactic nucleus (AGN) PG~1553+113 was first reported in the Palomar-Green catalog of UV-bright objects \\cite{green}. Its spectral characteristics are close to those of X-ray selected BL~Lacs \\cite{falomo90}. Despite several attempts, no emission or absorption lines have been found in the spectrum of PG 1553+113. Thus only indirect methods can be used to determine the redshift $z$ \\cite{sbarufatti05, sbarufatti06}. Recently the H.E.S.S. \\cite{hess1553} and MAGIC \\cite{magic1553} collaborations have presented $\\gamma$-ray signals at the $>4\\sigma$ and the $8.8\\sigma$ level, respectively, making this source a confirmed VHE $\\gamma$-ray emitter. Here we present the MAGIC observation results. ",
        "conclusions": "The MAGIC detection confirms the tentative signal seen by H.E.S.S. at a higher energy \\cite{hess1553}. The agreement between the H.E.S.S. and MAGIC energy spectra is reasonably good. While the spectral slope is consistent within errors, the absolute flux above 200~GeV in 2005 is by a factor 4 larger compared to H.E.S.S. This difference may in part be explained by the systematic errors of both measurements but also by variations in the flux level of the source (the observations with H.E.S.S. were commenced after MAGIC). The observed energy spectrum is steeper than that of any other known BL~Lac object. This may be an indication of a large redshift ($z > 0.3$), but can as well be attributed to intrinsic absorption at the AGN or, more naturally, to an inverse Compton peak position at lower energies. The spectrum can, however, be used to derive an upper limit on the source redshift from physical constrains on the intrinsic photon index ($\\Gamma_\\mathrm{int} > 1.5$) as discussed in \\cite{hess1553}. Using the lower limit on the evolving EBL density from \\cite{kneiske04} we derived a $2 \\sigma$ upper limit on the redshift of $z < 0.74$, which agrees on the limit found in \\cite{hess1553}. Other approaches to limit the redshift are given in \\cite{bc,fg}.  The spectral energy distribution (SED) of PG 1553+113 together with the results from a model calculation are shown in Fig.~\\ref{fig:spectrum_2}. The VHE data points correspond to the intrinsic spectrum of PG 1553+113 as derived for a redshift of $z = 0.3$. The black points at low energies denote the average optical and X-ray flux taken at the same time as the MAGIC observations. Non-simultaneous radio, optical and X-ray data were taken from \\cite{giommi02}. The solid line shows the result of a model fit to the simultaneously recorded data (black points) using a homogeneous, one-zone Synchrotron Self-Compton (SSC) model provided by \\cite{krawczynski04}. The $\\gamma$-ray, X-ray and optical data are well described by the model. This is not the case for the radio data, where intrinsic absorption requires a much larger emitting volume compared to X-rays and $\\gamma$-rays. Except for a somewhat smaller radius of the emitting region, identical model parameters as in \\cite{costamante02} have been used: Doppler factor $D=21$, magnetic field strength $B=0.7\\ \\mathrm{G}$, radius of the emitting region $R=1.16^{+0.62}_{-0.21} \\cdot 10^{16}\\ \\mathrm{cm}$, electron energy density $\\rho_e = 0.11^{+0.18}_{-0.06} \\ \\mathrm{erg / cm^3}$, slope of the electron distribution $\\alpha_e=-2.6$ for $8.2 < \\log \\left( E / \\mathrm{eV} \\right) < 9.8^{+0.2}_{-0.05}$ and $\\alpha_e=-3.6$ for $9.8^{+0.2}_{-0.05} < \\log \\left( E / \\mathrm{eV} \\right) <10.6^{+1.6}_{-0.0}$. The limits on some of these parameters indicate the change of the SED model parameters when varying the assumed redshift from $z=0.2$ up to $z=0.7$ (parameters without limits were kept constant for all fits). In the case of $z \\ge 0.56$ the SED model cannot accurately describe the data and, based on the obtained $\\chi^2$ value for the SED fit, $z>0.56$ is excluded on the $4.5\\sigma$ level.  \\begin{figure} \\includegraphics[width=\\linewidth]{icrc0072_fig04.eps} \\vspace*{-1.2cm} \\caption{SED of PG 1553+113. The solid lines are the result of a SSC model fit to the black data points using the code provided by \\cite{krawczynski04}. The gray points comprise non-simultaneous radio, optical and X-ray data from \\cite{giommi02}.} \\label{fig:spectrum_2} \\end{figure}  PG 1553+113 was in a high state in the optical in both years showing a strong flare at the end of March 2006. The high linear polarization of the optical emission ($8.3 \\pm 0.2 \\%$) indicates that a sizeable fraction of the optical flux is indeed synchrotron radiation. In $\\gamma$-rays only a significant change in the flux level from 2005 to 2006 is found while there is no evidence for variability in X-rays. As a result, a possible correlation between the different energy bands cannot be established. A possible connection between the $\\gamma$-ray detection and the optical high state can, however, not be excluded. The optical flare without X-ray or $\\gamma$-ray counterpart may still be explained by external-inverse-Compton models which predict a time lag of the X-rays and $\\gamma$-rays with respect to the optical emission."
    },
    "0711/0711.3310_arXiv.txt": {
        "abstract": "We present high spatial resolution ($\\sim 35$~parsec) $5-38\\mu$m spectra of the central region of M82, taken with the Spitzer Infrared Spectrograph. From these spectra we determined the fluxes and equivalent widths of key diagnostic features, such as the [NeII]12.8$\\mu$m, [NeIII]15.5$\\mu$m, and H$_2$ S(1)17.03$\\mu$m lines, and the broad mid-IR polycyclic aromatic hydrocarbon (PAH) emission features in six representative regions and analysed the spatial distribution of these lines and their ratios across the central region. We find a good correlation of the dust extinction with the CO 1-0 emission. The PAH emission follows closely the ionization structure along the galactic disk. The observed variations of the diagnostic PAH ratios across M82 can be explained by extinction effects, within systematic uncertainties. The $16-18\\mu$m PAH complex is very prominent, and its equivalent width is enhanced outwards from the galactic plane. We interpret this as a consequence of the variation of the UV radiation field. The EWs of the $11.3\\mu$m PAH feature and the H$_2$ S(1) line correlate closely, and we conclude that shocks in the outflow regions have no measurable influence on the H$_2$ emission. The [NeIII]/[NeII] ratio is on average low at $\\sim$0.18, and shows little variations across the plane, indicating that the dominant stellar population is evolved (5 - 6 Myr) and well distributed. There is a slight increase of the ratio with distance from the galactic plane of M82 which we attribute  to a decrease in gas density. Our observations indicate that the star formation rate has decreased significantly in the last 5 Myr. The quantities of dust and molecular gas in the central area of the galaxy argue against starvation and for negative feedback processes, observable through the strong extra-planar outflows. ",
        "introduction": "M82 (NGC 3034) is an irregular galaxy located at 3.3 Mpc \\citep{freedman88} in the M81 group. It is the closest starburst galaxy, seen nearly edge-on, with an inclination angle of about $80\\deg$. At infrared wavelengths it is the brightest galaxy on the sky, with a total infrared luminosity of $3.8\\times10^{10}L_{\\odot}$ \\citep{colbert99}. Most of its luminosity originates from the inner 500 pc hosting intense starburst activity presumably triggered by a tidal interaction with M81 \\citep[e.g.][]{yun93}. Evidence for this interaction comes from the HI streams which connect M81 to all three (M82, NGC 3077, and NGC 2976) peculiar members of the inner M81 group \\citep{apple81,yun94}. Recent deep optical images also revealed stars associated with the HI bridge between M81 and M82 \\citep{sun05}.  At the distance of M82, $1\\arcsec$ corresponds to 15 parsec, which allows spatially resolved studies of the starburst region. Evidence of a stellar bar $\\sim1$ kpc long is shown by near-infrared studies \\citep[e.g.][]{telesco91,larkin94}, mid-infrared [NeII]12.8$\\mu$m and millimetric CO emission studies \\citep{lo87}. According to \\citet{larkin94} and \\citet{achter95}, there is a rotating ring of ionized gas at a radius of $\\sim85$ pc, and on the inner side of a ring of molecular gas at $\\sim210$ pc. Two possible spiral arms were also identified by \\citet{shen95} and \\citet{mayya05}, at radii of $\\sim125$ pc and $\\sim400$ pc. The starburst of M82 drives a bipolar mass outflow out to several kiloparsecs perpendicular to the plane of the galaxy, especially evident in X-ray and H$\\alpha$ \\citep{bregman95,shop98,lehnert99,cappi99,strick04}. Dust has also been detected in the outflow region \\citep{alton99,heckman00,hoopes05,engel06}. The star forming regions of M82 are predominantly clustered in the volume enclosed by the molecular gas ring, indicated by the {\\sc HII} region tracers, like the [NeII]12.8$\\mu$m line and the mid- and far-infrared continuum emission \\citep{telesco91,walter02,larkin94,achter95,lipscy04}. Near-infrared hydrogen recombination lines also arise in these regions, but they are a more ambiguous tracer of star formation, as they can also be excited by shocks, though these are unlikely to dominate.  Near-infrared integral field spectroscopy and ISO-SWS mid-infrared spectroscopy by \\citet{schreiber01} allowed a detailed modelling of starburst activity in the central region \\citep{schreiber03}. These models are consistent with the occurrence of starburst activity in two successive episodes, about 10 and 5 Myr ago, each lasting a few million years. However, the spatial studies by \\citet{schreiber01} and others covered only near-infrared wavelengths.  The large aperture of the ISO-SWS provided a continuous $2.4-45 \\mu$m spectrum but covered the whole central region of M82.  ISOCAM-CVF data \\citep{schreiber02} provided better spatial resolution but the spectra only had a spatial resolution of $R \\sim 40$ and were shortward of $15\\mu$m.  \\citet{engel06} published Spitzer-IRS low-resolution spectra of a 1 arcmin wide strip along the minor axis of M82, intersecting the disk at the eastern side. The spectra, taken as part of the SINGS Legacy project, were combined with Spitzer 8 and 24 $\\mu$m images, and show that the emission by polycyclic aromatic hydrocarbons (PAH) and H$_2$ molecules extends far out from the disk (to 6 kpc) in both directions.  \\citet{engel06} suggest that the filamentary aromatic-dominated emission represents dust either expelled from the galaxy as a result of a powerful nuclear superwind, or that dust is in the halo being lit up by the starburst, perhaps coexisting with the extensive warm H$_2$ molecules. They suggest that this halo dust is probably a leftover from the interaction with M81.  In this paper we present mid-IR spectral maps at unsurpassed sensitivity and spatial resolution of the central $\\sim0.5$ kpc$^2$ of M82, covering the main contributors to the bolometric luminosity of the galaxy. Our goal is to provide a spatial and spectrally detailed description of the physical conditions within the central $\\sim500$ pc of M82, to help us to give an insight on the evolution of the starburst activity in this region. This involves the study of the distribution of the radiation field, gas density, and the physical properties of PAHs. Of particular interest are the spatial variations of the fine-structure lines, the excitation of the molecular hydrogen, and the distribution of the PAH molecules. In section 2 we describe the observations and data reduction, in section 3 we present the data, and in section 4 we discuss the scientific results, followed by our conclusions.   \\begin{figure} \\epsscale{1.15} \\plottwo{f1a.eps}{f1b.eps} \\caption{Left: Overlay of SL (red) and SH (blue) coverages and selected low-res extraction regions on an IRAC 8$\\mu$m image. Regions A and B, represented in yellow, are regions where LL (14 - 35 $\\mu$m) spectra were extracted. Right: Zoom-in of the IRAC 8$\\mu$m image with an overlay of the SH map area in blue and the selected regions (in green) from where the SH+SL spectra were extracted. The ISO-SWS aperture used in the $12-27 \\mu$m range by \\citet{schreiber01} is also overlayed in yellow on the image. Both figures are in logarithmic scaling.} \\label{figregions} \\end{figure} ",
        "conclusions": "We presented spatially resolved mid-infrared spectra of the central region of M82.  The spectra were taken with the Spitzer Infrared Spectrograph in both the $5-38\\mu$m low-resolution ($R\\sim 65-130$) and the $10-37\\mu$m high-resolution $R\\sim 600$ module. The high signal-to-noise and the continuous spatial and spectral coverage allowed us to study the nucleus of a starburst galaxy in unsurpassed detail. Our goal was to study the physical conditions of the interstellar medium and their spatial variations within the central kpc of M82.  Overall, the spectra show the typical features of a starburst: prominent PAH features, silicate absorption, fine-structure lines, and a steeply rising continuum. We built a spectral map with high-resolution spectra, selected six representative regions for spectral extraction, and studied the variations of the neon ionic lines and PAH feature emission among the regions.  We attempted to trace the structure of the ionizing radiation, and therefore the young stellar population of M82 through the diagnostic emission lines, [NeII]12.8$\\mu$m, and [NeIII]15.5$\\mu$m. The overall morphology of the ionizing radiation in the region as revealed by the spectral maps appears to be in agreement with previous ground-based observations \\citep{achter95}. There is surprisingly little spatial variation of the [NeIII]/[NeII] ratio across the disk, and it varies only by a factor of three throughout the central region, with these variations not corresponding to the position of the emission peaks. We observed an increase of the [NeIII]/[NeII] ratio by a factor of 2 outwards the galactic plane, which may be associated with the outflows. We suggest that the increase of the [NeIII]/[NeII] ratio with distance to the galactic plane is due to a decrease in gas density rather than a hardening of the field.  The [NeIII]/[NeII] ratio is low on average, which indicates that the dominant population consists of already evolved clusters ($>5$ Myrs). We cannot rule out the presence of ongoing star formation in the central region, but it must occur at a much reduced rate ($<5\\%$) compared to previous epochs. This drop is unlikely to be caused by starvation, as there is still a large amount of molecular gas present in the central region. It is more likely due to negative feedback processes causing a decrease in the star formation rate.  There are significant variations in the amount of dust extinction, which strongly correlate with the CO 1-0 emission across the central region. These variations are strong enough to affect the interpretation of PAH features, but due to limited spectral coverage the extinction estimates are uncertain.  The flux of the main PAH features correlates spatially with the flux of the neon ionic lines, and with previous IRAC observations. Variations in the PAH ratios such as $6.2/11.3\\mu$m were observed across the disk. However, they are strongly affected by the silicate feature at $10\\mu$m. We studied the variations the $6.2/7.7\\mu$m and $11.3/7.7\\mu$m PAH ratios, which are diagnostics for the size and the degree of ionization of PAHs. We found no correlation between the $6.2/7.7\\mu$m ratio with radiation hardness or with any other resolved spatial structure in the central region. Since the variations are comparable to the uncertainties in the measurement the results are not (yet) significant enough to support PAH size variations at parsec scales. Similarly, the variations of the $11.3/7.7\\mu$m ratio are mostly due to extinction and show little support for variations in PAH ionization throughout the region.  The $17\\mu$m PAH complex is very prominent in the center of M82. We did not find any relative variations within the complex, which is in agreement with predictions. The variations of the $17/6.2\\mu$m ratio are most likely due to extinction effects. Due to the uncertainties on the determination of extinction, we did not consider the remaining variations as a clear indicator of PAH size variation.  We observed an enhancement of the EWs of the $11.3\\mu$m and the other PAH features longwards of $10\\mu$m outwards from the galactic plane. Several explanations exist for this, but we favor the variation of the UV radiation field, which excites differently PAHs and VSGs, given their different sizes.  The S(1) and S(2) rotational transition lines of H$_2$ have been detected in our spectra throughout the central region. H$_2$ and PAHs coincide at least on spatial scales of the resolution of the map, which is 35~pc. We conclude that large scale shocks in the outflow regions have no measurable influence on the H$_2$ emission. However, we cannot distinguish between UV- and local shock excitation as produced by supernovae or energetic outflows, on scales smaller than the resolution of our pixels (35 pc).  The Spitzer-IRS observations of the central region of M82 complements previous studies, not only in mid-infrared, but also in other wavelengths. Our results demonstrate the importance of spatially resolved spectroscopy in starburst studies. They helped to constrain the age of the starburst and confirm results from other studies and also stressed the importance of a thorough study of extinction to investigate possible variations of PAH properties. Further research of the starburst feedback and quenching processes will elucidate the sharp decrease in star formation in the last 5 Myr. It would be interesting to see if higher spatial resolution, as expected from the MIRI on board of JWST will discover a wider variation of the radiation field or signs of variation of PAH sizes an/or ionization. In addition, spatially more extended spectral studies of M82 are necessary to study the connection of the ionic line ratios with the outflows."
    },
    "0711/0711.1912_arXiv.txt": {
        "abstract": "The  EGRET excess of diffuse Galactic gamma rays shows all the key features of dark matter annihilation (DMA) for a WIMP mass in the range 50-100 GeV, especially the distribution of the excess is compatible with a standard halo profile with some additional ringlike substructures at 4 and 13 kpc from the Galactic centre. These substructures coincide with the gravitational potential well expected from the ring of dust at 4 kpc and the tidal stream of dark matter from the Canis Major satellite galaxy at 13 kpc, as deduced from N-body simulations fitting to  the Monoceros ring of stars. Strong independent support for this substructure is given by the gas flaring in our Galaxy. The gamma rays from DMA are originating predominantly from the hadronization of mono-energetic quarks, which should produce also a small, but known fraction of protons and antiprotons. Bergstr\\\"om et al. an antiproton flux far above the observed antiproton flux and they conclude that the DMA interpretation of the EGRET excess can therefore be excluded. However, they used an isotropic propagation model, i.e. the same diffusive propagation in the disk and the halo. It is shown that  an anisotropic propagation model is consistent with the EGRET gamma ray excess, the antiproton flux and the ratios of secondary/primary and unstable/stable cosmic ray particles. Such an anisotropic propagation is supported by  the large bulge/disk ratio of the positron annihilation line, as observed  by the INTEGRAL satellite. In this case no need for new sources specific to the bulge are needed, so the claimed evidence for strong DMA in the bulge from these observations is strongly propagation model dependent. In the framework of Supersymmetry cross section predictions for direct dark matter searches are presented taking into account the EGRET data, the WMAP data and other electroweak constraints. ",
        "introduction": "Cold Dark Matter (CDM) is well established by the high rotation speeds in galaxies and clusters of galaxies. Recent cosmological measurements yield a dark matter (DM) density of $22\\pm2\\%$ of the energy of the Universe \\cite{wmap}. If this DM is created thermally during the Big Bang the present relic density  is inversely proportional to $\\langle\\sigma v\\rangle$, the annihilation cross section $\\sigma$ of DM particles, usually called WIMPS (Weakly Interacting Massive Particles), times their relative velocity. The average is taken over these velocities. This inverse proportionality is obvious, if one considers that a higher annihilation rate, given by $\\langle\\sigma v\\rangle n_\\chi$, would have reduced the relic density before freeze-out, i.e. the time, when the expansion rate of the Universe, given by the Hubble constant, became equal to or larger than the annihilation rate. The relation can be written as: \\cite{jungman} \\begin{equation} \\Omega_\\chi h^2=\\frac{m_\\chi n_\\chi}{\\rho_c} h^2\\approx (\\frac{3\\cdot 10^{-27} cm^3 s^{-1}}{\\langle\\sigma v\\rangle})\\label{eq1}. \\end{equation} For the present value of $\\Omega h^2=0.105 \\pm 0.008$, as measured by WMAP \\cite{wmap}, the thermally averaged total cross section at the freeze-out temperature of $m_\\chi/22$ must have been around $3\\cdot 10^{-26} ~{\\rm cm^3s^{-1}}$. Note that $\\langle\\sigma v\\rangle$ as calculated from Eq.~\\ref{eq1} is independent of the WIMP mass (except for logarithmic corrections) as can be shown by  detailed calculations \\cite{kolb}. If the s-wave annihilation is dominant, as expected  in many supersymmetric models, then the annihilation cross section is energy independent, i.e. the cross section given above is also valid for the cold temperatures of the present universe \\cite{susy}. Such a large cross section will lead to a production rate of mono-energetic quarks\\footnote{The quarks are mono-energetic, since the kinetic energy of the cold dark matter particles is expected to be negligible with the mass of the particles, so the energy of the quarks equals the mass of the WIMP.} in our Galaxy, which is 40 orders of magnitude above the rate produced at any accelerator. The fragmentation of these mono-energetic quarks will lead to a large flux of neutrinos, photons, protons, antiprotons, electrons and positrons in the Galaxy. From these, the protons and electrons disappear in the sea of many matter particles, but the photons and antimatter particles may be detectable above the background, generated by  cosmic ray interactions with the gas in the Galaxy. Such searches for indirect Dark Matter detection have been actively pursued. Recent  reviews and references to earlier work can be found in Refs. \\cite{bergstrom,sumner,Bertone:2004pz}.  Gamma rays have the advantage that they point back to the source and do not suffer energy losses, so they are the ideal candidates to trace the dark matter density and have a spectral shape characteristic for mono-energetic quarks. The charged components interact with Galactic matter and are deflected by the Galactic magnetic field, so they do not point back to the source. Therefore the charged particle fluxes have large uncertainties from the propagation models, which determine how many of the produced particles arrive at the detector. For gamma rays the propagation is straightforward: only the ones pointing towards the detector will be observed.  An excess of diffuse gamma rays compatible with dark matter annihilation (DMA) has indeed been observed by the EGRET telescope on board of NASA's CGRO (Compton Gamma Ray Observatory)\\cite{hunter}. Below 1 GeV the CR interactions describe the data perfectly well, but above 1 GeV the data are up to a factor two above the expected background. The excess shows all the features of DMA annihilation for a WIMP mass between 50 and 70 GeV \\cite{us}. Masses up to 100 GeV are possible if one assumes the cosmic ray energy spectrum to vary in the Galaxy. The excess was observed in all sky directions, which would imply that DM is not dark anymore, but shining in gamma rays \\cite{us}. Of course, such an important observation needs to be scrutinized heavily. Among the most important criticism was a paper by Bergstr\\\"om et al.\\cite{bergstrom1} claiming that the antiproton flux from DMA, using the DM distribution from the EGRET excess, would be an order of magnitude higher than the observed antiproton flux. However, they use an isotropic propagation model with the {\\it same} propagation in the disk and the halo. For the expected anisotropic propagation models everything can be made consistent by having a faster propagation perpendicular to the disk than in the disk, in which case the antiprotons from DMA in the halo do not return to the disk. This demonstrates  the large uncertainties from propagation models for indirect dark matter searches.   It is interesting to note that  anisotropic propagation models can easily explain a large bulge/disk (B/D) ratio of  the positron annihilation, as observed from the intensity of the 511 keV annihilation line of thermalized positrons. These positrons are e.g. produced in the $\\beta$-decays of radioactive nuclei from supernova explosions. A large B/D ratio is simply due to the different geometries of the disk and the bulge: in the bulge the positrons and electrons can annihilate and radiate before escaping to the halo, in the disk they enter the halo much faster  without having time to radiate and/or annihilate.  In contrast, in isotropic propagation models one can only explain the large bulge/disk ratio by a new source for positrons for the bulge alone. Anisotropic propagation models have enough freedom to explain the results without new sources.  So the claimed evidence for DMA from these observations\\cite{boehm} are strongly propagation model dependent.   The paper is organized as follows: in  section \\ref{egret} we summarize the DMA interpretation of the EGRET excess.  In section \\ref{prop1} we discuss the problems with the isotropic propagation models. An anisotropic propagation model, which simultaneously describes the antiproton flux, the EGRET excess and the observations concerning primary and secondary cosmic rays and cosmic clocks is discussed in section \\ref{prop2}, while in section \\ref{integral}  the consequences for the INTEGRAL excess of the 511 keV line  in this anisotropic propagation model are discussed. In section \\ref{direct} the constraints from direct dark matter searches are discussed. Section \\ref{conclusion} summarizes the results. ",
        "conclusions": "With an anisotropic propagation model the amount of antiprotons expected from DMA annihilation can be reduced by one to two orders of magnitude. Therefore the claim by Bergstr\\\"om et al.\\cite{bergstrom1} that the DMA interpretation of the EGRET excess of diffuse Galactic gamma rays is excluded ``by a large margin'' is strongly pro\\-pagation model dependent. It only applies for their simplified propagation model, which assumes the same diffusion in the halo and in the disk. An anisotropic propagation model with different propagation in the halo and the disk can reconcile the EGRET excess with the antiproton flux and the ratios of secondary/primary and unstable/stable nuclei. In addition the difference in geometry between bulge and disk leads to much more radiation and annihilation of positrons and electrons in the bulge as compared to the disk, thus alleviating the need to introduce new sources of positrons and electrons for the bulge to explain the INTEGRAL excess of the 511 keV line in the bulge\\cite{boehm}. Direct dark matter search experiments reach the cross section range expected from the EGRET excess. However, the experimental limits are inversely proportional to these cross sections times the local relic density. Both are uncertain: for the cross section  one assumes the WIMPS are the lightest supersymmetric partners of the minimal supersymmetric model with gravity inspired breaking of supersymmetry, while for the local relic density one assumes the WIMPS are smoothly distributed instead of the expected distribution in clumps. The clumpy nature can drastically reduce the local density if we are not located inside a clump.    In summary we consider DMA is a viable explanation of the EGRET excess of diffuse Galactic gamma rays, especially since it is observed with the same shape of the fragmentation of mono-energetic quarks in all sky directions and the intensity distribution of the excess traces the DM profile, as shown independently by the rotation curve, the gas flaring and the N-body simulation of the disruption of the Canis-Major satellite galaxy."
    },
    "0711/0711.2316_arXiv.txt": {
        "abstract": "Although dark matter makes up 80\\% of the gravitational mass of our Galaxy, its composition is not known. One hypothesis is that dark matter consists of massive particles called WIMPs.  WIMPs are expected to accumulate and coannihilate in the cores of stars, but the only signature of this accumulation has been thought to be hard- to-observe high-energy neutrinos. Here we propose an entirely new observable signature.   WIMP coannihilations in the core of a very low-mass star, brown dwarf, or planetary-mass object should alter the star's chemical composition via spallation reactions.  Very close to the Galactic center, these stars may acquire extremely high lithium, beryllium, and boron abundances, even for models with otherwise- undetectable WIMP-nucleon cross sections.   These abundances should be measurable in certain stellar systems and phenomena. ",
        "introduction": "In this era of high precision cosmology, the mass-energy density of the Universe is well-constrained.  {From a combination of measurements of the microwave background, large-scale strucure, rotation curves, gravitational lensing, and galaxy cluster temperatures, the Universe is known to contain about 4\\% ordinary matter and 24\\% dark matter.}  Dark matter is known to move and gravitate just like massive baryonic matter, but without any particle-particle interactions at the electromagnetic or strong-force scales.  The Standard Model of particle physics contains no massive non-interacting particle like this, but some extended models do; such particles are called WIMPs.  Supersymmetry (SUSY) naturally includes a WIMP, the neutralino ($\\chi_0$), which is an excellent dark-matter candidate; Kaluza-Klein bosons, which arise in theories with large or warped extra dimensions, are another example.  In this work we refer to WIMPs as $\\chi_0$, but our results are not model-specific (except as discussed below).  In order for WIMPs to be produced in the Big Bang, there must be a sizeable cross section for reactions like $x\\bar{x} \\rightarrow \\chi_0\\chi_0$, where $x$ is some Standard Model particle.   This implies that the reverse reaction, $\\chi_0\\chi_0 \\rightarrow x\\bar{x}$, will occur today at a rate proportional to the square of the WIMP density.  The WIMP-nucleus elastic scattering cross section, $\\chi_0 ^aN_z \\rightarrow \\chi_0 ^aN_z$ is not constrained by astrophysics, and may be extremely small. Terrestrial experiments allow us to constrain this cross section to be $\\sigma_0/A^2 \\tilde{<} 10^{-43}$ cm$^2$ for spin-independent WIMP-nucleus scattering, and $\\sigma_p \\tilde{<} 10^{-36}$ cm$^2$ for spin-dependent WIMP-proton scattering.  Any nonzero cross section for $\\chi_0 ^aN_z \\rightarrow \\chi_0 ^aN_z$ {allows} a Galactic {$\\chi_0$ particle} to scatter off nuclei in stars, to enter bound orbits which cross the star repeatedly, and thus to thermalize and be captured inside the star.  As the $\\chi_0$ number density in the star's core increases, the coannihilation rate will increase until it equilibrates with the  capture rate.  Recent work\\cite{moskalenko_dark_2007} shows that, in regions near the Galactic center's SuperMassive Black Hole (SMBH) where the dark-matter density is highest, the capture-and-annihilation process may provide enough thermal power to change a star's luminosity and evolution.  We show that these coannihilations can change stellar \\emph{chemical} composition in unambiguous ways.  Most importantly, the high-energy annihilation products can convert carbon and oxygen into $^3$He, Li, Be, and B via spallation.  We restrict our analysis to stars with $M < 0.1$ M$_\\odot$ (encompassing {the lowest-mass} main-sequence stars, brown dwarfs, and planetary-mass objects) which {we will refer to loosely as} substellar objects (SSOs). ",
        "conclusions": ""
    },
    "0711/0711.2857_arXiv.txt": {
        "abstract": "We present a theoretical investigation of the light, radius and radial velocity variations of the prototype $\\delta$ Cephei. We find that the best fit model accounts for luminosity and velocity amplitudes with an accuracy better than $0.8\\sigma$, and for the radius amplitude with an accuracy of $1.7\\sigma$. The chemical composition of this model suggests a decrease in both helium (0.26 vs 0.28) and metal (0.01 vs 0.02) content in the solar neighborhood. Moreover, distance determinations based on the fit of light curves agree at the $0.8\\sigma$ level with the trigonometric parallax measured by the Hubble Space Telescope (HST). On the other hand, distance determinations based on angular diameter variations, that are independent of interstellar extinction and of the $p$-factor value, indicate an increase of the order of 5\\% in the HST parallax. ",
        "introduction": "Classical Cepheids are the most important primary distance indicators. One of the key issues concerning the use of Cepheids as distance indicators is their dependence on chemical composition. Empirical tests seem to suggest that metal-rich Cepheids are, at fixed period, brighter than metal-poor ones, either over the entire period range \\citep[][]{KSS98,KSF03,K03,SCG04,G04,S04,M06} or for periods shorter than $\\approx 25$ days  \\citep[][]{SBR97,Sa04}. On the other hand, distance estimates based on the Near-Infrared Surface Brightness method indicate that the slope of the Period-Luminosity ($PL$) relation might be the same for Magellanic and Galactic Cepheids \\citep{GSB05}. However, the impact of the projection factor $p$ on the Cepheid distances is still debated \\citep[][]{S06,N07}. Early theoretical predictions based on linear pulsation models \\citep[e.g.][]{SBT99,A99,BA01} predicted a mild metallicity effect, while nonlinear convective models \\citep[][]{BMS99,C02,F02,MMF05} showed that the  metallicity affects both the zero-point and the slope of the $PL$ relations. This effect is more evident in the optical bands and it is nonlinear when moving from metal-poor ($Z=0.004$) to super metal-rich ($Z=0.04$) structures. Current predictions also suggest a turnover across solar metallicity ($Z$ $\\sim$0.02) due to the correlated increase in He content \\citep[][]{F02,MMF05}.  Three approaches are adopted to calibrate the Cepheid PL relations: cluster distances, trigonometric parallaxes and Baade-Wesselink (BW) method. Parallaxes based on Cepheids in open clusters and associations have been widely adopted to calibrate the PL relations \\citep{TB02}, but \\citet{Fo07} found a systematic difference with other methods. Accurate Cepheid parallaxes have been provided by \\citet{BFM02} and by \\citet{BFM07} using the Fine Guide Sensor on board the Hubble Space Telescope (HST). They measured the distance of several Galactic Cepheids with average relative errors of $\\approx 8$\\%. The BW method appears to be a robust approach, but some assumptions need to be investigated \\citep[][]{G87,BCS94,GSB05}. The value of the projection factor ($p$), i.e. the parameter adopted to transform radial into pulsation velocity, is still controversial \\citep[][]{S95,MK062}. The typical value used in the literature is $p=1.36$ \\citep{BMB82}, but period-dependent values are also adopted \\citep[][]{GFG93,GSB05}. However, the $p$-factor of $\\delta$ Cephei, $p=1.27\\pm 0.06$, measured by \\citet{MK05}, using the HST parallax, agrees quite well with the predicted value \\citep{NFM04}.  Although, several Cepheids are binaries dynamical mass estimates are only available for a handful objects \\citep[][]{E05}. Current mass estimates rely either on evolutionary or on pulsation prescriptions \\citep[][]{DC00,KW02,CBFMM05} and we are still facing the long-standing problem of the ``Cepheid mass discrepancy'' \\citep[]{Cox80}. Evolutionary masses are systematically larger than pulsation masses \\citep[][]{B01,BBK01,CBFMM05}.  Theoretical fits of Cepheid light curves provide robust constraints on their intrinsic parameters and distances \\citep[][]{WAS97,BCM02,KW06}. This approach, up to now, has only been applied to Bump Cepheids. We present, for the first time, a simultaneous fit of optical/NIR light, radial velocity, and angular diameter variations of the prototype $\\delta$ Cephei. We selected this object because accurate measurements of light, radius, velocity curves, and of the $p$-factor are available and its geometric distance is known with an accuracy of 4\\% \\citep[273$^{+12}_{-11}$][]{BFM02}. ",
        "conclusions": "We constructed a sequence of models at fixed chemical composition ($Z=0.02$, $Y=0.28$) and by assuming steps in mass and in effective temperature of $\\Delta M = 0.5M_\\odot$ and $\\Delta T_e=100K$. The comparison between theory and observations indicate that a model with $M=5.5M_\\odot$, and $Te=5700K$ accounts for the shape of the light curves and for the secondary bump located on the decreasing branch of the radial velocity curve ($\\phi \\approx 0.2$). To further improve the agreement between predicted and empirical observables, we constructed a new sequence of models by adopting smaller steps in mass and in effective temperature, namely $\\Delta M=0.1M_\\odot$,$\\Delta T_e=25 K$. For every model in this sequence, the mixing length parameter was changed from $\\alpha_{ml}=0.3$ to $\\alpha_{ml}=0.4$. Note that an increase in $\\alpha_{ml}$ causes a decrease in the pulsation driving, and in turn a decrease in the limit cycle amplitudes. Data plotted in the left panels of Fig. 1 show that the best fit model ($\\alpha_{ml}=0.33$) accounts for the optical and NIR light curves, but both the radial velocity ($\\sigma(\\Delta V_r)=4.2$) and the radius ($\\sigma(\\Delta R)=3.3$) amplitudes are systematically smaller than the observed ones. This discrepancy can only be removed by decreasing the $p$-factor from 1.27 to the barely plausible value of $p \\sim 1$. To further constrain the impact of the adopted chemical composition on the limit cycle stability, we constructed a new sequence of models by using the same metal abundance and a lower He content ($Y=0.26$). The new best fit model (see Table~2 and the middle panels of Fig. 1) accounts for optical and NIR light curves and better reproduces radial velocity ($\\sigma(\\Delta V_r)=1.6$) and radius ($\\sigma(\\Delta R)=2.3$) variations. The predicted $V$-band light curve presents a mild discrepancy along the rising branch. However, the discrepancy on luminosity amplitudes is smaller than $0.4\\sigma$. In order to verify the occurrence of a degeneracy in the input parameters, we also computed a new series of models by using a lower He content, ($Y=0.24$, see Table~2 and right panels of Fig. 1). Interestingly enough, velocity and radius amplitudes predicted by new best fit-model are at least $2.5\\sigma$ smaller than observed. Once again, a reasonable fit between theory and observations would require a barely plausible $p$-factor ($p=1.03$).  In order to test how the fits depend on metallicity, we computed a new series of models by adopting $Z=0.015$, and He content of $Y=0.27$ (see Table~2). Data plotted in the left panels of Fig. 2 show that the new best fit model agrees reasonably well with the observations. However, the discrepancy in the radial velocity is $1.6\\sigma$ and in the radius amplitude is $2.5\\sigma$. This finding indicates that a decrease in the metal content improves the agreement between theory and observations. Therefore, we constructed two new sequences of pulsation models by adopting a more metal-poor chemical composition ($Z=0.01$) and two different He contents ($Y=0.26, 0.24$). The new best fit models (see Table~2, middle and right panels of Fig. 2) account quite well for luminosity, radial velocity, and radius variations. A marginal difference between theory and observations is still present. However, the discrepancy is on average smaller than $1\\sigma$ for $Y=0.26$ and systematically larger for $Y=0.24$ (see Table~2). The comparison between theory and observations brings forward several findings. {\\em i)} the more metal-poor ($Z=0.01$) best fit models account for the observed light and velocity amplitudes at a level better than $0.8\\sigma$ and the average discrepancy over the entire cycle is $< 0.2\\Sigma$ (see Table~2). {\\em ii)} The fit of the radial velocity curve is very sensitive to the input parameters. The small ``bump`` located along the decreasing branch ($\\phi=0.2$) mainly depends on the mass value, and it is independent of the chemical composition. Moreover, the radial velocity amplitude strongly depends on the He content. For example, an increase/decrease of $\\sim 8$\\% in He content causes, at Z=0.02, an increase on the amplitude discrepancy of at least $1.4\\sigma$. The more metal-rich best fit models (Y=0.28, 0.25) to properly fit the velocity amplitude would require $p$-factor values that are at least $\\sim 20$\\% smaller than the empirical measurement by \\citet{MK062}. Therefore, they do not pass the empirical validation. {\\em iii)} Predicted radius variations are at least $1.7\\sigma$ smaller than observed while the average discrepancy over the entire curve is never smaller than $0.2\\Sigma$ for all the adopted chemical compositions. The mismatch is mainly caused by the few points located across the phases of minimum radius. We do not have a firm explanation for this discrepancy. These phases are the most rapid ones along the pulsation cycle and the radial displacements of the outermost layers might be affected either by shock waves or by velocity gradients \\citep{Ma06}. This critical point deserves further investigations, since radius measurements depend on the adopted distance and on limb darkening correction, but they are affected neither by the interstellar extinction nor by the $p$-factor. Therefore, they can play a key role in constraining theoretical predictions. {\\em iv)}  The model with $Z = 0.01$ and $Y = 0.26$ provides the optimum fit to all observables (except for the only three magnitudes on the rising branch of the V light curve). This finding taken at face value indicates that the He content might be 8\\% smaller than typically assumed to construct Galactic Cepheid models. The metal content of $\\delta$ Cephei seems to be $0.1-0.2$ dex more metal-poor than suggested by spectroscopy, but this constraint is less robust. Note that the $p$-factor adopted for this model ($p=1.28$) agrees quite well with the measurement by M\\'erand et al. (2006). {\\em v)} The best fit models present a luminosity excess (see column 7 in Table~2) that ranges from $\\Delta Log L/L_\\odot=0.07$ (Z=0.01, Y=0.26) to 0.17 (Z=0.02, Y=0.24). These values are between the luminosities predicted by a canonical ML relation \\citep{BCCM00} and those predicted by mild core overshooting ML relation ($\\Delta Log L=0.25$, Chiosi et al. 1993).  As a further validation of the current approach, we also estimated the absolute distance to $\\delta$ Cephei using the best fit models that passed the empirical validations. In order to provide a plausible estimate of the errors affecting distance determinations we assumed the largest uncertainties on the adopted input parameters: $\\Delta M=0.1M_\\odot$, $\\Delta T=50 K$, $\\Delta LogL=0.02$, $\\Delta R=0.5R_\\odot$. By accounting for the entire error budget the intrinsic uncertainties on current distance determinations is about $3\\%$ for photometric distances and $\\approx 5$ pc for the distances derived from the fit of radius variations. Distance estimates based on the apparent distance moduli provided by the fit of both optical and NIR light curves and on the reddening law provided by \\citet{CCM89} are listed in columns 6 and 7 of Table~3. The comparison with the HST parallax shows that distances based on $V,I$-bands agree at the $0.2\\sigma$ level, while those based on the $V,K$-bands are within $0.8\\sigma$ level. Errors either in the photometric zero-points or in the adopted color-temperature transformations or in the extinction law account for the difference in the distance determinations based on light curves. Moreover, the occurrence of a circumstellar envelope (CSE) around $\\delta$ Cephei \\citep{MK062} could also affect the reddening corrections due to the different relative distribution of the absorbing grains around the star, compared to typical interstellar medium distribution. Independent distance determinations follow from comparing the observed angular diameter variations by \\citet{MK05} with the radius curves from four different models. The distances (Table~3, column 8) agree well with each other and are on average only 5\\% ($1.3 \\sigma$) larger than the trigonometric distance. The comparison of observed and model radii is in principle a powerful method to derive Cepheid distances, since it is independent of the assumption on the $p$-factor and on interstellar extinction. Note that an increase in the distance would also imply an increase in the $p$-factor, since at fixed radial velocity amplitude, $p\\propto \\Delta \\Theta\\cdot d$, where $\\Delta \\Theta$ is the angular diameter amplitude.  Finally we mention that the mean radii predicted by the best fit models match within $1\\sigma$ the empirical estimate based on the angular diameter provided by ME06 ($R=43\\pm2 R_\\odot$), by assuming the HST parallax."
    },
    "0711/0711.2066_arXiv.txt": {
        "abstract": "We present results from a systematic investigation of the X-ray properties of a sample of moderate redshift ($0.3<z<0.6$) galaxy groups. These groups were selected not by traditional X-ray or optical search methods, but rather by an association, either physical or along the line of sight, with a strong gravitational lens.  We calculate the properties of seven galaxy groups in the fields of six lens systems. Diffuse X-ray emission from the intragroup medium is detected in four of the groups.  All of the detected groups have X-ray luminosities greater than $10^{42} h^{-2}$ erg s$^{-1}$, and lie on the $L_X$--$\\sigma_v$ relations defined by local groups and clusters.  The upper limits for the non-detections are also consistent with the local $L_X$--$\\sigma_v$ relationships.  Although the sample size is small and deeper optical and X-ray data are needed, these results suggest that lens-selected groups are similar to X-ray selected samples and, thus, are more massive than the typical poor-group environments of local galaxies. ",
        "introduction": "To have a full picture of galaxy evolution and structure formation in the Universe, it is crucial to understand the properties of galaxy groups.  Most galaxies in the local Universe reside in galaxy groups \\citep[e.g.,][]{turnergott,gellerhuchra,eke04}. In addition, groups are vitally important in driving changes in star formation rates and galaxy morphologies since $z \\sim 1$, because the low velocity dispersions and high density of groups make them likely locations for interactions and mergers \\citep[e.g.,][]{aarsethfall,barnes85,merritt85}.  Furthermore, indications are that the dark matter distributions in groups represent a transition between the dark-matter dominated profiles seen on cluster scales and galaxy-sized halos that are strongly affected by baryon cooling \\citep[e.g.,][]{oguri05}.  In this paper we examine the properties of moderate-redshift groups.  Groups in the local Universe have been well studied \\citep[e.g.,][]{zm98,mz98,gemsxray,rasmussen06}. In a systematic survey of 60 groups in the Group Evolution Multiwavelength Study (GEMS), \\citet{gemsxray} find that groups in the local Universe obey similar $L_X$--$\\sigma_v$ and $L_X$--$T$ scaling relations to more massive clusters, although a more recent analysis \\citep{gems_refit} finds a steepening in the $L_X$--$T$ slope for these groups.  There is also a large non-statistical scatter of a factor of 30 in X-ray luminosity and a factor of 3--4 in X-ray temperature for this group sample.  This scatter includes a class of spiral-rich groups with little or no emission, as well as some groups with high X-ray luminosities but low velocity dispersions. Possible explanations for these outliers include unrelaxed dynamical states, uncertain velocity dispersion measures from small numbers of redshifts \\citep[typically fewer than 10, see also][]{zm98}, and point source contamination of the X-ray flux \\citep{helsdon05}. \\citet{gemsxray} also find a strong correlation between the detection of diffuse emission and the presence of a centrally-located, dominant early-type galaxy, as well as an anti-correlation between spiral fraction and gas temperature.  Additional outliers from the canonical scaling relations are found by \\citet{rasmussen06}, who are studying a redshift-selected, statistically-unbiased sample of galaxy groups at $z \\approx 0.06$ with deep X-ray data. The first galaxy groups detected from this survey are X-ray faint ($\\sim 5 \\times 10^{40}~h^{-2}$ erg s$^{-1}$), are underluminous for their measured velocity dispersions, and do not host a dominant, central early-type galaxy, suggesting that they are collapsing for the first time. Optically-selected groups such as these represent a different, perhaps more common, class of groups than those detected through their X-ray emission.  Detailed studies of groups at moderate-to-high redshifts have been limited because groups are difficult to find, given their modest galaxy overdensities and X-ray luminosities.  On the optical side, the situation has been alleviated somewhat by large redshift surveys such as the Sloan Digital Sky Survey \\citep[SDSS;][]{sdss}, the Canadian Network for Observational Cosmology Field Galaxy Survey \\citep[CNOC2;][]{cnoc2}, and the DEEP2 survey \\citep{deep2}, but many of the group candidates found in these surveys have only 3-4 members \\citep[e.g.,][]{cnoc2groups1,deep2groups} and may not represent physically bound structures.  Furthermore, the intensive spectroscopic followup required to confirm each group candidate and characterize its properties has so far limited the size of well-studied group samples from these surveys \\citep[e.g.,][]{wilmangrp1,balogh07,fang07}.  Intermediate-redshift ($0.2 < z < 0.6$) groups selected on the basis of their X-ray emission are also being studied \\citep{xmmlss_groups,xmmlss_erratum,mulchaey96,jeltema06,jeltema07}. Results for these surveys suggest that these groups follow the low-redshift scaling relations between luminosity, temperature, and velocity dispersion. Many of these systems have high fractions of early-type galaxies, suggesting that this population is already in place by $z \\sim 0.5$ \\citep{mulchaey96,jeltema07}. However, there are clear indications of dynamical evolution, including many groups with brightest group galaxies that show multiple components, dominant early-type galaxies that are not centered on the diffuse emission, or no dominant early-type galaxies at all. While similar systems have been identified in low-redshift samples (see above), the luminosities and temperatures of the moderate-redshift examples are significantly higher, implying a group downsizing where more massive groups are still in the process of collapsing and virializing at these redshifts. The effects of evolution appear to continue to higher redshifts where studies of some optically-selected groups suggest that they are substantially and systematically underluminous, relative to their local counterparts, for a given velocity dispersion \\citep{fang07}.  Note however, that the velocity dispersions in \\citet{fang07} are all based on a small number of measured redshifts (ranging from 3 to 6 redshifts per group) and thus are highly uncertain and may be significantly overestimated.  Significant differences between the properties of X-ray and optically-selected groups have been noted at both low and intermediate redshift \\citep[e.g.,][]{rasmussen06,fang07,rykoff08}, suggesting that each method is selecting a distinct class of groups. To avoid these biases, we are conducting a survey of moderate-redshift groups that have been selected via a non-traditional technique, namely through their association with strong gravitational lenses.  There is growing evidence that strong gravitational lenses, i.e., those forming multiple images of the background object, are typically located in groups of galaxies at intermediate redshift \\citep[e.g.,][]{kundic1115,kundic1422,tonrygroups,momcheva,1608groupdisc,auger_groups1,auger1520}. Thus, strong lenses can be used to identify and study the properties of distant groups, selected in a manner that is completely independent from alternative techniques such as deep X-ray integrations or color-magnitude diagrams.  The lensing probability depends only on the projected mass distribution and does not depend on its kinematics or on the properties of the galaxy population or intragroup medium (IGM). Therefore, lens-selected groups provide an excellent sample to determine the properties of galaxies and hot baryons, and to understand selection effects (e.g., mass concentration for lensing, IGM luminosity and temperature for X-ray selection, and the homogeneity of the galaxy population for red sequence selection) by contrasting different methods.  It is not unexpected that the lensing galaxies in strong lens systems should reside in overdense regions of the Universe.  Searches for gravitational lenses are biased toward high density regions because (1) higher mass systems have a larger cross-section for lensing and (2) most lensing galaxies are ellipticals, which are preferentially found in groups or clusters \\citep[e.g.,][]{morphdensity,zm98}. Theoretical studies have predicted that a significant number of lens systems should be associated with groups or clusters, albeit with a large spread of values \\citep[]{keeton00,holderschechter,oguri_etal05}.  Unbiased photometric surveys of lens fields indicate that lenses lie along overdense lines of sight \\citep[][Fassnacht et al.\\ in prep.]{williams06}, and spectroscopic observations have confirmed several lens-group associations that can affect the lensing potential at the level of 5\\% or more \\citep[e.g.,][]{kundic1115,kundic1422,tonrygroups,0712group,momcheva,1608groupdisc,auger_groups1,auger1520}.   In this paper, we present the results from multi-object spectroscopic and deep {\\em Chandra} observations of seven groups detected through their association with gravitational lenses.  We assume a cosmological model with $(\\Omega_{\\rm M},\\Omega_\\Lambda) = (0.3,0.7)$.  We will use $h = H_0 / 100$\\kmsmpc\\ to represent the Hubble Constant when we do not have to choose a value.  When we do have to assign a value to the Hubble Constant, we use $h = 0.7$. ",
        "conclusions": "\\subsection{Comparison to Previous X-ray Analyses}  The diffuse X-ray emission from two of the lens-selected groups, PG1115 and B1422, has been analyzed previously by \\citet{grant_chandra}.  Because our analysis uses different inputs and techniques, care should be taken when comparing the results of the two analyses.  As one example, each of these two systems was observed under two separate programs (Obs IDs), namely 363 and 1630 for PG1115, and 367 and 1631 for B1422.  To simplify the analysis of the effects of the PSF, we used only the longer of the two observations for each system.  In contrast, the \\citet{grant_chandra} analysis combined the two programs in each case.  Therefore, the germane basic quantity to use in the comparison of the results is not the net counts from the diffuse emission, but rather the net count rates.  In each case, we find a higher count rate than that measured in the previous analysis: $2.7 \\times 10^{-3}$ vs.\\ $1.3 \\times 10^{-3}$ for PG1115 and $3.5 \\times 10^{-3}$ vs.\\ $1.8 \\times 10^{-3}$ for B1422, where all values are in counts s$^{-1}$.  We believe that the cause of this discrepancy is the manner in which we masked out the lensed AGN emission for the lenses.  The \\citet{grant_chandra} approach was conservative, masking out regions with diameters of 14\\arcsec--16\\arcsec, while our 99\\% masking regions (Fig.~\\ref{fig_rawximg}) cover somewhat smaller areas, with typical sizes of $\\sim$10\\arcsec\\ across.  In both cases, a correction was then made to account for the diffuse flux within the masked region. Given the small angular extent of the groups, changes in the masking area may lead to differences in the size of the correction and therefore to significant changes in the measured flux of the diffuse emission.  A second difference in technique is that the luminosities that we use in our final analysis are those within $R_{500}$, rather than just using the luminosities calculated from the observed region of significant detection. To facilitate comparison with the \\citet{grant_chandra} results, the bolometric rest-frame luminosities in Table~\\ref{tab_xray} are given for both the detection region and the $R_{500}$ region.  Other differences arise from the X-ray temperature used for PG1115, where \\citet{grant_chandra} use 0.7~keV and we use 1.0~keV, and the assumed cosmology.  The \\citet{grant_chandra} results are computed using an Einstein-de Sitter cosmology with $h = 0.5$, while we use $(\\Omega_{\\rm M},\\Omega_\\Lambda) = (0.3,0.7)$ and set $h = 0.7$ when we have to fix its value.   It is important to consider the effects that the analysis techniques will have on the interpretation of the results.  In terms of many of the conclusions drawn from the measured X-ray luminosities, including the discussion of the $L_X$--$\\sigma_v$ relationship below, the factors of a few in $L_X$ that come out of the different techniques will not change the interpretation.  In particular, given the log-log nature of the $L_X$--$\\sigma_v$ plot and the scatter in the observed relationship, these factors of 2--3 do not significantly change the location of the lens-group points in Figure~\\ref{fig_lx_sig}.  On the other hand, the determination of the location of the X-ray centroid is more sensitive to how the lensed AGN emission is masked, and therefore conclusions about the possible offsets of the brightest group galaxies from the center of the X-ray emission are tentative at best.  \\subsection{The Nature of the X-ray Emission in B2108}  The X-ray emission associated with the B2108 system is shown in the last row of Figure~\\ref{fig_smoothximg}, where its morphology is seen to be clearly different from that of the other systems.  While the smoothed data in each of the other fields show a single, relatively compact region, the B2108 field contains two elongated regions, connected by a bridge of emission.  These regions are roughly centered on the positions of masked point sources; the western component (B2108-1) is the one centered on the lens system.  It is likely that the X-ray emission in B2108-1 is produced by the the group associated with the B2108+213 lens system.  The lensing galaxy is a massive elliptical, is the brightest galaxy in the group, and is found nearly at the center of the X-ray emission.  However, the source of emission in B2108-2, the eastern X-ray component, is much less clear. An examination of optical imaging of this field, obtained with the Keck telescope, shows a $R>23$ source coincident with the X-ray position of the point source that is roughly at the center of B2108-2. The faintness of this optical object compared to the confirmed group members ($\\sim 5$ magnitudes fainter than the primary lensing galaxy) suggests that it is background AGN rather than a source associated with a massive elliptical in the $z = 0.36$ group.  If this is the case, then it is not clear whether the diffuse X-ray emission in B2108-2 is coming from a massive system associated with a background AGN or whether it is instead due to a second group at the redshift of the lensing galaxy. The clear removal of the point-source emission for the B1600 and B1608 systems suggests that the diffuse emission is not wholly due to residual emission from the point source. The majority of the B2108-2 emission is outside the region for which spectroscopic data have been obtained, so there is no information on whether there may be an overdensity of higher-redshift objects in the area covered by this second system.  That being said, the available spectroscopy does provide some insights on this complex field. The velocity distribution of the $z = 0.36$ group members is non-Gaussian, even with nearly 50 redshifts (McKean et al., in prep), suggesting that the group is in a dynamically disturbed state. The velocity distribution in the B2108 field and the presence of multiple elongated diffuse components connected by an apparent bridge of emission suggest that system is undergoing some kind of merger. Interestingly, the east-west alignment of the two diffuse X-ray components is roughly in the same direction as the low surface brightness lobes seen in deep radio imaging of this system \\citep{more2108}, although the radio lobes cover a much smaller area than the X-ray emission.  \\subsection{Detection Rate and Group Luminosities}  One of the most basic quantities that emerges from the analysis of the lens-group sample is the rate at which these lens-selected groups are detected when observed at X-ray frequencies.  We have examined the fields of six lenses, in which at least seven galaxy groups have been discovered using optical spectroscopy.  Of those seven groups, four are detected with the {\\em Chandra} observations, giving a formal detection rate of $\\sim$60\\% $\\pm$30\\%. All of the detected lens-selected groups have luminosities within $R_{500}$ of greater than $10^{42}\\ h^{-2}$ erg s$^{-1}$. Because most previous investigations of the X-ray properties of galaxy groups have been conducted at low redshifts, we compare the rate at which X-ray emission from the IGM is detected in the lens sample to these samples.  \\citet{mahdavi00} selected groups and clusters based on the CfA redshift survey and found a detection rate of 23\\%, which they correct to a 40\\% rate based on their X-ray selection function.  The X-ray data used for this work were from the ROSAT All-Sky Survey \\citep[RASS;][]{rass} and thus are not highly sensitive, with a limiting luminosity of $\\sim$10$^{42}$ erg s$^{-1}$. \\citet{gemsxray} concentrate more specifically on groups and detect $\\sim$60\\%, based on deeper pointed X-ray observations with ROSAT.  For this sample, the limiting luminosity is $L_X \\sim 10^{41}$ erg s$^{-1}$.  However, analysis and sample-selection issues make a straight comparison of detection rates in these earlier samples somewhat problematic.  The shallow X-ray data used in \\citet{mahdavi00} make it difficult to disentangle IGM emission from emission due to a hot halo associated with one of the group galaxies \\citep[e.g.,][]{mulchaey_araa,gemsxray}.  The \\citet{gemsxray} sample avoids this problem by using deeper ROSAT pointings, but their sample may be biased.  Their groups are optically-selected, but they require their groups to have pre-existing deep ($t_{\\rm exp} > $10,000 s) ROSAT data, which were often available because the RASS data had shown an X-ray source at that location.  Perhaps the best comparison sample for our lens-selected groups is the work on local optically-selected groups by \\citet{rasmussen06}.  Here, the selection is unbiased, at least with respect to the X-ray properties, and the X-ray observations, which are made with XMM, are sensitive. The IGM has been detected in two out of the four groups, giving a similar detection rate to the lens-selected samples.  We note, however, that none of the groups in this sample have luminosities above $10^{42}\\ h^{-2}$ erg s$^{-1}$. Although both the lens-group and \\citet{rasmussen06} samples are small, we may be seeing a difference in the sample properties. The optically-selected sample may be more representative of typical low-redshift galaxy groups, which often have low levels of X-ray emission.  In contrast, the lens-selected sample is picking out groups that are more like the X-ray--bright groups detected in local samples, which tend to be on the high-mass side of the local group distributions \\citep[e.g.,][]{mahdavi00,mz98}.  \\subsection{$L_X$--$\\sigma_v$ relationship}  In order to compare the properties of the X-ray emission from the lens-selected moderate-redshift groups with other group samples, we have assumed $h = 0.7$ and $(\\Omega_M,\\Omega_\\Lambda) = (0.3,0.7)$. The lens-group points are plotted as the large stars on the $L_X$--$\\sigma_v$ plot in Figure~\\ref{fig_lx_sig}. For the purposes of this plot, we have assumed that all of the measured redshifts in the B2108 field are associated with the B2108-1 group, and plotted the derived velocity dispersion against the B2108-1 X-ray luminosity. Also included in the plot are data from the low-redshift X-ray--detected groups in the GEMS sample of \\citet{gemsxray}, and from X-ray--selected moderate-redshift group samples \\citep{jeltema06,xmmlss_groups}.  The filled squares represent data on the four groups from \\citet{rasmussen06}.  As mentioned above, these optically-selected groups appear to be less X-ray luminous than typical systems in the X-ray--selected samples (Fig.~\\ref{fig_lx_sig}). In contrast, the lens-selected groups appear to be consistent with the X-ray--selected samples, given the scatter in the data.  \\ifsubmode % \\clearpage \\begin{figure} \\else % \\begin{figure*} % \\fi % \\plotone{f8.eps} \\caption{ $L_X$--$\\sigma_v$ plot showing moderate-redshift lens-selected groups from this paper (large stars, with open stars representing upper limits), low-redshift X-ray--selected groups from GEMS \\citep[][error bars with no points]{gemsxray}, moderate-redshift X-ray--selected groups from \\citet[][open circles]{jeltema06} and \\citet[][open triangles]{xmmlss_groups}, and low-redshift optically-selected groups from 2dF \\citep[][filled squares]{rasmussen06}.  The solid line represents the fit to the GEMS sample \\citep{gems_refit}, while the dotted line is the fit to the low-redshift cluster data of \\citet{markevitch98} and \\citet{horner_thesis}.  The long-dashed line is a fit to a sample of intermediate-redshift groups and clusters \\citep{modz_lxsigfit}. } \\label{fig_lx_sig} \\ifsubmode % \\end{figure} \\clearpage \\else % \\end{figure*} % \\fi %  Also shown in Fig.~\\ref{fig_lx_sig} are fits to low- and intermediate-redshift samples.  The low-redshift scaling relations are shown for the GEMS group sample \\citep[solid line;][]{gems_refit} and a cluster sample \\citep[short-dashed line;][]{markevitch98,horner_thesis}.  The other scaling relation (long-dashed line) in the plot is a fit made by \\citet{modz_lxsigfit} to data from an intermediate-redshift group sample, the majority of which is comprised of the \\citet{jeltema06} and \\citet{xmmlss_groups} data points represented as the open circles and triangles in Fig.~\\ref{fig_lx_sig}.  It is clear that the lens-group detections are more consistent with the low-redshift scaling relations than the intermediate-redshift relation.  However,  \\citet{modz_lxsigfit} do point out that their fitted slope may be flatter than it should be because the velocity dispersion measurements may be biased low. Of course, it is difficult to draw strong conclusions from only four detections in the lens-group sample, but the upper limits, while not constraining the properties of the lens-selected sample in a meaningful way, are still perfectly consistent with the low-redshift scaling relations and marginally consistent with the intermediate-redshift scaling relation.  It should be noted that the three lens-selected groups that were not detected all lie at $z>0.4$, while the four that were detected all have $z < 0.4$.  Deeper X-ray observations of the B1608 and B1600 fields, which now each have only $\\sim$30~ks of {\\em Chandra} data, would enhance the conclusions that can be drawn from the lens group sample.  \\subsection{BGG offsets}  In low-redshift samples, the brightest group galaxy (BGG) nearly always sits at the spatial and dynamical center of the group \\citep[e.g.,][]{zm98}, and the diffuse X-ray emission from group is also centered on the BGG \\citep[e.g.,][]{mz98,helsdonponman00,gemsxray}.  The coincidence between the diffuse X-ray gas and the BGG becomes less prominent in samples of X-ray--selected moderate redshift groups, where some groups have negligible BGG offsets but others have offsets up to $\\sim160 h^{-1}$ kpc \\citep{jeltema06,jeltema07}.  The lens-selected systems presented in this paper span a redshift range similar to the \\citet{jeltema07} sample and thus provide an interesting comparison sample.  We have identified the brightest spectroscopically-confirmed member of each of the groups for which diffuse X-ray emission has been detected, either from our own imaging (for B0712 and B2108) or from the data in \\citet[][for PG1115 and B1422]{momcheva}. The X-ray centroids are derived from the Gaussian-smoothed images in the last column of Figure~\\ref{fig_smoothximg}.  For the B2108 field, we calculate the offset between the BGG and the centroid of B2108-1. The offsets between the BGGs and the X-ray centroids are given in Table~\\ref{tab_offsets}, and overlays of the X-ray contours on HST images are shown in Figure~\\ref{fig_xray_opt_overlay}.  For B2108 the BGG is also the lensing galaxy.  However, in B0712 (because the group lies at a lower redshift than the lens), PG1115, and B1422 the lens is not the BGG.  Therefore, Table~\\ref{tab_offsets} also gives the offsets between the lensing galaxies and the X-ray centroids.   \\ifsubmode % \\clearpage \\begin{deluxetable}{lllrrllrr} \\else % \\begin{deluxetable*}{lllrrllrr} % \\fi % \\tabletypesize{\\scriptsize} \\tablecolumns{7} \\tablewidth{0pc} \\tablecaption{Offsets from X-ray Centroid} \\tablehead{ \\colhead{} & \\colhead{BGG} & \\colhead{BGG} & \\colhead{Offset} & \\colhead{Offset} & \\colhead{Lens Galaxy} & \\colhead{Lens Galaxy} & \\colhead{Offset} & \\colhead{Offset} \\\\ \\colhead{Lens System} & \\colhead{RA} & \\colhead{Dec} & \\colhead{(arcsec)} & \\colhead{($h^{-1}$ kpc)} & \\colhead{RA} & \\colhead{Dec} & \\colhead{(arcsec)} & \\colhead{($h^{-1}$ kpc)} } \\startdata B0712   & 07:16:05.01 & +47:09:04.8 & 16 & 49 & 07:16:03.58 & +47:08:50.0 & 9\\tablenotemark{a} & \\nodata\\tablenotemark{a} \\\\ PG1115  & 11:18:15.52 & +07:45:47.7 & 16 & 51 & 11:18:17.00 & +07:45:57.7 & 10\\phantom{$^a$} & 31 \\\\ B1422   & 14:24:38.39 & +22:55:53.5 &  7 & 24 & 14:24:38.09 & +22:56:00.6 & 1\\phantom{$^a$} &  3 \\\\ B2108-1 & 21:10:54.03 & +21:31:00.4 &  7 & 25 & 21:10:54.03 & +21:31:00.4 & 7\\phantom{$^a$} & 25 \\\\ \\enddata \\label{tab_offsets} \\tablenotetext{a}{The lens is not physically associated with the group.} \\ifsubmode % \\end{deluxetable} \\clearpage \\else % \\end{deluxetable*} % \\fi %  \\ifsubmode % \\clearpage \\begin{figure} \\else % \\begin{figure*} % \\fi % \\plotone{f9.eps} \\caption{ Optical imaging from the HST of the X-ray--detected groups, with contours representing the diffuse X-ray emission overlaid.  For the B0712, PG1115, and B2108 images, the data were taken with the F814W filter. For the B1422 image, the data were taken with the F791W filter. In each image the BGG is marked with a circle and the lens system is marked with a box.  Note that the lens system in the B0712 group is not physically associated with the group. } \\label{fig_xray_opt_overlay} \\ifsubmode % \\end{figure} \\clearpage \\else % \\end{figure*} % \\fi %  The offsets between the centroids of the diffuse X-ray emission and the BGGs are small to moderate, with all of the offsets being $\\leq$20\\arcsec.  At the redshifts of these groups, $z \\sim 0.3 - 0.4$, these offsets correspond to physical distances of $\\sim$25--50~$h^{-1}$~kpc.  For comparison, five out of the seven groups in the sample of \\citet{jeltema07} had BGG offsets of $<$15~$h^{-1}$~kpc, while the other two had offsets of 110--160~$h^{-1}$~kpc. In spite of the BGG--X-ray offsets measured in the lens-group sample, we cannot yet conclude that this sample is similar in character to the \\citet{jeltema07} one.  This is because slight changes in analysis techniques can lead to significant changes in the centroid position of lens-selected groups, as can be seen by comparing our X-ray images of PG1115 and B1422 to those in \\citet{grant_chandra}.  Even with the excellent angular resolution of {\\em Chandra}, the determination of the morphology of faint diffuse emission in the presence of bright lensed AGN images is challenging.  Over- or under-subtracting the AGN emission can bias the centroid position. Furthermore, the choice of smoothing technique can also lead to shifts in the derived X-ray centroid. Deeper X-ray imaging of all of the systems would help to address these centroiding issues.  Another complexity that enters into the determination of the BGG--X-ray offsets is the difficulty for some of the systems in identifying the BGG.  For many of the groups there is no clearly dominant central galaxy.  Instead, in several of the groups (B0712, B1608-3, B2108, and possibly PG1115) the brightest and next brightest galaxies are separated by less than half a magnitude.  Other investigations of intermediate-redshift groups \\citep{mulchaey06,jeltema06} have found systems for which the IGM is detected in X-rays, but no dominant BGG was identified.  For our lens-group sample, the lack of a dominant BGG in some groups may be due to seeing these groups at an earlier stage of their evolution \\citep[e.g.,][]{jeltema07}, or it may just be due to incomplete spectral information where the BGG has not yet been identified as a group member.  Deeper optical data are needed to resolve this issue.  \\subsection{Implications for Gravitational Lensing}  When a gravitational lens system resides in an overdense environment, that environment can impact the lensing potential and have clear effect on cosmological and astrophysics measurements made using the lens \\citep[e.g.,][]{keeton_group}.  An advantage of X-ray observations over spectroscopic surveys of the lens environments is that, in theory, the X-ray data can produce clear determinations of the center of the group potential and the total group mass (through the temperature of the X-ray gas).  However, the existing {\\em Chandra} data presented here are not sensitive enough to make robust measurements of either the centroids or the temperatures of the X-ray gas distributions.  \\subsection{Galaxy Properties}  In low-redshift groups, correlations have been found between the morphological distribution of the member galaxies and the overall group properties \\citep[e.g.,][]{zm98,mz98,gemsxray}. In particular, the fraction of early-type galaxies shows a significant correlation with such properties as the group velocity dispersion and whether the group has detectable diffuse X-ray emission.  At the redshifts of the lens-selected sample, it is not possible to use ground-based imaging to robustly determine galaxy morphologies.  Therefore, we have used HST imaging to determine morphological properties of the group members, where possible.  For nearly all of the systems, we used archival WFPC2 or ACS imaging (GO-5699, PI: Impey; GO-5908, PI: Jackson; GO-6555, PI: Schechter; GO-6652, PI: Impey; GO-7495, PI: Falco; GO-8628, PI: Impey; GO-9133, PI: Falco; GO-9138, PI: Impey; GO-9744, PI: Kochanek).  Most of these archival data sets were obtained by the CfA-Arizona Space Telescope Lens Survey team.  For B1608 we used our own deep ACS imaging (GO-10158: PI: Fassnacht). Table~\\ref{tab_f_e} lists the proportion of galaxies identified by morphology as early-type (E or S0) in the lens-selected groups. All of the groups for which diffuse X-ray emission has been detected have relatively high early-type fractions, ranging from $\\sim$50--70\\%.  For galaxy groups in the local Universe, the groups for which a diffuse intragroup X-ray component is detected also have significant early-type fractions.  In contrast, the group associated with B1600, which has no detected X-ray emission, has only one early-type galaxy among its confirmed members.  The B1608 groups are mixed, with B1608-1 having a low early-type fraction and B1608-3 having a high fraction.  It may be that diffuse X-ray emission was not detected in the B1608-3 group simply because the group is at a high redshift and the observations were not particularly deep.  It should be noted that the results in Table~\\ref{tab_f_e} must be considered with caution.  First, the number of member galaxies that have high-resolution imaging is small because of the fields of view of the HST cameras are limited.  Secondly, the fractions may well be biased because the small fields are centered on the lens galaxy and may be probing special regions of the groups, such as the group centers.  \\ifsubmode % \\clearpage \\fi % \\begin{deluxetable}{lrrrl} \\tabletypesize{\\scriptsize} \\tablecolumns{5} \\tablewidth{0pc} \\tablecaption{Group Early-type Fractions} \\tablehead{ \\colhead{Lens System} & \\colhead{$N_{\\rm tot}$} & \\colhead{$N_{HST}$} & \\colhead{$N_{\\rm E/S0}$} & \\colhead{$f_e$} } \\startdata B0712   & 15 &  9 &  5 & 0.56 \\\\ PG1115  & 13 &  7 &  5 & 0.71 \\\\ B1422   & 16 &  7 &  5 & 0.71 \\\\ B1600   &  7 &  7 &  1 & 0.14 \\\\ B1608-1 & 10 &  9 &  2 & 0.22 \\\\ B1608-3 &  8 &  7 &  6 & 0.85 \\\\ B2108   & 47 & 24 & 15 & 0.62 \\\\ \\enddata \\label{tab_f_e} \\end{deluxetable} \\ifsubmode % \\clearpage \\fi %  \\subsection{The Nature of the Lens-selected Groups}  Although the lens-selected group sample is small, the trends from the sample are suggestive.  As discussed in the previous sections, the systems for which diffuse X-ray emission has been detected appear to be quite similar to local X-ray--loud samples in their luminosities, their locations on the $L_X$--$\\sigma$ plot, and their early-type fractions.  This may not be totally unexpected, as the lensing galaxies in strong lens systems tend to have early-type morphologies, and local groups containing massive ellipticals are more likely to have detectable X-ray emission \\citep[e.g.,][]{mulchaey96,mz98}. In fact, the one group in the lens-selected sample in which the lensing galaxy is clearly a spiral also has the lowest early-type fraction and is one of the X-ray non-detections.  Thus, while the B1600 group and possibly the B1608-1 group appear to be like the typical group found in the local Universe, with low masses and low early-type fractions \\citep[e.g.][]{gellerhuchra,eke04}, the majority of the lens-group sample is more similar to the local groups for which X-ray emission from the IGM has been detected, which are more massive and dominated by early-type galaxies.  The one lens group that does not fit this picture is the B1608-3 group, with its high early-type fraction and a relatively high velocity dispersion.  The non-detection of B1608-3 in the {\\em Chandra} observations may be due to a combination of its redshift and the short exposure time used in the observations.  We feel that this is the group most likely to be detected if deeper observations of the field are undertaken."
    },
    "0711/0711.3689_arXiv.txt": {
        "abstract": "Differences between the inferred star formation histories (SFHs) of star clusters and field stars seem to suggest distinct star formation processes for the two. The Small Magellanic Cloud (SMC) is an example of a galaxy where such a discrepancy is observed. We model the observed age distributions of the SMC clusters and field stars using a new population synthesis code, SPACE, that includes stellar evolution, infant mortality and cluster dissolution. We find that the two observed age distributions can be explained by a single SFH, thus eliminating the need to assume two separate mechanisms for star formation. ",
        "introduction": "We consider the relation between age distributions of clusters and field stars. It has been suggested that separate formation mechanisms for clusters and field stars exist in galaxies \\citep[e.g.,][]{meurer95,chandar05}, implying different formation histories for the two. However, this appears to contradict theory and observations, as most stars are supposed to be formed in clusters.  Observations of the Small Magellanic Cloud (SMC) show a significant difference between the age distributions of clusters and field stars \\citep{hodge87,chandar06}. The SMC therefore provides an ideal opportunity to investigate whether the observed distributions can be reproduced by one star formation history (SFH) if all stars were formed in clusters. ",
        "conclusions": ""
    },
    "0711/0711.3640_arXiv.txt": {
        "abstract": "We study the large-scale anisotropic two-point correlation function using 46,760 luminous red galaxies at redshifts 0.16 -- 0.47 from the Sloan Digital Sky Survey. We measure the correlation function as a function of separations parallel and perpendicular to the line of sight in order to take account of anisotropy of the large-scale structure in redshift space. We find a slight signal of baryonic features in the anisotropic correlation function, i.e., a ``baryon ridge'' corresponding to a baryon acoustic peak in the spherically averaged correlation function which has already been reported using the same sample. The baryon ridge has primarily a spherical structure with a known radius in comoving coordinates. It enables us to divide the redshift distortion effects into dynamical and geometrical components and provides further constraints on cosmological parameters, including the dark energy equation-of-state. With an assumption of a flat $\\Lambda$ cosmology, we find the best-fit values of $\\Omega_m = 0.218^{+0.047}_{-0.037}$ and $\\Omega_b = 0.047^{+0.016}_{-0.016}$ (68\\% CL) when we use the overall shape of the anisotropic correlation function of $40<s<200\\himpc$ including a scale of baryon acoustic oscillations. When an additional assumption of $\\Omega_bh^2=0.024$ is adopted, we obtain $\\Omega_{\\rm DE}=0.770^{+0.051}_{-0.040}$ and $w=-0.93^{+0.45}_{-0.35}$.  These constraints are estimated only from our data of the anisotropic correlation function, and they agree quite well with values both from the cosmic microwave background (CMB) anisotropies and from other complementary statistics using the LRG sample.  With the CMB prior from the 3 year \\wmap  results, we give stronger constraints on those parameters. ",
        "introduction": "Recently, baryon acoustic oscillations have been observed in the large-scale structure of the universe. These observations include an analysis of the two-point correlation function (2PCF) of the Sloan Digital Sky Survey (SDSS) Luminous Red Galaxy (LRG) spectroscopic sample \\citep[E05, hereafter]{E2005}, the power spectrum of the Two-Degree Field (2dF) Galaxy Redshift Survey \\citep{Cole2005} and the SDSS LRG \\citep{Hutsi2006a, Hutsi2006b, Tegmark2006, Percival2007a, Percival2007b}, and the angular power spectrum of the SDSS LRG sample with photometric redshifts \\citep{Padmanabhan2007, Blake2006}. There is a hint of acoustic oscillations in the SDSS quasar sample \\citep{Yahata2005}. These analyses have established the ability of the baryon acoustic oscillations to constrain cosmological parameters competitively and complementarily with the CMB \\citep[e.g.,][]{Spergel2007} and Type Ia supernovae \\citep{Riess1998, Perlmutter1999}.  These previous analyses of the baryon oscillations, however, use angle-averaged 2PCFs, $\\xi(s)$, or angle-averaged power spectra, $P(k)$, where $s$ and $k$ are the separation and wavenumber in redshift space, respectively. A certain amount of information is lost when anisotropies of structure are ignored. In their pioneering work, \\citet{AP1979} proposed that geometrical anisotropies in redshift space of the high-z universe can be used as a probe of the cosmological constant.  \\citet{MS1996} and \\citet{Ballinger1996} pointed out that the anisotropy of the 2PCF and the power spectrum in redshift surveys can constrain the dark energy components. Recently, methods which directly use anisotropy of the baryon acoustic oscillations have been theoretically developed for both the power spectrum \\citep{Hu2003, Seo2003, GB2005, Seo2007} and the 2PCF \\citep{M2004}. These approaches use functions of two variables: separations parallel and perpendicular to the line of sight. Observationally, estimations of such two-variable functions are noisier than one-variable functions. The baryon acoustic signature in the large-scale structure is weak even in one-variable statistics when presently available samples of galaxies are used (see, e.g., E05). Therefore, methods that directly treat the anisotropy of the baryon acoustic feature require very large survey volume \\citep{EHT1999,M2001}. \\citet{Peacock2001} and \\citet{Hawkins2003} measured the 2PCF with two variables from the 2dFGRS \\citep{Colless2001} and detected the detailed signature of large-scale coherent infall, and as a result were able to constrain a value of $\\beta \\simeq \\Omega_m^{0.6}/b$ which parameterizes linear redshift distortions. However, they used information from scales much smaller than the acoustic scale because of limited survey volume and did not analyze the effect of geometrical distortion. On the other hand, the correlation analyses of the 2dF QSO survey placed constraints on the cosmological constant (Hoyle et al.~2002; Outram et al.~2004; da \\^Angela et al.~2005; see also Ross et al.~2007). Their analyses still focused on smaller scales than the baryon acoustic scale.   In this paper we analyze the anisotropic 2PCF, including the baryon acoustic peak in the large-scale structure and constrain relevant cosmological parameters. We use a spectroscopic sample of the SDSS LRG, which is the most useful sample for our purpose. Our analysis differs from previous studies of baryon acoustic oscillations in the LRG sample in that, owing to a theoretical development by \\citet{M2004}, we take into consideration the fully two-dimensional feature in the 2PCF to detect a geometrical distortion effect; this is the first cosmological application of the two-dimensional acoustic peaks.  Before proceeding to the next section, we take note of nonlinearity on large scales and the scale dependence of the galaxy biasing. The nonlinear effects on the baryonic features that appear around $100\\himpc$ play an important role in taking account of percent-level cosmology \\citep[e.g.,][]{Meiksin1999, Seo2005, Jeong2006}. Recent work also suggests that the scale-dependent biasing poses a serious problem in analyzing galaxy surveys \\citep[e.g.,][]{Blanton2006, Percival2007b, Smith2007, Coles2007,Sanchez2008}. However, for the sake of simplicity, in this paper we consider only the large-scale clustering and assume the biasing to be scale-independent and linear.    The structure of this paper is as follows. In \\S~\\ref{sec:sdss} we describe the SDSS LRG sample used in our analysis. We then, in \\S~\\ref{sec:2pcf}, measure the anisotropic 2PCF in redshift space and estimate its covariance matrix. In \\S~\\ref{sec:theory} we outline the cosmological parameter dependence on the modeled 2PCF, including dynamical and geometrical distortions. Cosmological parameters are constrained by the measured anisotropic 2PCF in \\S~\\ref{sec:parameter}. In \\S~\\ref{sec:conclusion} our conclusions are given. ",
        "conclusions": "\\label{sec:conclusion}  We have presented the 2PCF in redshift-space for the SDSS Luminous Red Galaxy sample considering the anisotropy in 2D redshift space. In particular, we have focused on the distorted features of the 2PCF in redshift-space from both peculiar velocities of galaxies and geometrical effect. The distorted features of the Kaiser and finger-of-God effects were clearly detected. The baryon ridges, which are the baryonic acoustic features in the anisotropic 2PCF, are a nearly spherical object in comoving space. We found indications of baryon ridges in the measured 2PCF. Beyond qualitative comparison between data and theory, evaluation of the covariance matrix is needed for cosmological parameter estimation. We constructed the matrix by generating mock samples using the second-order Lagrangian perturbation theory with an artificial biasing scheme. We have constrained the cosmological parameters by comparing the observed 2PCF with linear theory.  We have obtained constraints on fundamental cosmological parameters, $\\Omega_m = 0.218^{+0.047}_{-0.037}$, $\\Omega_b = 0.0473^{+0.0157}_{-0.0160}$, $h = 0.702^{+0.187}_{-0.117}$ , and $n_s=1.122^{+0.152}_{-0.183}$ when we have used the data of $40<s<200\\himpc$.  The constraint on $\\Omega_m$ was better mainly because of the clear detection of the Kaiser effect, which directly depends on $\\Omega_m$ through $\\beta$. We have also obtained the constraints on the dark energy as $\\Omega_{\\rm DE}=0.770^{+0.051}_{-0.040} $ and $w=-0.93^{+0.45}_{-0.35}$ when we fix $\\Omega_bh^2=0.024$ and the other parameters, $h$, $n_{\\rm s}$, $\\sigma_8$, and $b$ are marginalized over.  These constraints are mainly due to the overall shape of the anisotropic 2PCF and the information from geometrical distortions including the scale of the baryon ridge.  We have demonstrated that a pure LRG analysis can constrain $w$ by considering the anisotropy of the structure accurately.  As for the parameters related to the clustering amplitude, we have obtained $\\sigma_8=0.66^{+0.289}_{-0.216}$ and $b=1.55^{+1.42}_{-0.75}$.  While these two parameters are strongly coupled, the degeneracy was alleviated from anisotropy of the 2PCF through the redshift distortion factor, $\\beta$.  In addition, stronger constraints on the cosmological parameters above were obtained by the CMB prior from the 3 year \\wmap results.  All the constraints summarized above agree with the previous studies in literature.  The current analysis can be improved by considering two issues below. The first issue is theoretical improvement accounting for the nonlinearity of the gravitational evolution, the redshift distortions and the galaxy biasing. Although the baryonic signature emerges on very large scales, the width of the baryon peak in the 2PCF is an order of $10~$Mpc. The nonlinearities nontrivially affect such a feature. In fact, such effects have already been investigated using $N$-body simulations, higher-order perturbation theories, and renormalization perturbation theory \\citep{Meiksin1999, Seo2005, Springel2005, Jeong2006, Crocce2007, Nishimichi2008, M2007}. The degradation of the acoustic signature was well modeled; it was shown that the acoustic peak of the linear density field in the 2PCF can be reconstructed \\citep{E2007a, E2007b}. In addition, the overall shape of the redshift-space 2PCF is also affected by nonlinear dynamics \\citep{Scoccimarro2004}. According to his result, the redshift-space 2PCF for pairs parallel to the line of sight in a random Gaussian field deviates from the prediction of standard linear theory even on fairly large scales. In this work we do not use the data along the line of sight; however, we shall include these issues in the future analysis. The theoretical and numerical studies also suggest that the biasing is potentially scale dependent even on large scales \\citep[e.g.,][]{Schulz2006, Smith2007, Coles2007}, which poses a serious problem for estimating cosmological parameters from galaxy surveys \\citep{Blanton2006, Percival2007b, Sanchez2008}.  The second issue is the calculation of the covariance matrix for the measured 2PCF. As described in \\S~\\ref{sec:covariance}, we have constructed a covariance matrix by the second-order perturbation theory. Using N-body simulations which fully include nonlinearity provides better estimation of the covariance matrix. However, this approach is too computationally expensive to produce a large number of independent realizations. The approximation by the second-order perturbation theory is valid on scales which we consider in this work. However, in order to utilize the information at smaller scales for more accurate cosmological parameter estimation, we must estimate the nonlinearities more accurately using $N$-body simulations or more sophisticated methods such as a halo occupation model from the second-order Lagrangian perturbation theory \\citep[e.g.,][]{Scoccimarro2002}.  The most important point of our analysis is that we directly include anisotropies of the structure. The baryonic features enable to divide the effect of the redshift distortions into dynamical and geometrical components. The anisotropy due to the geometric distortion, in particular, contributes to better estimation of the equation-of-state parameter for the dark energy. Various methods using the scale of the oscillations as a standard ruler have been considered for both the power spectrum and the 2PCF \\citep{EH1998, BG2003, Hu2003, Seo2003, M2004, GB2005}. This work is the first application of the anisotropy in the 2PCF with baryon acoustic features to observational data, which was proposed by \\citet{M2004}. Direct measurement of the growth function from the Kaiser's effect and the two-dimensional acoustic scales which depend on $D_A(z)$ and $H(z)$ is also an attractive challenge of the analysis using the anisotropic 2PCF, so these topics will be definitely pursued in future work with an improved LRG sample.  The baryonic signature from the redshift range of SDSS LRGs is not strong, because their number density is relatively small and nonlinear effects weaken the baryonic feature. There are many plans for constraining the dark energy by future wide-field, deep galaxy surveys: the Fiber Multiobject Spectrograph \\citep[FMOS;][]{Kimura2003}, Wide-Field Multiobject Spectrograph \\citep[WFMOS; ][]{GB2005, Bassett2005}, Baryon Oscillation Probe \\citep[BOP; ][]{Glazebrook2005}, and the Hobby-Eberly Dark Energy Experiment \\citep[HETDEX; ][]{Hill2004}, and so on. When the baryonic signature is detected with high accuracy from future redshift surveys, the analysis of the anisotropic 2PCF as in this work will be an important ingredient for stringently constraining properties of the dark energy."
    },
    "0711/0711.4256_arXiv.txt": {
        "abstract": "Window functions describe, as a function of orbital period, the probability that an existing planetary transit is detectable in one's data for a given observing strategy. We show the dependence of this probability upon several strategy and astrophysical parameters, such as length of observing run, observing cadence, length of night, and transit duration.  The ability to detect a transit is directly related to the intrinsic noise of the observations. In our simulations of the window function, we explicitly address non-correlated (white) noise and correlated (red) noise and discuss how these two different noise components affect window functions in different manners. ",
        "introduction": "\\label{introduction}  The signal-to-noise ratio (SNR) of a planetary transit detection in a given photometric light curve (here defined as simply time versus magnitude, irrespective of presence of absence of variability) can, in the simplest case, be approximated by:  \\begin{equation} SNR = \\frac{depth}{\\sigma}\\sqrt{n}, \\label{simple_equation} \\end{equation}  \\noindent where {\\it depth} is the transit depth in magnitudes, $\\sigma$ represents the photometric measurement uncertainty in magnitudes per data point (assumed here to be the same for all data points), and {\\it n} equals the number of data points observed during transit. The fundamental assumption in this equation is the absence of any statistically correlated (red) noise, i.e., the only source of noise is random (white) noise.  White noise is defined as noise that is uncorrelated from data point to data point; typical sources are photon noise and sky background noise. White noise decreases with increasing brightness of the observed target.  Red noise is defined as noise that is correlated from data point to data point; it is not necessarily removed through standard differential or ensemble photometry techniques. Typical sources of red noise may be weather, seeing changes, or tracking errors. It does not change as a function of target magnitude.  As \\cite{p06} and \\cite{pzq06} initially pointed out, the assumption of the presence of only white noise in equation \\ref{simple_equation} is incorrect, and one needs to account for the presence of red noise in the calculation of the SNR and the corresponding expected yields for transit surveys, given survey and astrophysical parameters. A more detailed description of the transit detection SNR which includes both uncorrelated (white) and correlated (red) noise components is given by \\cite{pzq06}:  \\begin{equation} SNR = \\frac{ depth }{\\sqrt{\\frac{1}{n^2} \\sum_{i,j}C_{ij} }} = \\frac{ depth }{\\sqrt{\\frac{\\sigma^2}{n} + \\frac{1}{n^2}\\sum_{i\\neq j}C_{ij} }}, \\label{cov_equation} \\end{equation}  \\noindent where $C_{ij}$ is the covariance matrix. In equation \\ref{cov_equation}, the elements $C_{ij}$ represent the correlation coefficents between the $i$-th and $j$-th measurement obtained during transit. All diagonal elements $C_{ii}=\\sigma^2_i$ are not correlated with other measurements and thus represent the uncorrelated or white noise uncertainties in the $i$-th measurement. It is furthermore assumed that $\\sigma_i = \\sigma$ for all values of $i$.  In order to make the above equation more practically calculable, \\cite{pzq06} assume that statistical correlation among data points from different transits will be much weaker than among data points observed during the same transit. They furthermore separate the total noise into a purely uncorrelated (white) component $\\sigma_w$ and a purely correlated (red) component $\\sigma_r$ and use these to derive an approximation of equation \\ref{cov_equation}:  \\begin{equation} SNR = \\sqrt{ \\frac{ \\left ( depth \\cdot n \\right )^{2}}{\\sum^{N_{tr}}_{k=1}\\left[ n^2_k \\left ( \\frac{\\sigma^2_w}{n_k} + \\sigma^2_r \\right ) \\right ] } }, \\label{approximation_equation} \\end{equation}  \\noindent where $n$ is the total number of data points observed during all transits, $N_{tr}$ is the total number of transits observed, $n_{k}$ is the number of data points observed during the $k$-th transit, and $\\sigma_w$ and $\\sigma_r$ are the white and red noise components, respectively.  By means of equation \\ref{approximation_equation}, the SNR can be regarded as a function of transit survey strategy and astrophysical parameters (see \\S \\ref{algorithm}). If it exceeds a certain threshold value, then an existing transiting planet is detectable in the data. The window function indicates the probability, as a function of orbital period, that the SNR ratio exceeds this threshold and thus produces a detectable planet transit.  We briefly outline our algorithm and methods in \\S 2, show respective influences of varying white and red noise components in \\S 3, and examine the effects of various survey strategy and astrophysical parameters in \\S4. We summarize and conclude in \\S 5. ",
        "conclusions": "In this presentation, we illustrate the influence of several parameters on the probability that an existing planetary transit is detectable in a data set with given noise properties. Red noise dominates in the regime in which planets are typically found (the brightest stars in one's sample). In order to beat down red noise effects and improve detection efficiency, detrending is a vital instrument.  Transit survey strategy can be employed to maximize the projected yield of a given survey. We examine how much the detection efficiency for different orbital periods would suffer when changing one's observing strategy, e.g., to increase the number of monitored stars. We show examples involving the number of consecutive observing nights, typical night length, as well as the observing cadence, and we find that sacrificing full nights or parts of nights can significantly lower transit detection probability. Finally, we show the effects of non-central transits and the associated change of transit duration upon detection efficiency. Our results indicate that, for $\\sigma_r =$ 3 mmag (typically found in some of the major transit surveys before detrending is applied), small changes in inclination angle and associated transit duration do not greatly affect detection efficiency.  In all simulations presented here, the presence of red noise (for $\\sigma_r$ = 3 mmag) dominates all other effects. However, the variation of some of the parameters examined in this work (in particular the observing cadence and the transit duration) appears to have a bigger effect on white noise dominated (e.g., detrended) data."
    },
    "0711/0711.4583_arXiv.txt": {
        "abstract": "The fluorescence detection of ultra high energy cosmic rays requires a detailed knowledge of the fluorescence light emission from nitrogen molecules over a wide range of atmospheric parameters, corresponding to altitudes typical of the cosmic ray shower development in the atmosphere. We have studied the temperature and humidity  dependence of the fluorescence light spectrum excited by MeV electrons in air. Results for the 313.6 nm, 337.1 nm, 353.7 nm and 391.4 nm bands are reported in this paper. We found that the temperature and humidity dependence of the quenching process changes the fluorescence yield  by a sizeable amount (up to 20\\%) and its effect must be included for a precise estimation of the energy of ultra high energy cosmic rays. ",
        "introduction": "\\label{intro} The detection of ultra high energy ($\\gtrapprox 10^{18}$ eV) cosmic rays using nitrogen fluorescence light emission from extensive air showers (EAS) is a well established technique, used by the Fly's Eye \\cite{flyseye}, HiRes \\cite{hires}, and Pierre Auger Observatory \\cite{auger} experiments, and for the Telescope Array \\cite{telaray} under construction.  It has also been proposed for the satellite-based EUSO and OWL projects. Excitation of atmospheric nitrogen by EAS charged particles induces fluorescence emission, mostly in the wavelength range between 300 to 430 nm. Information on the longitudinal EAS development can be obtained by fluorescence telescopes by recording the light intensity as a function of time and incoming direction. However, the fluorescence light yield from EAS charged particles must be well known at each point within the shower, and corrections applied for atmospheric effects between the shower and the telescope for an accurate primary energy determination. Thus, the intensities of the fluorescence bands should be measured over a range of air pressure and temperature corresponding to altitudes up to about 16 km, the typical elevation of EAS development in the atmosphere. The presence of humidity will also affect the fluorescence yield, the effect being more important for satellite experiments which will detect showers over the oceans.  The AIRFLY (AIR FLuorescence Yield) collaboration is pursuing an extensive measurement program of the fluorescence light yield with significantly improved precision with respect to previous experiments \\cite{flbib}. A precise measurement of the pressure dependence of the air fluorescence band intensities has been published \\cite{airflyAP}, and is summarized in \\cite{obermeier}. Measurements of the fluorescence yield dependence on the electron kinetic energy from keV to GeV are presented in \\cite{maximo}. The status of AIRFLY measurement of the absolute fluorscence yield is presented in \\cite{absolute}. The data reported here address the measurement of the fluorescence yield spectrum dependence on air temperature and humidity.  The fluorescence yield of a band of wavelength $\\lambda$ at a given pressure $p$ and temperature $T$ can be written as \\cite{airflyAP}: \\begin{eqnarray} Y_{air}(\\lambda,p,T)&=& Y_{air}(337,p_0,T_0) \\nonumber \\cdot  I_{\\lambda}(p_0,T_0)  \\nonumber \\\\ && \\cdot \\frac{1+\\frac{p_0}{p'_{air}(\\lambda,T_0)}}{1+\\frac{p}{p'_{air}(\\lambda,T_0)\\sqrt{\\frac{T}{T_0}}\\frac{H_\\lambda(T_0)}{H_\\lambda(T)}}} , \\label{eq:phish} \\end{eqnarray} where  $Y_{air}(337,p_0,T_0)$ is the absolute yield of the 337 nm band at pressure $p_0$ and temperature $T_0$ (in photons emitted per MeV of energy deposited), $I_{\\lambda}(p_0,T_0)$ is the $\\lambda$ band intensity relative to the 337 nm band and $p'_{air}(\\lambda,T_0)$  is the band quenching reference pressure. $H_\\lambda(T)$ has been introduced to take into account a possible temperature dependence of the collisional cross sections.  The effect of humidity in the fluorescence yield can be introduced by substituting in Eq. (\\ref{eq:phish}): \\begin{eqnarray} \\frac{1}{p'_{air}} &\\rightarrow& \\frac{1}{p'_{air}} \\left( 1-\\frac{p_{h}}{p} \\right)  + \\frac{1}{p'_{\\rm{H_2O}}}\\frac{p_{h}}{p}, \\label{eq:hum} \\end{eqnarray} where $p_{h}$ is the water vapour partial pressure and $p'_{\\rm{H_2O}}$ is the water vapour collisional quenching pressure.  The experimental method used for the measurement of the temperature and humidity dependence is described in Section \\ref{sec:expmethod}. The results obtained for the humidity and temperature dependence of four major fluorescence bands (313.6 nm, 337.1 nm, 353.7 nm and 391.4 nm) are presented in Sections \\ref{sec:temperature} and \\ref{sec:humidity}, respectively. Implications for the energy determination of ultra high energy cosmic rays of the AIRFLY measurements reported in this paper are discussed in Section \\ref{sec:implications}. Conclusions are given in Section \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions} The measurement of the EAS energy with the fluorescence detection technique requires a detailed knowledge of the air fluorescence emission as a function of pressure, temperature and humidity. The AIRFLY experiment has performed precise measurements of the fluorescence light spectrum excited by MeV electrons in dry air. The relative intensities of 34 fluorescence bands as well as their pressure dependence have been fully reported \\cite{airflyAP}, and are summarized in a contribution to this workshop \\cite{obermeier}. In this paper, we reported measurements of the temperature and humidity dependence of the 313.6 nm, 337.1 nm, 353.7 nm and 391.4 nm fluorescence bands in air. The temperature dependent measurements of the fluorescence yield, performed in the range between 240~K and 310~K, are the first reported in the literature. Our data show that collisional cross sections are temperature dependent, contrary to the assumption so far made in the UHECR community. We also measured the quenching effect of water vapour, resulting in about 20\\% decrease of the fluorescence yield for rH=100\\%. The effect of the temperature and humidity dependence of the fluorescence yield on the energy estimation of EAS was studied, and we found that it should not be neglected. A full analysis of the AIRFLY data, with the measurement of the temperature and humidity dependence of 34 fluorescence bands, is in preparation."
    },
    "0711/0711.4310_arXiv.txt": {
        "abstract": "{ The times of maximum brightness collected in the GEOS RR Lyr database allowed us to trace the period variations of a sample of 123 galactic RRab variables. These data span a time baseline exceeding 100~years. Clear evidence of period increases or decreases at constant rates has been found, suggesting evolutionary effects. The observed rates are slightly larger than those predicted by theoretical models; moreover, there is an unexpected large percentage of RRab stars showing a period decrease. The new possibilities offered by the use of robotic telecopes (TAROTs, REM) and of data from satellite (CoRoT) are expected to speed up the project to measure stellar evolution in real time. ",
        "introduction": "Horizontal branch stars cross the instability strip in different stages of their evolution; in these phases they become RR Lyr stars. The crossing can take place in both directions; as a consequence, the periods will be either increasing if the stars evolve from blue to red or decreasing if they evolve from red to blue. Despite its importance as a test for the stellar evolution theory, the observed rate of the period changes is still an unknown quantity.  The amateur/professional association GEOS (Groupe Europ\\'een d'Observations Stellaires) has built a database aiming to put together all possible RR Lyr light maximum times published in the literature. The publications from the end of the XIX$^{\\rm th}$ century up to today were scanned for this purpose and recent observations from amateur astronomers of the European groups GEOS and BAV (Bundesdeutsche Arbeitsgemeinschaft f\\\"ur Ver\\\"anderliche Sterne) fed the database in a continuous way. To date, it contains about 50000 maximum times from more than 3000 RR Lyr stars. The GEOS database is freely accessible on the internet at the address http://dbrr.ast.obs-mip.fr/. Therefore, this database is the adequate tool to search for secular period variations in the galactic RR Lyr  stars, since  most of them have been studied for tens of years, several of them since the end of the XIX$^{\\rm th}$ century. The first result of the project, {\\it Stellar evolution in real time}, have been presented by \\citet{flb}. It is noteworthy that its outlines have been sketched during several GEOS meetings, where the different knowledge of amateur and professional astronomers found a very profitable synthesis.   \\begin{figure} \\includegraphics[width=1.0\\columnwidth,height=0.80\\columnwidth]{EPoretti.fig1.ps} \\caption{\\footnotesize Histogram of the ratios between the modulus of the $a$ coefficient and its formal error. } \\label{isto} \\end{figure} \\begin{figure*} \\begin{center} \\includegraphics[width=0.666666\\columnwidth,height=0.6666\\columnwidth]{EPoretti.fig2a.ps} \\includegraphics[width=0.666666\\columnwidth,height=0.6666\\columnwidth]{EPoretti.fig2b.ps} \\includegraphics[width=0.666666\\columnwidth,height=0.6666\\columnwidth]{EPoretti.fig2c.ps} \\includegraphics[width=0.666666\\columnwidth,height=0.6666\\columnwidth]{EPoretti.fig2d.ps} \\includegraphics[width=0.666666\\columnwidth,height=0.6666\\columnwidth]{EPoretti.fig2e.ps} \\includegraphics[width=0.666666\\columnwidth,height=0.6666\\columnwidth]{EPoretti.fig2f.ps} \\caption{\\footnotesize For each star the upper panel shows the O--C values as obtained from linear elements, the lower panel from a parabolic fit. Measure units are elapsed cycles (E/10000) on the $x$--axis and days on the $y$-axis.} \\label{oc} \\end{center} \\end{figure*} ",
        "conclusions": "The most remarkable result is the detection of a large number of O--C plots showing a parabolic pattern when linear elements are used. The cause of this pattern is a regular period variation, having the same order of magnitude of the expected evolutionary changes. On the other hand, in 54 cases (i.e., a number of cases similar to that of RR Lyr stars showing period variations) we did not detect a significant period variation and linear elements satisfy the available times of maximum brightness. The upper panel in Fig.~\\ref{istobeta} shows the distribution of the $\\beta$ values in the sample of the galactic RRab stars. The central peak is given by the stars showing a constant period, but the number of stars both in the negative and in the positive planes are relevant.  Indeed, the comparison between RRab stars showing negative and positive rates yields us the most important scientific feedbacks. They  can be evaluated in an easier way by removing the central peak (lower panel in Fig.~\\ref{istobeta}). It is quite evident that the blueward evolution (decreasing periods, negative $\\beta$ values; 21 cases) is as common as the redward one (increasing periods, positive $\\beta$ values; 27 cases). \\begin{figure} \\includegraphics[width=1.00\\columnwidth,height=1.00\\columnwidth]{EPoretti.fig4.ps} \\caption{\\footnotesize Histograms of the distributions of the $\\beta$ values of the whole sample (upper panel) and restricted to the stars shown an appreciable linearly variable period (lower panel). The observed $\\beta$ values are slightly larger than those predicted by the evolutionary models. } \\label{istobeta} \\end{figure} By analyzing and comparing these two samples, \\citet{flb}  could stress some well--established observational facts: \\begin{enumerate} \\item RRab stars experiencing blueward evolution are quite common, only slightly less than RRab stars experiencing redward evolution. The period range covered by the two groups are very similar and the average and median periods are nearly coincident; \\item the absolute values of period changes are larger than expected. The theoretical evolutionary rates calculated by \\citet{lee} predict $\\beta$ values smaller than 0.30 for almost all stars. $\\beta$ values greater than 0.30 are  reported in exceptional cases. Since in our sample we have 18 stars with $0.00<\\beta<0.30$ and 9 stars with $\\beta>$0.30, the 2:1 ratio seems too small to be caused by the exceptional conditions. \\item Blazhko effect is often superimposed on secular changes, but the monotonic trend due to evolutionary variations still remains visible; \\end{enumerate} In particular, the rates of the period changes are difficult to associate with canonical horizontal--branch evolution. The large observed $\\beta$ values could be ascribed to a sort of ``noise\" superimposed on the canonical evolutionary tracks, irrespective of where the star is in the course of its horizontal--branch evolution. However, this is still surprising since the main sources of this ``noise\" had been invoked to act toward the exhaustion of the helium in the core, but not before \\citep{sweig}. At this purpose, we also note that there is a numerous group of RRab stars showing erratic O--C behaviours, with abrupt or continuous period variations, never monotonic. The regular variations caused by light--time effect (and hence duplicity), often invoked to explain large O--C excursions, are not convincingly observed in any case (except maybe for RZ Cet)  of our sample."
    },
    "0711/0711.2646_arXiv.txt": {
        "abstract": "A cosmological model is proposed for the current Universe consisted of non-interacting baryonic matter and interacting dark components. The dark energy and dark matter are coupled through their effective barotropic indexes, which are considered as functions of the ratio between their energy densities. It is investigated two cases where the ratio is asymptotically stable and their parameters are adjusted by considering best fits to Hubble function data. It is shown that the deceleration parameter, the densities parameters,  and the luminosity distance have the correct behavior which is expected for a viable present scenario of the Universe. ",
        "introduction": "Recent observations of Supernova Ia (SNIa) suggest that the Universe has entered into a stage of an accelerated expansion with a redshift $z \\lesssim 1$, \\cite{Riess:1998cb}, \\cite{Perlmutter:1998np}, \\cite{Astier:2005qq}.  This has been confirmed by precise measurements of the spectrum of the Cosmic Microwave Background (CMB) anisotropies \\cite{Spergel:2003cb}, \\cite{Spergel:2006hy} as well as the baryon acoustic oscillations (BAO) in the Sloan Digital Sky Survey (SDSS) luminous galaxy sample \\cite{Eisenstein:2005su}. All usual types of matter with positive pressure generate attractive forces and decelerate the expansion of the universe. For that reason, a dark energy component with negative pressure was suggested to account for a fluid that drives the current accelerated expansion (see \\cite{Sahni:1999gb} for reviews).  The simplest explanation of dark energy is provided by a cosmological constant, but the scenario is plagued by a severe fine tuning problem associated with its energy scale. The vacuum energy density falls below the value predicted by any sensible quantum field theory by many orders of magnitude \\cite{Weinberg:1988cp}, and it unavoidably leads to the coincidence problem, i.e.,``Why are the vacuum and matter energy densities of precisely the same order today?\" \\cite {Steinhardt}. More sophisticated models replace the cosmological constant by a dynamical dark energy that may be a scalar field (quintessence), tachyon field, phantom field or exotic equations of state. These models fit the observational data but it is doubtful that they can solve the coincidence problem \\cite{Chimento:2000kq},\\cite{Chimento:2003iea}.  Besides, in a viable dark energy scenario we require that the energy density of the dark fluid remains subdominant during radiation and matter dominant eras and that it becomes important only at late times to account for the current acceleration of the Universe.  In several papers, it has been proposed that dark matter and dark energy are coupled and do not evolve separately \\cite{Binder:2006yc},\\cite{Amendola:1999er}. The coupling between matter and quintessence is either motivated by high energy particle physics considerations \\cite{Amendola:1999er} or is constructed by requiring the final matter to dark energy ratio to be stable against perturbations \\cite{Chimento:2003iea},\\cite{Zimdahl:2001ar}.  In the same sense, we will use two arbitrary interacting fluids in the dark sector, with constant barotropic indexes, and will show that there is a stable solution for the ratio dark matter-dark energy, which satisfies the expected behavior from the matter dominated era until today. We explain in section II the general features of the model for two different interactions and find the stability conditions of the ratio dark matter-dark energy. The section III is devoted to obtain the values of parameters that best fit the experimental data using the recently published Hubble function $H(z)$ data \\cite{Simon:2004tf}, extracted from differential ages of passively evolving galaxies. The allowed regions of probability $1\\sigma$ and $2\\sigma$ for the parameters are showed. Also, deceleration parameter, densities parameters, ratio dark matter-dark energy, effective equations of state and luminosity distance are exhibited in section IV for the best fit parameters, showing the correct behavior expected for a viable scenario. Finally, in section V we present our conclusions. ",
        "conclusions": "To sum up: in this work a cosmological model for the present Universe was proposed whose constituents were the non-interacting baryonic matter and interacting dark components. The interaction between the dark energy and dark matter was related with their effective barotropic indexes, which were considered as functions of the ratio between their energy densities. Furthermore, the parameters of the functions were adjusted by considering best fits to Hubble function data. Two functions were analyzed, namely, a constant  and a variable one. It was shown: (i) for the proposed functions the ratio $r$ is asymptotically stable; (ii) the energy transfer from the dark energy to the dark matter is more efficient for the variable case; (iii) for both functions the dark energy density predominates in the future; (iv) for both functions the present values of the deceleration parameter and the values of the transition red-shift -- where the regime changes from a decelerated to an accelerated -- are of the same order as the experimental values; (v) for the variable function, the effective ratio between the pressures and energy densities for the dark components indicates that the dark energy behaves as quintessence whereas the dark matter as practically a pressureless fluid (dust) when $z>0$;  (vi) the behavior of the parameter $\\mu_0$ -- related with the luminosity distance -- with the red-shift does not have a sensible difference for the two functions and practically coincide with the one of the $\\Lambda$CDM model, indicating a good fit with the experimental values."
    },
    "0711/0711.0643_arXiv.txt": {
        "abstract": "We describe source code level parallelization for the {\\tt kira} direct gravitational $N$-body integrator, the workhorse of the {\\tt starlab} production environment for simulating dense stellar systems. The parallelization strategy, called ``j-parallelization'', involves the partition of the computational domain by distributing all particles in the system among the available processors.  Partial forces on the particles to be advanced are calculated in parallel by their parent processors, and are then summed in a final global operation. Once total forces are obtained, the computing elements proceed to the computation of their particle trajectories.  We report the results of timing measurements on four different parallel computers, and compare them with theoretical predictions.  The computers employ either a high-speed interconnect, a NUMA architecture to minimize the communication overhead or are distributed in a grid. The code scales well in the domain tested, which ranges from 1024 - 65536 stars on 1 - 128 processors, providing satisfactory speedup. Running the production environment on a grid becomes inefficient for more than 60 processors distributed across three sites. ",
        "introduction": "Numerical studies of gravitational $N$-body problem have made enormous advances during the last few decades due to algorithmic design \\citep{1963ZA.....57...47V,1964ApNr....9..313A,1968BAN....19..479V,1975ARA&A..13....1A,1999PASP..111.1333A}, developments in hardware \\citep{1986LNP...267...86A,1996IAUS..174..141T,1998sssp.book.....M,2001ASPC..228...87M,2003PASJ...55.1163M}, and refinements to approximate methods \\citep{1971Ap&SS..13..284H,1979ApJ...234.1036C,1984ApJ...283..801M,1985ApJ...299....1B,1986Natur.324..446B}. A number of interesting books have been written on the subject \\citep{Newton:1687,1987gady.book.....B,1987degc.book.....S,Hockney1988,2003gmbp.book.....H,Aarseth2003}.  The main reasons for the popularity of the direct evaluation techniques are its very high accuracy and lack of prior assumptions.  Approximate methods, by contrast, often deal with idealized circumstances, which, regardless of their advances in performance, may be hard to justify. Regardless of all the advances over the last 40 years the evaluation of the mutual forces between all stars remains the bottleneck in the calculation. Without the use of hybrid force evaluation techniques, as pioneered by \\cite{1993ApJ...414..200M,1997AGAb...13...53H}, the $\\mathcal{O}(N^2)$ performance complexity will remain a problem for the forseeable future. Nevertheless, it is to be expected that the direct integration method for gravitational $N$-body simulations will remain popular for the next decade.  The main technique adopted to increase the performance (and accuracy) of direct $N$-body methods in stellar dynamics applications is individual time stepping, providing orders of magnitude increase in computational speed.  In the early 1990s this technique enabled astronomers to study star clusters with up to $\\sim 1000$ stars \\cite{1990ApJ...362..522M,1991ApJ...372..111M,1994ApJ...427..793M}. The main development in the 1990s was the development of special-purpose hardware to provide the required teraflop/s raw performance. In particular GRAPE-4 \\citep{1996IAUS..174..141T}, which enables simulation of several tens of thousands of stars, and GRAPE-6 which raised the bar above $10^5$ stars for the first time \\citep{2003PASJ...55.1163M}. For the forseeable future, special purpose hardware may still take the lead in advancing the study of self gravitating systems, particularly if that hardware is multi-purpose \\citep{2005PASJ...57..799H} and therefore able to serve a larger group of users and investors.  The use of graphical processor units (GPU), is an interesing parallel development with a promising performance/price ratio \\citep{2007NewA...12..641P,2007arXiv0707.0438B,2007astro.ph..3182H}. Another likely developmental pathway is the more general use of large parallel computers.  The latter, however, requires efficient parallel kernels for force evaluation.  In the last decade the number of processors per computer has increased dramatically. Even today's personal computers are equipped with multiple cores, and large supercomputers often carry many thousands of processors, for example the IBM BlueGene/L which has 131072 processors.  Efficient interprocess communication for such large clusters is far from trivial, and requires significant development efforts.  The communication overhead and development challenges increase even more once large calculations are performed on the grid, as shown in the pioneering study of \\cite{2007PARCO..33..3G,2007arXiv0709.4552G}.  The parallelization of direct $N$-body kernels, however, has not kept pace with the increased speed of multiprocessor computers.  Whereas the long-range of the gravitational interaction introduces the need to perform full $N^2$ force evaluations, it also requres all-to-all memory access between processors in the computer, giving a large communication overhead. This overhead increases linearly with the number of processors \\cite{2007NewA...12..357H}, although algorithms with reduced communication characteristics have been proposed \\cite{2002NewA....7..373M}. In fact, many algorithmic advances, such as individual time steps \\citep{1991ApJ...369..200M,1993ApJ...414..200M} and neighbor schemes \\cite{1973ApJ...179..885A}, appear to hamper the parallelization of direct $N$-body codes.  The first efficiently parallelized production $N$-body code was NBODY6++ \\citep{2007MNRAS.374..703K}, for which the systolic-ring algorithm was adopted \\citep{Dorband,2007PARCO..33..3G,2007NewA...12..357H}.  In this research note we report on our endevour to parallelize the {\\tt kira} direct gravitational $N$-body integrator, the workhorse of the {\\tt starlab} environment. The source code of the package, including the parallelized version, here reported, is available online at {\\tt http://www.manybody.org/}. We now briefly report on the implementation of the gravitational $N$-body methodology (\\S\\,\\ref{Sect:Integrator}) and the adopted parallelization topology (\\S\\,\\ref{Sect:parallelization}). In \\S\\,\\ref{Sect:performance} we report the performance of the package when ported to several parallel systems. ",
        "conclusions": "We report the results of our parallelization on the {\\tt kira} $N$-body integrator in the {\\tt starlab} package. In order to maintain a platform independent source code we opted for the Message Passing Interface standard.  The parallelization scheme is standard $j$-parallelization with a copy of the entire system on each node to guarantees that the perturber lists are consistent with the serial version of the code.  In this way we achieve that the numerical results are independent of the number of processor units.  We had to include an additional broadcast step to guarantee that the neighbour list is consistent between the sequential and the parallel code. This extra overhead, is particularly noticeable for small numbers of particles, and does not seriously hamper the scaling of the code with the number of processors. For simulations with a large number of binaries or when running on clusters with more than 128 processors we also found a drop in performance.  The code scales well in the domain tested, which ranges from 1024 to 65536 stars on 1 to 128 processors.  A super linear speed-up was achieved in several cases for 4096--32768 stars on 16--64 processors.  The overall performance for clusters with single stars is rather satisfactory, irrespective of the density profile of the star cluster. But for primordial binaries the performance drops considerably. With 100\\% binaries the performance of a large simulation on many processors hardly results in a reduction of the wall-clock time compared to a few processors. Runs with a large number of primordial binaries still benefit from parallelization for $\\aplt 16$ processors.  This performance hit in the presence of binaries is mainly the results of the additional communication required to synchronize the perturber list and of the sequential part in the binary treatment."
    },
    "0711/0711.2756_arXiv.txt": {
        "abstract": "We present a general scheme for constructing Monte Carlo realizations of equilibrium, collisionless galaxy models with known distribution function (DF) $f_0$.  Our method uses importance sampling to find the sampling DF $f_s$ that minimizes the mean-square formal errors in a given set of projections of the DF $f_0$.  The result is a multi-mass $N$-body realization of the galaxy model in which ``interesting'' regions of phase-space are densely populated by lots of low-mass particles, increasing the effective $N$ there, and less interesting regions by fewer, higher-mass particles.  As a simple application, we consider the case of minimizing the shot noise in estimates of the acceleration field for an $N$-body model of a spherical Hernquist model.  Models constructed using our scheme easily yield a factor $\\sim100$ reduction in the variance in the central acceleration field when compared to a traditional equal-mass model with the same number of particles.  When evolving both models with a real $N$-body code, the diffusion coefficients in our model are reduced by a similar factor.  Therefore, for certain types of problems, our scheme is a practical method for reducing the two-body relaxation effects, thereby bringing the $N$-body simulations closer to the collisionless ideal. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.0423_arXiv.txt": {
        "abstract": "Radiatively-driven transfer flow perpendicular to a luminous disk is examined in the relativistic regime of $(v/c)^2$, taking into account the gravity of the central object. The flow is assumed to be vertical, and the gas pressure as well as the magnetic field are ignored. Using a velocity-dependent variable Eddington factor, we can solve the rigorous equations of the relativistic radiative flow accelerated up to the {\\it relativistic} speed. For sufficiently luminous cases, the flow resembles the case without gravity. For less-luminous or small initial radius cases, however, the flow velocity decreases due to gravity. Application to a supercritical accretion disk with mass loss is briefly discussed. ",
        "introduction": "Mass outflow from a luminous disk is a clue to the formation mechanism of astrophysical jets and winds in the active objects. In particular, in a supercritical accretion disk, the disk local luminosity exceeds the Eddington one, and mass loss from a disk surface driven by radiation pressure would take place (see Kato et al. 1998 for a review of accretion disks).  So far, radiatively driven outflows from a luminous disk have been extensively studied by many researchers (Bisnovatyi-Kogan, Blinnikov 1977; Katz 1980; Icke 1980; Melia, K\\\"onigl 1989; Misra, Melia 1993; Tajima, Fukue 1996, 1998; Watarai, Fukue 1999; Hirai, Fukue 2001; Fukue et al. 2001; Orihara, Fukue 2003), and by numerical simulations (Eggum et al. 1985, 1988; Okuda 2002; Ohsuga et al. 2005; Ohsuga 2006). In almost all of these studies, however, the luminous disk was treated as an external radiation source, and the radiation transfer in the flow was not solved.   Radiation transfer in the disk, on the other hand, was investigated in relation to the structure of a static disk atmosphere and the spectral energy distribution from the disk surface (e.g., Meyer, Meyer-Hofmeister 1982; Cannizzo, Wheeler 1984; Shaviv, Wehrse 1986; Adam et al. 1988; Hubeny 1990; Ross et al. 1992; Artemova et al. 1996; Hubeny, Hubeny 1997, 1998; Hubeny et al. 2000, 2001; Davis et al. 2005; Hui et al. 2005). In these studies, however, the vertical movement and the mass loss were not considered.  Recently, mass outflow as well as radiation transfer have been examined for the first time in the subrelativistic (Fukue 2005a, 2006a) and fully relativistic cases (Fukue 2005b). In the latter case, it is pointed out some singular behavior inherent in the relativistic radiative flow (e.g., Turolla, Nobili 1988; Turolla et al. 1995; Dullemond 1999). When the gaseous flow is radiatively accelerated up to the relativistic regime, the velocity gradient becomes very large in the direction of the flow. As a result, the radiative diffusion may become {\\it anisotropic} in the comoving frame of the gas. Hence, in a flow that is accelerated from subrelativistic to relativistic regimes, the Eddington factor should be different from $1/3$ even in the optically thick diffusion limit.  In order to avoid the singular behavior of such a relativistic regime, for a plane-parallel case, Fukue (2006b) proposed a {\\it velocity-dependent Eddington factor}, which depends on the flow velocity $v$: \\begin{equation} f(\\beta) = \\frac{1+2\\beta}{3}, \\end{equation} where $\\beta=v/c$. In Fukue (2006b) this form (1) was adopted as the simplest one among various forms, which satisfy several necessary conditions to avoid the singularity. Physically speaking, this form (1) can be interpreted as follows. In the high velocity regime, where the radiative diffusion may become anisotropic in the comoving frame, the `apparent' optical depth $\\tau$ would be of the order of \\begin{equation} 1+\\tau = \\frac{1}{\\beta}. \\end{equation} That is, as the flow is accelerated and approaches to the speed of light, the optical depth becomes zero (outward {\\it peaking}). In this case, the form (1) can be read as \\begin{equation} f(\\tau) = \\frac{3+\\tau}{3+3\\tau}, \\end{equation} which recovers a similar form of a usual variable Eddington factor (see, e.g., Tamazawa et al. 1975). Hence, the applicability and accuracy of the form (1) from a low speed regime to a high speed one would be similar to those of a variable Eddington factor from an optically thick regime to an optically thin one.  Similarly, for a spherically symmetric case, Akizuki and Fukue (2006) proposed a variable Eddington factor, which depends on the flow velocity $\\beta$ as well as the optical depth $\\tau$: \\begin{equation} f(\\tau, \\beta) = \\frac{\\gamma(1+\\beta)+\\tau}{\\gamma(1+\\beta)+3\\tau}, \\end{equation} where $\\gamma$ ($=1/\\sqrt{1-\\beta^2}$) is the Lorentz factor.  In Fukue (2006b), the plane-parallel case was examined as an example, although the gravity of the central object was ignored for simplicity. In this paper, we thus consider the radiatively driven vertical outflow -- {\\it moving photosphere} -- in a luminous flat disk within the framework of radiation transfer in the relativistic regime using $f(\\beta)$, while taking into account the gravity of the central object, although the gas pressure is ignored.  In the next section we describe the basic equations in the vertical direction. In section 3 we show our numerical examination of the radiative flow. In section 4 we briefly apply the present model to the case of a supercritial accretion disk. The final section is devoted to concluding remarks. ",
        "conclusions": "In this paper we have examined the relativistic radiative transfer flow in a luminous disk in the relativistic regime of $(v/c)^2$, taking account of gravity of the central object. In such a relativistic regime, we adopt the velocity-dependent variable Eddington factor. The flow is assumed to be vertical, and the gas pressure is ignored for simplicity. The basic equations are numerically solved as a function of the optical depth $\\tau$, and the flow velocity $v$, the height $z$, the radiative flux $F$, and the radiation pressure $P$ are obtained for a given radius $r$, the initial optical depth $\\tau_0$, and the initial coditions at the flow base (disk ``inside''), whereas the mass-loss rate $J$ is determined as an eigenvalue of the boundary condition at the flow top (disk ``surface''). For sufficiently luminous cases, the flow resembles the case without gravity. For less-luminous cases, however, the flow velocity decreases.  Application to the critical accretion disk was also examined. If the disk thickness becomes so large, the present plane-parallel approximation violates and other treatment, such as numerical simulations, should be needed.   The radiative flow investigaed in the present paper must be a quite {\\it fundamental problem} for accretion-disk physics and astrophysical jet formation, although there are many simplifications at the present stage. For example, we have ignored the gas pressure. In general cases, where the gas pressure is considered, there usually appears sonic points, and the flow is accelerated from subsonic to supersonic. In this paper we consider a purely vertical flow, and the cross section of the flow is constant. If the cross section of the flow increases along the flow, the flow properties such as a transonic nature would be influenced. Moreover, we do not consider the rotation of the gas. In accretion disks around a black hole, the gas usually rotates around the hole at the relativistic speed. In the vicinity of the equator, the vertical flow approximation safely holds, since the radial gravity is balanced with the centrifugal force. When the flow is accelerated to be lift up the large height, the streamline would be curved outward, since the centrifugal force dominates. This would violate the vetical flow approximation.  There remain many problems to be solved.     \\vspace*{1pc}  The authors thank an anonymous referee for useful comments, which greatly improved the original manuscript. This work has been supported in part by a Grant-in-Aid for the Scientific Research (18540240 J.F.) of the Ministry of Education, Culture, Sports, Science and Technology."
    },
    "0711/0711.3616_arXiv.txt": {
        "abstract": "{ The motivation for considering distributed large scale dynamos in the solar context is reviewed in connection with the magnetic helicity constraint. Preliminary accounts of 3-dimensional direct numerical simulations (in spherical shell segments) and simulations of 2-dimensional mean field models (in spherical shells) are presented. Interesting similarities as well as some differences are noted. ",
        "introduction": " ",
        "conclusions": "We have briefly presented preliminary results comparing direct simulations in spherical shell segments with their Cartesian analogues, on the one hand, and mean field models in spherical shells, on the other.  Simulations in Cartesian boxes and spherical shell segments both show qualitatively similar features in the mean magnetic field as well as similar saturation behaviour in the normalized mean magnetic field. On the other hand the saturation value is lower in mean field models than in the three-dimensional simulations.  Clearly more work is needed to clarify further the relation between these models, before we can proceed to the arguably more realistic case with a penetrable boundary. Work is in progress in this direction and will be reported elsewhere."
    },
    "0711/0711.4807_arXiv.txt": {
        "abstract": "{Rapidly rotating, chemically homogeneously evolving massive stars are considered to be progenitors of long gamma-ray bursts.} {We present numerical simulations of the evolution of the circumstellar medium around a rapidly rotating 20~$\\mso$ star at a metallicity of $Z=0.001$. Its rotation is fast enough to produce quasi-chemically homogeneous evolution. While conventionally, a star of 20~$\\mso$ would not evolve into a Wolf-Rayet stage, the considered model evolves from the main sequence directly to the helium main sequence.} {We use the time-dependent wind parameters, such as mass loss rate, wind velocity and rotation-induced wind anisotropy from the evolution model as input for a 2D hydrodynamical simulation.} {While the outer edge of the pressure-driven circumstellar bubble is spherical, the circumstellar medium close to the star shows strong non-spherical features during and after the periods of near-critical rotation.} {We conclude that the circumstellar medium around rapidly rotating massive stars differs considerably from the surrounding material of non-rotating stars of similar mass. Multiple blue-shifted high velocity absorption components in gamma-ray burst afterglow spectra are predicted. As a consequence of near critical rotation and short stellar evolution time scales during the last few thousand years of the star's life, we find a strong deviation of the circumstellar density profile in the polar direction from the $1/R^2$ density profile normally associated with stellar winds close to the star.} ",
        "introduction": "Within the collapsar picture (Woosley~\\cite{W93}), long gamma-ray bursts (GRB) are thought to be produced at the end of the evolution of massive stars, if their rapidly rotating iron cores collapse to form black holes. The circumstellar medium (CSM) around GRB progenitors is interesting because it determines the GRB afterglow evolution (Chevalier \\& Li \\cite{CL00}, Panaitesku \\& Kumar \\cite{PK01, PK02}, Chevalier et al. \\cite{CLF04}, Ramirez-Ruiz et al. \\cite{RGSP05}, Nakar \\& Granot \\cite{NG06}), and because it may produce distinct absorption features in GRB afterglow spectra (Mirabal et al. \\cite{Metal03}, Schaefer et al. \\cite{Setal03}, Fiore et al. \\cite{Fietal05}, Starling et al. \\cite{Setal05}). In previous papers, van Marle et al. (\\cite{MLG05b}, \\cite{MLAG06}, \\cite{MLG07}) have computed the CSM evolution of various types of supernova and GRB progenitors without including the effect of rotation of the central star.  It has been recently suggested that rapidly rotating, chemically homogeneous massive stars evolve into long GRBs (Yoon \\& Langer \\cite{YL05}; Woosley \\& Heger \\cite{WH06}; Yoon et al. \\cite{YLN06}; Cantiello et al. \\cite{CYLL07}). This type of evolution is thought to occur at low metallicity, where stellar winds are weaker and angular momentum loss by winds (Langer~\\cite{L98}) is therefore limited. Since models of chemically homogeneously evolving stars avoid a redward evolution in the HR diagram, but rather evolve from the main sequence directly to the helium main sequence, the structure of their circumstellar medium is expected to be different from that of Galactic Wolf-Rayet stars, which are thought to be the descendants of red supergiants or Luminous Blue Variables (Garcia-Segura et al. \\cite{GML96, GLM96}). Additionally, the chemically homogeneous stellar models remain very rapid rotators throughout their lives, and encounter various stages of near-critical rotation (Yoon \\& Langer \\cite{YL05}). Wind anisotropies may therefore play an important role in shaping the CSM of these stars.  Here, we undertake an effort to understand the evolution of the circumstellar medium around such stars. We perform detailed numerical simulations of its time evolution, were we use a detailed stellar evolution model as input. We employ a prescription of the anisotropy of the mass outflow and wind velocity as a function of the rotation of the central star, and adopt this, together with the time dependent global wind properties and ionizing photon fluxes, as the inner boundary condition for our CSM evolution model. As result, we obtain the density, velocity and temperature structure of the CSM of our stellar model, from the main sequence to the pre-SN stage.  In Sect.~\\ref{sec-evol} we describe the evolution of the properties of the chemically-homogeneously evolving stellar model which is used as input for our hydrodynamic calculations. In Sect.~\\ref{sec-num} we discuss the numerical method used to simulate the hydrodynamical evolution of the CSM. The results of our simulation are presented in Sect.~\\ref{sec-result1}. In Sect.~\\ref{sec-spectrum}, we show the absorption profiles that would result from the morphology of the circumstellar medium when the star explodes. Section~\\ref{sec-dprof} deals with the density profile of the circumstellar medium, which is important for the lightcurve of a GRB afterglow. In Sect.~\\ref{sec-concl} we discuss our results.  \\begin{figure} \\centering \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90]{8802f01.eps}} \\caption{Evolution the rapidly rotating 20 $\\mso$ model at $Z=0.001$ which is used as input for the CSM simulations in the Hertzsprung-Russel diagram. The model is similar to the ones used in Yoon et al. (\\cite{YLN06}). Its initial rotational velocity is 414~km/s, which is about 60\\% of the critical velocity. The track starts at the zero age main sequence and ends at the pre-supernova stage. } \\label{fig:HRdiag} \\end{figure}  \\begin{figure} \\centering \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90]{8802f02.eps}} \\caption{Latitudinally averaged mass loss rate, and $\\Omega$ ($=\\vrot/\\vcrit$) for our 20 $\\mso$ sequence, as a function of time. The horizontal axis shows the logarithm of the time before core collapse ($t_{cc}$). The gray area highlights the time interval which is explored in Figs.~\\ref{fig:2d1} to \\ref{fig:2d3} below. } \\label{fig:mloss1} \\end{figure}   \\begin{figure} \\centering \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90]{8802f03.eps}} \\caption{Number of high energy photons, and terminal wind velocity of our 20 $\\mso$ sequence star, as a function of time. The horizontal axis shows the logarithm of the time before core collapse ($t_{cc}$). The gray area highlights the time interval explored in Figs.~\\ref{fig:2d1} to \\ref{fig:2d3} below. } \\label{fig:mloss2} \\end{figure}  \\begin{figure} \\centering \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90]{8802f04.eps}} \\caption{The result of a 1D hydrodynamic simulation of the CSM around our 20 $\\mso$ star, just before the star reaches critical rotation for the first time ($4.644\\times 10^5$yr before core collapse; cf. Fig.~2). The solid line shows the mass density, the dashed line the internal energy density. Moving outward from the star, we first encounter the free streaming wind, then the shocked wind material ($6~\\lesssim~R~\\lesssim~30$~pc), the photo-ionized ISM ($30~\\lesssim~R~\\lesssim~43$~pc), the shell of shocked ISM ($R~\\simeq~45$~pc) and finally, the unperturbed ISM. The dotted line marks the Str{\\a\"o}mgren radius. We use this point in time to start the 2D simulation. } \\label{fig:ms} \\end{figure} ",
        "conclusions": "\\label{sec-concl} The time evolution of the mass loss rate and wind velocity of a rapidly rotating, chemically homogeneous massive star differs radically from the evolution of a non-rotating star. This difference is reflected in the evolution of the CSM. As was shown in Sects. \\ref{sec-result1} and \\ref{sec-result2}, our rapidly rotating 20~$\\mso$ star produces a large bubble, which resembles that of non-rotating stars. However, during the two periods of near-critical rotation, large-scale deviations from spherical symmetry are produced. While the outer edge of the bubble remains spherical, closer to the star the asphericity of the wind creates unique morphologies. As the second phase of near-critical rotation occurs immediately before the supernova explosion, the supernova itself, and a possible GRB, will expand into such an aspherical medium. This will clearly influence the supernova remnant evolution, and the GRB afterglow light curve.  For a GRB, the effect of a non-spherical CSM can be felt in two ways. First, the relativistic jet will not expand in a normal, free streaming wind, but in a medium that has a far more complex structure (Sect.~\\ref{sec-result2}). In particular, the density slope in the innermost few parsec, well inside the wind termination shock, may differ significantly from $R^{-2}$, and may also show significant changes as function of radius. bf As the changes of the stellar wind during the last few 10000 yr depend on the internal structure of the star at this stage, it requires the consideration of a large grid of stellar evolution models to find out what the typical final density profiles would be, and how large their expected variation is  Second, the medium in front of the jet does not have a constant velocity. Assuming that not all of this material is ionized by the GRB, it may be observable in the form of blue-shifted absorption lines (Sect. \\ref{sec-spectrum}). Our investigation shows that the wind-wind interactions during the final stages of the stellar evolution produce multiple absorption lines at high velocity. Such lines were observed in the afterglow of \\object{GRB~021004}. While it is debated as to whether these absorption lines are in fact circumstellar in origin (Prochaska et al. \\cite{PCB06}, Chen et al. \\cite{Cetal07}), the possibility of producing them in this way is shown here for the first time. Furthermore, such lines might also be observed in early-time spectra of nearby Type~Ib/c supernovae at low metallicity."
    },
    "0711/0711.1755_arXiv.txt": {
        "abstract": "Kinematics combined with detailed element abundances provide a method of analysis of stellar populations that uses as much available information as possible, in contrast to other methods. Here we employ this technique on local A-type horizontal branch stars in an ongoing programme to search for $\\omega$ Centauri debris among these objects. This led us to discover another group of stars with very similar kinematics and abundances, the {\\bf C}ometary {\\bf O}rbit {\\bf G}roup (COG). We also comment on the future of this kind of undertaking in the Gaia era. ",
        "introduction": "There are quite a few approaches when it comes to the study of the stellar component of our Galaxy. Each of them has its merits and caveats, and in an overall picture these different methods complement each other. There are several important items to consider when planning a project aiming at analysing Galactic structure: \\begin{itemize} \\item What part of the Galaxy is the aim? Is it young or old? What has previously been known about it? How can previous results be improved on by adding new data or better defined and larger samples? \\item What is the best approach to achieve the goal, i.e. is it better to have a large sample at the expense of not having all information and/or tolerating less accurate information, or study a limited sample with more detailed and better information available? Obviously, it would be ideal to have as large a sample as possible with all information; however, in most cases this is not feasible. High-quality information such as (especially high-resolution) spectroscopy is expensive in terms of observing time and might even be beyond the reach of current telescopes for faint objects. The fortunate situation today is that more and more ambitious programs have become feasible. \\item Selection effects and statistical ambiguities: every sample, however carefully composed, is subject to selection effects. Owing to the Sun's location within the Milky Way, the pitfalls of selection effects are nowhere so common as in Galactic astronomy. While many of these can be accounted for in a statistical sense, some can only be described in a qualitative way, and need to be kept in mind when interpreting the results. Minimising such unwanted systematic side effects is one of the most important constraints when undertaking a study of Galactic structure. \\end{itemize}  Several classical methods exist, some of them having been in use for almost a century. The simplest one is star counts, especially in pencil beam fields \\citep[see e.g.][]{Altmann2007}. Here the spatial distribution of all stars or stars of a given temperature or colour range in a well defined field in the sky is studied. For this one needs to determine distances by means of photometric or spectroscopic parallaxes. Star counts can also be accomplished on the whole sky or parts of the sky, which allows for the detection of smaller-scale anomalies such as streams and moving groups.  \\begin{figure} \\begin{center} \\includegraphics[scale=0.6]{altmann_fig1.ps} \\end{center} \\caption{Toomre diagram for the HBA stars in our sample \\citep[Fig. 1 of][]{Altmann2005}. The concentric circles show the absolute peculiar velocity ($v_{\\rm pec}$), i.e. the total deviation from circular velocity (shown as an open star; the Sun's position in this diagram is shown as an open circle). The hexagons depict the stars belonging to the members of the Cometary Orbit Group (COG) discussed in this paper, the triangles denote putative $\\omega$ Cen debris candidates (identified as retrograde moving stars broadly conforming to the abundance range of $\\omega$~Cen giants), the other HBAs are represented by squares.} \\end{figure} \\vfill   Adding kinematics gives access to more parameters, such as the movement of objects. While a purely spatial analysis allows for the description of populations such as Galactic disks in a statistical sense, knowledge of the kinematics allows to assign particular objects to stellar groups or populations. It would be ideal to have the complete kinematical information, i.e. radial velocities {\\it and} proper motions, but significant progress has been made using one component alone, such as radial velocities in the case of the discovery of the Sagittarius dwarf \\citep{Ibata94}. Target objects can be stars in general or specific tracers, like white dwarfs \\citep[see, e.g.,][]{Pauli2003,Pauli2006} or horizontal branch (HB)-like stars.  For example, blue subdwarf (sdB) stars have been in the focus of several such studies \\citep[see, e.g.,][]{deBoer1997,Altmann2004}. The study of \\citet{Altmann2000} dealt with the complete HB blueward of the instability strip, while \\citet{Kaempf2005} studied red HB (RHB) and \\citet{Maintz2005} RR Lyrae stars \\citep[see][for a description of the several HB components]{mc05}. Caveats of using kinematics include variability of radial velocities caused by close binarity or pulsations. Proper motions are often difficult to find in the literature and take a long time to obtain, since one needs data from two epochs with as large as possible a time baseline. In recent years, several whole-sky catalogues \\citep[e.g.,][]{USNO} have appeared, improving the situation dramatically. Space-based astrometry from Hipparcos \\citep{Hipparcos} provided excellent proper motions for about 120,000 stars.  A further step in complexity is to add abundances to the kinematic analysis. It is known that the various components of the Galaxy have different abundance patterns, i.e. ratios between various element groups, such as iron peak and $\\alpha$-capture elements \\citep[see, e.g.,][and references therein]{Fuhrmann1998,mc07}. Differential abundances add important information about memberships of stars to certain groups and also the evolution of the Galaxy as a whole, and thus of galaxies in general. Since every population of stellar objects has its distinct star formation history as revealed by its abundance pattern, knowing kinematics {\\it and} differential abundances allows us to distinguish objects of different origin from within the Galaxy or from its outside, i.e. in dwarf galaxies that have been accreted by the Milky Way. A typical example for such a chemodynamical study is given by \\citet{Altmann2005}, which will also be the focus of the remainder of this paper. ",
        "conclusions": ""
    },
    "0711/0711.4941_arXiv.txt": {
        "abstract": "{We present high angular resolution ($\\sim$ 3$''$) and sensitive 1.3 mm continuum, cyanogen (CN) and vinyl cyanide (C$_2$H$_3$CN) line observations made with the Submillimeter Array (SMA) toward one of  most highly obscured objects of the W51 IRS2 region, W51 North.  We find that the CN line exhibits a pronounced inverse P-Cygni profile indicating that the molecular gas is infalling inwards this object with a mass accretion rate between 4 and 7 $\\times$ 10$^{-2}$ M$_\\odot$ yr$^{-1}$. The C$_2$H$_3$CN traces an east-west rotating molecular envelope that surrounds either a single obscured (proto)star with a kinematic mass of 40 M$_{\\odot}$ or a small central cluster of B-type stars and that is associated with a compact high velocity bipolar outflow traced by H$_2$O masers and SiO molecular emission. We thus confirm that the W51 North region is part of the growing list of young massive star forming regions that have been associated with infalling motions and with large mass accretion rates ($\\sim$ 10$^{-2}$ -- 10$^{-4}$  ), strengthening the evidence for massive stars forming with very high accretion rates sufficient to quench the formation of an UCHII region. } ",
        "introduction": "One of the main questions related to star formation is whether massive stars ($>$ 10 M$_{\\odot}$) are formed through gas accretion via a circumstellar disk/torus or whether other mechanims play a role. It was believed that the powerful radiation fields and stellar winds produced at the very beginning of nuclear burning will increasingly inhibit farther accretion of material thereby limiting the maximum stellar mass to about $10$ M$_{\\odot}$ \\citep{Kahn1974, LarsonStarrfield1971, YorkeKruegel1977}. Several theoretical models have since been proposed to solve this puzzle: the formation of massive stars through dense disks with jets/outflows \\citep{Nakano1989, JijinaAdams1996}, through merging of smaller stars \\citep{Bonnelletal1998}, through turbulent accretion \\citep{MckeeandTan2003},  through competitive accretion \\citep{Bonnellandbate2001}, and through ionized accretion flows \\citep{KetoWood2006}. However, due to the lack of good observational guidance, these alternatives have remained controversial \\citep{ZinneckerandYorke2007}.  With a total bolometric luminosity of about 3 $\\times$ 10$^6$ L$_{\\odot}$ the W51-IRS2 region is one of the most luminous massive star forming regions in our Galaxy, \\citep{EricksonTokunaga1980}. It is located 6--7 kpc away in  the Sagittarius spiral arm \\citep{Genzeletal1981, Imaietal2002}. We adopt here a distance to the W51 region of 7 kpc. The W51 IRS2 region comprises a complex group of highly obscured young objects (with not mid-infrared counterparts, see please \\citet{Kraemeretal2001,Okamotoetal2001}) called ``W51 North'' and ``W51d2'', and a cluster of massive strong infrared ZAMS stars \\citep{Kraemeretal2001,Okamotoetal2001, Lacyetal2007}, which the most prominent member being the source  ``IRS2d'' associated with an extended edge-brightened cometary HII region called ``W51d'' \\citep{Gaumeetal1993, Lacyetal2007}.  The W51 North object shows strong thermal dust emission at (sub)millimeter wavelengths, molecular emission from a large ($\\sim$ 4 $\\times$ 10$^4$ AU) {\\it hot core} at an excitation temperature of 100-200 K \\citep{Zhangetal1998}, and very faint centimeter free-free emission \\citep{Gaumeetal1993, Zhangetal1998, Eisneretal2002}, suggesting that it is forming an extremely young massive star. This object, in addition, contains a group of strong H$_2$O, OH and SiO masers that are located in the center of this molecular and dusty structure (called {\\it \"The Dominant Center\"}) \\citep{Schnepsetal1981, GaumeMutel1987, Hasegawaetal1986, Ukitaetal1987, Moritaetal1992}. Observations of the proper motions of the H$_2$O masers (which traces shocks in dust-laden gas close to the exciting protostars, Elitzur 1992) revealed the presence of a compact ($\\sim$ 7000 AU) northwest-southeast high velocity ($>$ 100 km s$^{-1}$) outflow \\citep{Schnepsetal1981, Eisneretal2002, Imaietal2002}. Moreover, high angular resolution observations showed that the SiO masers seem to be tracing the innermost parts of this powerful outflow \\citep{Eisneretal2002}. Finally, this object has previously been identified  with spectroscopic signatures of dynamical collapse using emission from HCO$^{+}$ \\citep{Rudolphetal1990}, and SO$_{2}$ \\citep{Sollinsetal2004}.  Here we present 1.3 mm continuum, cyanogen and vinyl cyanide line observations toward the W51 North region made with the SMA. We report the presence of molecular gas accretion onto a 40 M$_{\\odot}$ (proto)star or a small central cluster of B stars located in the center of W51 North region, and with an accretion rate between 4 and $7\\times$ 10$^{-2}$ M$_\\odot$ yr$^{-1}$ .  \\begin{figure*}[ht] \\begin{center} \\includegraphics[scale=0.44]{f1.eps}\\vspace{-0.8cm} \\caption{\\scriptsize Integrated molecular emission of the lines CN N=2-1, $J = 5/2-3/2, F = 5/2-3/2$ and  N=2-1, $J = 5/2-3/2, F =7/2-5/2$ observed with the SMA is shown as a color image of the W51 IRS2 region. The contours are -14, -13, -12, -11, -10, -9, -8, -7, -6, -5, -4, 4, 5, 6, 7, 8, 10, 11, 12, 13, and 14 times 970 mJy beam$^{-1}$, the rms noise of the image. The integration is over a velocity range from 50 to 80 km s$^{-1}$. The synthesized beam is 3.4 $''$ $\\times$ 3.2$''$ with a P.A. = -87$^\\circ$, and is shown in bottom left corner. The scale bar indicates the molecular line emission and absorption in Jy Beam$^{-1}$. The white diamonds indicate the positions of the infrared sources KJD3 and IRS2d \\citep{Kraemeretal2001, Okamotoetal2001, Lacyetal2007} and the radio source W51d2 \\citep{Gaumeetal1993}. The yellow line indicates the orientation of the SiO(5-4) high velocity bipolar molecular outflow centered on the water masers (Zapata et al. in prep.). The green circle indicates the position of the center of W51 North \\citep{Schnepsetal1981}.} \\label{fig1} \\end{center} \\end{figure*} ",
        "conclusions": "\\subsection{The Molecular and Millimeter Continuum Emission}  In Figure \\ref{fig1}, we present the integrated CN line emission image of the W51 IRS2 region, and the positions of the H$_2$O masers \\citep{Imaietal2002}, the infrared sources W51 IRS2d and KJD3/OKYM1 \\citep{Kraemeretal2001,Okamotoetal2001, Lacyetal2007}, and the radio source W51d2 \\citep{Gaumeetal1993}, all of them in the neighborhood of the W51 North object.  Furthermore, we mark the position and orientation of the strong high velocity molecular bipolar outflow traced by the SiO \\textit{v}$ =0; J=5 - 4$ line found by Zapata et al. (in prep.), which is centered on the cluster of masers.  In this image we can see two components of the molecular line distribution, one in unresolved absorption toward the W51 North object and the other also quite compact, but resolved in emission surrounding this region. We interpret this appearance the signature of molecular absorption against the strong compact millimeter continuum source associated with W51 North and shown in Figure 3.   Figure \\ref{fig2} shows the spectrum of the CN line towards the center of the W51 North object. The line shows a pronounced inverse P Cygni profile.  The emission feature appears at a V$_{LSR}$=55 km s$^{-1}$ and absorption at V$_{LSR}$=65 km s$^{-1}$, which is consistent with the HCO$^+$ molecular observations of \\citet{Rudolphetal1990}. Given that the core's  systemic velocity is 59 km s$^{-1}$ \\citep{Zhangetal1998}, the location of the redshifted absorption projected against the bright continuum emission of the central highly obscured object implies inward motions away from the observer. The brightness temperature of the emission from the infalling material is lower than $\\sim$ 10 K, the brightness temperature of the continuum emission from the core.  From Figure 6 of \\citet{Rudolphetal1990} and Figure 2 shown here, we estimate that the velocity of infall ($V_{infall}$) is about 4 km s$^{-1}$.  With this information and taking the values of the density ($\\rho$=2$\\times$10$^{6}$ cm$^{-3}$), linear radius (r=1.4 $\\times$10$^4$ AU) reported for the compact continuum source located in W51 North \\citep{Zhangetal1997, Zhangetal1998} and a radius of the hot core of 2 $\\times$ 10$^4$ AU, and following Beltran et al.  (2006), we calculate that the mass infall rate (\\.{M}=$4 \\pi r^2 \\rho V_{infall}$) is between 4 and $7 \\times 10^{-2}$ M$_\\odot$ yr$^{-1}$. This value have large uncertainties, due to the uncertainty on the density and on the radius at which V$_{infall}$ is measured.  The first moment map of the C$_2$H$_3$CN {\\bf $J_{K_a, K_c}= 23_{2,22}-22_{2,21}$}  line is shown in Figure 3. The emission is tracing an unresolved east-west rotating molecular ``envelope'' or ``core'' with a total velocity shift of 1.5 km s$^{-1}$ and a size of 4 $\\times$ 10$^4$ AU. Moreover, the integrated emission from this molecule is well centered on the millimeter continuum source, suggesting that this species is tracing high density gas close to the (proto)star (see Figure 3). Finally, our bandpasses contained lines from other molecules (e.g. HCOOCH$_3$ and CH$_3$OH) associated with W51 North. However, they were very much contaminated by the emission from the hot molecular core associated with W51d2, not allowing us to search for similar east-west velocity gradients associated with these line molecular tracers. The W51 North source is associated with the extended {\\it hot core} found by \\citet{Hoetal1983} and \\citet{Zhangetal1998}, and it is centered on the cluster of masers as reported in other observations e.g. NH$_3$; \\citet{Hoetal1983,Eisneretal2002}, CH$_3$CN; \\citet{Zhangetal1998}, SO$_2$; \\citet{Sollinsetal2004}.  If we assume that the molecular gas is rotating as a rigid body (i.e. the dynamical mass is M$_{dyn}$=v$^2$rsin$^2$({\\it i})/G, where v is the rotation velocity, r is the radius of the envelope, {\\it i} is the inclination angle of the envelope assumed to be 90$^\\circ$ and G is the gravitational constant), we estimate a mass for the central object(s) of 40 M$_\\odot$.  This central object might be associated with a single central O-type (proto)star or with a small group of B-type (proto)stars. However, as there is a strong and compact bipolar outflow in the center of the core (see Figure 3), it seems to be dominated by one central massive star.  From Figure 3 and assuming that at a wavelength of 1.3 mm we are observing isothermal optically thin dust emission with a dust mass opacity coefficient that varies with frequency as $\\kappa \\propto \\nu^{\\beta}$, with $\\beta=1$, (the size of the source suggest that we are observing emission from the envelope; hence, we adopt $\\beta$=1, however, this value is uncertain, see \\citet{Beckwithetal1990} that observe how this value varies in pre-main-sequence stars), a gas-to-dust ratio of 100 (which may not be the most adequate to use for protostellar sources since erosion of the circumstellar envelope by photoevaporation from near OB stars may decrease the gas-to-dust ratio, see \\citet{Williamsetal2005, ThroopBally2005}), an adopted value of $\\kappa_{1.3 mm}$ = 1.5 cm$^2$ g$^{-1}$ and a dust temperature value of  about 100 K (with uncertainties of a factor of about 1.5, \\citet{Zhangetal1998}), we estimate an enclosed mass of the molecular core W51 North of 90 M$_\\odot$, very close to the 100 M$_\\odot$ estimated by \\citet{Zhangetal1997}. Due to the uncertainties referred to above, the values of the derived masses are good within a factor of 2.  \\begin{figure}[ht] \\begin{center} \\includegraphics[width=7.8cm, height=8.5cm]{f3.eps}\\\\ \\vspace{-0.8cm} \\includegraphics[width=7.83cm, height=8.9cm]{f4.eps} \\caption{\\scriptsize {\\it \\bf Upper:} SMA 1.3 mm continuum emission color image overlayed with the  moment zero distribution of the C$_2$H$_3$CN line (pink contours) of the W51 North region. The integration is over a velocity range of 55 to 65 km s$^{-1}$.  The synthesized beam with a FWHM 3.4 $''$ $\\times$ 3.2$''$ and a P.A. of  $-87^\\circ$ is shown in the bottom left corner.  The contours are -4, 4, 8, 12, 16, 20, 24, 30, 40, 50, 60, 80, 90, 100, 120, 150, 170 and 200 times 40 mJy beam$^{-1}$ km s$^{-1}$, the rms noise of the image.  The scale bar indicates the continuum peak flux density in Jy beam$^{-1}$.  {\\it \\bf Lower:} First moment color image of the C$_2$H$_3$CN emission toward the W51 North region.  The integration is over a velocity range of 55 to 65 km s$^{-1}$.  The scale bar indicates the velocity shift in km s$^{-1}$.  The white diamonds indicate the positions of the infrared sources KJD3 and IRS2d \\citep{Kraemeretal2001, Okamotoetal2001, Lacyetal2007} and the radio source W51d2 \\citep{Gaumeetal1993}. The blue and red crosses indicate the position of the blue- and red-shifted strong H$_2$0 masers spots, respectively, reported by \\cite{Imaietal2002}. Note that the central cluster of masers is tracing a high velocity outflow with a northwest-southeast orientation \\citep{Imaietal2002, Eisneretal2002}. The yellow line indicates the orientation of the SiO(5-4)  high velocity bipolar outflow centered on the water masers (Zapata et al. in prep.). The black circle indicates the position of the center of W51 North \\citep{Schnepsetal1981}.} \\end{center} \\label{fig3} \\end{figure}  \\subsection{Forming an Early O-type star in W51 North?}  The combined 1.3 mm continuum, C$_2$H$_3$CN and CN data from W51 North suggest that this object is forming a massive O5-type (proto)star in its center through molecular gas accretion and with a very large accretion rate between 4 and $7\\times$ 10$^{-2}$ M$_\\odot$ yr$^{-1}$.  Moreover, the powerful compact bipolar high velocity ($>$ 100 km s$^{-1}$) outflow traced by the H$_2$O masers pinpoints the position of this putative central massive object.  Zapata et al. (in prep.) in addition found a compact high velocity SiO bipolar outflow with both unresolved-lobes spatially separating for less than one arcsecond and forming a P. A. of 150$^\\circ$ $\\pm$ 30$^\\circ$. This is centered at the position of the compact H$_2$O outflow which has a P.A. of 110-145$^\\circ$ \\citep{Schnepsetal1981, Imaietal2002, Eisneretal2002}.  We propose that these two outflows might be manifestations of a single powerful outflow.  The masers appear to be tracing the innermost regions of it  (as observed in other outflows, e.g. IRAS 20126+4104, Moscadelli et al. 2000), while the SiO is tracing the more extended shocked molecular gas. From this point of view, the southeastern cluster of H$_2$O masers (see Figure 3) maybe tracing another older ejected bow-shock. However, higher angular resolution SiO molecular observations are neccesary to confirm this picture.  It is interesting to note that the P.A. of the velocity gradient across the C$_2$H$_3$CN envelope ($\\sim$ 90$^\\circ$) is not exactly perpendicular to the orientation of the molecular outflow (P.A. $\\sim$ 140$^\\circ$), as might be expected. This may indicate that outflow could be precessing due the presence of binary system or that the C$_2$H$_3$CN  emission is contaminated by the outflow. This physical phenomenon of precessing outflows has been reported in other outflows: IRAS 20126+4104 \\citep{Shepherdetal2000}, L115{\\bf 7} \\citep{Bachilleretal2001}, and NGC7538IRS1 \\citep{Krausetal2006}.  At present, there is a list of early massive (proto)stars that have been associated with possible infalling motions and with large mass accretion rates, e.g.  W51e2: a gas mass of 200 M$_{\\odot}$ and an accretion rate of 10$^{-3}$ M$_\\odot$ yr$^{-1}$ \\citep{Zhangetal1997,Hoandyoung1996}; NGC7538-IRS9: a gas mass of 100-300 M$_{\\odot}$ and an accretion rate of 10$^{-3}$ M$_\\odot$ yr$^{-1}$ \\citep{Sandelletal2005}; G24.78+0.08: a (proto)stellar mass of $\\sim$ 20 M$_\\odot$ and an accretion rate between 10$^{-2}$ to 10$^{-4}$ M$_\\odot$ yr$^{-1}$ \\citep{Beltranetal2006}; IRAS 16547-4247: associated with an O-type (proto)star and an accretion rate of about 10$^{-2}$ M$_\\odot$ yr$^{-1}$ \\citep{Garayetal2007} and W51 North, a (proto)stellar mass of $\\sim$ 40 M$_\\odot$ and an accretion rate between 4-7 $\\times$ 10$^{-2}$ M$_\\odot$ yr$^{-1}$  (these results). This suggests that the very massive stars (O-type) might form starting with very high accretion rates, sufficient to quench the formation of an UCHII region.  We note that this hypothesis has already been proposed by recent numerical simulations \\citep{Banerjeeetal2007}.  We thank the anonymous referee for many valuable suggestions. R.G. is grateful to the {\\it Alexander von Humboldt Foundation} for a Humboldt Research Fellowship."
    },
    "0711/0711.4988_arXiv.txt": {
        "abstract": "We study the geometry and the physical conditions in the inner (AU-scale) circumstellar region around the young Herbig~Be star MWC~147 using long-baseline spectro-interferometry in the near-infrared (NIR $K$-band, VLTI/AMBER observations and PTI archive data) as well as the mid-infrared (MIR $N$-band, VLTI/MIDI observations). The emission from MWC~147 is clearly resolved and has a characteristic physical size of $\\sim1.3$~AU and $\\sim9$~AU at $2.2~\\mu$m and $11~\\mu$m respectively (Gaussian diameter). The MIR emission reveals asymmetry consistent with a disk structure seen under intermediate inclination. The spectrally dispersed AMBER and MIDI interferograms both show a strong increase in the characteristic size towards longer wavelengths, much steeper than predicted by analytic disk models assuming power-law radial temperature distributions.  We model the interferometric data and the spectral energy distribution of MWC~147 with 2-D, frequency-dependent radiation transfer simulations. This analysis shows that models of spherical envelopes or passive irradiated Keplerian disks (with vertical or curved puffed-up inner rim) can easily fit the SED, but predict much lower visibilities than observed; the angular size predicted by such models is 2 to 4 times larger than the size derived from the interferometric data, so these models can clearly be ruled out. Models of a Keplerian disk with optically thick gas emission from an active gaseous disk (inside the dust sublimation zone), however, yield a good fit of the SED and simultaneously reproduce the absolute level and the spectral dependence of the NIR and MIR visibilities. We conclude that the NIR continuum emission from MWC~147 is dominated by accretion luminosity emerging from an optically thick inner gaseous disk, while the MIR emission also contains contributions from the outer, irradiated dust disk. ",
        "introduction": "\\label{sec:intro}  The spatial structure of the circumstellar material around Herbig~Ae/Be (HAeBe) stars, i.e.\\ intermediate-mass, pre-main sequence stars, is still a matter of debate.  Until recently, the spatial scales of the inner circumstellar environment (a few AU, corresponding to $\\lesssim $0\\farcs1) were not accessible to optical and infrared imaging observations, and conclusions drawn on the spatial distribution of the circumstellar material were, in most cases, entirely based on the modeling of the spectral energy distribution (SED). However, fits to the observed SEDs are highly ambiguous, and very different models such as spherical envelopes \\citep{mir97}, geometrically thin disks \\citep{hil92}, disks surrounded by a spherical envelope \\citep{nat95, nat01, mir99}, flared outer disks, puffed-up inner disk rims \\citep{dul01}, and disk plus inner halo models \\citep{vin06} have been used to successfully fit the observed SEDs of HAeBes. It was also unclear how the accretion of circumstellar material onto the star affects the structure and emission of the inner circumstellar environment.  In the last couple of years, long-baseline interferometry in the near- (NIR) and mid-infared (MIR) spectral range provided, for the first time, direct spatial information on the {\\em inner} circumstellar regions of young stars at scales of $\\lesssim 10$~milli-arcseconds \\citep[mas; e.g.,][]{mal98, mil99, ake00, mil01, wil03, eis04, lei04, mon05,mal07}. The interferometric data were usually interpreted by fitting simple analytic models for the brightness distribution to the measured visibilities. The characteristic sizes of the emitting regions derived in this way were found to be correlated to the stellar luminosity, consistent with the idea that the NIR continuum emission mainly traces hot dust close to the dust sublimation radius, i.e.~where the grains are heated above their sublimation temperature ($T_{\\rm sub} \\sim 1\\,500$~K) and destroyed by the (stellar) radiation field \\citep{tut02, mon02}. While the expected simple $R_{\\star}\\propto L_{\\star}^{1/2}$ scaling law between stellar luminosity $L_{\\star}$ and NIR size $R_{\\star}$ appears to hold throughout the low- to medium-luminosity part of the observed stellar sample, some very luminous early B-type stars exhibited considerably smaller NIR sizes than predicted by this relation \\citep{mon02,eis04,mon05}. \\citet{mon02} suggested that this might be due to the presence of an inner gaseous disk, which shields the dust disk from the strong stellar ultraviolet (UV) radiation. Since this shielding would be most efficient for hot stars, it would allow the inner rim of the dust disk around B-type stars to exist closer to the star. Several subsequent studies favour ``classical'' accretion disk models \\citep{eis04,mon05,vin07} or ``two-ring'' models \\citep{eis07}, in which the infrared emission contains contributions from the thermal emission of optically thick gas in the innermost disk regions.  Following initial attempts by \\citet{hin01}, who used single-dish nulling interferometry to put upper limits of $\\lesssim 20$~AU on the MIR size of some Herbig~Ae stars, a large sample of HAeBes disks could be resolved with the VLTI/MIDI long-baseline interferometer at MIR wavelengths. In contrast to the NIR flux, which originates from hot dust in the innermost disk regions close to the dust sublimation radius, the MIR emission also traces cooler ($\\sim 200-300$~K) dust further out in the disk. \\citet{lei04} determined characteristic dimensions of the $10~\\mu$m emitting regions for a sample of HAeBes, which ranged from 1~AU to 10~AU.  Due to the quite limited $uv$-plane coverage of most existing infrared interferometric data sets, most published studies were only able to derive estimates of characteristic sizes and, in some cases, to look for possible elongation of the emitting region, but could not investigate the geometry of individual sources in detail. A more comprehensive  interferometric study of one Herbig~Ae star, HR~5999, was recently performed by our group \\citep{pre06}. Based on a set of ten MIDI measurements at different projected baseline lengths and position angles (PAs), modeling with 2-D frequency-dependent radiation transfer simulations provided relatively detailed information on the disk size and inclination.\\newline  In this study, we combine, for the first time, NIR and MIR interferometric observations to constrain the spatial structure of the dust and gas environment around a Herbig~Be star. Besides the wide spectral coverage (from $\\sim 2$ to $13~\\mu$m), our interferometric data is also dispersed into several spectral channels (resolution $R=\\lambda/\\Delta\\lambda \\approx 30$), allowing us to measure the wavelength-dependence of the visibility over the $K$- and $N$-band and within spectral features such as the silicate emission feature around $10~\\mu$m. In combination with detailed 2-D radiative transfer modeling, this spectro-interferometric data set provides unique information about the inner circumstellar structures of this star.  \\medskip  \\begin{deluxetable}{rcl} \\tabletypesize{\\scriptsize} \\tablecaption{Stellar parameters for MWC~147 from \\citet{her04b}, assuming $R_V=3.1$. \\label{tab:stellarproperties}} \\tablewidth{0pt} \\centering  \\tablehead{ \\colhead{Parameter} & \\colhead{} & \\colhead{Value} } \\startdata Spectral Type    &             & B6\\\\ Effective Temperature & $T_{\\mathrm{eff}}$ & 14\\,125~K\\\\ Luminosity       & $L_{\\star}$ & 1\\,550~$L_{\\sun}$\\\\ Mass             & $M_{\\star}$ & 6.6~$M_{\\sun}$\\\\ Age              & $t$         & 0.32~Myr\\\\ Distance         & $d$         & 800~pc\\\\ Extinction       & $A_V$       & 1.2~mag\\\\ Stellar Radius   & $R_{\\star}$ & 6.63~$R_{\\sun}$\\\\ \\enddata \\end{deluxetable}  MWC~147 (alias HD\\,259431, BD+10\\,1172, HBC 529, V700\\,Mon) is  a well-studied Herbig~Be star in Monoceros. There is some uncertainty concerning the distance and the physical parameters of this star. From the analysis of the Hipparcos parallax data by \\citet{van98}, a lower limit on the distance of $>\\!130$~pc was derived, while a reanalysis suggested a distance of $290^{+200}_{-84}$~pc \\citep{ber99}. This distance estimate, however, is in conflict with the apparent location of MWC~147 in the NGC\\,2247 dark cloud, which is a part of the cloud complex in the Monoceros OB1 association at a distance of $\\sim 800-900$~pc \\citep{oli96}. We assume a distance of 800~pc for MWC~147 (consistent with most other recent studies) and use the main stellar parameters as listed in Tab.~\\ref{tab:stellarproperties}, which were taken from \\citet{her04b}.  Numerous observational results strongly suggest the presence of a circumstellar disk around MWC~147. The object shows a strong infrared excess of about 6~mag at MIR wavelengths, clearly demonstrating the presence of circumstellar material. \\citet{hil92} fitted the SED of MWC~147 with a model assuming an accretion disk and estimated an accretion rate of $\\dot{M}_{\\rm acc} = 1.01 \\times 10^{-5}~M_{\\sun} $yr$^{-1}$. MIR (10~$\\mu$m and 18~$\\mu$m) imaging observations revealed an elongated diffuse emission component around MWC~147 along PA $\\sim$50\\degr, extending out to $\\sim$6\\arcsec\\ and contributing $\\sim 34$\\% to the total flux \\citep{pol02}. \\citet{man94} determined the 1.1~mm flux of MWC~147 and estimated the mass in the circumstellar disk/envelope to be $< 0.09~M_{\\sun}$. The study of the far-UV spectrum of MWC~147 by \\citet{bou03} also suggested the presence of a flared circumstellar disk. \\citet{pol02} imaged MWC~147 in the MIR and concluded that the star is surrounded by a moderately flared disk and probably an extended envelope. Measurements by \\citet{jai90} showed a significant amount of linear polarization ($\\sim$1\\% along PA $\\sim 106\\degr$) but no wavelength-dependence of the polarization. The high observed rotational velocity of $v \\sin i = 90$~km\\,s$^{-1}$ \\citep{boe95} suggests a high inclination of the rotation axis of MWC~147 with respect to the line-of-sight. This implies that the orientation of the circumstellar disk should be closer to edge-on than to face-on.  Recently, \\citet{bri07} presented a high-resolution NIR spectrum of MWC~147, showing a strong Br$\\gamma$ emission line.  Using an empirical relation between the Br$\\gamma$ luminosity and the accretion rate \\citep[as derived from UV veiling;~][]{van05}, % they derive a mass accretion rate of $\\dot{M}_{\\rm acc} = 4.1 \\times 10^{-7}~M_{\\sun} $yr$^{-1}$.  As discussed in their Section 5.2, it is well possible that this method underestimates the true mass accretion rate.  Adopting the stellar parameters used in our study (Tab.~\\ref{tab:stellarproperties}) will also result in a larger value for the derived accretion rate.  Evidence for a strong stellar wind from MWC~147 comes from the observed P~Cygni profiles in several emission lines \\citep{bou03}. A quantitative modeling of FUSE spectra revealed multiple absorption components with different temperatures, consistent with a flared disk interpretation \\citep{bou03}. Based on the intensity ratio of infrared hydrogen lines, \\citet{nis95} estimated a mass loss rate of $2.0 \\pm 0.4 \\times 10^{-7}~M_{\\sun}$yr$^{-1}$, which is slightly higher than the mass loss rates determined from radio observations \\citep[$0.68 \\times 10^{-7}~M_{\\sun}$yr$^{-1}$, ][]{ski93}.  The star has a faint visual companion at a projected separation of 3\\farcs1~\\citep[$\\sim 2\\,500$~AU, $\\Delta R=6.82$, ][]{bai06}. While \\citet{vie94} classified MWC~147 as a spectroscopic binary with a period of about one year, this claim could not be confirmed in more recent observations \\citep{cor99}.  First interferometric measurements on MWC~147 were presented by \\citet{mil01}, providing an upper limit on the $H$-band size. \\citet{ake00} observed MWC~147 with the Palomar Testbed Interferometer (PTI) and resolved its emission in the $K$-band at baselines around 100~m. They derived a best-fit Gaussian FWHM diameter of 1.38~mas (=1.1~AU) in the $K$-band. ",
        "conclusions": "\\label{cha:conclusions}  We have presented infrared long-baseline interferometric observations of MWC~147, constraining, for the first time, the inner circumstellar environment around a Herbig~Be star over the wavelength range from 2 to 13~$\\mu$m.  The interferometric data obtained from the PTI archive and with VLTI/AMBER suggest a characteristic diameter of just $\\sim 1.3$~AU (Gaussian FWHM) for the NIR emitting region, while the MIR structure is about a factor of 7 more extended (9~AU at 11~$\\mu$m). Within the $K$-band, we measure a significant increase of size with wavelength. Comparing the measured wavelength-dependence of the characteristic size with the class of analytic accretion disk models, which are currently most commonly applied for the modeling of HAeBe SEDs \\citep{hil92} and NIR interferometric data \\citep[e.g.][]{vin07} yields that these models cannot well reproduce the measured significant increase of the apparent size with wavelength. To test whether more realistic physical models of the circumstellar dust environment yield better agreement, we employed 2-D radiative transfer modeling.  The radiative transfer models were constructed to fit the SED from 0.3 to 450~$\\mu$m, including an extended envelope and an inner dust shell/disk geometry. While models of passive irradiated disks with or without puffed-up rims are able to reproduce the SED, they are in conflict with the interferometric observables, significantly overestimating the size of both the NIR and MIR emission. This is also the case, to an even bigger extend, for spherical shell geometries. Therefore, in our radiative transfer models, we incorporated additional accretion luminosity emitted from an inner gaseous disk, yielding significantly better agreement with the interferometric data. The best-fit was obtained with a flared Keplerian disk seen under an inclination of $\\sim 50$\\degr, extending out to 100~AU and exhibiting a mass accretion rate of $7 \\times 10^{-6}~M_{\\sun}$yr$^{-1}$.  Since MWC~147 belongs to the group of ``undersized'' Herbig~Be star disks (which have measured NIR diameters smaller than expected from the size-luminosity relation; see \\citealt{mon05}), our detailed spectro-interferometric study confirms earlier speculations that the NIR emission of Herbig~Be stars in this group might be dominated by gas emission from inside the dust sublimation radius.  Furthermore, our study demonstrates that the spectro-interferometric capabilities of the latest generation of long-baseline interferometric instruments are particularly well suited  to reveal the contributions from active accretion processes taking place close to the star. To achieve further progress on MWC~147, it seems promising for future spectro-interferometric observations not only to extend the spectral coverage (e.g.\\ to the $J$- and $H$-band) but to employ also a higher spectral resolution in order to study the spatial origin of accretion-related emission lines \\citep[e.g.\\ Br$\\gamma$, ][]{gar06}, maybe tracing ionized, optically thin gas in the disk or magnetospheric accretion \\citep{har94}. In addition, future long-baseline interferometric observations will measure the closure phase relation, which is a sensitive measure to asymmetries in the source brightness distribution, and which provides also the key to direct aperture synthesis imaging of the sub-AU circumstellar environment."
    },
    "0711/0711.3085_arXiv.txt": {
        "abstract": "We show that TeV \\gr\\ emission produced via interactions of high-energy particles with anisotropic radiation field of a massive star in binary systems should have a characteristic rotating hollow cone anisotropy pattern. The hollow cone, whose axis is directed away from the massive star, rotates with the period equal to the orbital period of the system. We note that the two maxima pattern of the TeV energy band lightcurve of the \\gr\\ loud binary LS 5039 can be interpreted in terms of this rotating hollow cone model. Adopting such an interpretation, we are able to constrain the geometry of the system -- either the inclination angle of the binary orbit, or the elevation of the \\gr\\ emission region above the orbital plane. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.4616_arXiv.txt": {
        "abstract": "This paper presents Spitzer-IRS $\\lambda/\\Delta\\lambda\\sim 600$ spectroscopy of the CO$_2$ 15.2\\,$\\mu$m bending mode toward a sample of 50 embedded low-mass stars in nearby star-forming clouds, taken mostly from the ``Cores to Disks (c2d)'' Legacy program. The average abundance of solid CO$_2$ relative to water in low-mass protostellar envelopes is $0.32\\pm 0.02$, significantly higher than that found in quiescent molecular clouds and in massive star forming regions. It is found that a decomposition of all the observed CO$_2$ bending mode profiles requires a minimum of five unique components. In general, roughly 2/3 of the CO$_2$ ice is found in a water-rich environment, while most of the remaining 1/3 is found in a CO environment with strongly varying relative concentrations of CO$_2$ to CO along each line of sight. Ground-based observations of solid CO toward a large subset of the c2d sample are used to further constrain the CO$_2$:CO component and suggest a model in which low-density clouds form the CO$_2$:H$_2$O component and higher density clouds form the CO$_2$:CO ice during and after the freeze-out of gas-phase CO. The abundance of the CO$_2$:CO component is consistent with cosmic ray processing of the CO-rich part of the ice mantles, although a more quiescent formation mechanism is not ruled out. It is suggested that the subsequent evolution of the CO$_2$ and CO profiles toward low-mass protostars, in particular the appearance of the splitting of the CO$_2$ bending mode due to pure, crystalline CO$_2$, is first caused by distillation of the CO$_2$:CO component through evaporation of CO due to thermal processing to $\\sim 20-30\\,$K in the inner regions of infalling envelopes. The formation of pure CO$_2$ via segregation from the H$_2$O rich mantle may contribute to the band splitting at higher levels of thermal processing ($\\gtrsim 50\\,$K), but is harder to reconcile with the physical structure of protostellar envelopes around low-luminosity objects. ",
        "introduction": "Although CO$_2$ is not an abundant gas-phase molecule in molecular clouds, it is one of a small number of molecular species consistently found in very high abundances inside ice mantles on dust grains \\citep[$>10^{-5}$ with respect to H$_2$][]{Gerakines99,Whittet07}. The generally high abundance of solid CO$_2$ became apparent with the spectroscopic surveys conducted with the Infrared Space Observatory (ISO) \\citep{Gerakines99, Nummelin01}. Other species known to belong to this class of very abundant molecules are CO and H$_2$O. In less than 10\\% of dark cloud regions surveyed, methanol (CH$_3$OH) is also found with similar abundances \\citep{Pontoppidan03b,Boogert07}. Depending on the density and temperature of a cloud, the CO is found partly in the gas-phase and partly frozen onto grain surfaces, while the CO$_2$ and H$_2$O are completely frozen as ice mantles \\citep{Bergin95}, except in very hot or shocked regions \\citep{Boonman03,Nomura04,Lahuis07}. The system of CO, CO$_2$, H$_2$O and, under some conditions, CH$_3$OH therefore represents the bulk of solid state volatiles in dense star forming clouds, and interactions between these four species can be expected to account for most of the solid state observables. Other species with abundances of less than 5\\% relative to water, such as CH$_4$, NH$_3$, OCN$^-$, HCOOH and OCS will be good tracers of chemistry and their local molecular environment, but are unlikely to strongly affect the molecular environments, and therefore the band profiles, of the four major species.  The formation mechanism of solid CO$_2$ in the cold interstellar medium is still not understood, although a number of plausible scenarios have been proposed. Since the direct surface route, $\\rm CO+O\\rightarrow CO_2$, is thought to possess a significant activation barrier, it was initially suggested that strong UV irradiation was needed to produce the observed CO$_2$ ice \\citep{dhendecourt85}. Laboratory simulations of interstellar ice mixtures of H$_2$O and CO confirmed that CO$_2$ is indeed readily formed during strong UV photolysis \\citep{dhendecourt86}, and initial detection of abundant CO$_2$ ice around UV-luminous massive young stars seemed to confirm this picture. However, recent detections of similar abundances of CO$_2$ in dark clouds observed along lines of sight toward background stars, far away from any ionizing source \\citep{Bergin05,Knez05,Whittet07}, argue against a UV irradiation route to CO$_2$, at least through enhanced UV from nearby protostars. \\cite{Pontoppidan06} showed evidence for an increasing abundance of CO$_2$ ice with gas density in at least one low-mass star forming core. Furthermore, the original premise of a barrier to oxygenation of CO is now in doubt \\citep{Roser01}. Consequently, both theoretical and laboratory efforts to understand the formation of CO$_2$ are still very active.   Extensive surveys of the 3.1 and 4.67\\,$\\mu$m stretching mode of H$_2$O and CO ices have been carried out in a range of different star-forming environments \\citep{Whittet88, Chiar95, Pontoppidan03}. However, CO$_2$ can only be observed from space, and surveys have until recently been limited to a small sample of very luminous young stars \\citep{Gerakines99}. While the 4.27\\,$\\mu$m stretching mode of CO$_2$ was detected toward a few background field stars by ISO \\citep{Whittet98,Nummelin01}, recent Spitzer observations of the 15.2\\,$\\mu$m bending mode of CO$_2$ have extended the sample of CO$_2$ along quiescent lines of sight considerably \\citep{Knez05, Bergin05, Whittet07}.  The profiles of the CO$_2$ ice bands observed in quiescent regions and along lines of sight toward luminous protostars show intriguing differences. In particular, ice in massive star-forming regions, which presumably traces more processed material, tends to show a splitting of the 15.2\\,$\\mu$m bending mode. This splitting has been identified as a general property of crystalline pure CO$_2$ and is readily reproduced in laboratory simulations of interstellar ices \\citep[e.g.][]{Ehrenfreund97, Broekhuizen06}. It seems implausible that this pure CO$_2$ layer would form directly through gas-phase deposition and subsequent surface reactions, and it has been suggested that the CO$_2$ segregates from a CO$_2$:H$_2$O:CH$_3$OH=1:1:1 mixture upon strong heating \\citep{Gerakines99}. Annealing in a laboratory setting produces the same effect  \\citep{Ehrenfreund97}, and strong heating is not unreasonable in the envelopes of massive young stars with luminosities in excess of $10^3\\,\\rm L_{\\odot}$.   In this paper, a survey of the 15.2\\,$\\mu$m CO$_2$ bending mode toward $\\sim$50 young {\\it low mass} stars using the high resolution mode of the Infrared Spectrograph on board the Spitzer Space Telescope (Spitzer-IRS) is presented. The sample stars have typical luminosities in the range $0.1-10\\,\\rm L_{\\odot}$, thus bridging the observational gap between the background stars and the massive protostars from the ISO sample. The Spitzer spectra are complemented by ground-based observations of H$_2$O and CO ices, where available, as well as the archival spectra from the Infrared Space Observatory used in \\cite{Gerakines99}. While this paper concentrates on the region around 15\\,$\\mu$m, \\cite{Boogert07} discusses the ices causing the 5-8\\,$\\mu$m absorption complex (henceforth referred to as Paper I). The 7.7\\,$\\rm\\mu$m CH$_4$ and 9.0\\,$\\mu$m NH$_3$ bands are described in two separate papers ({\\\"O}berg et al. 2008, submitted and Bottinelli et al., in prep, respectively).  The central questions that are addressed using the new Spitzer spectra of the CO$_2$ bending mode are:  \\begin{itemize} \\item What are the differences, if any, between CO$_2$ ice in massive and low-mass star-forming environments? \\item What is the average abundance of CO$_2$ in low-mass protostellar envelopes compared to lower density quiescent clouds? \\item In which molecular environments can solid CO$_2$ be found, and what are their relative abundances? \\item Which process forms CO$_2$ in CO-dominated environments when the CO accretes from the gas-phase at high densities? \\item How does the component of pure CO$_2$ (as measured by the well-known splitting of the bending mode) form in low-mass protostellar envelopes? \\item What is the evolution of the CO$_2$ ice and what does it tell us about protostellar evolution? \\end{itemize} ",
        "conclusions": "In this paper, a picture is presented in which the dominant ice components, H$_2$O, CO$_2$, CO and, in a limited number of cases, CH$_3$OH appear to constitute an intimately connected system; the observational characteristics of each species are directly affected by the presence of the others. Other molecules act as trace constituents with abundances that are too low to significantly affect the profiles of the dominant species, i.e. other species did not have to be invoked to explain the observed band profiles. This is in contrast to the very complex 5-8\\,$\\mu$m region discussed in Paper I.  The survey has shown that the CO$_2$ bending mode profiles toward low-luminosity ($0.1<L<100\\,\\rm L_{\\odot}$), solar-type YSOs do not differ strongly in their basic structure from those observed toward massive, luminous YSOs ($\\rm L> 1000\\,L_{\\odot}$). All protostellar envelopes as well as quiescent molecular clouds probed by lines of sight toward background stars are dominated by CO$_2$ mixed with water with abundances of 20\\% relative to water ice, but with a significant additional contribution of 10\\% of CO$_2$ mixed with CO. The other components needed to explain the CO$_2$ bending mode are present with much smaller abundances, or a few \\% each with respect to water ice.  While the CO$_2$:H$_2$O component might form from CO and OH via reaction (\\ref{OHroute}), it is probable that, because of the lack of residual atomic hydrogen, the CO$_2$:CO component forms through a different chemical route. Additionally, since the band profiles indicate that the CO$_2$ in this component can have concentrations of anything from 2:1 to less than 1:100 relative to CO, a variable effect highly sensitive to either the environment or age of the ice must play a role. Based on observations of a sample of background stars, \\cite{Bergin05} suggest that the formation of the CO$_2$:CO component is related to the abundance of atomic oxygen relative to atomic hydrogen in the gas-phase and predict that it forms in low-density regions of the cloud. Observationally, this would produce CO and CO$_2$ profiles dominated by the CO$_2$:CO$\\sim$1:1 component in low-density regions. It would also result in the CO$_2$:CO layer being placed below the water-rich layer on each grain. In this case, the formation of pure CO$_2$ by the desorption of CO as proposed here would not work.  A different scenario is suggested in which the CO$_2$:CO component is connected to the rapid freeze-out of pure CO at higher densities. This scenario is practically independent on the gas-phase chemistry, but requires a mechanism for forming CO$_2$ directly from CO. Such a mechanism will most likely involve a highly energetic input from cosmic rays; UV irradiation with $\\lambda > 1200\\,\\AA$ will not work in the absence of water \\citep{Oberg07_letter}. While the energy input from a standard cosmic ray field is likely sufficient, this process would also tend to form more complex carbon oxides in the ice. While these are are not detected in our Spitzer spectra, their absence is not strongly constraining, given that the molecular properties of their bands, such as strengths, are not well known.  It is confirmed that the CO$_2$ ice profile is an excellent ice temperature tracer. Comparison with the stretching vibration band of solid CO shows that the CO$_2$ correlates well with the pure CO to CO:H$_2$O ratio, an established ice temperature tracer. The observed differences between CO$_2$ bending mode profiles are consistent with differences in the fraction of the sight lines that have been heated above a certain threshold temperature, $T_{\\rm crit}$, to produce a double-peaked profile. More luminous YSOs have generally had a larger fraction of their ice column densities at temperatures above $T_{\\rm crit}$, but the prevalence of the characteristic double-peak shows that even the low-luminosity YSOs have thermally processed inner envelopes.  There are several possibilities for the value of $T_{\\rm crit}$. For the massive YSOs, strong annealing of methanol-rich ices to temperatures higher than 100\\,K in a laboratory setting was invoked to explain the observed abundance of pure CO$_2$ ice by \\cite{Gerakines99}. It is suggested that this corresponds to $T_{\\rm crit}=50-80\\,$K under the conditions in collapsing protostellar envelopes. These are high temperatures that puts restrictions on the density structure of the protostellar envelopes in order to reproduce the observed abundances of pure CO$_2$. Laboratory experiments measuring the kinetics of the annealing process in methanol-poor CO$_2$:H$_2$O ice mixtures will be needed to use this process in quantitative modeling of ice mantle processing.  It is therefore argued that the presence of a significant fraction of pure CO$_2$ toward so many of the low-mass YSOs in the survey indicates that another, lower temperature process may play a role. Our survey, as well as others, measure a significant fraction of the CO$_2$ ice mixed with CO, rather than with water. This mantle component will form a layer of pure CO$_2$ upon very moderate heating to $T_{\\rm crit}=$20-30\\,K as the CO desorbs, leaving the CO$_2$ behind. This process has therefore the potential to create pure CO$_2$ by {\\it distillation} rather than {\\it segregation}. Note that the new observations presented here do not rule out that the segregation process is responsible for a significant fraction of the pure CO$_2$ in the sample of {\\it massive} YSOs originally discussed in \\cite{Gerakines99}, in particular since less CO will be frozen out in the warm envelopes surrounding massive YSOs. The simple protostellar models presented here seem to indicate that both mechanisms for producing CO$_2$ in general contribute to the double-peak in low-mass young stellar objects."
    },
    "0711/0711.1519_arXiv.txt": {
        "abstract": "We present an analysis of the X-ray variability of three symbiotic X-ray binaries, GX 1+4, \\17, and \\src, using observations made with the Swift Burst Alert Telescope (BAT) and the Rossi X-ray Timing Explorer (RXTE) All-Sky Monitor (ASM).  Observations of \\src\\ with the Swift BAT show modulation at a period near 5 hours. Models to explain this modulation are discussed including the presence of an exceptionally slow X-ray pulsar in the system and accretion instabilities.  We conclude that the most likely interpretation is that \\src\\ contains one of the slowest known X-ray pulsars. Unlike \\src, neither GX 1+4 nor \\17\\ show any evidence for modulation on a timescale of hours.  An analysis of the RXTE ASM light curves of GX 1+4, \\17, and \\src\\ does not show the presence of periodic modulation in any source, although there is considerable variability on long timescales for all three sources.  There is no modulation in GX 1+4 on either the optical 1161 day orbital period or a previously reported 304 day X-ray period. For \\17\\ we do not confirm the 404 day X-ray period previously proposed for this source from a shorter duration ASM light curve. We conclude that all three sources have substantial low-frequency noise in their power spectra that may give the appearance of periodic modulation if this noise is not properly accounted for, particularly if short duration light curves are examined. ",
        "introduction": "A symbiotic binary is a system which contains a hot object accreting from an M giant companion, either from the wind of the M star or via Roche lobe overflow (e.g. Iben \\& Tutukov 1996). In most symbiotic stars the accreting object is a white dwarf, and these systems can be modest X-ray emitters (e.g. M\\\"urset et al. 1997). In far fewer systems is the accreting object thought to be a neutron star. There are presently only five sources for which there is strong evidence that they are ``symbiotic X-ray binaries'' with a neutron star component. This paper discusses three of these sources:  GX 1+4, \\src, and \\17. GX 1+4 has been known as a symbiotic source containing an X-ray pulsar for some time (e.g. Davidsen et al. 1977). \\17\\ was suspected to be a symbiotic source and this is now confirmed and \\src, which was initially thought to be a high-mass X-ray binary, has also recently been found to have an M giant counterpart (Masetti et al. 2006). In addition to these sources, Sct X-1 IGR J16194-2810, and 1RXS J180431.1-273932 have also recently been suggested to be symbiotic X-ray binaries (Kaplan et al. 2007, Masetti et al. 2007b, Nucita et al. 2007).  We present an analysis of the X-ray variability of GX 1+4, \\src, and \\17 using long term observations made with the Swift Burst Alert Telescope (BAT) and the Rossi X-ray Timing Explorer (RXTE) All Sky Monitor (ASM). Sct X-1, IGR J16194-2810, and 1RXS J180431.1-273932 are not considered here as these sources were not strongly detected by the BAT and ASM light curves are not available for IGR J16194-2810 and 1RXS J180431.1-273932. These observations enable a study of variability on timescales from hours to years. The BAT observations show the presence of a strong modulation with period of about 5 hours in \\src, first reported by Corbet et al. (2006). We consider possible mechanisms that could cause this modulation. Modulation on timescales of hours is not seen in either \\gx\\ or \\17. We find no evidence for orbital modulation in any of these three systems. ",
        "conclusions": "The three symbiotic X-ray binaries considered here show substantial X-ray variability on long time scales, dominated by the presence of slow flares. In none of the systems does the long term X-ray flux show detectable orbital modulation and there is no evidence for modulation at any of the previously reported periods. Any accretion rate variations caused by modulation in an eccentric orbit appear to be much less than that which occurs during the non-periodic flares.   In contrast, on shorter timescales two of the three sources, \\gx\\ and \\src, exhibit periodic modulation that appears to come from a neutron star rotation period. This, together with the pulsations shown by Sct X-1 and 1RXS J180431.1-273932, suggests that further observations of \\17, preferably during an outburst, might also reveal a pulsation period in this source. Additional observations of \\src\\ should help to better constrain the nature of the 5 hour modulation by showing whether this is a persistent feature of the light curve and, if so, how the period changes on long timescales in response to varying accretion rates."
    },
    "0711/0711.4420_arXiv.txt": {
        "abstract": "X-ray timing of neutron stars in low-mass X-ray binaries (LMXBs) with the {\\it Rossi X-ray Timing Explorer}\\/ has since 1996 revealed several distinct high-frequency phenomena. Among these are oscillations during thermonuclear (type-I) bursts, which (in addition to persistent X-ray pulsations) are thought to trace the neutron star spin. The recent discoveries of 294~Hz burst oscillations in IGR~J17191$-$2821, and 182~Hz pulsations in Swift~J1756.9$-$2508, brings the total number of measured LMXB spin rates to 22. An open question is why the majority of the $\\approx100$ known neutron stars in LMXBs show neither pulsations nor burst oscillations.  Recent observations suggest that persistent pulsations may be more common than previously thought, although detectable intermittently, and in some cases at very low duty cycles. For example, the 377.3~Hz pulsations in HETE~J1900.1$-$2455 were only present in the first few months of it's outburst, and have been absent since (although X-ray activity continues). Intermittent (persistent) pulsations have since been detected in a further two sources. In two of these three systems the pulsations appear to be related to the thermonuclear burst activity, but in the third (Aql~X-1) they are not. This phenomenon offers new opportunities for spin measurements in known systems.  Such measurements can constrain the poorly-known neutron star equation of state, and neutron stars in LMXBs offer observational advantages over rotation-powered pulsars which make the detection of more rapidly-spinning examples more likely. Even so, spin rates of at least 50\\% faster than the present maximum appear necessary to give constraints stringent enough to discriminate between the various models. Although the future prospects for such rapidly-spinning objects do not appear optimistic, several additional observational approaches are possible for LMXBs. The recent study of EXO~0748$-$676 is an example. ",
        "introduction": "The equation of state (EOS) of cold matter at super-nuclear densities remains one of the foremost outstanding problems for fundamental physics (e.g. \\citep{lp07}). The major uncertainty in high-density quantum chromo-dynamics theory (which has otherwise been so successful in explaining the properties and behaviour of subatomic particles) is in the regime where the density is at or above that reached in the atomic nucleus. Cold matter beyond nuclear density may be composed primarily of neutrons, as is normally thought, or it could be dominated by exotic components such as hyperons, pion or kaon condensates, or quark matter (e.g. \\citep{witten84}). Such states of matter are {\\em purely theoretical} at the present time, and their detection --- whether it be via accelerator experiments, or in the astrophysical ``laboratories'' available to astronomers --- would represent a significant step forward for modern physics.  Particle accelerators probe the conditions in matter at extreme temperatures and densities (up to a factor of ten higher than nuclear). Matter within neutron stars is also expected to reach super-nuclear densities, but at comparatively ``cool'' temperatures (no more than $10^9$~K!). Neutron stars thus play an important complementary role for studies of condensed matter, and measurements which may constrain the EOS are a high priority for observers.  Since the EOS affects the bulk properties (mass and radius) of neutron stars, simultaneous measurement of these parameters with moderate precision in an individual object would in some cases be sufficient to identify the EOS. However, such measurements have proved surprisingly elusive. The masses of neutron stars in binary (rotation-powered) pulsars can be measured in some cases to a fraction of a percent (e.g. \\citep{tc99}) although simultaneous radius measurements are generally not available. While the maximum neutron star mass also provides a constraint on the EOS, most of the measured masses cluster around $1.4\\,M_\\odot$, which is not useful in distinguishing between different models.  Measurement of the spin rate in rapidly-rotating neutron stars provides a relatively model-independent way to constrain the EOS. The maximum spin rate of a neutron star (above which it will break up due to centrifugal forces) can be expressed in terms of the neutron star mass $M$ and radius $R$, roughly independent of the EOS \\cite[]{lp07}:  \\begin{equation} \\nu_{\\rm max} = 1045(M/M_\\odot)^{1/2}(10\\,{\\rm km}/R)^{3/2}~{\\rm Hz} \\end{equation} where $M$ and $R$ are the neutron star mass and radius. Constraining the possible candidates for the neutron star EOS thus requires detection of ever-more rapidly spinning neutron stars. The fastest-spinning neutron star presently known is the radio pulsar PSR~J1748$-$2446ad, at 716~Hz \\cite[]{hessels06}. Although it's mass is unknown, assuming a value consistent with the measurements from other radio pulsars leads to an upper radius limit of 14.4~km. Without a mass measurement, this limit does not yet allow us to reject any individual EOS.  A compelling observational goal, then, is to detect ever-more rapidly spinning neutron stars.  Despite much effort, progress in this regard has been slow; the detection of PSR~J1748$-$2446ad represented the first increase in the known maximum neutron-star spin rate in 23 years. In the radio band, selection effects make more rapidly-spinning rotation powered pulsars significantly harder to detect. Although these selection effects do not affect spin measurements via X-ray observations of accretion-powered neutron stars, faster-spinning systems have not yet been convincingly detected, suggesting perhaps that some physical mechanism prevents further spin-up.  Regardless, measurement of neutron-star spins in LMXBs remains a high observational priority. In this paper I will discuss the phenomenology of the various types of high-frequency timing phenomena detected to date, and assess the prospects for future detections which may provide the first strong constraints on the neutron-star equation of state. ",
        "conclusions": "Having presented in some detail the phenomenology of measuring accreting neutron-star spins, I turn to the prospects for the future potential for stringent constraints on the neutron-star equation of state. The combined spin frequency distribution for the \\numns\\ burst oscillation sources and millisecond pulsars is approximately flat between 45--620~Hz (Fig. \\ref{spins}). The spin distribution for rotation-powered pulsars, in contrast, is subject to significant selection effects which mask the true distribution. Radio pulsars are subject to pulse smearing due to dispersion in the interstellar medium, which makes previously unknown examples harder to detect. This problem becomes worse for very fast pulsars, and the increasing effect of eclipses by the outflowing pulsar wind further decreases the sensitivity for detection. X-ray pulsations, on the other hand, are not subject to either effect, and thanks to {\\it RXTE}'s timing capability well above 1~kHz, the sensitivity to X-ray pulsars rotating at frequencies well above the current maximum is effectively flat. Thus, accretion-powered pulsars may offer the best chance to detect maximally-rotating neutron stars, and thus provide future constraints on the neutron star EOS.  \\begin{figure} \\includegraphics[width=0.45\\textwidth]{ns_spins} \\caption{ The neutron star spin frequency distribution, plotted separately for different types of systems: radio (rotation-powered) pulsars (from the ATNF Pulsar Catalogue, as of June 2007), accretion-powered millisecond pulsars, and burst oscillation sources. The overwhelming majority of rotation-powered pulsars spin slowly; the bar at 0--50~Hz is cut off by the $y$-axis range and includes 1480 sources. The distribution for the accreting sources is much flatter, and is not affected by any known selection effects which make detection of more rapidly-spinning systems less likely. \\label{spins} } \\end{figure}  In the absence of any known selection effects, the present lack of accretion-powered neutron stars spinning faster than 620~Hz strongly suggests that such systems are rare, if they exist at all. One possible explanation for this lack of faster-spinning objects is the increasing role of gravitational radiation which may prevent further spin-up (e.g. \\citep{bil98c}). Regardless of the mechanism which apparently prevents further spin-up, this paucity of more rapidly spinning systems has serious implications for the prospects for future detections and corresponding constraints on the neutron star EOS.  To explore these implications further, it is worthwhile to consider how much faster a neutron star need be discovered before significant constraints on the possible EOS are achieved. In this respect the candidate 1122~Hz burst oscillation, even though it has not been confirmed, has prompted a timely exploration of the consequences for the EOS. According to \\cite[]{lavagetto06}, this result leads for the first time to ``strong, model-independent observational constraints'' to the neutron star EOS (see also \\citep{bhz07}).  For such a rapidly spinning neutron star to be comprised of nucleonic matter would require a rather large mass of $\\gtrsim2~M_\\odot$ (perhaps providing an alternative explanation of why extremely rapidly-spinning neutron stars have been so hard to find). The rotational mass-radius limit for an 1122~Hz neutron star just intersects the M-R trajectories for several plausible equations of state at the highest possible mass (\\citep{lp07}, Figure 2). This suggests that this spin rate is a convenient empirical target for observers; neutron stars spinning slower than this rate likely cannot significantly constrain the EOS (unless other, complementary constraints are available) while neutron stars spinning at even higher rates have a good chance to constrain the EOS.  It is also worth noting that the prospect for access to an X-ray timing mission in the near future is far from guaranteed. The present \\xte\\/ cycle 12, through February 2009 at the latest, may be the last observing cycle for the instrument\\footnote{URL \\url{ http://heasarc.gsfc.nasa.gov/docs/xte/cycle12.html}}. Efforts are underway to continue the mission through 2009 and beyond, but if these efforts are unsuccessful, no alternative timing mission is currently planned for the near future by ESA or NASA. The best chance for a replacement high-sensitivity, high time resolution X-ray instrument is the Large-Area Xenon Proportional Counter (LAXPC) onboard the Indian multiwavelength satellite {\\it ASTROSAT}\\footnote{URL \\url{http://meghnad.iucaa.ernet.in/~astrosat}}, currently scheduled for launch in December 2009. The LAXPC has comparable spectral and timing resolution to the \\xte/PCA, with improved high-energy sensitivity; in addition, the satellite will also feature soft- and hard-X-ray imaging telescopes, an all-sky monitor, and a UV telescope for broadband coverage.  In the unfortunate event that \\xte\\/ ceases operation before {\\it ASTROSAT}\\/ is launched, there will be no further spin measurements for rapidly-rotating accreting neutron stars in the meantime. Even if this situation is avoided, if some phenomenon prevents the spin-up of neutron stars to spin rates much in excess of 750~Hz (as is suggested by the present distribution of measured spins), the prospects for strong constraints on the neutron star EOS by measurement of rapid spins alone appear poor. However, the prospects for constraints via multiple observational measurements remain promising. For the accretion-powered neutron stars, this approach is illustrated by the recent results from EXO~0748$-$676, which combined the spin rate with measurements of the surface gravitational redshift, the peak flux of radius-expansion thermonuclear bursts (the Eddington limit) and the (apparent) blackbody radius of the star from the X-ray flux in the burst tail \\cite[]{ozel06}. Although the spin rate in this system is the slowest measured in any LMXB at 45~Hz, and so cannot alone give any useful constraints on the EOS, the combination of other measurements allowed those authors to rule out all the ``soft'' equations of state for this system.  Although this result is not without it's own caveats (see e.g. \\citep{alford06}), many of the issues appear resolvable. The energetics of both thermonuclear bursts and carbon-burning ``superbursts'' may also allow complementary measurements of the heat flux from the neutron star crust, which also constrains the interior properties and hence the EOS (e.g. \\citep{pc05}).      \\begin{table} \\begin{tabular}{lcccccl} \\hline &      & $\\nu_{\\rm spin}$ & $P_{\\rm orb}$ & $d$ & $\\dot{M}$ & \\\\ Source & Type\\tablenote{B = burster, T = transient, L = long-duration ($>1$~yr) transient, D = dipping source. Adapted from \\cite[]{lmxb07}.} &  (Hz) & (min)        & (kpc)\\tablenote{For the sources with radius-expansion bursts, the distance is determined from the mean peak flux of those bursts \\cite[]{bcatalog}. For sources with only non-radius expansion bursts, the maximum peak flux of the bursts leads to an upper limit on the distance. For sources with no bursts, the ratio of the long-term average X-ray flux and the expected accretion rate driven by gravitational radiation leads to a lower limit on the distance \\cite[]{gal06b} } & ($10^{-11}\\ M_\\odot\\,{\\rm yr^{-1}}$)\\tablenote{The inferred mass accretion rate, based on the time-averaged broadband X-ray flux and the distance estimate, and assuming a $1.4M_\\odot$ neutron star with radius 10~km. For transients, the peak rate observed during outburst is indicated in parentheses. For the long-duration transients, the rate is measured while the source is X-ray active. For sources without bursts, the estimated rate is a lower limit based on the minimum companion mass and the orbital period (following \\citep{bc01}).} & Ref.\\\\ \\hline \\multicolumn{7}{c}{Burst oscillation sources}\\\\ \\hline EXO 0748$-$676\t& BDT\t& 45\t& 229 & 7.5 & 20--45 (120) & \\cite{villarreal04} \\\\ 4U 1916$-$05\t& BD\t& 270\t& 50 & 8.9 & 22--110 & \\cite{1916burst} \\\\ IGR J17191$-$2821\t& BT\t& 294\t& $\\ldots$ & $<11$ & $\\ldots$ & \\cite{mark07a} \\\\ 4U 1702$-$429\t& B\t& 329\t& 1320 & 5.5 & 35--80 & \\cite{mss99} \\\\ 4U 1728$-$34\t& B\t& 363\t& $\\ldots$ & 5.2 & 45--210 & \\cite{stroh96} \\\\ KS 1731$-$26\t& BL\t& 524\t& $\\ldots$ & 7.2 & 85--350 & \\cite{muno00} \\\\ A 1744$-$361\t& BT\t& 530\t& $\\ldots$ & $<9$ & ($<230$) & \\cite{bhatt06b} \\\\ MXB 1658$-$298\t& BDL\t& 567\t& 427\t& 12 & 60--200 & \\cite{wij01} \\\\ 4U 1636$-$536\t& B\t& 581\t& 228\t& 6.0 & 55--330 & \\cite{zhang97} \\\\ GRS 1741.9$-$2853\\tablenote{The source of the bursts with 589~Hz oscillations assumed by \\cite{bcatalog}. } & BT\t& 589\t& $\\ldots$ & 8 & 0.01 (1) & \\cite{stroh97,muno03b} \\\\ SAX J1750.8$-$2900 & BT\t& 601\t& $\\ldots$ & 6.8 & 20 (180) & \\cite{kaaret02} \\\\ 4U 1608$-$52\\tablenote{Hartman et al. 2008, in preparation}\t& BT\t& 620\t& 773\t& 4.1 & 100 (530) & \\\\ \\hline \\multicolumn{7}{c}{Accretion-powered millisecond pulsars} \\\\ \\hline Swift J1756.9$-$2508 & T & 182 & 54.0\t& 8? & $>0.14$ (200) & \\cite{krimm07} \\\\ XTE J0929$-$314 & T & 185 & 43.6 & $>4$\t& $>0.4$ (30) & \\cite{gal02d} \\\\ XTE J1807$-$294 & T & 191& 40.1 & $>5$ & $>0.3$ (40) & \\cite{markwardt03b} \\\\ XTE J1814$-$338\\tablenote{Also a burst oscillation source} & BT & 314 & 257 & $<8$ & 1.4--1.6 ($<60$) & \\cite{stroh03a} \\\\ SAX J1808.4$-$3658$^\\mathparagraph$ & BT\t& 401 & 120.9 & 3.4--3.6 & 0.6--1.3 (75) & \\cite{wij98b,chak98d,gal06c} \\\\ XTE J1751-305 & T & 435\t& 42.4\t& $>8$ & $>1.2$ (190) & \\cite{markwardt02} \\\\ IGR 00291+5934 & T & 599 & 147.4 & 5--6 & 1.3 (90) & \\cite{gal05a,gal06b} \\\\ \\hline \\multicolumn{7}{c}{Intermittent pulsars}\\\\ \\hline HETE J1900.1-2455 & BL & 377 & 83.3 & 5.0 & 7--70 & \\cite{kaaret05b,gal07a} \\\\ SAX J1748.9$-$2021 & BT & 442 & 520 & 8.1 & 8 (420) & \\cite{gavriil07,altamirano07} \\\\ Aql X-1$^\\mathparagraph$ & BT & 550 & 1137 & 5.0 & 105 (510) & \\cite{casella07} \\\\ \\hline \\end{tabular} \\caption{Rapidly-rotating accreting neutron stars} \\label{tab:a} \\end{table}   \\begin{theacknowledgments} Rudy Wijnands provided timely and insightful comments on this paper. I would also like to thank the conference organizers for a highly entertaining and enlightening meeting. \\end{theacknowledgments}"
    },
    "0711/0711.1802_arXiv.txt": {
        "abstract": "The cosmological constant $\\L$ affects gravitational lensing phenomena. The contribution of $\\L$ to the observable angular positions of multiple images and to their amplification and time delay is here computed through a study in the weak deflection limit of the equations of motion in the Schwarzschild-de Sitter metric. Due to $\\L$ the unresolved images are slightly demagnified, the radius of the Einstein ring decreases and the time delay increases. The effect is however negligible for near lenses. In the case of null cosmological constant, we provide some updated results on lensing by a Schwarzschild black hole. ",
        "introduction": "The interpretation of the cosmological constant $\\L$ is a very fascinating and traditional topic in theoretical physics. On the observational side, large scale structure observations have made a strong case for $\\L$ as a possible choice for dark energy. In fact, a very small value of $\\L  \\sim 10^{-52}\\mathrm{m}^{-2}$, together with dark matter, can provide a suitable framework for observational cosmology \\cite{pe+ra03}.  Since the cosmological constant should take part in all kinds of gravitational phenomena, investigations have been performed on very different scale-lengths. Despite no convincing method for constraining $\\Lambda$ in an Earth's laboratory has been proposed \\cite{je+st05}, local astronomical phenomena seem to be more promising. The cosmological constant can influence the motion of massive bodies  \\cite{isl83,wri98,ker+al03}. Perihelion precession of solar system planets together with other solar and stellar tests has been considered to put an upper bound of $\\L \\ls 10^{-42}\\mathrm{m}^{-2}$ \\cite[and references therein]{wri98,je+se05,ior06,se+je06}. The cosmological constant  also affects the gravitational equilibrium of large astrophysical structures \\cite{ba+no05,bal+al06,bam07} and produces a lower velocity dispersion around the Hubble flow on the scale of the Local Volume \\cite{tee+al05}.  Recently, \\citet{ri+is07} discussed how the cosmological constant takes part in gravitational lensing. Taking into account $\\L$ through the Schwarzschild-de Sitter (SdS) metric, they showed that even if the exact differential equation for a light path in the coordinate space can be written in a form that does not involve $\\L$ \\cite{isl83}, the cosmological constant contributes to the bending of light through the metric itself, which determines the actual observations that can be made on the orbit equation. In fact, one must consider not only the null geodesic equation but also the process of measurement \\cite{bak+al07,lak07}.  Following this correction of the long-standing misconception that $\\L$ does not affect the observed deflection angle, in this paper I further investigate the effect of the cosmological constant in gravitational lensing observations in near systems. The weak deflection limit considered throughout the paper allows to have a clear insight on the effect of $\\L$ but it is to be remarked that a gravitational lens equation without approximations can be written in generic spherically symmetric and static spacetimes \\cite{per04}. In this paper, the lens equation is derived from the lightlike null geodesics of the SdS metric. Results concerning observable quantities are expressed in terms of the invariants of the light ray, avoiding ambiguities connected to coordinate-dependent quantities \\cite{ke+pe05,bo+wi03}.  The paper is organized as follows. In Sec.~\\ref{sec:geod}, the null orbits are solved in the weak deflection limit. In Sec.~\\ref{sec:lens} the lens equation is first written in terms of the observed image position angle and then solved with a perturbation method. Image amplification and time delay are discussed in Sec.~\\ref{sec:magn} and~\\ref{sec:time}, respectively. Some quantitative estimates of the effect of $\\L$ are illustrated in Sec.~\\ref{sec:loca} whereas Sec.~\\ref{sec:conc} is devoted to some final considerations. ",
        "conclusions": "\\label{sec:conc}  The stagnant theoretical affair between the cosmological constant and the bending of light rays took an hit recently when \\citet{ri+is07} pointed out how the study of the orbit equation in the coordinate space is not enough to describe the observations of lensing phenomena. This noteworthy criticism has then stimulated some new interest on the subject \\citep{lak07}. In this paper, I have performed an analysis of lensing phenomena in the framework of the SdS metric, which allows a full treatment for systems much smaller than the Hubble radius. I have based my results on a perturbation expansion of the equation of motions, from which I have derived a lens equation accounting for $\\L$. The analysis has also showed that the usual argument against $\\L$, i.e. that the cosmological constant is dropped out from the exact differential equation for a light path, does not apply to the time delay. It is also to be remarked that the degeneracy between the orbital differential equation in the Schwarzschild metric and that in the SdS spacetime breaks down in presence of a non null angular momentum of the lens.  The argument that $\\L$ affects lensing through the metric itself at the observer position is not restricted to the weak deflection limit and applies as well to light rays passing very near to the photon horizon of a black hole. Since SdS null geodesics are formally identical to the Schwarzschild case, the calculation of the deflection angle should be performed as usual but the relation between the constant of motion and the observed angle should be revised. However, since the angular separations of the relativistic images are very small with respect to the splitting of the primary images, in the strong deflection limit it is customary to neglect higher order corrections.  Even though important on a theoretical point of view, the effect of $\\L$ on near lenses, such as the Sun, the supermassive black hole in the Galactic center or compact objects in the halo of near galaxies, is quantitatively very small. \\citet{ish+al07} tried to extend the result obtained in the framework of the SdS metric to a cosmological scenario where the distances between lens, source and observer are comparable with the Hubble radius. Some caution should be however used in such an extrapolation. The cosmological lens equation is usually derived combining local results on the light deflection, which are based on an asymptotically flat metric, with considerations on global geometry and  angular diameter distances, which are on turn based on the global Friedmann-Lema\\^{\\i}tre-Robertson-Walker spacetime  in which the system is embedded \\cite{sch+al92,sei+al94}. As shown in \\citep{ri+is07} and in the present analysis, both based on the SdS metric, the main contribution of $\\L$ to lensing observables comes through the value of the metric at the observer position, which lies in a region of spacetime curved by the cosmological constant. In the classical reasoning at the basis of the cosmological lens equation, local effects are related to a small region in the neighborhood of the lens whereas global effects are connected to the large regions of spacetime between source, lens and observer. Then, the main contribution of $\\L$ to gravitational lensing should be seen as global in the sense that it is connected to the observer radial distance. This view is also supported by the fact that, as shown in Sec.~\\ref{sec:lens}, the effect of $\\L$ on the lensing equation can be already considered through the angular diameter distances of the background smooth spacetime, which express global relations.  As far as distances are small with respect to the de Sitter horizon, we can safely apply the expressions obtained in the present analysis and neglect higher order correction connected to the coupling between $\\L$ and the black hole mass, but if distances are comparable to the Hubble length then the results should be likely revisited. This will the subject of a future analysis.  A further consideration is that if we are assuming that a constant energy background as the one provided by $\\L$ affects lensing, then every other background, such as that provided by dark matter, should have a similar effects. The McVittie metric, which accounts for the presence of a generic cosmological fluid around the central mass and the related expansion of the spacetime, should be used instead of the more specific SdS spacetime and the effect of all the contributions to the cosmological energy budget should be considered even on a small scale \\cite{se+je07}.  Then, even though $\\L$ changes in some ways the expression for the bending angle, the dark matter, whose uniform distribution counteracts the cosmological constant, should nearly compensate in the opposite direction.  In the case of $\\L =0$ the results in this paper updates previous studies for lensing in the Schwarzschild spacetime that were based either on an approximate lens equation, differently from the present analysis which is based on a perturbation analysis of an exact lens equation, or neglected the effect of the metric in the relation between the impact parameter and the observed image position angle. The present study is also relevant to lensing in extended theories of gravity \\citep{rug07}, in which the SdS metric provides an exact solution suitable to evaluate the effects of the non linearity of the gravity Lagrangian."
    },
    "0711/0711.3346_arXiv.txt": {
        "abstract": "A recent article from the Pierre Auger Collaboration links the direction of charged cosmic rays to possible extragalactic sites of emission. The correlation of the direction of such particles with the direction of the emitter allows constraining the value of large-scale magnetic fields B. Assuming for ${\\rm B}$ a coherence length $\\lambda$ in the range  between 1 Mpc and 10 Mpc, we find values of ${\\rm B}$ between 0.3 and 0.9 nG. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.1043.txt": {
        "abstract": "{}{To study cosmological models in scalar tensor theories of gravity with power law potentials as models of an accelerating universe.}{We consider cosmological models in scalar tensor theories of gravity that describe an accelerating universe, and we study a family of inverse power law potentials, for which exact solutions of the Einstein equations are known. We also compare theoretical predictions of our models with observations. For this we use the following data: the publicly available catalogs of type Ia supernovae and high redshift Gamma Ray Bursts, the parameters of large scale structure determined by the 2-degree Field Galaxy Redshift Survey (2dFGRS), and measurements of cosmological distances based on the Sunyaev-Zel'dovich effect, among others.}{ We present a class of cosmological models which describe evolution of homogeneous and isotropic universe filled in with dust like matter and a scalar field that is non minimally coupled to gravity. We show that this class of models depends on three parameters: $V_0$ - the amplitude of the scalar field potential, $\\widehat{H}_0$ - the present value of the Hubble constant, and a real parameter $s$ which determines the overall evolution of the universe. It turns out that these models have a very interesting feature of producing in a natural way an epoch of accelerated expansion. We fix the values of these parameters by comparing predictions of our model with observational data. It turns out that our model is compatible with the presently available observational data.}{} ",
        "introduction": "Recent observations of the type Ia supernovae and CMB anisotropy indicate that the total matter-energy density of the universe is now dominated by some kind of dark energy causing an accelerated expansion of the Universe (\\cite{perl,rie+al98, Riess04,spergel}). The origin and nature of this dark energy remains unknown~(\\cite{zel67,weinberg2}).\\\\ Prompted by this discovery recently a new class of cosmological models has been proposed. In these models the standard cosmological constant $\\Lambda$-term is replaced by a dynamical, time-dependent component - quintessence or dark energy - that is added to baryons, cold dark matter (CDM), photons and neutrinos. The equation of state of the dark energy is assumed to be of a hydrodynamical type $w_{\\phi} \\equiv \\rho_{\\phi} /p_{\\phi}$, where $\\rho_{\\phi}$ and $p_{\\phi}$ are, respectively, the energy density and pressure, and $-1 \\leq w_{\\phi}<0$, what implies a negative contribution to the total pressure of the cosmic fluid. When $w_{\\phi} =-1$, we recover the standard cosmological constant term. One of the possible physical realizations of quintessence is a cosmic scalar field (\\cite{cal+al98}), which induces dynamically a repulsive gravitational force, that is responsible for the observed now accelerated expansion of the universe. \\\\ The existence of dark energy that now dominates the overall energy density in the universe is posing several theoretical problems. Firstly, it is natural to ask, why do we observe the universe at exactly the time when the dark energy dominates over the matter (\\textit{cosmic coincidence} problem). The second issue, a \\textit{fine tuning problem}, arises from the fact that if the dark energy is constant, like in the standard cosmological constant scenario, then at the beginning of the radiation era its energy density should have been vanishingly small in comparison with the radiation and matter component. This poses a problem, since in order to explain the inflationary behaviour of the early universe and the late time dark energy dominated regime, the dark energy should evolve and cannot simply be a \\textit{constant}. All these circumstances stimulated a renewed interest in the generalized gravity theories, and prompted consideration of a variable $\\Lambda $ term in more general classes of theories, such as the scalar tensor theories of gravity.\\\\  In our earlier paper (\\cite{nmc}) we have analyzed extended quintessence models, for which exact solutions of the Einstein equations are known, and discussed how in such models it is possible to treat the fine tuning problem in an alternative way. We applied our consideration to a special model, based on one of the most commonly used quintessence potentials $V(\\phi)=\\lambda \\phi^{4}$, corresponding to the coupling $F(\\phi)=(3/32)\\phi^2$ (so called induced gravity). We showed that this model corresponds to a special, and physically significant case, which emerged by requiring the existence of a Noether symmetry in the {\\em pointlike} Lagrangian. In this paper we analyze a new and wider class of theories derived from the Noether symmetry requirement. One of the main advantages of such models is the fact that they exhibit power law couplings and potentials, and admit a {\\it tracker behaviour}. In some sense we complete and generalize the analysis initiated in (\\cite{alma, rug1}), where the attention was focused on the mechanism of obtaining an effective cosmological constant through the {\\it cosmological no hair theorem}, and the analysis of the solution was restricted to the asymptotical $t\\rightarrow \\infty$ regime only. Extending our analysis to the whole time evolution, we are able not only to clarify the properties of such solutions, but also to compare predictions of such models with observations. We concentrate on the following data: the publicly available data on type Ia supernovae and Gamma Ray Bursts, the parameters of large scale structure determined by the 2-degree Field Galaxy Redshift Survey (2dFGRS), and  the measurements of cosmological distance with the Sunyaev-Zel'dovich effect. ",
        "conclusions": "In this paper we have extended the analysis that we performed in Paper I (where  we have analyzed a special extended quintessence model, based on one of the most commonly used quintessence potentials $V(\\phi)=\\lambda \\phi^{4}$, corresponding to the coupling $F(\\phi)=(3/32)\\phi^2$ ), considering a new and wider class of theories for which exact solutions of the Einstein equations are known, and discussed how in such models it is possible to treat the fine tuning problem in an alternative way. We have shown that in the family of such models selected by requiring that their corresponding point like Lagrangian admits a Noether symmetry an epoch of accelerated expansion appears in a natural way. In the non minimally coupled scalar tensor theory of gravity it is possible to perform an appropriate conformal transformation and to move from the Jordan picture to the standard Einstein one but then matter becomes coupled to the scalar field. We have explored both descriptions and also considered the neutrino mass varying model as a possible example of non minimally coupled scalar tensor theory.   It turns out that the imposed requirement of existence of a Noether symmetry is quite restrictive and we obtained a family of models that is fully specified by 3 parameters: a parameter $s$ that determines the strength of the non minimal coupling and the potential of the scalar field, $H_0$ the Hubble constant, and a parameter $V_0$ that determines the scale of the potential. To determine the values of these parameters we compared predictions of our model with several independent observational data. The results of this parameter determination procedure are presented in Table 2. We see that with our average value of $s= - 1.46$ the scale factor, for small $t$, is changing as $a\\sim t^{0.66}\\sim t^{2/3}$, and for large $t$, as $a\\sim t^{1.25}\\sim t^{5/4}$ while $\\phi\\sim t^{0.0032}$ and $\\phi \\sim t^{0.083}$ in corresponding asymptotic regimes. The potential $V$ is decaying to zero for large $t$, after reaching a maximum value, (see Fig. (\\ref{Vvariation1})). Similarly, the effective gravitational coupling $G_{eff}=-{1\\over {2F}}$ is decreasing for  large $t$, until it becomes zero for $t\\rightarrow \\infty $ ( we have a sort of \\textit{ asymptotic freedom } at $ t\\rightarrow \\infty $). It turns out that in our model the observational constraints on the variation of the effective gravitational constant are respected. Actually a new analysis of the Big Bang Nucleosynthesis~(\\cite{bbnG}) restricts the variations of $G$ to \\begin{equation} \\Delta G_{BBN}/G_0 = 0.01^{+0.20}_{-0.16} \\end{equation} at $68\\%$ confidence level, where  $\\Delta G_{BBN}/G_0 ={G_{BBN}-G_{0}\\over G_0} $, and $G_0\\equiv G_{Newton}$. The combined analysis of the new $^4$He and WMAP data implies that \\begin{equation} -0.10 < \\Delta G_{BBN}/G_0 < 0.13\\,. \\end{equation} A recent analysis of the secular variation of the period of nonradial pulsations of the white dwarf G117-B15A shows~(\\cite{gwd}) that $0<\\dot G/G < 4.0\\times 10^{-11}\\,{\\rm yr}^{-1}$ at 2$\\sigma$, which is of the same order of magnitude as previous independent bounds (see also \\cite{gwd2}). With our unit of time this becomes $0<\\,\\dot G/G < 4.0 \\,\\tau \\times 10^{-2}$, where  $\\tau = 13.73^{+ 0.16}_{-0.15}$  as given by the WMAP team (\\cite{spergel}). \\begin{figure} \\begin{minipage}{13 cm} \\centering{ \\includegraphics[width=6 cm, height=5 cm]{Vvariation3.eps}} \\caption{\\small The potential $V$ as a function of the redshift $z$.} \\label{Vvariation1} \\end{minipage} \\end{figure}    Initially, for small $t$, the matter energy density is larger than the energy density of the scalar field.  In Table (\\ref{tab1}) we present results of our analysis, they show that predictions of our model are fully compatible with the recent observational data. Comparing results of this paper with our previous analysis of minimally coupled scalar field models (see Paper I) we conclude that the present day observational data connected with the post recombination evolution of the universe can be fitted by several different models of quintessence. More data on high redshift supernovae of type Ia and GRBs are needed as well as more information on the early phase of structure formation in order to place stronger restrictions on the allowed type of dark energy. \\begin{table*}[ht]\\label{tab1} \\begin{center} \\caption{The basic cosmological parameters derived from our model are compared with observational data. } \\begin{tabular*}{12.15 cm}{|c|c|c|c|c|} % after \\\\: \\hline or \\cline{col1-col2} \\cline{col3-col4} ... \\cline{1-5} \\bf{dataset} & $\\mathbf{s}$ & $\\mathbf{{\\widehat H}_0}$  & $\\mathbf{\\Omega_{\\Lambda_{eff}}}$ & $w_{\\phi}$ \\\\ \\hline high redshift SNIa  &  $ -1.46^{+0.04}_{-0.03}$ & $ 0.96^{+0.05}_{-0.05}$& $0.69^{+0.04}_{-0.06}$ & $-1.01\\pm 0.02$  \\\\ \\hline dimensionless coordinates test &  $-1.49^{+ 0.02}_{-0.04}$ & $ 0.97^{+0.04}_{-0.03}$ & $ 0.73\\pm 0.08 $ &$-0.98 \\pm 0.03$ \\\\ \\hline relative galaxy ages &  $-1.49^{+0.03}_{-0.09}$ & $ 1.01^{+0.01}_{-0.03}$ & $ 0.80 \\pm 0.07$ & $-0.99 \\pm 0.04$\\\\ \\hline   sze data &     $-1.49^{+0.03}_{-0.05}$&$ 1^{+0.01}_{-0.03}$ &$ 0.70\\pm  0.05$&$ -1 \\pm 0.04$  \\\\ \\hline  GRBs dataset &$-1.43^{+0.02}_{-0.04}$ & $0.98^{+0.03}_{-0.03}$ & $ 0.76 \\pm 0.06 $  &$-0.9\\pm 0.04 $  \\\\ \\hline  fraction of gas in clusters & $-1.39^{+0.04}_{-0.01}$&$ 0.98\\pm 0.04$ &$ 0.77^{+0.05}_{-0.03}$&$-0.87 \\pm 0.05$  \\\\ \\hline galaxies peculiar velocity &   $-1.4\\pm 0.1$&$0.94^{+0.05}_{-0.02}$ &$ 0.65\\pm 0.08 $&$-0.85 \\pm 0.04$  \\\\ \\hline average &   $\\mathbf{-1.46 \\pm 0.09}$&$\\mathbf{ 0.97 \\pm 0.02}$ &$ \\mathbf{0.74\\pm 0.03}$&$\\mathbf{-1.01\\pm 0.02}$ \\\\ \\hline \\end{tabular*} \\end{center} \\end{table*}"
    },
    "0711/0711.4572_arXiv.txt": {
        "abstract": "Homan \\& Lister (2006) have recently published circular-polarizatio (CP) detections for 34 objects in the MOJAVE sample -- a set of bright, compact AGN being monitored by the Very Long Baseline Array at 15~GHz. We report the detection of 15-GHz parsec-scale CP in two more AGN (3C345 and 2231+114), and confirm the MOJAVE detection of CP in 1633+382. It is generally believed that the most likely mechanism for the generation of this CP is Faraday conversion of linear polarization to CP. A helical jet magnetic-field ({\\bf B}-field) geometry can facilitate this process -- linearly polarized emission from the far side of the jet is converted to CP as it passes through the magnetised plasma at the front side of the jet on its way toward the observer. In this case, the sign of the generated CP is essentially determined by the pitch angle and helicity of the helical {\\bf B} field. We have determined the pitch-angle regimes and helicities of the helical jet {\\bf B} fields in 8 AGN for which parsec-scale CP has been detected, and used them to predict the expected CP signs for these AGN if the CP is generated via conversion in these helical fields. We have obtained the intriguing result that our predictions agree with the observed signs in all eight cases, provided that the longitudinal {\\bf B}-field components in the jets correspond to South magnetic poles. This clearly non-random pattern demonstrates that the observed CP in AGN is directly associated with the presence of helical jet {\\bf B} fields. These results suggest that helical {\\bf B} fields are ubiquitous in AGN jets. ",
        "introduction": "The radio emission of core-dominated, radio-loud Active Galactic Nuclei (AGN) is synchrotron radiation generated in the relativistic jets that emerge from the nucleus of the galaxy, presumably along the rotational axis of a central supermassive black hole. Synchrotron radiation can be highly linearly polarized, to $\\simeq 75\\%$ in the case of a uniform magnetic ({\\bf B}) field (Pacholczyk 1970), and linear polarization (LP) observations can yield unique information about the orientation and degree of order of the {\\bf B} field in the synchrotron source, as well as the distribution of thermal electrons and the {\\bf B}-field geometry in the immediate vicinity of the AGN (e.g., via Faraday rotation of the plane of polarization).  After a number of early circular-polarization (CP) studies for various samples of AGN in the 1970's and 1980's (Gilbert \\& Conway 1970; Conway et al. 1971; Roberts et al. 1975; Weiler \\& Wilson 1977; Weiler \\& de Pater 1983; Komesaroff et al. 1984), little observational work was done in this area until the late 1990's, when the Australia Telescope Compact Array was optimised for CP measurements by Rayner et al. (2000), and techniques for deriving CP information on parsec scales were developed by Homan and his collaborators (Homan \\& Wardle 1999; Homan, Attridge \\& Wardle 2001) using the NRAO\\footnote{The National Radio Astronomy Observatory of the USA is operated by Associated Universities, Inc., under co-operative agreement with the US NSF.} Very Long Baseline Array (VLBA). The high resolution provided by these latter observations makes it possible to more precisely localize the regions of CP within the source, and thereby more accurately estimate the degree of CP in these regions. Certain tendencies became clear from the very first VLBA CP measurements: the CP was nearly always coincident with the VLBI core, with the degrees of CP typically being tenths of a percent. Comparisons with integrated data obtained earlier indicated that the sign of the CP for a given source usually remained constant on time scales of decades. In addition, it became clear that CP could occasionally be detected in compact components in the inner VLBI jets.  CP measurements of AGN have recently made another major step forward with the publication of first-epoch CP results for the 133 sources in the MOJAVE sample of radio-bright AGN, which are regularly monitored by the VLBA at 15~GHz (Homan \\& Lister 2006).  Of these sources, 34 displayed CP at the positions of their total-intensity peaks (nearly always the VLBI core) at a level of $2\\sigma$ or greater, with the typical degrees of CP being a few tenths of a percent. Comparison with 20 different optical, radio and intrinsic parameters of the AGN yielded virtually no evidence for any correlation between the degree of CP and these other parameters. Another exciting result of the first-epoch MOJAVE CP measurements was that 5 of the sample sources clearly displayed appreciable CP in their VLBI {\\em jets}, well resolved from the core, indicating that a CP-generation mechanism capable of operating effectively in optically thin regions was involved.  We consider here the hypothesis that the parsec-scale CP is generated via Faraday conversion in helical jet {\\bf B} fields in these AGN, focusing on the {\\em sign} of the observed CP, which has not been investigated in previous studies. We first examine the sign of CP generated by Faraday conversion (Section~2), then turn to the special case of conversion in a helical magnetic-field geometry (Section~3). In Section~3, we also describe our approach for using multi-frequency VLBI polarization data to determine the expected CP sign for particular AGN in a simple helical jet B-field model.  Sections~4 and 5 discuss the data used for our analysis and our results: new detections of CP and transverse jet Faraday-rotation gradients in several AGN, and our remarkable result that the CP signs expected in our simple model agree with the observed signs in all eight AGN for which the comparison was possible, provided their longitudinal {\\bf B}-field components correspond to South poles. The robustness and implications of this result are discussed in Section~6, and our conclusions summarized in Section~7. ",
        "conclusions": "Despite extensive analyses of the degree of CP in previous studies, we are the first to consider the {\\em sign} of the parsec-scale CP of AGN. The new transverse Faraday-rotation gradient and parsec-scale CP detections reported here have enabled us to compile a list of 8 AGN with both CP detections and direct evidence for helical jet {\\bf B} fields. The perfect agreement between the observed CP signs and the CP signs expected if the CP is generated by Faraday conversion in helical jet fields whose poloidal component is oriented opposite to the jet conclusively demonstrates that the observed CP is intrinsically related to the helical jet {\\bf B} fields in these objects. Although the possibility of conversion-generated CP has been considered in a number of theoretical studies, this represents the first direct \\emph{observational} evidence pointing toward this specific CP-generation mechanism in the more than three decades since the earliest studies of CP in AGN.  This result has far-reaching implications for our understanding of the processes occurring in AGN jets, since it suggests that helical jet {\\bf B} fields are common in AGN, possibly even ubiquitous. Indeed, such helical {\\bf B} fields can be produced very naturally through the ``winding up'' of a seed field by rotation of the central black hole and accretion disc, combined with the jet outflow (e.g. Meier et al. 2001; Kato et al. 2004; McKinney 2006). The generation of a helical field requires the presence of a significant poloidal field, which can come about due to the jet outflow. The presence of helical jet {\\bf B} fields provides a means to collimate the jets, and demands that the jets carry current (e.g. Lovelace et al. 2002, Tsinganos \\& Bogovalov 2002). Thus, it is clear that we must view AGN jets as fundamentally electromagnetic structures, although it is less clear from the available observations whether or not the jets are actually dynamically dominated by the associated electromagnetic forces.  Our results also appear to provide evidence that the initial ``seed'' fields of the central black holes are typically quadrupolar rather than dipolar, with a predominance of quadrupolar fields with two South poles. Numerical simulations of the joint action of rotation and outflow on an initial quadrupolar field would be helpful in determining if there is a theoretical basis for this possible asymmetry between these two field configurations. An interesting observational distinction between dipolar and quadrupolar field configurations for the central black holes was recently pointed out by Blandford (2007): in the case of a helical jet {\\bf B} field generated from a dipolar seed field, the transverse Faraday-rotation gradients for the approaching jet and receding counterjet should have \\emph{opposite} directions, while these two gradients will be in the \\emph{same} direction if the seed field that is ``wound up'' by the rotation of the system is quadrupolar. The observational challenge is therefore to identify AGN in which both total intensity and polarization are detected in both the jet and counterjet  on scales on which the transverse Faraday-rotation gradients due to the helical jet {\\bf B} fields are reliably detected, and to carry out sensitive Faraday-rotation measurements for such objects (the apparent absence of clear RM gradients in FRI jets on kiloparsec scales (e.g. Laing \\& Bridle 2007) may indicate that the systematic transverse RM gradients are disrupted by chaotic or turbulent motions on these scales).  We plan to analyze multi-wavelength VLBA polarization observations for several additional AGN with detected parsec-scale CP in order to search for transverse Faraday-rotation gradients in these objects, with the aim of further investigating the relationship between the linear-polarization structures, {\\bf B}-field helicities and observed CP in AGN.  We also plan to consider the possible influence of modest amounts of Faraday rotation on the observed CP in a future paper."
    },
    "0711/0711.1950_arXiv.txt": {
        "abstract": "The contribution of cosmological perturbations to the time drift of the cosmological redshift is derived. It is shown that the dominant correction arises from the local acceleration of both the emitter and the observer. The amplitude of this effect is estimated to be of the order of 1\\% of the drift signal at $z=2-4$, but can easily be lowered down to 0.1\\% by using many absorption lines and quasars. ",
        "introduction": "The increasing observational evidences that the expansion of the universe is accelerating (see Ref.~\\cite{refsn,refdata} and reference therein) has stimulated a rising interest to the reconstruction of its expansion history. An important outcome of these theoretical studies is to clarify the sensitivity of observational tests to dark energy properties and to assess how each could be corrupted by extra-noise from other cosmological effects. Comprehensive investigations of these nuisances have been carried out on ``standard\" tests, like SNIa, weak lensing, BAO, CMB, ISW or clusters of galaxies.  In contrast, very few focussed on the time drift effect that changes the observed redshift  of an object as function of time. The recent claim that it  may drive the conceptual design of instrument for next generation giant telescopes raised the need that similar attention should be paid to the theoretical ground of this novel technique. Interestingly, such an observation may lead to a better understanding of the physical origin of the recent acceleration~\\cite{refde,refde2} and to a determination of the dark energy equation of states~\\cite{Lake07} as well as constraints on dark energy models~\\cite{demodels} or tests of the variation of fundamental constants~\\cite{constant,Molaro}.  As first pointed out by Sandage~\\cite{Sandage}, in a homogeneous and isotropic spacetime, the time drift of the observed redshift is directly related to the Hubble function by \\begin{equation}\\label{SLformula} \\dot z = (1+z)H_0-H(z)\\equiv \\dot{\\bar z}(\\eta_0,z)\\ . \\end{equation} Given the  most likely ranges of cosmological parameter values derived from observations, in a $\\Lambda$CDM model the typical amplitude of the redshift drift is of order $\\delta z\\sim -4\\times10^{-10}$ on a  time scale of $\\delta t= 10$~yr, for  a source at redshift $z=4$. This corresponds to a spectroscopic tiny velocity shift, $\\delta v \\equiv c\\delta z/(1+z)$, of $\\delta v\\sim 2.5\\ {\\rm cm/s}$. Fig.~\\ref{AmpError} (left panel) shows the time drift as function of redshift for the standard $\\Lambda$CDM model and a dark energy models with an equation of state changed by only 10\\% ($w=-0.9$), all other parameters being kept constant. Both curves have similar shape but the difference the drifts between a standard $\\Lambda$CDM model and cosmological models tends to zero at hight redshift. Fig.~\\ref{AmpError} (right panel) depicts this difference for two models with either $w = -0.95$ or $w = -0.98$.  The feasibility of this measurement is most challenging and impossible with present-day astronomical facilities. However, it was  recently revisited~\\cite{Loeb} in the context of the new generation of Extremely Large Telescopes\\footnote{ {\\tt http://www.eso.org/projects/e-elt/Publications/ ELT\\_SWG\\_apr30\\_1.pdf}} (ELT), arguing that with such outstanding collecting areas one could measure velocity shifts of order $\\delta v\\sim 1-10\\ {\\rm cm/s}$ over a 10 year period from the observation of the Lyman-alpha forest on QSO absorption spectra. In particular, it is one of the main science driver to design the COsmic Dynamics EXperiment (CODEX) spectrograph~\\cite{Pasquini1,Pasquini2} for the future European ELT (E-ELT).  The performances of CODEX and its capability to measure a time drift of very distant objects were estimated using Monte-Carlo simulations of quasar absorption spectra. The expected velocity accuracy of this experiment can be written as follows (see Ref.~\\cite{Pasquini1}) $$ \\sigma_v=1.4\\left(\\frac{S/N}{2350}\\right)^{-1} \\left(\\frac{N_{\\rm QSO}}{30}\\right)^{-1/2} \\left(\\frac{1+z}{5}\\right)^{-1.8}\\,{\\rm cm/s}\\ , $$ provided the absorption lines are resolved. $S/N$ denotes the signal-to-noise ratio, for a pixel scale of $0.0125\\, {\\rm\\AA}$ and $N_{\\rm QSO}$ is the number of quasars. Thus, spectroscopic measurements of about 40 quasars with $S/N\\sim2000$ ten years apart can reach a $1.5\\,{\\rm cm/s}$ accuracy.  This is within the reach of a CODEX instrument mounted on a $60-80$ meter ELT by observing a 16.5$^{\\rm th}$ magnitude QSO during  2000 hrs~\\cite{Pasquini1}.  Many systematic effects that may spoil the time drift signal, such as Earth rotation, proper motion of the source, relativistic corrections etc., are discussed in Ref.~\\cite{Pasquini1}. The acceleration of the Sun in the Galaxy seems more a serious problem because its amplitude  may be of the same order than the cosmic signal. However, it has not been measured yet, so its nuisance is still unknown. On the other hand, subtle contaminations like accelerations produced by large scale structures have never been estimated in the error budget. The purpose of this work is to address this issue and to estimate whether it may hamper the cosmological interpretation of the time drift. ",
        "conclusions": "In order to measure the cosmological time drift of the redshift, many systematic effects will have to be understood. Besides the systematic errors of astrophysical origin that may affect the observation of the Lyman-alpha forest, large scale structures will induce a dispersion of $\\dot z$. This work addresses this issue.  First, we have derived the expression of the time drift of cosmological redshift at first order in the perturbation. This was then used to estimate its variance and then to demonstrate it depends on two  main effects, the accelerations and the local gravitational potential at both the source and observer positions.  The contributions at the observer position have not been discussed further. High precision astrometric observations with GAIA will soon provide exquisite  knowledge of the motion of the Earth in the Milky Way. It  will pinpoint all local acceleration terms with enough accuracy to remove this contribution easily~\\cite{gaiaref}.  In contrast, the contributions at the source position are much more difficult to subtract. In the linear regime, we have shown that the gravitational potential contribution is negligible while the acceleration of the source is typically of the order of 1\\% at $z=2-4$. We argue a dominant contribution to this term is the acceleration of galaxy clusters near the source.  In order to understand whether the amplitude of this variance can be reduced,  we have estimated the effect of averaging the signal over several absorption lines. One can either profit from the total spectral range covered by the spectrograph  to measure the drift from all lines detected along a line of sight, provided correlated acceleration contributions are taken into account, or use the mean  drift over many randomly selected lines of sight. The first option reduces the variance by the square root of the number of uncorrelated clouds systems along a line of sight, the second by the square root of the number of independent lines of sight. For an instrument having the current specifications of the CODEX spectrograph, it is then easy to drop the contribution of large-scale structures to the total variance of the time drift down to a 0.1\\% level.   \\noindent{\\bf Acknowledgements:} We thank Jochen Liske and Luca Pasquini for providing Ref.~\\cite{Pasquini1} before publication and for their comments, Patrick Petitjean for useful discussions, and Eric Linder his useful comments. We also thank St\\'ephane Charlot and Jean-Gabriel Cuby for organizing the Programme National de Cosmologie discussion on ELT which triggered these questions."
    },
    "0711/0711.1357_arXiv.txt": {
        "abstract": "The major theoretical limitation for extracting cosmological parameters from the CMB sky lies in the precision with which we can calculate the cosmological recombination process. Uncertainty in the details of hydrogen and helium recombination could effectively increase the errors or bias the values of the cosmological parameters derived from the {\\sl Planck} satellite, for example.  Here we modify the cosmological recombination code R{\\sc ecfast} by introducing one more parameter to reproduce the recent numerical results for the speed-up of the helium recombination.  Together with the existing hydrogen fudge factor, we vary these two parameters to account for the remaining dominant uncertainties in cosmological recombination. By using the C{\\sc osmo}MC code with {\\sl Planck} forecast data, we find that we need to determine the parameters to better than ten per cent for He\\,{\\sc i} and one per cent for H, in order to obtain negligible effects on the cosmological parameters.  For helium recombination, if the existing studies have calculated the ionization fraction to the 0.1 per cent level by properly including the relevant physical processes, then we already have numerical calculations which are accurate enough for {\\sl Planck}.  For hydrogen, setting the fudge factor to speed up low redshift recombination by 14 per cent appears to be sufficient for {\\sl Planck}.  However, more work still needs to be done to carry out comprehensive numerical calculations of all the relevant effects for hydrogen, as well as to check for effects which couple hydrogen and helium recombinaton through the radiation field. ",
        "introduction": "{\\sl Planck}~\\citep{Planck:2006}, the third generation Cosmic Microwave Background~(CMB) satellite will be launched in 2008; it will measure the CMB temperature and polarization anisotropies $C_\\ell$ at multipoles $\\ell=1$ to $\\simeq 2500$ at much higher precision than has been possible before.  In order to interpret these high fidelity experimental data, we need to have a correspondingly high precision theory. Understanding precise details of the recombination history is the major limiting factor in calculating the $C_\\ell$ to better than 1 per cent accuracy.  An assessment of the level of this uncertainty, in the context of the expected {\\sl Planck} capabilities, will be the subject of this paper.  The general physical picture of cosmological recombination was first given by \\citet{Peebles:1968} and \\citet{Zeldovich:1968}. They adopted a simple three-level atom model for hydrogen~(H), with a consideration of the Ly\\,$\\alpha$ and lowest order 2s--1s two-photon rates.  Thirty years later, \\citet{Seager:1999km} performed a detailed calculation by following all the resonant transitions and the lowest two-photon transition in multi-level atoms for both hydrogen and helium in a blackbody radiation background. \\citet{Lewis:2006ym} first discussed how the uncertainties in recombination might bias the constraints on cosmological parameters coming from {\\sl Planck}; this study was mainly motivated by the effect of including the semi-forbidden and high-order two-photon transitions~\\citep{Dubrovich:2005fc}, which had been ignored in earlier calculations.  There have been many updates and improvements in the modelling of recombination since then.  \\citet{Switzer:2007sn} presented a multi-level calculation for neutral helium~(He\\,{\\sc i}) recombination including evolution of the radiation field, which had been assumed to be a perfect blackbody in previous studies. Other issues discussed recently include the continuum opacity due to neutral hydrogen (H)~\\citep[see also][]{Kholupenko:2007qs}, the semi-forbidden transition $2^3$p--$1^1$s ~\\citep[the possible importance of which was first proposed by][]{Dubrovich:2005fc}, the feedback from spectral distortions between $2^1$p--$1^1$s and $2^3$p--$1^1$s lines, and the radiative line transfer. In particular, continuum absorption of the $2^1$p--$1^1$s line photons by neutral hydrogen\\,(H\\,{\\sc i}) causes helium recombination to end earlier than previously estimated~(see Fig.~\\ref{fig1}). \\citet{Hirata:2007sp} also found that the high order two-photon rates have a  negligible effect on He\\,{\\sc i}, and the same conclusion was made by other groups for hydrogen as well~\\citep{Wong:2006iv,Chluba:2007qk}, largely because the approximate rates adopted by~\\citet{Dubrovich:2005fc} had been overestimated.  The biggest remaining uncertainty in He\\,{\\sc i} recombination is the rate of the $2^3$p--$1^1$s transition, which causes a variation equal to about 0.1 per cent in the ionization fraction $x_{\\rm e}$~\\citep{Switzer:2007sq}.  For hydrogen, \\citet{Chluba:2006bc} improved the multi-level calculation by considering seperate angular momentum $\\ell$ states.  This brings about a 0.6 per cent change in $x_{\\rm e}$ at the peak of the visibility function, and about 1 per cent at redshifts $z < 900$.  The effect of the induced 2s--1s two-photon rate due to the radiation background~\\citep{Chluba:2005uz} is partially compensated by the feedback of the Ly\\,$\\alpha$ photons ~\\citep{Kholupenko:2006jm}, and the net maximum effect on $x_{\\rm e}$ is only 0.55 per cent at $z \\simeq 900$. The high-order two photon transitions bring about a 0.4 per cent change in $x_{\\rm e}$ at $z \\simeq 1160$~\\citep{Wong:2006iv,Chluba:2007qk}. There are also 0.22 per cent changes in $x_{\\rm e}$ at $z \\simeq 1050$ when one considers the Lyman series feedback up to $n=30$, and there is additionally possibility of direct recombination, although this has only a roughly $10^{-4}$ per cent effect~\\citep{Chluba:2007yp}.  The list of suggested updates on $x_{\\rm e}$ is certainly not complete yet, since some additional effects, such as the convergence of including higher excited states and the feedback-induced corrections due to the He\\,{\\sc i} spectral distortions,  may enhance or cancel other effects.  In general we still need to develop a complete multi-level code for hydrogen, with detailed interactions between the atoms and the radiation field.  However, what is really important here is establishing how these effects propagate into possible systematic uncertainties in the estimation of cosmological parameters.  Since the uncertainties in cosmological recombination discussed in the \\citet{Lewis:2006ym} paper have been reduced or updated, it is time to revisit the topic on how the new effects or remaining uncertainties might affect the constraints on cosmological parameters in future experiments. The recent full version of the He\\,{\\sc i} recombination calculation~\\citep{Switzer:2007sn,Switzer:2007sq,Hirata:2007sp} takes too long to run to be included within the current Boltzmann codes for $C_\\ell$.  So instead, in this paper, we try to reproduce the updated ionization history by modifying R{\\sc ecfast}~\\citep{Seager:1999bc} using a simple parametrization based on the fitting formulae provided by \\citet{Kholupenko:2007qs}. We then use the C{\\sc osmo}MC~\\citep{Lewis:2002ah} code to investigate how much this impacts the constraints on cosmological parameters for an experiment like {\\sl Planck}.  \\begin{figure} \\centering \\vspace*{6cm} \\leavevmode \\special{psfile=gxe_topbottom.ps angle=0 hoffset=-115 voffset=-50 hscale=37 vscale=31} \\caption{\\small Top panel: Ionization fraction $x_{\\rm e}$ as a function of redshift $z$. The dotted\\,(red) line is calculated using the original RECFAST code. The solid\\,(black) line is the numerical result from \\citet{Switzer:2007sq}, while the dashed\\,(blue) and long-dashed\\,(green) lines are both evaluated based on the modification given by \\citet{Kholupenko:2007qs} -- the dashed one has $b_{\\rm He}=0.97$ (the value used in the original paper) and the long-dashed one has $b_{\\rm He}=0.86$. Bottom panel: The visibility function $g(z)$ versus redshift $z$. The two curves calculated~(dotted and long-dashed) correspond to the same recombination models in the upper panel. The cosmological parameters used for these two graphs are, $\\Omega_{\\rm B} = 0.04592$, $\\Omega_{\\rm M} = 0.27$, $\\Omega_{\\rm CDM}= 0.22408$, $T_{\\rm CMB}=2.728$\\,K, $H_0=71$ \\,kms$^{-1}$ Mpc$^{-1}$ and $f_{\\rm He}=0.079$.} \\label{fig1} \\end{figure} \\begin{figure} \\centering \\vspace*{6cm} \\leavevmode \\special{psfile=gvaryb.ps angle=0 hoffset=-110 voffset=-50 hscale=35 vscale=30} \\caption{\\small Ionization fraction $x_{\\rm e}$ as a function of redshift $z$ calculated based on the modified He\\,{\\sc i} recombination discussed here with different values of the helium fitting parameter $b_{\\rm He}$. The curve with $b_{\\rm He}=0$~(dot-dashed, blue) overlaps the line using Saha equilibrium recombination~(solid, black). The cosmological parameters used in this graph are the same as for Fig.~\\ref{fig1}.} \\label{fig3} \\end{figure} \\begin{figure} \\centering \\vspace*{6cm} \\leavevmode \\special{psfile=gvaryfudge03.ps angle=0 hoffset=-110 voffset=-50 hscale=35 vscale=30} \\caption{\\small The ionization fraction $x_{\\rm e}$ as a function of redshift $z$ calculated with different values of the hydrogen fudge factor $F_{\\rm H}$.  The cosmological parameters used in this graph are the same as in Fig.~\\ref{fig1}.} \\label{fig2} \\end{figure} ",
        "conclusions": "In this paper, we modify R{\\sc ecfast} by introducing one more parameter $b_{\\rm He}$ (besides the hydrogen fudge factor $F_{\\rm H}$) to mimic the recent numerical results for the speed-up of He\\,{\\sc i} recombination. By using the C{\\sc osmo}MC code with forecast {\\sl Planck} data, we examine the variation of these two factors to account for the remaining dominant uncertainties in the cosmological recombination calculation.  For He\\,{\\sc i}, the main uncertainty comes from the $2^3$p--$1^1$s rate~\\citep{Switzer:2007sq}, which corresponds to about a 1 per cent change in $b_{\\rm He}$.  We find that this level of variation has a negligible effect on the determination of the cosmological parameters.  Therefore, based on this simple model, if the existing studies have properly considered all the relevant physical radiative processes in order to provide $x_{\\rm e}$ to 0.1 per cent accuracy during He\\,{\\sc i} recombination, then we already have numerical calculations which are accurate enough for {\\sl Planck}.  For H, since there is still no comprehensive model which considers all the interactions between the atomic transitions and the radiation background, we consider the size of the updates as an indication of the existing level of uncertainty. We represent this uncertainty by varying the fudge factor $F_{\\rm H}$, because the largest update on $x_{\\rm e}$ occurs at $z \\lesssim 900$, and comes from a consideration of the separate angular momentum states~\\citep{Chluba:2006bc}.  We find that $F_{\\rm H}$ needs to be determined to better than 1 per cent accuracy in order to have negligible effect on the determination of cosmological parameters with {\\sl Planck}.  Hydrogen recombination is of course important for the formation of the CMB anisotropies $C_{\\ell}$, since it determines the detailed profile around the peak of the visibility function $g(z)$. A comprehensive numerical calculation of the recombination of H (similar to He\\,{\\sc i}) to include at least all the recent suggestions for updates on the evolution of $x_{\\rm e}$ is an urgent task.  We need to determine that the phenomenological parameters $F_{\\rm H}$ and $b_{\\rm He}$ are fully understood at the $\\lesssim 1$ per cent level before we can be confident that the uncertainties in the details of recombination will have no significant effect on the determination of cosmological parameters from {\\sl Planck}."
    },
    "0711/0711.3022_arXiv.txt": {
        "abstract": "A search for diffuse neutrinos with energies in excess of 10$^{5}$~GeV is conducted with AMANDA-II data recorded between 2000 and 2002. Above 10$^{7}$~GeV, the Earth is essentially opaque to neutrinos. This fact, combined with the limited overburden of the AMANDA-II detector (roughly 1.5 km), concentrates these ultra high-energy neutrinos at the horizon. The primary background for this analysis is bundles of downgoing, high-energy muons from the interaction of cosmic rays in the atmosphere. No statistically significant excess above the expected background is seen in the data, and an upper limit is set on the diffuse all-flavor neutrino flux of E$^{2}$ $\\Phi$$_{\\mathrm{90\\% CL}}$ $<$ 2.7 $\\times$ 10$^{-7}$ GeV cm$^{-2}$ s$^{-1}$ sr$^{-1}$ valid over the energy range of 2 $\\times$ 10$^{5}$~GeV to 10$^{9}$~GeV. A number of models which predict neutrino fluxes from active galactic nuclei are excluded at the 90\\% confidence level. ",
        "introduction": "AMANDA-II (Antarctic Muon and Neutrino Detector Array), a neutrino telescope at the geographical South Pole designed to detect Cherenkov light from secondary particles produced in collisions between neutrinos and Antarctic ice, has placed limits on the flux from point-like and diffuse sources of astrophysical neutrinos \\citep{ach07, ack04, ack05, ahr03a, ahr03b}. This work describes a search for neutrinos with energies above 10$^{5}$~GeV, which we define as ultra high-energy (UHE) neutrinos. These neutrinos are of interest because they could be associated with the potential acceleration of hadrons by active galactic nuclei \\citep{man95, mpr00, hal97, pro96, ste92}; they could potentially be produced by exotic phenomena such as the decay of topological defects \\citep{sig98} or possibly associated with the Z-burst mechanism \\citep{yos98}; and they are guaranteed by-products of the interactions of high-energy cosmic rays with the cosmic microwave background \\citep{eng01}.  This analysis is sensitive to all three flavors of neutrinos. Leptons and cascades from UHE electron, muon and tau neutrinos create bright, energetic events (Fig. \\ref{figevt}) which can be identified by AMANDA-II as far as 450~m from the center of the array (Fig. \\ref{figdist}). The sensitivity of this analysis starts at energies roughly coincident with the highest energy threshold of other diffuse analyses conducted with AMANDA-II \\citep{ach07, ack04}.  At UHE energies, the interaction length of neutrinos in rock is shorter than the diameter of the Earth \\citep{gan98}, so neutrinos from the Northern Hemisphere will interact before reaching AMANDA-II. Combined with the limited overburden above the AMANDA-II detector, this concentrates UHE events at the horizon. This contrasts with the majority of other astrophysical neutrino analyses completed using data from the AMANDA-II detector, which search for neutrinos from the Northern Hemisphere with energies below 10$^{5}$~GeV.  The flux of atmospheric neutrinos is negligible at UHE energies, with fewer than 10 events in three years expected from the model in \\citet{lip93} after intermediate UHE selection criteria have been applied. This drops to 0.1 events after application of all selection criteria. Similarly, there are fewer than 0.6 events expected in three years at the final selection level from prompt neutrinos from the decay of charmed particles produced in the atmosphere (using the ``C'' model from \\citet{zas93}). Therefore, the primary background for the UHE analysis is composed of many lower energy processes that mimic higher energy signal events. Cosmic ray collisions in the upper atmosphere that generate large numbers of nearly parallel muons (or ``muon bundles'') can generate high-energy signatures even though the individual muons have much lower energy than single leptons or cascades from UHE neutrinos. Signal and background events spread light over roughly equivalent areas in the detector, but UHE neutrino events are distinguishable because they have higher energy and higher light density than background events. Specialized selection criteria which use these properties, as well as differences in reconstruction variables, separate the UHE neutrinos from the background of muon bundles from atmospheric cosmic rays.  Limits have been placed on the all-flavor neutrino flux in the ultra high-energy range by other experiments (Fig. \\ref{figrep}). Additionally, a previous analysis using an earlier configuration of the AMANDA detector called AMANDA-B10, consisting of 302 optical modules \\citep{ack05}, has placed limits on the all-flavor UHE neutrino flux (Fig. \\ref{figrep}). This analysis uses 677 optical modules (OMs) of the AMANDA-II detector and gives a combined result using data from three years (2000-2002) with a livetime of 456.8 days.  A description of the AMANDA-II detector is given in section 2. Sections 3 and 4 discuss possible sources of astrophysical neutrinos and background, and the simulation of both. The selection criteria used to separate UHE neutrino signals from background are discussed in section  5. A study of systematic uncertainties is presented in section 6, and the results are shown in section 7. ",
        "conclusions": "The diffuse neutrino flux limit for a 1:1:1 $\\nu_e:\\nu_{\\mu}:\\nu_{\\tau}$ flavor ratio set by this analysis of \\begin{equation} \\mathrm{E^{2}\\Phi_{90\\% CL} \\le 2.7 \\times 10^{-7} GeV\\ cm^{-2}\\ s^{-1}\\ sr^{-1},} \\end{equation} is the most stringent to date above 10$^{5}$~GeV. A number of models for neutrino production have been rejected (see Table \\ref{tbl-mrf} for a full list). AMANDA-II hardware upgrades which were completed in 2003 should lead to an improvement of the sensitivity at ultra high-energies \\citep{sil05}. Additionally, AMANDA-II is now surrounded by the next-generation IceCube detector which is currently under construction. The sensitivity to UHE muon neutrinos for 1 year is expected to increase by roughly an order of magnitude as the IceCube detector approaches its final size of 1~km$^{3}$ \\citep{ahr04b}."
    },
    "0711/0711.1027_arXiv.txt": {
        "abstract": "{Large-scale accretion shocks around massive clusters of galaxies, generically expected in hierarchical scenarios of cosmological structure formation, are shown to be potential sources of the observed ultrahigh energy cosmic rays (UHECRs) by accelerating a mixture of heavy nuclei including the iron group elements. Current observations can be explained if the source composition at injection for the heavier nuclei is somewhat enhanced from simple expectations for the accreting gas. The proposed picture should be testable by current and upcoming facilities in the near future through characteristic features in the UHECR spectrum, composition and anisotropy. The associated X-ray and gamma-ray signatures are also briefly discussed.}  \\begin{document} ",
        "introduction": "The origin of UHECRs with energies $10^{18}$-$10^{20}$ eV and above remains one of the biggest mysteries in physics and astrophysics \\cite{nw00}. Only a few types of astrophysical objects appear capable of accelerating UHECRs to the highest observed energies, such as the jets of radio-loud active galactic nuclei or gamma-ray bursts \\cite{hil84, ino07}. However, no unambiguous identification with any kind of source has been achieved so far.  In the currently favored picture of hierarchical structure formation in the CDM cosmology, all massive clusters of galaxies should be surrounded by strong accretion shocks, as a consequence of continuing infall of dark matter and baryonic gas \\cite{ryu03}. Such shocks should be interesting sites of particle acceleration, and have also been proposed as sources of UHECRs \\cite{nma95}. Here we summarize our recent work on this subject invoking UHECR nuclei; more details can be found in Ref. \\cite{isma07}. ",
        "conclusions": "\\begin{figure}[htb!] \\begin{center} \\includegraphics [width=0.47\\textwidth]{fig2a.eps}\\\\ \\includegraphics [width=0.47\\textwidth]{fig2b.eps} \\end{center} \\caption{Observed UHECR spectrum (top) and mean mass composition (bottom) versus energy $E$ (1 EeV $\\equiv 10^{18}$ eV) from cluster accretion shocks for $\\alpha=1.7$ and $\\beta=0.5$, compared with the current data for HiRes (bars) and Auger (asterisks). The histograms are the average result over different model realizations for the cases with (thick) and without (thin) EGMF, and the thin curves outline the median deviations due to cosmic variance for the former case only. The straight line in the top panel denotes the injection spectrum.} \\label{fig:speccomp} \\end{figure}  Fig.\\ref{fig:speccomp} shows our results for the observed spectrum and composition for $\\alpha=1.7$ and $\\beta=0.5$, which are consistent with the current data for HiRes \\cite{abb04} and Auger \\cite{yam07} within about 2 sigma of combined cosmic and statistical variance, for both the differential and integral fluxes (and possibly AGASA \\cite{tak98} as well). Values of $\\alpha<2$ are naturally expected at the high energy end for nonlinear shock acceleration that accounts for the dynamical back reaction from CRs \\cite{md01}. The spectral steepening at $\\ga 10^{20}$ eV is due both to propagation losses and the $E_{\\max}$ limit at the source. Normalization to the observed flux and comparison with the available accretion power for $M > 10^{15} M_\\odot$ fixes $f_{\\rm CR}$, which is $\\simeq 0.005-0.3$ for cases with EGMF and $\\simeq 0.002$ for the case without. Low values of $f_{CR}$ may reflect inefficient escape of CRs from the system, which is conceivable in view of the converging nature of the accretion flow. CR escape may be mediated by diffusion in directions away from the filaments, or possibly by advection in outflows during merging events.  The mass composition at $\\la 3 \\times 10^{19}$ eV is predominantly light and consistent with HiRes reports \\cite{abb05}, while the rapid increase of the average mass at higher energies is a clear prediction of the scenario to be tested by the new generation experiments (and is in line with the latest Auger results on the shower elongation rate \\cite{ung07}).  Despite the relative rarity of massive clusters in the local universe, strong deflections of the highly charged nuclei in EGMF allow consistency with the currently observed global isotropy (Fig. \\ref{fig:aniso}). On the other hand, with a sufficient number of accumulated events, clear anistropies toward a small number of individual sources should appear, although this prediction is subject to uncertainties in the EGMF and Galactic fields.  \\begin{figure}[htb!] \\begin{center} \\includegraphics [width=0.42\\textwidth]{fig3.eps} \\end{center} \\caption{Angular power spectrum $C(l)$ of UHECR arrival directions above $4 \\times 10^{19}$ eV versus multipole $l$, for a realization with EGMF and a single, dominant cluster at $D \\sim$ 50 Mpc. The crosses are for 100 events with AGASA + SUGAR exposure and diamonds for 1000 events with Auger North + South exposure. Vertical bars indicate statistical errors.} \\label{fig:aniso} \\end{figure}  An aspect of this scenario that warrants further study is the spectral domain $< 10^{19}$ eV and the implications for the Galactic-extragalactic transition region \\cite{all07}."
    },
    "0711/0711.4602_arXiv.txt": {
        "abstract": "Motivated by the string gas cosmological model, which predicts a blue tilt of the primordial gravitational wave spectrum, we examine the constraints imposed by current and planned observations on a blue tilted tensor spectrum. Starting from an expression for the primordial gravitational wave spectrum normalised using cosmic microwave background observations, pulsar timing, direct detection and nucleosynthesis bounds are examined. If we assume a tensor to scalar ratio on scales of the CMB which equals the current observational upper bound, we obtain from these current observations constraints on the tensor spectral index of $n_{T} \\lesssim 0.79$, $n_{T} \\lesssim 0.53$, and $n_{T} \\lesssim 0.15$ respectively. ",
        "introduction": "A stochastic background of primordial gravitational waves represents valuable information about the very early universe as well as a way to discriminate between the myriad of cosmological models currently proposed. Although a stochastic background of gravitational waves has yet to be directly detected, efforts are being undertaken to do so by many current and future experiments, some of the largest being LIGO, GEO600, TAMA, and LISA \\cite{Allen:1996vm}. A stochastic background of gravitational waves is characterised by the gravitational wave spectrum \\begin{equation} \\Omega_{gw} \\, = \\, \\frac{1}{\\rho_{c}}\\frac{d\\rho_{gw}}{d\\ln f}\\,, \\end{equation} where $\\rho_{c}$ is the critical density of the universe, $\\rho_{gw}$ is the energy density of the background gravitational waves and $f$ is frequency \\cite{Allen:1996vm}. In the above, the energy density in gravitational waves is written as an integral over $\\ln f$, and the derivative picks out the integrand. In practice, the gravitational wave spectrum is commonly assumed to depend on frequency as a power of $f$.  Within the framework of the inflationary universe scenario the primordial gravitational wave spectrum is predicted to be nearly scale invariant with a slight red tilt \\cite{Spergel:2006hy}, i.e. more power at large scales. The reason for the red tilt is that the amplitude of the gravitational wave spectrum on a fixed scale $k$ is set by the Hubble constant $H$ at the time $t_i(k)$ when the scale $k$ exits the Hubble radius during the period of inflation. Smaller scales exit the Hubble radius later when the Hubble constant is smaller, leading to the red tilt. However, there do exist alternative cosmological models to the standard inflationary scenario which predict a blue tilt of the gravitational wave spectrum, i.e. more power at small scales, which have not yet been ruled out by observations. One cosmological model that predicts such a tilt is string gas cosmology \\cite{BV}.  String gas cosmology \\cite{BV} is an approach to string cosmology which starts from the new degrees of freedom and symmetries which string theory contains, but particle physics-based models lack, namely string winding modes, string oscillatory modes and T-duality symmetry, and uses them to develop a new cosmological model. The claim is that by making these crucial stringy additions, one obtains a new cosmological model which is singularity free \\cite{BV}, generates nearly scale-invariant scalar metric perturbations from initial string thermodynamic fluctuations \\cite{NBV,BNPV1} and provides a natural explanation for the observed dimensionality of space \\cite{BV} (see \\cite{RHBrev} for an overview of the string gas cosmology structure formation scenario). A key result which emerges \\cite{BNPV2} from string gas cosmology is that the gravitational wave spectrum has a slight blue tilt, giving rise to a testable prediction different from that of the inflationary universe paradigm.  Using string gas cosmology as a motivation, in this work we will focus on how a tensor spectrum with a blue tilt can be constrained by current observational results, and what the prospects for improved constraints from some planned experiments are.  The starting point of our analysis will be an expression for the primordial gravitational wave spectrum normalised by cosmic microwave background (CMB) observations. Chongchitnan and Efstathiou \\cite{Chongchitnan:2006pe} have derived just such an expression with a pivot scale $k_{0}=0.002\\:\\mbox{Mpc}^{-1}$ using the combined results from multiple surveys. Using the value $\\mathcal{P}_{S}(k_{0})\\simeq2.21\\times10^{-9}$ for the amplitude of the scalar power spectrum evaluated at the pivot scale, they found $\\Omega_{gw}(f)$ can be written as \\begin{equation}\\label{chong} h^{2}\\Omega_{gw}(f) \\, \\simeq \\, 4.36\\times10^{-15}r\\left(\\frac{f}{f_{0}}\\right)^{n_{T}}\\,, \\end{equation} where $f_{0}=3.10\\times10^{-18}$ Hz. Solving this result for the tilt of the gravitational wave spectrum, $n_{T}$, we get the explicit expression \\begin{equation}\\label{tilt} n_{T} \\, \\simeq \\, \\frac{1}{\\ln(f)-\\ln(f_{0})}\\ln\\left(2.29\\times10^{14}\\frac{h^2\\Omega_{gw}(f)}{r}\\right)\\,. \\end{equation} In the above equations $r$ is the tensor-to-scalar ratio evaluated at the pivot scale, \\begin{equation} r \\, = \\, \\frac{\\mathcal{P}_{T}(k_{0})}{\\mathcal{P}_{S}(k_{0})}\\,. \\end{equation}  The calculation of the tensor-to-scalar ratio depends quite sensitively on the parameters of the cosmological model under consideration. For that reason we choose to leave $r$ as a free parameter in the main expressions calculated in this work. However, for the sake of examining some numerical values of the constraints derived here we will insert a value of $r$ corresponding to the current upper bound into our results. Note, however, that the bounds we derive depend only logarithmically on $r$. In each case we choose to use the value of the tensor-to-scalar ratio given by the combined three-year Wilkinson Microwave Anisotropy Probe (WMAP) and lensing normalised Sloan Digital Sky Survey (SDSS) data applied to a typical $\\Lambda\\mbox{CDM}$ model including tensors \\cite{Spergel:2006hy}.  In Sections II and III we will use the expression for the tilt of the gravitational wave spectrum \\eqref{tilt} along with several direct detection constraints to calculate the bounds on the required blue tilt of the spectrum. In Section IV we calculate the bound on the required blue tilt of the spectrum arising indirectly from the theory of big bang nucleosynthesis. In Section V we will investigate whether the CMB observations are compatible with the calculated bounds, or if they will offer even tighter constraints. We conclude with a discussion of our findings along with other issues related to the method used. ",
        "conclusions": "A stochastic background of primordial gravitational waves is a prediction of many cosmological models. Assuming the gravitational wave spectrum depends as a power on frequency, then this spectrum can be characterised by its tilt and amplitude. Most models predict the tilt to be nearly scale-invariant but slightly red, while some models, like string gas cosmology, predict a slight blue tilt. Although a gravitational wave background has yet to be directly detected, observational results can already be used to constrain it.  Using the current results from pulsar timing observations, direct detection observations and the theory of nucleosynthesis, we have placed bounds on the possible blue tilt of the gravitational wave background. After completion of this work we became aware of \\cite{Boyle} in which a master equation was derived which relates the short wavelength observable $\\Omega_{gw}(f)$ to the tensor to scalar ratio measured with the CMB. The goal of that work was to develop a formulation which is as general as possible. In particular, the equation of state parameter and the tensor spectral index are taken to be arbitrary functions of the scale factor and wavenumber respectively, not constants as is often assumed. Their master equation is thus a more general version of our ``master'' equation \\eqref{chong} and we could have just as easily used it as our starting point. In fact, we have confirmed that by choosing a constant value $w=1/3$ for the equation of state parameter (which is described as the most logical value in \\cite{Boyle}) and numerical values for other cosmological parameters that match our choices above, we can indeed re-derive all of our results from the master equation of \\cite{Boyle}. We consider this a good consistency check for the constraints derived in this work. The authors of \\cite{Boyle} do include a discussion of some of the same types of observations mentioned here, namely laser interferometer and pulsar timing, however, we stress that they do not use the current numerical constraint from any particular observatories to compute actual upper bounds on $n_T$, which was the purpose of this work. The authors of \\cite{Boyle} also discuss a constraint on $n_T$ coming from BBN, but they take the constraint on $\\Omega_{gw}$ from BBN to be a constant across all frequencies rather than integrating their master equation as was done in this work. In the end they find a weaker bound on the tilt, $n_T\\lesssim0.36$, than the one obtained here.  By far the tightest constraint on the tilt comes from big-bang nucleosynthesis. If we take the tensor to scalar ratio on CMB scales to be given by the current observational upper bound, and if we take the ultraviolet cutoff scale in the spectrum of gravitational radiation to be the Planck scale, then the bound is $n_{T}\\lesssim 0.15$, tighter than even Advanced LIGO, the future PPTA, and LISA can hope to achieve. It is not surprising that nucleosynthesis provides the tightest bounds on a blue tilt of the gravitational wave spectrum since nucleosynthesis probes physics on scales much smaller than the other experiments we analysed, and spectra with blue tilts have more power on the smallest scales. That is, BBN gives us the largest ``lever arm'' to probe gravitational waves in conjunction with CMB observations, a point also made in \\cite{Peiris}. From our results we can clearly see the trend that the bound on the blue tilt of the tensor spectral index tightens as the length scale probed by the given experiment decreases.  Simulations of the angular power spectrum of the temperature anisotropies in the CMB did not offer any tighter constraints on the tensor spectral index, with each of the constraints calculated in this paper producing a temperature power spectrum that was within the cosmic variance error of one calculated for a standard $\\Lambda\\mbox{CDM}$ model with tensor modes included.  As mentioned in the Introduction, Equation \\eqref{chong} has been normalised at the scale of cosmic microwave background observations. However, those experiments probe scales that are approximately ten orders of magnitude larger than those probed by the PPTA and approximately nineteen orders of magnitude larger than those probed by LIGO, with LISA probing between the two. Extrapolating between such a large difference in scales is not straightforward and we should note that in \\cite{Chongchitnan:2006pe} the authors conclude from their analysis that even within the framework of the inflationary universe paradigm the formula for the primordial gravitational wave spectrum \\eqref{chong} is too restrictive, and they believe it is indeed not possible to extrapolate reliably over such a large difference in scales. Whether or not this is the case in the string gas cosmology model should perhaps be examined more carefully in future work.  Continuing with string gas cosmology, we conclude that the current bounds on the tilt of the gravitational wave spectrum are weak. The predicted magnitude of the blue tilt of the gravity wave spectrum is thought to be comparable to the magnitude of the red tilt of the spectrum of scalar metric fluctuations \\cite{BNPV2}. If the latter is taken to agree with the current bounds, we predict a blue tilt of less than $n_T = 0.1$ which will not be easy to detect. There may, however, be models similar to string gas cosmology in which the scalar and tensor tilts are not related, and for which planned experiments could set valuable constraints on the model parameter space."
    },
    "0711/0711.2024_arXiv.txt": {
        "abstract": "We present the first high spatial resolution X-ray study of NGC~2244, the 2~Myr old stellar cluster immersed in the Rosette Nebula, using the {\\it Chandra X-ray Observatory}.  Over 900 X-ray sources are detected; 77\\% have optical or FLAMINGOS near-infrared (NIR) stellar counterparts and are mostly previously uncatalogued young stellar cluster members.  All known OB stars with spectral type earlier than B1 are detected and the X-ray selected stellar population is estimated to be nearly complete between 0.5 and 3~M$_{\\odot}$. The X-ray luminosity function (XLF) ranges from $29.4 < \\log L_x < 32.0$ ergs s$^{-1}$ in the hard ($2-8$~keV) band. By comparing the NGC~2244 and Orion Nebula Cluster XLFs, we estimate a total population of $\\sim 2000$ stars in NGC~2244. A number of further results emerge from our analysis: 1) The XLF and the associated $K$-band luminosity function indicate a normal Salpeter initial mass function (IMF) for NGC~2244. This is inconsistent with the top-heavy IMF reported from earlier optical studies that lacked a good census of $<4$M$_{\\odot}$ stars. 2) The spatial distribution of X-ray stars is strongly concentrated around the central O5 star, HD 46150. The other early O star, HD 46223, has few companions. The cluster's stellar radial density profile shows two distinctive structures: a power-law cusp around HD~46150 that extends to $\\sim 0.7$~pc, surrounded by an isothermal sphere extending out to 4~pc with core radius 1.2~pc. This double structure, combined with the absence of mass segregation, indicates that this 2~Myr old cluster is not in dynamical equilibrium. Our results will strongly constrain models of the cluster formation process. The spatial distribution of X-ray selected $K$-excess disk stars and embedded stars is asymmetric with an apparent deficit towards the north. 3) The fraction of X-ray-selected cluster members with $K$-band excesses caused by inner protoplanetary disks is 6\\%, slightly lower than the 10\\% disk fraction estimated from the FLAMINGOS study based on the NIR-selected sample. This is due to the high efficiency of X-ray surveys in locating disk-free weak-lined T Tauri stars. 4) X-ray luminosities for 24 stars earlier than B4 confirm the long-standing $\\log (L_x/L_{bol})\\sim -7$ relation. The Rosette OB X-ray spectra are soft and consistent with the standard model of small-scale shocks in the inner wind of a single massive star. 5) About 50 intermediate-mass ($2 < M < 8$~M$_{\\odot}$) cluster members are identified; they exhibit a wide range of X-ray luminosities consistent with previously studied samples of Herbig Ae/Be stars.  Their $K$-excess disk fraction is $\\sim$10\\%, indicating that the Herbig Ae/Be phenomenon is rare or short-lived. ",
        "introduction": "During their evolution from Class I protostars to zero-age-main-sequence (ZAMS) stars, young stellar objects are readily identified in X-rays due to their highly elevated X-ray emission compared to the older Galactic stellar population \\citep[see reviews by][]{Feigelson99,Favata03,Feigelson07}. High spatial resolution {\\em Chandra} observations of well-known Galactic star forming regions (e.g., the {\\em Chandra} Orion Ultradeep Project, hereafter COUP, Getman et al.\\ 2005; RCW 38, Wolk et al.\\ 2006; Cepheus B, Getman et al.\\ 2006), more distant molecular cloud and HII region complexes (e.g., NGC 6334, Ezoe et al.\\ 2006; NGC 6357, J. Wang et al.\\ 2007; M~16, Linsky et al.\\ 2007; M~17, Broos et al.\\ 2007), the Galactic Center \\citep[e.g., the  Arches and Quintuplet clusters,][]{Wangqd06,Muno06}, and extragalactic star forming regions \\citep[e.g., 30 Dor,][]{Townsley06,Townsley06b}, have greatly advanced our knowledge of star formation processes in these regions.  Moreover, these studies demonstrate the unique power of studying star formation in the X-ray band. Besides the high energy phenomena, other new information about the young clusters such as population, membership, and star formation environs can be obtained when combining X-ray detections with knowledge obtained from longer wavelength studies. For example, although X-ray-selected samples contain a very small fraction of extragalactic sources and Galactic field stars, X-ray sources that have optical and infrared (IR) counterparts are mostly cluster members. This is in contrast to the high percentage of non-members in the optical/IR images where, except for those sources with massive dusty disks, membership for individual stars generally has to be ascertained via spectroscopy. In initial mass function (IMF) studies of more distant, high mass star forming regions, traditional optical measurements are significantly encumbered by large reddening and membership confusion. In comparison to the massive stars, the lower mass pre-main-sequence (PMS) stellar populations are much less accessible and thus much harder to evaluate. In the past, $H_{\\alpha}$ emission was used in general as a youth and membership indicator. However, in addition to the observational challenges due to prevalent bright $H_{\\alpha}$ nebulosity in HII regions, the emission itself requires accretion activity in protoplanetary disks \\citep{Muzerolle98,Muzerolle01} and thus is susceptible to the disk evolutionary stages.  It has long been recognized that X-ray emission circumvents these problems and is very effective for securing PMS membership of young clusters \\citep{Feigelson99}. Modern X-ray observatories like {\\em Chandra} and {\\em XMM-Newton} enable identification of hundreds of individual members from their X-ray emission. In this work and subsequent papers, we report {\\em Chandra} studies of a well-known star forming complex, concentrating here on a new census of the low mass cluster members and new knowledge of the IMF arising from the X-ray perspective.  The Rosette star-forming complex, situated in a large star formation site in the Perseus spiral arm, provides an ideal testbed for studying sequential formation of clusters due to the favorable orientation of its morphological components, consisting of an expanding blister HII region on the edge of a giant molecular cloud oriented perpendicular to the line-of-sight.  Thorough review of past and present research on this popular and important star formation region appear in \\citet{TFM03} (hereafter TFM03) and \\citet{RL07} (hereafter RL07). Here we highlight what is most relevant to this work.  The Rosette Nebula (= Sharpless 275 = W 16 = NGC 2237--2239, NGC 2244, and NGC 2246) is a large HII region at the tip of the Rosette Molecular Cloud (RMC). In both radio and optical images, it shows a prominent ring-like morphology, with a cluster of ionizing young stars located in the central hole \\citep{Celnik85,TFM03}.  The large-scale IRAS data \\citep{Cox90} and CO emission map \\citep{Heyer06} clearly show a similar annular morphology extending into the molecular cloud (Figure~\\ref{fig:large_view}a).  Extending to the southeast of the Rosette Nebula, the RMC is an elongated giant molecular cloud with $\\sim 10^5 M_{\\odot}$ of gas and dust \\citep{Blitz80}. Multiwavelength observations in mid-IR and radio show clumpy structure in the RMC \\citep{Williams98}.  Embedded star clusters have been revealed in the densest parts through near-IR imaging surveys \\citep{Phelps97,Roman05,Roman07}. Figure~\\ref{fig:large_view}a demonstrates the association between \\citet{Phelps97} IR clusters and the molecular clumps where the CO emission peaks. \\citet{Lada03} have shown that embedded clusters are physically associated with the most massive and dense cores in molecular clouds, based on systematic and coordinated surveys \\citep[e.g.,][]{Lada92}.  In the X-ray band, previous imaging study of the Rosette Complex was hampered by low spatial resolution. An early {\\it Einstein Observatory} study detected a few individual O stars and extended $\\sim 2$~keV X-ray emission at the center of the nebula \\citep{Leahy85}. \\citet{BC02} considered the integrated contribution from X-ray emitting low mass PMS stars in the cluster and concluded that apparently diffuse emission seen by {\\em ROSAT} could be explained by unresolved point sources. \\citet{Chen04} analyzed the same {\\em ROSAT} data set and attributed the brightest X-ray sources to massive stars, active T Tauri stars, and foreground stars. \\citet{Gregorio98}, studying a {\\em ROSAT} observation of the RMC, reported faint X-ray point sources associated with T Tauri stars and Herbig Ae/Be (HAeBe) stars and X-ray ``hot spots'' from unresolved embedded low mass star clusters. In TFM03, the first high resolution {\\em Chandra} X-ray image mosaic of this high-mass star forming region was presented and soft diffuse X-ray plasma ($kT\\simeq 0.06$ and $\\simeq 0.8$ keV) with luminosity $L_x\\simeq 6\\times 10^{32}$ ergs s$^{-1}$ was detected in the HII region. It was attributed to a combination of fast O star winds and unresolved T Tauri stars.  The young stars powering the Rosette HII region are members of the massive open cluster NGC 2244.  Figure~\\ref{fig:large_view}b shows a Digital Sky Survey image ($59^{\\prime}\\times 59^{\\prime}$) of this region with a few bright stars and interesting objects labeled. Despite its apparent low concentration of stars in the optical, NGC 2244 contains $\\sim 30$ early-type stars between O4V and B3V (Table 6 in TFM03), whose cluster memberships are secured from deep photometric study along with proper motion data and spectroscopy \\citep{Verschueren91,PS02}. No other known massive young stellar cluster within 2 kpc, other than RCW 38 \\citep{Wolk06} and M 17 (Broos et al.\\ 2007), is comparably rich. The PMS members were largely unknown from optical studies; only a handful of $H_{\\alpha}$ emission objects have been identified as young PMS members \\citep{PS02,Li04}. \\citet{BC02} presented 138 sources selected from {\\it ROSAT} PSPC and HRI detections. Although incomplete, they revealed previously unknown PMS stars to K spectral types.  The distance to NGC 2244 has been measured in many visual photometry studies and ranges between 1.4 kpc and 1.7 kpc.  \\citet{Hensberge00} derive a distance of $1.39\\pm 0.1$ kpc to an eclipsing binary member V578 Mon, using a novel Fourier spectral disentangling technique. The main sequence (MS) turn-off age estimated by \\citet{PS02} is 1.9 Myr, which is consistent with the inferred age from V578 Mon, $2.3\\pm 0.2$ Myr \\citep{Hensberge00}. This makes NGC~2244 the youngest cluster within the larger Mon OB2 association \\citep{Hensberge00}. We adopt a distance of 1.4~kpc and a cluster age for NGC~2244 of 2~Myr throughout our studies; these are consistent with TFM03. Note that a larger distance value $d=1.6$ kpc has been used recently by other researchers \\citep{PS02,Roman07}. This discrepancy in distance affects our derived X-ray luminosities ($\\log L_x$) by only 0.1 dex. At the distance of 1.4~kpc, 1$^{\\prime}$ corresponds to 0.4 pc.  Whereas the TFM03 paper was dedicated to study of the diffuse X-ray emission in the Rosette HII region, we present here an X-ray point source study of the Rosette complex based on TFM03 data and a new 75 ks {\\em Chandra} observation centered on NGC 2244. The field of view of our mosaic {\\em Chandra} fields are outlined by the polygons in Figure~\\ref{fig:large_view}a. We separate our study into a series of four papers with different astrophysical emphasis. In this work (Paper I), we report {\\em Chandra} observations of the NGC 2244 cluster and the Rosette HII region and study the young stellar population in detail.  In an upcoming paper (Wang et al., in preparation; Paper II), we describe {\\em Chandra} observations of the embedded clusters in the RMC, aiming to investigate cluster formation in a sequential manner and to test whether molecular clumps preferentially forming embedded clusters of low-mass stars make up the fundamental building blocks of star formation in molecular clouds. The westernmost {\\em Chandra} field, designed to study triggered star formation and the X-ray detection of a twin cluster to NGC 2244 \\citep{Li05,Roman07}, will be presented in Paper III (Wang et al., in preparation). A detailed analysis of the diffuse X-ray emission in the HII region will appear in Paper IV (Townsley et al., in preparation).  New mid-IR observations are also contributing to our knowledge of the Rosette complex.  A shallow {\\em Spitzer Space Telescope} survey of NGC 2244 has been reported \\citep{Balog07} and a deep {\\em Spitzer} survey for disk emission from {\\em Chandra} stars in NGC 2244 is also underway (PI: Bouwman). Together with the {\\em Spitzer} MIPS coverage of the massive cores in the RMC (PI: Bonnell) and a new {\\em Spitzer} program (PI: Rieke), virtually the entire nebula and the molecular cloud will be completely mapped.  This paper is organized as follows. First, we describe the {\\em Chandra} observations and data reduction in \\S2. In \\S3, we identify the X-ray sources with optical and infrared counterparts, and evaluate the fraction of contaminants through simulations utilizing stellar population synthesis model and the $\\log N$--$\\log S$ distribution for extragalactic X-ray sources. \\S4 is devoted to global properties of the NGC 2244 cluster such as the X-ray luminosity function, the initial mass function, the $K$-band luminosity function, spatial structures, mass segregation, and the $K$-band excess disk fraction among the X-ray detected stars. We present collective properties of interesting X-ray sources in \\S5, ending with a summary in \\S6. ",
        "conclusions": "}  We present a high spatial resolution X-ray study of the NGC 2244 cluster in the Rosette Nebula obtained via deep {\\it Chandra} observations. Our main findings follow:  1. We detect 919 X-ray sources with a limiting X-ray sensitivity of $L_{t,c}\\sim 1\\times 10^{29}$ ergs s$^{-1}$.  Positional coincidence matching yields a total of 712 ONIR counterparts. We estimate 8\\% extragalactic and Galactic contamination. The rest of the X-ray sources without ONIR counterparts are likely new NGC~2244 members clustered around the massive star HD~46150 or deeply embedded in the cloud. The X-ray detected population provides the first deep probe of the rich low mass population in this massive cluster.  2. The locations of most ACIS sources in the color-magnitude plot indicate a large population of 2 Myr old PMS low mass stars ($M \\la 2 M_{\\odot}$) subject to a visual extinction of $1\\la A_V \\la 2$ at 1.4 kpc. We derive an overall $K$-excess disk frequency of $\\sim$6\\% for stars with mass $M \\ga 0.5M_{\\odot}$ using the X-ray selected sample, slightly lower than the 10\\% $K$-excess disk fraction using a FLAMINGOS selected sample. Both fractions are significantly lower than the 45\\% mid-IR disk fraction in a {\\em Spitzer} sample that is more sensitive to disks. We emphasize that the combination of young stars identified in X-rays (mostly Class III stars) and those selected through IR colors will provide the best census to date for the young stellar population of this region. Three objects have Class I colors.  3. The derived XLF ($L_{h,c}$) for NGC 2244 is compared to the XLFs of the ONC, M17, Cep B, and NGC 6357; this exercise indicates that the unobscured population in NGC 2244 is 1.2 times larger than that of the ONC, or $\\sim 1000$ stars detectable in COUP-sensitivity X-ray observations. Taking into account the obscured population, the total stellar population in NGC 2244 is $\\sim 2000$, in good agreement with the estimated population from the spatially complete 2MASS study. The XLF and KLF suggest a normal Salpeter IMF for NGC 2244; we do not confirm a top heavy IMF reported from earlier optical studies.  4. We examine the spatial distribution of the X-ray identified NGC 2244 cluster members; the stellar surface density map suggests a spherical cluster with substructure. We confirm the existence of a subcluster around HD 46150 with $\\sim 50$ members in a 1~pc region; a second small subcluster consisting of a number of late type stars is also found.  The O4 star HD 46223 has few companions. The radial density profile of NGC 2244 shows a larger relaxed structure around the central subcluster. Similar structure is seen in NGC 6357. No evidence for significant mass segregation is found in this cluster. Altogether we suggest that this 2 Myr cluster is not dynamically evolved and has a complex star formation history. Our results will strongly constrain models of the cluster formation process.  5. We detected all 9 OB stars with spectral types B0.5 or earlier, but only 6 out of 14 B stars with spectral types B1--B3 in our field of view. X-ray spectra for the massive stars in NGC 2244 all show soft emission. We confirm the long-standing $\\log(L_x/L_{bol})\\sim -7$ relation for the NGC 2244 O stars. Large scatter around this correlation was found for the B stars.  6. We report X-ray emission detected from a few interesting individual objects, including the ionizing source of the optical jet Rosette HH1, binary systems, a magnetic star, and a possible X-ray luminous uncataloged massive star."
    },
    "0711/0711.2481_arXiv.txt": {
        "abstract": "The TeV blazar PKS 2155--304 was the subject of an intensive 2 week optical and near-infrared observing campaign in 2004 August with the CTIO 0.9m telescope. During this time, simultaneous X-ray data from \\emph{RXTE} were also obtained. We compare the results of our observations to the results from two previous simultaneous multiwavelength campaigns on PKS 2155--304. We conclude that the correlation between the X-ray and UV/optical variability is strongest and the time lag is shortest (only a few hours) when the object is brightest. As the object becomes fainter, the correlations are weaker and the lags longer, increasing to a few days. Based on the results of four campaigns, we find evidence for a linear relationship between the mean optical brightness and lag time of X-ray and UV/optical events. Furthermore, we assert that this behavior, along with the different multiwavelength flare lag times across different flux states is consistent with a highly relativistic shock propagating down the jet producing the flares observed during a high state. In a quiescent state, the variability is likely to be due to a number of factors including both the jet and contributions outside of the jet, such as the accretion disk. ",
        "introduction": "In the broad class of extragalactic objects known as active galactic nuclei (AGNs), blazars distinguish themselves in many ways. Blazars are a radio-loud subclass of AGNs, typically featuring very core-dominated radio morphologies. This class of objects is characterized by (1) a strong, featureless continuum, (2) strong (up to $\\sim$20\\%) and variable polarization, and (3) the presence of large-amplitude, rapid flux variability observed on timescales of less than an hour to several years at all wavelengths. The most extreme and unique property of blazars is their highly beamed continuum, most likely produced by a jet of relativistic material aimed close to the observer's line of sight. BL Lac objects have particularly weak spectral lines, so their redshifts, and hence luminosities, are difficult to determine. Blazars exhibit the most extreme variability observed for any class of AGN. See \\citet{urr95} and the references therein for further details.  The spectral energy distribution (SED) of blazars is distinguished by two peaks: one in the radio/UV range and the other in the X-ray/$\\gamma$-ray range. The spectrum in the radio to UV range is generally agreed to arise from synchrotron emission from relativistic electrons spiralling around the jet's magnetic field lines \\citep{ulr97}. The X-ray/$\\gamma$-ray spectrum is most likely due to inverse Compton (IC) emission in low-frequency peaked BL Lac objects (LBLs) and flat-spectrum radio quasars (FSRQs), while most of the X-rays are the high-energy tail of the synchrotron emission in high-frequency peaked BL Lac objects (HBLs) \\citep{ulr97}. Any $\\gamma$-ray emission is most likely explained as inverse Compton scattering of photons by relativistic electrons in the jet (e.g. \\citet{mar85}).  Multiwavelength campaigns are a very useful tool for investigating the underlying structure and physics of blazar jets. Assuming different areas of the jet emit different wavelengths of light, simultaneous multiwavelength campaigns can examine the physical processes operating in the jet and indicate spatial relationships between different emission regions. Correlations in multiwavelength flares can be used to constrain models predicting the source of photons which are upscattered to produce the IC emission observed in blazars. In many blazars, the X-ray variability is observed to be more rapid and of larger amplitude than the variability at longer wavelengths, implying that this emission originates from a smaller region than the optical and radio emitting regions. We must understand the physics of the region near the central engine in order to fully understand blazars, since this is where jet particles are collimated and accelerated to relativistic speeds.  PKS 2155--304 is one of the best-studied blazars in the southern hemisphere. It was first observed in the radio as part of the Parkes survey \\citep{shi74}. \\citet{gri78} noted a blue stellar object with a featureless continuum at the same location. \\citet{sch79} and \\citet{hew80} first identified PKS 2155--304 as a BL Lac object. \\citet{sch79} report the first X-ray observations of this object, using \\emph{HEAO 1}. Significant optical variability ($\\sim$1.5 mag) and polarization ($\\sim$5\\%) was observed by \\citet{gri79}. A redshift of 0.117 $\\pm$ 0.002 was determined by \\citet{bow84}; the current accepted value is 0.116 \\citep[e.g., ][]{fal93}. The long-term R-band light-curve is displayed in Figure 1. Recently, it has been detected at TeV energies \\citep{rob99,dja03,emm06}, thus opening up new possibilities for multiwavelength campaigns. During the past 15 years, PKS 2155--304 has been the subject of numerous campaigns examining possible correlations between X-ray, UV, optical, and radio variability.  The two most prominent and intensive of these earlier campaigns occurred in 1991 November \\citep{ede95, urr96} and 1994 May \\citep{urr96, urr97}. In both campaigns, at least one set of temporally correlated flares was observed in X-ray and ultraviolet wave bands. The 1991 campaign showed correlated optical flares as well. However, the time lags between flares in different wave bands differed in each campaign.  Since PKS 2155--304 is a HBL, the radio through X-ray emission are all thought to be due to the synchrotron process in the jet, with higher energy emission occuring closer to the central engine than lower energy emission \\citep{ulr97}. Hence, flares at lower energies should lag flares at higher energies. The flaring activity in both the 1994 and 1991 campaigns follow this trend. The results from these campaigns will be discussed further in conjunction with the results from the new 2004 campaign on PKS 2155--304. ",
        "conclusions": "We present the comparison of our observations to the results from two previous simultaneous multiwavelength campaigns on PKS 2155--304. We conclude that the correlation between the X-ray and UV/optical variability is strongest and the lag is shortest when the object is brightest. As the object becomes fainter, the correlations are weaker and the are lags longer. Furthermore, we assert that the different lag times across different flux states indicate that a relativistic shock propagating down the jet causes the flares during a high state. During a low state, the flaring could be due to something other than the jet, for example the accretion disk (e.g. Mangalam \\& Wiita 1993), explaining the observed lack of clearly correlated multiwavelength activity in 2004. These differences in correlated flaring activity could also come from a ``swinging'' jet which has a constant Lorentz factor, but small changes over time in the angle between the jet and the observer \\citep{gop92}. All of the observed flaring activity was consistent with the synchrotron process. An interesting future prospect would be to conduct more simultaneous multiwavelength campaigns on other blazars to see if similar patterns emerge.  During a recent campaign on fellow XBL PG 1553+11, clear flares were observed at X-ray and optical wavelengths \\citep{ost06}. The peaks of these flares appeared to be separated by about 10 days. These flares may be correlated, with the X-ray event leading the optical. If, as has been argued, PKS 2155--304 and PG 1553+11 are inherently similar objects, then the flux state of PG 1553+11 may also be related to the correlation of multiwavelength flares. Following the patterns observed in PKS 2155--304, a lag between X-ray and optical flares of as long as 10 days would indicate a relatively low flux state in PG 1553+11 during the 2003 campaign. The long-term R-band data on PKS 2155--304 and PG 1553+11 agree with these assessments of the relative flux states of each object at the observed epochs. PG 1553+11 was observed around R$=$13.5 to 13.7 during the 2003 campaign, a low state compared to almost all other observations from 2000--2006 \\citep{awc04}. The longterm PKS 2155--304 light-curve also indicates that this object was brighter during the 1991 campaign than in the 1994 campaign, and the 1994 state was brighter than the state during the 2004 campaign.  Additional simultaneous campaigns are needed on other blazars to see if similar patterns emerge in the multiwavelength flaring activity. Also, future campaigns should be performed on PKS 2155--304 and PG 1553+11 during different flux states to see if the pattern found for PKS 2155--304 continues to be observed in these objects. These campaigns would benefit from broader wavelength coverage, particularly in the TeV regime, which would include IC emission from PKS 2155--304, PG 1553+11, and similar objects. The \\emph{GLAST} observatory, scheduled to launch at the end of 2007, would provide valuable $\\gamma$-ray regime coverage to similar campaigns on LBLs and FSRQs. Subsequent campaigns would also benefit from more extensive temporal coverage."
    },
    "0711/0711.2448_arXiv.txt": {
        "abstract": "A standard disk around an accreting black hole may become effectively optically-thin and scattering dominated in the inner region, for high accretion rates (as already predicted by the Shakura-Sunyaev model). Radiative emission from that region is less efficient than blackbody emission, leading to an increase of the colour temperature in the inner region, by an order of magnitude above the effective temperature. We show that the integrated spectrum has a power-law-like shape in the $\\sim 1$--$5$ keV band, with a soft excess at lower energies and a downward curvature or break at higher energies, in agreement with the observed spectra of many ULXs. This scenario offers a physical explanation for the phenomenological ``dense corona'', successfully used to fit those spectra. It also relates their fit parameters to physical properties of the accreting black hole, such as mass and accretion rate. ",
        "introduction": "Fitting colour temperatures to the X-ray spectra of accreting black holes (BHs) has traditionally been a useful, indirect method to estimate BH masses, typically within a factor of 2. This method relies on the assumption that the accretion disk extends down to the innermost stable circular orbit and is optically-thick, emitting a multicolour-blackbody spectrum (Shakura \\& Sunyaev 1973; Frank, King \\& Raine 2002). If so, the fitted peak temperature and luminosity give a characteristic inner-disk area and hence BH mass. These conditions are usually verified in Galactic BHs in their high-soft state ($0.01 \\la L_{\\rm X}/L_{\\rm Edd} \\la 1$).  For ultraluminous X-ray sources (ULXs), the scenario that the disk is optically thick and reaches the innermost stable circular orbit could be a plausible initial assumption, in the absence of dynamic mass determinations. From a characteristic feature in the soft spectrum (``soft excess''), peak disk temperatures $\\approx 0.15$--$0.20$ keV have been inferred and used to suggest masses $\\sim 1000 M_{\\odot}$ (Miller, Fabian \\& Miller 2004). A strong objection against this argument is that most of the X-ray flux in ULXs comes out in a harder, power-law-like component, more consistent with inverse-Compton scattering (Stobbart, Roberts \\& Wilms~2006; Winter, Mushotzky \\& Reynolds~2006). This suggests that the soft thermal component only traces emission from the outer, optically-thick disk, at radii $R \\ga 10 R_{\\rm ISCO}$; if so, its temperature and luminosity cannot be directly used to estimate the BH mass. An additional parameter is needed to model the location of the transition radius between the outer, standard disk and the inner, scattering-dominated region. In general, based on a variety of physical arguments, we expect that the transition radius $R_{\\rm c} \\propto \\dot{m}^\\beta R_{\\rm ISCO}$ for $\\dot{m} > 1$, where $\\dot{m} \\equiv \\dot{M}/\\dot{M}_{\\rm Edd}$ and $\\beta \\sim 1$. ",
        "conclusions": "We have shown that the accretion disk in BHs with $M \\ga 10 M_{\\odot}$, $\\dot{M} \\sim$ a few $\\dot{M}_{\\rm Edd}$ is expected to develop a scattering-dominated, effectively-thin inner region ($R \\la 10 \\, R_{\\rm ISCO}$), with a corresponding increse in the colour temperature, to compensate for the decreased emission efficiency. The predicted spectrum is very similar to what is observed from many ULXs at $L_{\\rm X} \\sim 10^{40}$ erg s$^{-1}$. Phenomenologically, those spectra have been successfully fitted with an optically-thick disk-blackbody emission from the outer disk, plus a dense scattering corona in the inner region, ``covering'' the inner disk. We have shown here that the characteristic scattering opacity, size and temperature observationally inferred for those ``coronae'' are the same as predicted for a radiation-pressure dominated standard disk, for accretion rates and BH masses consistent with ULX parameters. So, we argue that there is no need to postulate an ad-hoc corona covering or replacing the inner disk. It is the standard disk itself that becomes a hot, scattering-dominated environment at small radii, and produces the observed spectra. The advantage of seeing it as the inner disk rather than a corona is that we do not have to worry about how energy is transferred from one to the other, or how their interface looks like. Moreover, in the scattering disk scenario, the peak temperature, optical depth and emitted flux are not just empirical parameters: they are predictable functions of $M$, $\\dot{M}$, $\\alpha$, known from standard-disk models. This provides the tools to determine the physical properties of the accreting BHs in ULXs from their fitted spectral parameters, and relate their spectral evolution to changes in the accretion rate. In this scenario, both the cool ($T \\sim 0.2$ keV) and the hot component ($T \\ga 5$ keV) come from the disk."
    },
    "0711/0711.4164_arXiv.txt": {
        "abstract": "We present Spitzer IRAC images that indicate the presence of cavities cut into the dense outer envelope surrounding very young pre-main sequence stars. These young stellar objects (YSOs) characterized by an outflow represent the earliest stages of star formation. Mid-infrared photons thermally created by the central protostar/disk are scattered by dust particles within the outflow cavity itself into the line of sight. We observed this scattered light from 27 nearby, cavity-resolved YSOs, and quantified the shape of the outflow cavities. Using the grid models of \\citet{rob06}, we matched model spectral energy distributions (SEDs) to the observed SEDs of the 27 cataloged YSOs using photometry from IRAC, MIPS, and IRAS. This allows for the estimation of geometric and physical properties such as inclination angle, cavity density, and accretion rate. By using the relative parameter estimates determined by the models, we are able to deduce an evolutionary picture for outflows. Our work supports the concept that cavities widen with time, beginning as a thin jet-like outflow that widens to reveal the central protostar and disk until the protostellar envelope is completely dispersed by outflow and accretion. ",
        "introduction": "In even their earliest stages of formation, young stellar objects (YSOs) are known to produce outflows that perturb their surrounding parental gas cloud \\citep[e.g.][]{shuetal87}. Due to extreme obscuration, for the youngest, most embedded sources, these outflows may be one of the best means of probing the physical conditions of star formation. An outflow can have a profound effect on its parent cloud as jets and prestellar winds carve cavities and inject energy and momentum into the surrounding medium \\citep[e.g.][]{shu00,bal07}. The presence of an outflow can be influential enough to disrupt neighboring cloud cores through turbulent motions \\citep{bal96}. With protostellar winds that originate close to the surface of the forming star \\citep[e.g.][]{shu95,sha07,pud07}, outflows may not only have an effect on the surrounding material but on the final properties of the newborn star as well. Outflows clear out gas and dust from the protostellar surroundings, resulting in the termination of the infall stage \\citep{vel98} and affecting the star formation efficiency of the cloud \\citep{mat00}.  A general evolutionary picture of star formation has been formed over the years to be consistent with observations of young stars \\citep{lad84,ada86,ada87}. The current picture of star formation begins with the densest parts of a molecular cloud gravitationally and quasi-statically contracting to form a dense prestellar core. Collapse can only take place once gravity has overtaken internal turbulent and/or magnetic support within the collapsing cloud region. As the core continues to gravitationally collapse, a central protostar forms surrounded by an infalling envelope and self-supported protostellar disk. Initially, the central object is completely shielded by the dense envelope, which functions as the primary mass reservoir for the accreting star. Over time, the surrounding envelope loses mass to both infall onto the central object/disk and outflow into the surrounding interstellar medium. Eventually, the central young star and prestellar disk are revealed once the envelope is completely dispersed by both outflow and accretion.  Because this evolutionary sequence is fluid, identifying specific evolutionary stages can be difficult. The most popular categorization system is a class scheme (Class 0, I, II, III) \\citep{lad87,and93}. Class 0, the earliest evolutionary stage, is characterized by the densest infalling envelope and an outflow. This outflow is often times jet-like or collimated in nature $-$ having a very small opening angle. More evolved sources (Class I) have larger outflows possibly consisting of a jet and a broader protostellar wind outflow component \\citep{arc06,shu00}. The surrounding envelope's density begins to decrease as mass is lost to the central protostar or is swept away by the outflow. By Class II, the protostellar disk has grown in mass, and the outflow cavity has greatly widened and may no longer have a definite structure \\citep{bac96}. The most evolved stage of YSO, a Class III object consists solely of a T Tauri, pre-main sequence, star and either an optically thin disk or no disk at all. Some Class III objects may still possess an optical bipolar jet \\citep[e.g.][]{bal07}.  \\subsection{YSO Outflows}  The presence of an outflow is a fundamental property of young stellar objects; some of the earliest observations revealed that outflows are common around young forming stars \\citep[e.g.][]{bal83,edw82,edw83,edw84}. YSOs develop outflows during the earliest stages of evolution; indeed, Class 0 sources are characterized by outflows \\citep{and93}. The cavities carved-out by outflows are bipolar \\citep{shu87}, signaling the presence of bipolar winds or jets. Bipolar jets or collimated outflows are a ubiquitous phenomenon of astronomical objects involving accretion, rotation, and magnetic fields \\citep[e.g.][]{bal07}. Initially during the infall phase, the inflowing material acts to suppress the stellar outflow. Because the total column is the most depleted at the poles \\citep{ter84} and the envelope is rotationally flattened, breakout will first occur at the poles \\citep{shu87}. Combined with the presence of a thick protostellar disk, these physical characteristics lead to outflows with a symmetry about the poles.  Although never conclusively demonstrated with a very large sample of sources, it is expected that the cavity opening angles must widen over time \\citep{shu87,arc06,tob07}. A forming star's evolution, which begins with a heavily embedded protostar and an outflow, must ultimately result in the presence of a protostar and disk without the embedding envelope. Today there is no singular generally accepted explanation for outflow cavity widening. It may be that jet axis precession could produce wider cavities over time \\citep{mas93}. However, despite ample evidence of axis wandering \\citep{rei97,ter99}, in most cases, the angle of the precession is too small to account for the full extent of the cavity widening \\citep{yu99,rei00,arc04}. It has also been suggested that the outflow cavity is formed by both a collimated wind or jet and a broader wide-angle protostellar wind \\citep{kwa95,kon00,shu00,mat03}. The X-wind model predicts a high density protostellar wind along the axis, with a decreasing density with increasing angle from the axis \\citep{sha98}. According to this model, in very early stages (young Class 0 sources), only the dense, high-velocity, collimated component of the outflow can puncture the surrounding high density envelope \\citep{wil03}. Over time, the envelope loses mass, allowing the weaker, less dense, and lower velocity broad-angle component to carve a wider cavity into the envelope. Combined with the entraining of the cavity's edge material out of the envelope, this two-component wind can lead to the widening of the outflow. Additionally, since the accretion rate lessens over time, the decrease in infall eventually allows the outflow to dominate over the infall. This could allow the wind to carve out wider cavities as the system evolves \\citep{del00}.  \\subsection{Procedure}  YSOs emit most of their radiation at wavelengths longward of the near-infrared \\citep[e.g.][]{ada87}. The primary radiation source is thermal in nature, emanating from the envelope, circumstellar disk, and central protostar. Due to intense scattering and extinction in the dense envelope, direct imaging of the envelope is only possible in the millimeter regime. On the other hand, the incredible sensitivity of the Spitzer Space Telescope allows us to observe through the dense envelope enshrouding the young star to the scattered light within the outflow cavities. In the infrared, the outflow cavity is the most prominent feature \\citep[e.g.][]{tob07}. Dust particles within the cavity scatter stellar and disk radiation into the line of sight, releasing light that would be completely obscured without the presence of the cavity. The shape and structure of the outflow can be determined by imaging this scattered light in the mid-IR. In this paper, we use archival \\textit{Spitzer Space Telescope} data from 27 nearby YSOs to determine the opening angle of the outflow cavities. In order to apply the measured opening angles to an evolutionary scheme, we match observed photometry data to pre-computed model Spectral Energy Distributions (SEDs) of 20,000 YSOs at varying stages of prestellar evolution. We inserted mid-IR and far-IR data into a web-based\\footnote{http://caravan.astro.wisc.edu/protostars/} SED fitting program \\citep{rob07} to find the best fit models for the 27 nearby YSOs. Using the best-fit models, we identified evolutionary trends in the physical structure of young stellar objects and their surrounding environments. ",
        "conclusions": "Our thorough search through nearby star forming regions has resulted in a sample of resolved YSOs and their companion outflow cavities that enables a detailed study of the morphological evolutionary effects to outflows. Estimated ages of the catalog sources have been explored through the use of IR-color and SED fitting to pre-computed models. The latter allows an extrapolation of relative physical parameters such as cavity density and accretion rate.  The mid-IR images provided by the IRAC camera give evidence to the presence of a cavity that has been cleared of its dense envelope material by an outflow. There appears to be an evolutionary trend in the morphology of the outflow cavities. The sources estimated to be the youngest by SED fitting and color possess the narrowest cavities. More evolved sources generally have larger, more open cavitites, indicating a widening of outflow cavities with time. Should our conclusions be correct that outlfow cavity angles widen over time, our work should provdie confindence that the 2-D radiation tranasfer models do, to first order, accuratley depict the structure and physics associated with star formation. Outflows are clearly affected by the aging of young pre-main sequence stars, and it is likely they play a major role in the star formation process. A detailed knowledge of the morphological changes to a cavity should produce insights into the outflow's creation and evolution, ultimately providing more accurate theories of star formation."
    },
    "0711/0711.4352_arXiv.txt": {
        "abstract": "As galaxy surveys become larger and more complex, keeping track of the completeness, magnitude limit, and other survey parameters as a function of direction on the sky becomes an increasingly challenging computational task. For example, typical angular masks of the Sloan Digital Sky Survey contain about $N=300{,}000$ distinct spherical polygons. Managing masks with such large numbers of polygons becomes intractably slow, particularly for tasks that run in time $\\mathcal{O}\\left(N^{2}\\right)$ with a naive algorithm, such as finding which polygons overlap each other. Here we present a {}``divide-and-conquer'' solution to this challenge: we first split the angular mask into predefined regions called {}``pixels,'' such that each polygon is in only one pixel, and then perform further computations, such as checking for overlap, on the polygons within each pixel separately. This reduces $\\mathcal{O}\\left(N^{2}\\right)$ tasks to $\\mathcal{O}\\left(N\\right)$, and also reduces the important task of determining in which polygon(s) a point on the sky lies from $\\mathcal{O}\\left(N\\right)$ to $\\mathcal{O}\\left(1\\right)$, resulting in significant computational speedup. Additionally, we present a method to efficiently convert any angular mask to and from the popular \\emph{HEALPix} format. This method can be generically applied to convert to and from any desired spherical pixelization. We have implemented these techniques in a new version of the \\noun{mangle} software package, which is freely available at \\texttt{\\url{http://space.mit.edu/home/tegmark/mangle/}}, along with complete documentation and example applications. These new methods should prove quite useful to the astronomical community, and since \\noun{mangle} is a generic tool for managing angular masks on a sphere, it has the potential to benefit terrestrial mapmaking applications as well. ",
        "introduction": "\\begin{figure*} \\includegraphics{fig1}   \\caption{\\label{fig:speed_fractions}Speed trials for a series of portions of the SDSS DR5 mask with and without pixelization. (a) Time required for pixelization, snapping, balkanization and unification of the mask. (b) Time required to identify in which polygon each of the ${\\sim}400{,}000$ SDSS DR5 galaxies lies. Each set of trials is fitted with a power law to show how the processing time scales with the number of polygons $N$. Also shown on the x-axis are the number of polygons in the 2dFGRS mask, the SDSS DR5 mask, and conservative estimates for the Pan-STARRS and LSST large-scale structure masks based on scaling up 2dFGRS.} \\end{figure*} Over the past few decades, galaxy surveys have provided a wealth of information about the large-scale structure of our Universe, and the next generation of surveys currently being planned promises to provide even more insight. In order to realize the full potential of upcoming surveys, it is essential to avoid unnecessary errors and approximations in the way they are analysed. The tremendous volumes of data produced by these new surveys will shrink statistical uncertainty to unprecedented levels, and in order to take advantage of this we must ensure that the systematic uncertainties can keep pace. The purpose of this paper is to maximize the scientific utility of next-generation surveys by providing methods for processing angular masks as rapidly and accurately as possible.  Angular masks of a galaxy survey are functions of direction on the sky that model the survey completeness, magnitude limit, seeing, dust extinction, or other parameters that vary across the sky. The earliest galaxy redshift surveys -- the first Center for Astrophysics redshift survey (CfA1; \\citealt{1983ApJS...52...89H}) and the first Southern Sky Redshift Survey (SSRS1; \\citealt{1991ApJS...75..935D}) -- had simple angular masks defined by boundaries in declination and Galactic latitude. The next generation of surveys -- IRAS \\citep{1992ApJS...83...29S} and PCSz \\citep{2000MNRAS.317...55S} -- had somewhat more complex masks, with some regions of high contamination excluded from the survey.  The present generation of surveys -- the Two Degree Field Galaxy Redshift Survey (2dFGRS; \\citealt{2001MNRAS.328.1039C,2003astro.ph..6581C}) and the Sloan Digital Sky Survey (SDSS; \\citealt{2000AJ....120.1579Y}) -- consist of photometric surveys that identify galaxies and measure their angular positions combined with spectroscopic surveys that measure a redshift for each galaxy to determine its distance from us. Angular masks are useful for describing parameters for both photometric and spectroscopic surveys -- for example, seeing and magnitude limit are key parameters to model in photometric surveys, and the survey completeness -- i.e., the fraction of photometrically selected target galaxies for which a spectrum has been measured -- is vital for analysing spectroscopic surveys.  The angular masks of SDSS and 2dFGRS consist of circular fields defined by the spectroscopic plates of the redshift survey superimposed on an angular mask of the parent photometric survey. 2dFGRS uses the Automatic Plate Measurement (APM) survey (Maddox et al. 1990a,b, 1996) \\nocite{1990MNRAS.243..692M,1990MNRAS.246..433M,1996MNRAS.283.1227M}% as its parent photometric survey and covers approximately 1500 square degrees. The APM angular mask consists of 269 $5^{\\circ}\\times5^{\\circ}$ photographic plates which are drilled with numerous holes to avoid bright stars, satellite trails, plate defects, and so forth. Combining the photographic plates, holes, and spectroscopic fields gives a total of 3525 polygons that define the spectroscopic angular mask of 2dFGRS.  The SDSS covers a larger area on the sky -- 5740 square degrees in Data Release 5 (DR\\%; \\citealt{2007ApJS..172..634A}) -- and has a yet more complicated angular mask than 2dFGRS. The SDSS photometric survey is done by drift-scanning: each scan across the sky covers six long, narrow scanlines, and the gaps between these lines are filled in with a second scan slightly offset from the first, producing a {}``stripe'' about $2.5^{\\circ}$ wide assembled from 12 scanlines. In addition to the fairly intricate pattern produced by this scanning strategy, there are nearly 250,000 holes masked out of the photometric survey for various reasons, plus the circular $3^{\\circ}$ spectroscopic fields. Combining all of these elements produces an angular mask for the spectroscopic survey that contains 340,351 polygons.  To accurately manage the 2dFGRS and SDSS angular masks, \\citet{2004MNRAS.349..115H} developed a suite of general-purpose software called \\noun{mangle}, which performs several important procedures on angular masks using computational methods detailed in \\citet{1993ApJ...406L..47H,1993ApJ...417...19H}. This software has proved to be a valuable resource to the astronomical community: it has been used in several analyses of galaxy survey data \\citep{2002MNRAS.335..887T,2004ApJ...606..702T,2005MNRAS.361.1287M,2005PASJ...57..709H,2005ApJ...633...11P,2005ApJ...635..990C,2007ApJ...658..898P,2007PASJ...59...93N,2007AJ....133.2222S,2007ApJ...664..608W,2007arXiv0707.3445T}. Additionally, it was used extensively in the preparation of the New York University Value-Added Galaxy Catalog (NYU-VAGC; \\citealt{2005AJ....129.2562B}), which has been used as the basis for almost all publications on large-scale structure by the SDSS collaboration.  However, many of functions in the original version of \\noun{mangle} run in $\\mathcal{O}\\left(N^{2}\\right)$ time, which becomes quite computationally challenging as the size and complexity of surveys continues to increase -- computations involving the SDSS mask can take several months of CPU time. In this paper we present new algorithms that can process complicated angular masks such as SDSS dramatically faster with no loss of accuracy. Our method is based on splitting an angular mask into pixels, reducing the processing time to $\\mathcal{O}\\left(N\\right)$ by adding an $\\mathcal{O}\\left(N\\log N\\right)$ preprocessing step. Similar methods based on hierarchical spatial subdivisions have been found to be useful in the field of computational geometry (see, e.g., \\citealt{bigbluebook}), but have not previously been applied to angular masks in an astronomy context.  This will be especially useful for next-generation surveys, such as the Dark Energy Survey (DES; \\citealt{2005astro.ph.10346T}) and surveys done with the Wide-Field Multi-Object Spectrograph (WFMOS; \\citealt{2006PhRvD..74f3525Y,wfmos}), the Panoramic Survey Telescope and Rapid Response System (Pan-STARRS; \\citealt{2004SPIE.5489...11K}), and the Large Synoptic Survey Telescope (LSST; \\citealt{2002SPIE.4836...10T,2004AAS...20510802S,2006AIPC..870...44T,LSST}). DES, Pan-STARRS, and LSST will perform photometric surveys and use techniques for estimating redshifts based on the photometric information; WFMOS will perform spectroscopic surveys using one of the upcoming photometric surveys for target selection. The methods we present here are useful for both photometric and spectroscopic surveys -- keeping track of factors such as seeing and dust extinction could prove particularly important for photometric redshift determinations \\citep{2007MNRAS.375...68C,2008ApJ...674..768O,2008MNRAS.386.1219B}.  \\begin{figure*} \\includegraphics{fig2a}\\includegraphics{fig2b}   \\caption{\\label{fig:examplemask}\\label{fig:examplemask_final}A portion of the SDSS DR5 angular mask \\citep{2005AJ....129.2562B,2007ApJS..172..634A}. Left: Polygons defining the mask: spectroscopic plates and lines delineating different scans and spectroscopic plates are shown in black, and holes in the mask are shown in blue/gray. Right: Processed version of the mask, shaded according to survey completeness.} \\end{figure*} \\begin{table*} \\begin{minipage}[c][70mm]{175mm}%  \\caption{\\label{tab:mangleterms}Definitions of Terms, in Alphabetical Order}  \\begin{tabular}{@{}lp{155mm}} Term & Definition \\tabularnewline \\hline boundary & A set of edges bounding a polygon. \\tabularnewline cap & A spherical disk, a region above a circle on the unit sphere. \\tabularnewline circle & A line of constant latitude with respect to some arbitrary polar axis on the unit sphere. \\tabularnewline edge & An edge is part of a circle. A polygon is enclosed by its edges. \\tabularnewline great circle & A line of zero latitude with respect to some arbitrary polar axis on the unit sphere. A great circle is a circle, but a circle is not necessarily a great circle. \\tabularnewline mask & The union of an arbitrary number of weighted polygons. \\tabularnewline pixel& A special polygon that specifies some predefined region on the sky as part of a scheme for discretizing the unit sphere. Once a mask has been pixelized, each polygon in the mask is guaranteed to overlap with exactly one pixel.\\tabularnewline polygon & The intersection of an arbitrary number of caps. \\tabularnewline rectangle & A special kind of polygon, a rectangular polygon bounded by lines of constant longitude and latitude. \\tabularnewline vertex & A point of intersection of two circles. A vertex of a polygon is a point where two of its edges meet. \\tabularnewline weight & The weight assigned to a polygon. The spherical harmonics of a mask are the sum of the spherical harmonics of its polygons, each weighted according to its weight. A weight of 1 is the usual weight. A weight of 0 signifies an empty polygon, a hole. In general the weight may be some arbitrary positive or negative real number. \\tabularnewline \\hline \\end{tabular}% \\end{minipage}%  \\end{table*} The proposed large-scale structure survey to be produced by Pan-STARRS will cover ${\\sim}30{,}000$ square degrees in 5 wavelength bands, with each field being observed ${\\sim}50$ times such that the images can be co-added. A naive scaling up of the 2dFGRS area and number of polygons gives an estimate of ${\\sim}2\\times10^7$ polygons for the final Pan-STARRS mask. Similarly, the LSST large-scale structure survey will cover ${\\sim}20{,}000$ square degrees in 6 bands, with ${\\sim}200$ co-added images, suggesting ${\\sim}8\\times10^7$ polygons for the LSST mask. The need for an improvement in mask processing speed is clearly illustrated in Fig.~\\ref{fig:speed_fractions}: with the old algorithms, the projected processing time for the LSST mask would be over 6000 years. With our new method, this time is reduced by a factor of ${\\sim}24{,}000$ to just ten days.  \\begin{figure*} \\includegraphics{fig3}   \\caption{\\label{fig:pix_concept}A cartoon illustrating the process of balkanization (a) with no pixelization, (b) with pixelization.} \\end{figure*} In addition to developing faster algorithms for processing angular masks, we have also integrated the \\noun{mangle} utilities with \\emph{HEALPix}, a widely used tool for discretizing the celestial sphere \\citep{2005ApJ...622..759G}. The methods used by \\noun{mangle} are complementary to \\emph{HEALPix}: \\noun{mangle} is best used for functions that are piecewise-constant in distinct regions of the sky, such as the completeness of a galaxy survey. In contrast, \\emph{HEALPix} is optimal for describing functions that are continuously varying across the sky, such as the cosmic microwave background (CMB) or the amount of extinction due to Galactic dust. The ability to convert rapidly between these two formats allows for easy comparison of these two types of data without the unnecessary approximation inherent in discretizing an angular mask. Furthermore, converting a mask into \\emph{HEALPix} format allows users to take advantage of pre-existing \\emph{HEALPix} tools for rapidly computing spherical harmonics.  The spectacular surveys on the horizon are preparing to generate massive, powerful datasets that will be made publically available -- this in turn necessitates powerful and intuitive general-purpose tools that assist the community to do science with this avalanche of data. We provide a such tools with this new generation of the \\noun{mangle} software and describe these new tools here. However, this paper is not a software manual (a manual is provided on the \\noun{mangle} website) but rather a description of the underlying algorithms. These tools have been utilized in recent analyses of SDSS data \\citep{2006PhRvD..74l3507T,2008MNRAS.385.1635S}; we are now making them public so others can use them as well.  The outline of this paper is as follows: in \\S\\ref{sec:Mangle-terminology} we give an overview of the terminology we use to describe angular masks and the basic tasks we wish to perform, and in \\S\\ref{sec:Speedup:-pixelization} we detail our algorithms for accelerating these tasks and quantify their speed. We describe our methods for integrating \\noun{mangle} with \\emph{HEALPix} in \\S\\ref{sec:Unification-with-HEALPix} and summarize in \\S\\ref{sec:Summary}. ",
        "conclusions": "\\label{sec:Summary}  As technologies for surveying the sky continue to improve, the process of managing the angular masks of galaxy surveys grows ever more complicated. The primary purpose of this paper has been to present a set of dramatically faster algorithms for completing these tasks.  These algorithms are based on dividing the sky into regions called {}``pixels'' and performing key operations only within each pixel rather than across the entire sky. The pixelization is based on a hierarchical subdivision of the sky -- this produces a quadtree data structure that keeps track of which polygons are nearby each other. The preprocessing step of pixelization is $\\mathcal{O}\\left(N\\log N\\right)$ for a mask with $N$ polygons and ${\\sim}N$ pixels, and it reduces the mask processing time from $\\mathcal{O}\\left(N^{2}\\right)$ to $\\mathcal{O}\\left(N\\right)$. Furthermore, it reduces the time required to locate a point within a polygon from $\\mathcal{O}\\left(N\\right)$ to $\\mathcal{O}\\left(1\\right)$.  This method is exact, i.e., it does not make a discrete approximation to the mask, and it takes only a tiny fraction of the computation time. It accelerates the processing of the SDSS mask by a factor of about 1200 and reduces the time to locate all of the galaxies in the SDSS mask by a factor of nearly 500. It will provide even more dramatic gains for future surveys: processing time for the LSST large-scale structure mask could be reduced by a factor of about 24,000.  We have also described a method for converting between masks described by spherical polygons and sky maps in the \\emph{HEALPix} format commonly used by CMB and large-scale structure experiments. This provides a convenient way to work with both piecewise-constant functions on the sky such as the completeness of a galaxy survey and continuously varying sky maps such as the CMB temperature. Converting angular masks into \\emph{HEALPix} format also allows users to take advantage of existing \\emph{HEALPix} tools for rapidly computing spherical harmonics.  All of the new algorithms and features detailed here have been integrated into the \\noun{mangle} software suite, which is available for free download at \\texttt{\\url{http://space.mit.edu/home/tegmark/mangle/}}. This updated software package should prove increasingly useful in the coming years, especially as next-generation surveys such as DES, WFMOS, Pan-STARRS, and LSST get underway."
    },
    "0711/0711.0601_arXiv.txt": {
        "abstract": "{Photometry of Procyon obtained by the MOST satellite in 2004 has been searched for p~modes by several groups, with sometimes contradictory interpretations.} {We explore two possible factors that complicate the analysis and may lead to erroneous reports of p~modes in these data.} {Two methods are used to illustrate the role of subtle instrumental effects in the photometry: time-frequency analysis, and a search for regularly spaced peaks in a Fourier spectrum based on the echelle diagramme approach.} {We find no convincing evidence of a p-mode signal in the MOST Procyon data. We can account for an apparent excess of power close to the p-mode frequency range and signs of structure in an echelle diagramme in terms of % instrumental effects.} {} ",
        "introduction": "The search for acoustic oscillations in the star Procyon A (which we call hereafter simply Procyon) has been a key part of the young field of stellar seismology for two decades, since the first claim of p-mode detection by \\citet{bernard86}. In recent years, interest grew rapidly with the announcement by \\citet{martic99} of global acoustic modes in Procyon detected in high-resolution, time-resolved spectroscopy, confirmed by \\citet{caro}.  Interest turned to controversy with the non-detection of intensity oscillations by the MOST (Microvariability \\& Oscillations of STars) space-borne photometer \\citep{jaymie}, while contemporaneous groundbased spectroscopic campaigns by \\citet{martic04}, \\citet{eggen04} and \\citet{bouchy04} appeared to confirm the presence of stellar oscillations. This non-detection triggered considerable speculation about non-stellar noise sources in the MOST data \\citep{jcd04,bedding05} and independent analyses of those data claiming possible evidence for p~modes, but using different analysis methods \\citep{clara,garcia05}. The contrast between the reported p-mode signal in luminosity and integrated radial velocity also inspired theorists to refine models of convection and turbulence in Procyon \\citep[e.g.,][]{robinson2005}.  The MOST team has described the nature of the Fabry Imaging data on which the Procyon photometry is based \\citep{reegen2006}, including the recognised effects of stray light due to Earthshine modulated by the MOST satellite orbit and a beating effect of the science and startracker CCD timing in the early stages of the mission.  We present an independent analysis of the MOST Procyon data compared to simulations, calling attention to how the timing effects and orbital artifacts can be misinterpreted as evidence for p-mode oscillations.  \\begin{figure*} \\centering \\hbox{\\includegraphics[height=9cm,angle=90]{procyon_to.ps} \\includegraphics[height=9cm,angle=90]{t_ccd.ps}} \\caption{Left: Time/frequency diagramme of Procyon data, showing excess power between 500 and  1200~$\\mu$Hz during the first eight days of observation. Right: same analysis of the CCD temperature, showing the same excess of power} \\label{fig:tempsfreq} \\end{figure*} ",
        "conclusions": "From the two approaches presented here, one can draw two conclusions. First, the Procyon MOST light curve obtained in 2004 shows some changes in the spectral power content after about seven days of observation (which were the first seven days of observation of the MOST mission) that covers the frequency range relevant to the p-mode search and thus requires special care in this search. A counterpart to this change in the CCD temperature telemetry indicates that it is not intrinsic to Procyon.\\\\ Second, the echelle-diagramme technique must be applied to Procyon MOST data with care. Despite an excellent observation window, the orbital artifacts in the data introduce apparent spacings close to the p-mode spacing expected for Procyon even though the gaps are very short and their signature in the window spectrum is very weak. However, they are unexpectedly not negligible when searching for regularly--spaced peaks in the Fourier spectrum, reinforcing the need for a careful analysis when looking for regular spacings in the signal spectrum even in the case of a very good observation window.\\\\ In conclusion, we do not find any evidence for a stellar oscillation signal in the MOST photometry, and the issue of the nature of luminosity oscillations in Procyon remains open."
    },
    "0711/0711.2929_arXiv.txt": {
        "abstract": "Astrophysical observations originate from matter that interacts with radiation or transported particles. We develop a pragmatic approximation in order to enable multi-dimensional simulations with basic spectral radiative transfer when the available computational resources are not sufficient to solve the complete Boltzmann transport equation. The distribution function of the transported particles is decomposed into a trapped particle component and a streaming particle component. Their separate evolution equations are coupled by a source term that converts trapped particles into streaming particles. We determine this source term by requiring the correct diffusion limit for the evolution of the trapped particle component. For a smooth transition to the free streaming regime, this 'diffusion source' is limited by the matter emissivity. The resulting streaming particle emission rates are integrated over space to obtain the streaming particle flux. Finally, a geometric estimate of the flux factor is used to convert the particle flux to the streaming particle density, which enters the evaluation of streaming particle-matter interactions. The efficiency of the scheme results from the freedom to use different approximations for each particle component. In supernovae, for example, reactions with trapped particles on fast time scales establish equilibria that reduce the number of primitive variables required to evolve the trapped particle component. On the other hand, a stationary-state approximation considerably facilitates the treatment of the streaming particle component. Different approximations may apply in applications to stellar atmospheres, star formation, or cosmological radiative transfer. We compare the isotropic diffusion source approximation with Boltzmann neutrino transport of electron flavour neutrinos in spherically symmetric supernova models and find good agreement. An extension of the scheme to the multi-dimensional case is also discussed. ",
        "introduction": "Most applications in computational astrophysics involve matter that is in thermal or reactive equilibrium. Astrophysical observations, however, require information to propagate from the astrophysical site to a terrestrial observer. Hence, the observationally interesting events involve additional particle species that are not trapped in the fluid and must therefore be treated by radiative transfer. In many traditional numerical models, the equilibrated matter is treated in the hydrodynamic limit, while transported particle species are treated by a suitable algorithm of radiative transfer. If this splitting in the numerical method is based on the particle species, it can become extremely challenging and inefficient if the range of conditions in the astrophysical scenario is large. In situations where radiative particles are in thermal or reactive equilibrium, the non-local algorithm of radiative transfer has to be perfectly consistent with the hydrodynamics scheme and capable of handling very stiff source terms in order to evolve the applicable equilibria.  For example core collapse supernova events have challenged computer models for several decades \\citep{Colgate.White:1966,Arnett:1966,Bowers.Wilson:1982,Bethe.Wilson:1985,Bruenn:1985,Myra.Bludman:1989}. These stellar explosions occur at the end of the life of a massive star when the growing iron core, as the end product of nuclear fusion, becomes unstable against gravitational collapse. The collapse is halted only after nuclear density is reached at the centre. The repulsive strong interaction and the neutron degeneracy pressure reduce the compressibility of matter so that a sound wave travels outward through the causally connected inner core. It turns into a shock wave when it reaches supersonically accreting matter from the outer layers and, due to dissociation and neutrino losses, stalls within few milliseconds to an initially hydrostatic accretion front, which is thought to slowly expand over several hundreds of milliseconds before the explosion sets in. The new-born protoneutron star (PNS) at the centre of the event has a larger density by many orders of magnitude than the shock-heated and dissociated matter that continues to accumulate behind the accretion front. This accreted matter is subject to a variety of fluid instabilities and asymmetric flow patterns that couple to the energetic neutrino radiation field produced within or close to the surface of the PNS. The emergence of an explosion and the detailed interaction between the multi-dimensional flow and the neutrino field is subject to a long-standing and ongoing debate. See for example \\citep{Bethe:1990} for a theoretical description of the scenario or \\citep{Bruenn.Hix.ea:2006,Dessart.Burrows.ea:2007,Marek.Janka:2007} for recent axisymmetric computer simulations with spectral neutrino transport.  Several aspects of the core collapse supernova scenario make computer models difficult to perform: The relevant densities in the computational domain range over ten decades from about \\( 10^{15} \\) g/cm\\( ^{3} \\) at the centre of the PNS down to about \\( 10^{5} \\) g/cm\\( ^{3} \\) in the outermost layers (the star is much larger but does not dynamically react to the changes at the centre during the first few hundred milliseconds after bounce). This density range implies very disparate dynamical time scales. At high densities, the weak interaction rates are much faster than the dynamical time scale, while at low density, they are even slower than the evolution time scale. This requires implicit finite differencing of many processes to perform simulations in a reasonable amount of time. Moreover, the gravitational and stored internal energy scales are of order \\( 10^{53} \\) erg, while typical kinetic explosion energies are two orders of magnitude smaller. The presence of a strong shock with near-relativistic infall velocities further challenges the resolution and stability of numerical models. In this very dynamical environment with multi-dimensional fluid instabilities and differential rotation, a computer model has to quantify the exchange of energy between matter and the emitted neutrino field in order to allow statements about the explosion dynamics to be made. As the neutrino cross sections at given matter conditions scale with the square of the neutrino energy, the neutrino transport problem must be considered as a superposition of several single-energy (monochromatic) transport problems, where the geometry of the scattering sphere depends on the neutrino energy and type. This leads to multi-group transport schemes where all propagation quantities are solved in separate energy groups \\citep{Mazurek:1975,Arnett:1977}. Moreover, stationary-state studies with spherically symmetric Boltzmann neutrino transport showed that the angular distribution of the neutrino propagation directions have also to be taken into acount to obtain accurate neutrino heating and cooling rates \\citep{Wilson:1971,Mezzacappa.Bruenn:1993a,Messer.Mezzacappa.ea:1998,Yamada.Janka.ea:1999,Burrows.Young.ea:2000}.  In the mean time, the spherically symmetric Boltzmann transport equation is routinely solved for all neutrino flavours in dynamical models of stellar core collapse and postbounce evolution \\citep{Liebendoerfer.Mezzacappa.ea:2001,Rampp.Janka:2002,Thompson.Burrows.Pinto:2003,Sumiyoshi.Yamada.ea:2005}. These comprehensive and sophisticated codes treat disparate time scales by implicit finite differencing and require large computational resources. However, due to the enormous density range in the scenario, a considerable fraction of the computational domain is either neutrino-opaque or neutrino-transparent. In a fully neutrino-opaque regime, a numerical solution to the Boltzmann transport equation is rarely more accurate than the solution of the simpler diffusion equation that describes the correct physical limit. Similarly, the angular distribution of the neutrino propagation directions in the far-field free streaming regime is precisely determined by the geometry and emissivity of the distant sources, while an angular discretisation of the Boltzmann equation may lead to discretisation errors. Hence, we suggest that a combination of different approaches to radiative transfer is applied in the same astrophysical simulation. This leads to the concept of an 'adaptive algorithm', which should not be regarded as an inconsistency or drawback. On the contrary, it saves computation power that can be invested for better resolution, more input physics, or the coverage of a larger parameter space. Current state-of-the-art supernova models with spectral neutrino transport are performed under the assumption of axisymmetry. As a 2D solution of the Boltzmann equation is computationally inefficient for the reasons discussed above, different groups either apply 1D Boltzmann solutions in separate angular segments in combination with 2D neutrino advection in opaque regimes (ray-by-ray) \\citep{Marek.Janka:2007}, resort to the multi-group flux-limited diffusion (MGFLD) approximation \\citep{Walder.Burrows.ea:2004,Ott.Burrows.ea:2008}, or use a combination of both \\citep{Bruenn.Hix.ea:2006}.  Three-dimensional supernova models were hitherto only affordable with 'grey' neutrino transport \\citep{Fryer.Warren:2004} or with other simple ad hoc approximations to the neutrino physics \\citep{Scheck.Plewa.ea:2004,Ott.Dimmelmeier.ea:2007,Scheidegger.Fischer.ea:2008}. But many features of the supernova should be studied in three spatial dimensions, for example the pattern of the accretion flow \\citep{Herant.Benz.ea:1994}, the standing accretion shock instability \\citep{Blondin.Mezzacappa.DeMarino:2003}, the excitation of PNS g-modes \\citep{Burrows.Livne.ea:2006}, and the evolution of magnetic fields (see e.g. \\citet{Kotake.Sato.Takahashi:2006} and references therein). In order to enable three-dimensional supernova models with spectral neutrino transport and in order to accelerate parameter studies with simulations of lower dimensionality, we construct the isotropic diffusion source approximation based on several ideas scattered across the literature. Our algorithm is not meant to compete in accuracy with more detailed solutions of the transport equation in the semi-transparent regime. As the transition from the opaque to the transparent regime will be handled by interpolation and basic geometric considerations, it is important that the approach is verified by an accurate transport solution for every new field of application. Our approach shares this limitation with the well-known flux-limited diffusion approximation. The main advantage of the isotropic diffusion source approximation over the latter is that the fluxes and flux factors in the transparent regime are determined by the non-local distribution of sources rather than the local intensity gradient, which may give an incorrect flux direction. Moreover, the new approach tries to limit the computationally expensive solution of the multi-dimensional diffusion equation to the opaque regime where the diffusion approximation is adequate. The goal is to create a flexible algorithm that efficiently implements only the dominant features of radiative transfer in one consistent framework.  The most obvious method of avoiding the inefficient global application of algorithms is a decomposition of the problem in space, so that one algorithm is used in one subdomain while another algorithm is used in another subdomain \\citep{Chick.Cassen:1997}. Another approach is the decomposition in the momentum phase space into particles with thermal velocities and particles with supra-thermal velocities. If one treats the thermal particles by an efficient fluid model, a hybrid kinetic/fluid model is obtained \\citep{Crouseilles.Degond.Lemou:2004}. The isotropic diffusion source approximation described in this article is more similar to the so-called \\( \\delta f \\)-method \\citep{Parker.Lee:1993,Brunner.Valeo.Krommes:1999}. We also decompose the distribution function of the transported particles into a thermal component and a perturbation that is allowed to overlap the thermal part in the entire computational domain and particle phase space. However, in our case, the perturbation is not considered to be small. Inspired by flux-limited diffusion \\citep{Levermore.Pomraning:1981}, the evolution of the radiation component is guided by the diffusion limit at large opacities and by the free streaming limit at low opacities. In contrast to flux-limited diffusion, we build our scheme conceptually on approximations of the collision integral rather than particle fluxes. We assume that the free streaming particle flux is dominated by the flux emerging from the diffusive domain so that the sources for the far field can easily be integrated. In the stationary-state limit, the determination of the particle fluxes reduces to the solution of a Poisson equation \\citep{Gnedin.Abel:2001}. Since the matter interacts according to the local particle abundances and not fluxes, we have to convert the particle fluxes to local particle densities. This is achieved by a geometric estimate of the flux factor as suggested and evaluated by Bruenn in \\citep{Liebendoerfer.Messer.ea:2004}.  In Sect. \\ref{sec:idsa} we describe in detail how these concepts enter the framework of the isotropic diffusion source approximation, which we design for the transport of massless fermions through a compressible gas. Its connection to the well-known diffusion limit is made in Appendix A. In Sect. \\ref{sec:verification}, we evaluate the performance of this approximation in comparison with Boltzmann neutrino transport in spherical symmetry. Finally, in Sect. \\ref{sec:multid}, we discuss the extension to multi-dimensional simulations. Details of the finite differencing and implementation are given in Appendix B. ",
        "conclusions": "The isotropic diffusion source approximation (IDSA) has been designed to treat radiative transfer most efficiently in astrophysical applications where particles diffuse out of an opaque domain that is subject to multi-dimensional dynamics. This scenario occurs frequently in new-born gravitationally bound objects, accretion discs, or dynamical atmospheres. We have developed and tested the IDSA in the context of spherically symmetric supernova models. The neutrino transport in core collapse supernovae provides a very challenging application where comprehensive solutions of the Boltzmann transport equation have been studied in great detail (e.g. \\citet{Mezzacappa:2005} and references therein). Our approximation treats only the most important features of radiative transfer. In a supernova model, these are (i) the thermodynamics of trapped neutrinos (for example the \\( dE_{\\nu} = -p_{\\nu}dV \\)-term), (ii) the correct diffusion limit, (iii) the spectrum of transported neutrinos, and (iv) the angular focussing of the neutrino propagation directions with increasing distance from the neutrinospheres.  We implemented the IDSA in the hydrodynamics code Agile \\citep{Liebendoerfer.Rosswog.Thielemann:2002} to compare its performance with a more sophisticated transport model. We compare the postbounce evolution of a \\( 13 \\) M\\( _{\\odot } \\) star \\citep{Nomoto.Hashimoto:1988} in Newtonian gravity with the corresponding model N13 that is based on Boltzmann neutrino transport \\citep{Liebendoerfer.Rampp.ea:2005}. We find good agreement between the results of the IDSA and the reference model during the early postbounce phase when the neutrino burst is launched and the accretion shock expands to its maximum radius. The IDSA leads to a somewhat larger maximum shock radius and a faster shock retraction thereafter, i.e. a slightly more pessimistic model with respect to the possibility of a delayed supernova explosion. We conclude that the concept of the IDSA accurately captures the feedback of neutrino transport onto the hydrodynamics evolution in spherically symmetric supernova models. The neutrino fluxes and spectra are in nice agreement if small time steps are taken, but unphysical fluctuations in the stationary-state neutrino luminosities develop when the time steps are increased much beyond the neutrino propagation time scale. We will continue to improve the numerical stability of the scheme in future work.  With respect to the choice of physical assumptions and with respect to the quality of the results, the IDSA is most closely related to the well-known flux-limited diffusion (FLD) approximation. Like the latter, the IDSA is guided by the diffusion limit at high opacities and the free streaming limit at low opacities. The FLD approximation is based on diffusive fluxes that are limited by the physical particle speed at the transition to the free streaming regime. This limiting is well-defined in a spherically symmetric setup where the limiting can be applied to the radial velocity direction. In a multi-dimensional simulation, however, FLD fails to determine the correct direction of the net flux outside the diffusive regime. In the free streaming regime there is no physical connection between the local radiation intensity and the flux direction. A related point is the efficiency of the FLD approach. Since the implemented numerical solution strategies are optimised for the diffusive domain, they easily become inefficient in the free streaming regime, which represents often the larger part of the computational domain. These are the two main shortcomings of FLD which we address with the IDSA. The IDSA is based on an overlapping two-component description. The trapped particle component dominates in the diffusive regime while the streaming component dominates in the free streaming regime. In the diffusive regime, the IDSA is equivalent to FLD. In the semi-transparent regime, the IDSA is based on a different interpolation between the diffusive and the free streaming regimes. In the IDSA, the limiter is applied to the sources instead of the fluxes. In the free streaming limit of a multi-dimensional application, the direction of fluxes is well-defined by the non-local position of the sources. The distinction between the trapped and streaming particle components allows us to separately test more efficient solution algorithms for both regimes.  For example, we assumed that the trapped particle component is close to thermal and weak equilibrium while the streaming particle component carries a non-equilibrium spectrum. In the framework of the IDSA we have compare a grey and a spectral (i.e. energy-averaged and energy-dependent) treatment of the evolution of  the trapped particle component. We found visible differences of several percent. We believe that the differences are large enough to value the spectral approach, but at the same time they are small enough to be ignored if the efficiency of the algorithm in 3D simulations becomes critical. However, we emphasise that this comparison must not be confused with the comparison of a globally grey scheme and a globally spectral approach. We believe that the {\\em streaming} particle component, which represents the lepton and energy flux in a supernova application, must be treated in a spectral manner. The location of the neutrinospheres is highly energy-dependent. Hence, the geometry of the radiative transfer problem is different for each neutrino energy. An energy-averaged local quantity can only be expected to be accurately represented if one averages over solutions of the energy-dependent transport problem. A grey scheme, that solves an energy-averaged transport problem instead, is unlikely to produce accurate results in the supernova context.  Aside from the field of neutrino-driven supernova models, FLD has been used in fields ranging from planet migration \\citep{Paardekooper.Mellema:2006} to cosmology \\citep{Bonanno.Romano:1994}, including areas such as accretion discs, both around stars \\citep{Kley:1989} and more energetic objects where the accretion disc is dominated by radiation (e.g. \\citep{Turner.Stone.Sano:2002,Hujeirat.Camenzind:2000}), and star formation \\citep{Whitehouse.Bate:2006}. The IDSA is potentially applicable to these and other areas of astrophysics where FLD is commonly used. However, neither the IDSA nor the FLDA can replace the development of more reliable transport schemes because, for each application, it has to be ascertained that no important physical ingredients are missed. We hope that the IDSA is well-suited to accompany the application of sophisticated radiative transfer algorithms in order to enable an efficient exploration of input physics, parameter space, or dimensionality in computer models where a comprehensive solution of the Boltzmann transport equation is not affordable."
    },
    "0711/0711.0437_arXiv.txt": {
        "abstract": "We report on resolved interferometric observations with VLTI/MIDI of the massive young stellar object (MYSO) W33A. The MIDI observations deliver spectrally dispersed visibilities with values between 0.03 and 0.06, for a baseline of 45m over the wavelength range 8-13$\\mu$m. The visibilities indicate that W33A has a FWHM size of approximately 120\\,AU (0.030'') at 8$\\mu$m which increases to 240\\,AU at 13$\\mu$m, scales previously unexplored among MYSOs. This observed trend is consistent with the temperature falling off with distance. 1D dust radiative transfer models are simultaneously fit to the visibility spectrum, the strong silicate feature and the shape of the mid infrared spectral energy distribution (SED). For any powerlaw density distribution, we find that the sizes (as implied by the visibilities) and the  stellar luminosity are incompatible. A reduction to a third of W33A's previously adopted luminosity is required to match the visibilities; such a reduction is consistent with new high resolution 70$\\mu$m data from Spitzer's MIPSGAL survey. We obtain best fits for models with shallow dust density distributions of $r^{-0.5}$ and $r^{-1.0}$ and for increased optical depth in the silicate feature produced by decreasing the ISM ratio of graphite to silicates and using optical grain properties by Ossenkopf et al. (1992). ",
        "introduction": "Massive young stellar objects are luminous, embedded infrared (IR) sources that show many signs that they are still actively accreting mass. Their luminosity ($ L>10^{4}L_{\\odot}$) is such that they are expected to be ionizing their surroundings to produce an H II region, yet they only have weak radio emission due to ionized winds or jets (Hoare 2002\\nocite{2002ASPC..267..137H}). Most appear to be driving bipolar molecular flows and (sub-)millimetre interferometry is beginning to reveal evidence of rotating, disc-like structures on scales of hundreds of AU (Patel et al. 2005\\nocite{2005Natur.437..109P}; Beltr\\'{a}n et al. 2006\\nocite{2006Natur.443..427B}; Torrelles et al. 2007\\nocite{2007ApJ...666L..37T}). There is great interest in knowing the distribution of the infalling and outflowing material on smaller scales to provide clues to which physical processes are controlling the dynamics and setting the final mass of the star (Beuther et al. 2007\\nocite{2007prpl.conf..165B}; Hoare et al. 2007\\nocite{2007prpl.conf..181H}).  The mid-IR (8-13$\\mu$m) emission from MYSOs is thought to arise in the warm ($\\sim$300\\,K) dust in the envelope heated by the young star. Previous modelling (e.g. Churchwell et al. 1990\\nocite{1990ApJ...354..247C}) has indicated that the size of the mid-IR emission region should be unresolved at the typical distances of MYSOs by single-dishes and this is borne out by observations (Kraemer et al. 2001\\nocite{2001ApJ...561..282K}; Mottram et al. 2007\\nocite{2007arXiv0709.2040M}). Exceptions occur when MYSOs are viewed close to edge-on when a dense torus can completely obscure the bright central region and only the warm dust in the outflow cavities is seen (de Buizer 2006\\nocite{2006ApJ...642L..57D}).  Here we present new results from mid-IR interferometric observations of the MYSO W33A using the VLTI MIDI instrument. W33A has a kinematic distance of 3.8 kpc and a luminosity as derived from IRAS fluxes of $\\rm 1 \\times 10^{5} L_{\\odot}$ (Fa\\'{u}ndez et al. 2004\\nocite{2004A&A...426...97F}). It has weak, compact, optically thick radio continuum emission (Rengarajan \\& Ho 1996\\nocite{1996ApJ...465..363R}; van der Tak \\& Menten 2005\\nocite{2005A&A...437..947V}) and broad ($\\sim$100 km s$^{-1}$), single-peaked H I recombination emission lines (Bunn et al. 1995\\nocite{1995MNRAS.272..346B}) consistent with an ionized stellar wind origin. IR images from 2MASS and Spitzer's GLIMPSE survey clearly show a large scale monopolar nebula emerging to the SE, which is most likely to demarcate the blue-shifted lobe of a bipolar outflow.  The cold dust envelope has been studied at a few arcsecond resolution at millimetre wavelengths (van der Tak et al. 2000\\nocite{2000ApJ...537..283V}), probing the imprint of the star formation process on the circumstellar dust distribution. Many different scenarios have been put forward predicting this distribution. In case of W33A, G\\\"{u}rtler et al. (1991\\nocite{}) modelled the SED with a constant density spherical envelope and found an inner cavity radius of 135 AU (35 mas at 3.8 kpc) with the dust close to the sublimation temperature. MIDI provides data on comparable size scales of 50 mas, and since the mid-IR size will be somewhat larger than the sublimation zone we expect the emission to be well-resolved. ",
        "conclusions": "We presented mid-IR high-resolution dispersed interferometric observations of the forming massive star W33A. The visibility spectrum indicates an equivalent Gaussian FWHM of the emitting region of $\\sim 120$AU, at 8$\\mu$m, increasing to $\\sim 240$AU, at 13$\\mu$m. We interpreted the interferometric data with simple spherically symmetric DUSTY models  representative of a (unresolved) star embedded in a dusty envelope, aiming for a simultaneous fit to the SED, silicate profile and visibility spectrum. For any radial dust distribution we found that the canonical value of W33A's luminosity is not compatible with the visibilities. This is supported by MIPSGAL data. We found that even for a reduced luminosity, the model produces too large emitting regions causing the visibilities and SED to be mutually incompatible. Changing the dust composition improves the situation for an increased graphite to silicate ratio and using warm silicate optical constants from Ossenkopf et al. (1992). We resorted to a substantial lowering of the $T_{\\rm eff}$ in order to obtain a satisfactorily match between observables and models.  Further coverage of the uv-plane will be very rewarding, constraining more appropriate, higher dimensionality models, which will eventually lead in a proper description of the circumstellar environment of an accreting massive star. A full 2D axi-symmetric treatment of this and other MIDI data will be the subject of a future paper."
    },
    "0711/0711.0747_arXiv.txt": {
        "abstract": "Detecting the parity-odd, or B-mode, polarization pattern in the cosmic microwave background radiation due to primordial gravity waves is considered to be the final observational key to confirming the inflationary paradigm. The search for viable models of inflation from particle physics and string theory has (re)discovered another source for B-modes: cosmic strings. Strings naturally generate as much vector mode perturbation as they do scalar, producing B-mode polarization with a spectrum distinct from that expected from inflation itself. In a large set of models, B-modes arising from cosmic strings are more prominent than those expected from primordial gravity waves. In light of this, we study the physical underpinnings of string-sourced B-modes and the model dependence of the amplitude and shape of the $C_l^{BB}$ power spectrum. Observational detection of a string-sourced B-mode spectrum would be a direct probe of post-inflationary physics near the GUT scale. Conversely, non-detection would put an upper limit on a possible cosmic string tension of $G\\mu \\lesssim 10^{-7}$ within the next three years. ",
        "introduction": "\\vspace{-0.1in}  In the past few years, observations of the cosmic microwave background (CMB) and large scale structure data have largely confirmed the predictions of the inflationary paradigm. Now, observational cosmologists have begun to turn their instruments towards the final unconfirmed signature of inflation: primordial B-mode, or odd-parity, polarization in the CMB \\cite{tfcmb}. Meanwhile, string theorists have worked diligently towards realizing observationally viable models of inflation within reasonably well-understood string compactifications (see, e.g.,  \\cite{HenryTye:2006uv, McAllister:2007bg, Burgess:2007pz} for reviews). In parallel, viable models of hybrid inflation continue to be explored in the context of particle physics (e.g., \\cite{smooth1,cop,rc,linde}). Brane and hybrid inflation models often predict the production of cosmic strings \\cite{HenryTye:2006uv, kofman,tkachev,jeannerot,mairi,cop,rc, Jeannerot:2006dy,jst,costring,DV03,Polchinski03,mcgill,Firouzjahi:2005dh}, linear defects arising from the breaking of U(1) symmetries towards the end of the inflationary epoch. These relics of the inflationary period provide a rich phenomenology, but here we focus on their contribution to the polarization of the CMB.  The production of cosmic strings in such models seems at first unimportant, since they cannot source the majority of the CMB anisotropy \\cite{ABR97}. However, it now appears that strings may be the most prominent source of observable B-mode polarization in many of these braneworld or hybrid inflationary models \\cite{Seljak:2006hi,Battye:2007si}. Strings, as extended objects, source vector mode perturbations in the CMB as readily as they do scalar mode perturbations. This is in contrast with the inflationary picture, where gravity waves generated by inflation source tensor mode perturbations that are generally much smaller than the adiabatic scalar mode perturbations. Defect-sourced B-mode polarization was first calculated in \\cite{pen_pol} for global defects, and for local strings in \\cite{battye_98, Benabed:2000tr,PTWW03}. Recently, B-mode predictions from field theoretical simulations of local strings were also reported in \\cite{Bevis:2007qz}. It is true that cosmic strings are subdominant sources of CMB perturbation, limited to producing less than about 10\\% of the primordial anisotropy \\cite{contaldi,battye,bouchet,bevis,Wyman:2005tu,Seljak:2006bg,Bevis:2007gh}\\footnote{It was shown in \\cite{Bevis:2007gh,Battye:2007si} that CMB data can actually favor a contribution from strings if the inflationary spectrum is either exactly Harrison-Zeldovich ($n_s=1$) or, more generally, has $n_s \\ge 1$. The degree of preference is somewhat dependent on the shape of the spectrum. String models that predict a pronounced peak at $\\ell \\sim 400$ in the string contribution to the CMB spectrum  are less likely to improve the fit to the data than those with no peak}. But since the inflationary tensor-to-scalar ratio - the parameter that governs the strength of inflationary gravity waves -- is expected also to be below 10\\%, cosmic strings and primordial gravity waves may well be observables of similar detectability.  Models of hybrid and brane inflation \\cite{Spergel:2006hy}, tend to produce a low tensor-to-scalar ratio \\cite{Martineau:2007dj}. Hybrid inflationary models from particle physics, for instance, typically have unobservably small ($r \\lesssim 10^{-4}$) tensor to scalar ratios \\cite{Linde:1993cn, Dufaux:2007pt}. Most models of inflation that have been realized in string theory are of the hybrid type, and also have  small scalar-to-tensor ratios, typically $r \\sim 10^{-10}$ (e.g. \\cite{Brax:2006yq, Avgoustidis:2006zp, Bond:2006nc}). On the other hand, many of these models produce cosmic strings near the GUT scale \\cite{Lin:2006xta, Clauwens:2007wc,Jeannerot:2006jj, Cui:2007js}, which implies a cosmic string tension ($G\\mu$) close to the observational upper bound of $G\\mu \\lesssim 2 \\times 10^{-7}$ \\cite{Wyman:2005tu,Seljak:2006bg,Bevis:2007gh}.  Among models of inflation motivated by string theory, brane inflation (reviewed recently in Refs. \\cite{Burgess:2007pz, HenryTye:2006uv}) can create the widest range of cosmic sting tensions, with  $10^{-13} < G\\mu < 10^{-6}$ possible. In certain warped throat ``slow-roll'' scenarios, for instance, $G\\mu \\sim 10^{-9}$ is expected, while in other ``fast-roll'' realizations $G\\mu \\gtrsim 10^{-7}$ can easily be achieved \\cite{Firouzjahi:2005dh}. In the low-multipole region ($2 < \\ell < 100$), experimental sensitivity to primordial gravity waves can be converted to sensitivity to strings via\\footnote{This conversion covers only the region of $2 < \\ell < 100$, and was calculated by comparing the total BB power in strings versus the total BB power from tensor modes in this $\\ell$ range. This particular relation is dependent on the parameters of the string model, and a more complete conversion is suggested in  \\S\\ref{ddetection}.} $G\\mu\\leftrightarrow 1.4 \\times 10^{-6}\\sqrt{r}$; hence, the promise that near-term experiments can observe $r\\sim10^{-2}$ in this multipole range implies sensitivity to strings with tensions as low as $G\\mu \\sim 10^{-7}$. Brane inflation, however, may not be able \\footnote{More complicated set-ups involving ``wrapped'' branes might loosen such bounds \\cite{Krause:2007jr, Becker:2007ui}.} to create a tensor-to-scalar ratio above $r \\sim \\mbox{(few)} \\times 10^{-3}$ \\cite{Baumann:2006cd}. For high $\\ell > 100$, the situation is different: the dominant B-modes are expected to come from the lensing of primordial-E modes into B-modes by large scale structure. However, the string sourced B-mode spectrum peaks in this high $\\ell$ range, so if strings are present they will show up as an excess over the expected signal. An apparent excess of 10\\%, for instance, would be caused by strings with a tension of $G\\mu \\simeq 6 \\times 10^{-8}$ (assuming standard string network parameters).   To summarize: in models of inflation that are motivated either by string theory or particle physics, the more promising avenue for creation of B-mode polarization is via cosmic strings, not tensor modes. Hence, we might hope that if B-modes from inflationary tensor modes are not observed, B-modes from cosmic strings produced after inflation might be. Inflationary models with large tensor modes (like chaotic inflation) tend not to create defects. Distinguishing the source of the polarization will be a challenge. To distinguish string B-modes from primordial gravity waves requires fine scale resolution in the range $\\ell < 100$. Distinguishing string sourced polarization at $\\ell > 100$ requires an accurate calculation of the amplitude of the B-mode polarization that comes from the conversion of E-mode polarization to B-mode polarization by gravitational lensing. This is because the peak BB contribution from strings falls at $\\ell \\sim 1000$, overlapping heavily with the spectrum expected from lensing of E. Hence, the presence of a string-induced B-mode would induce an excess in power throughout this region of $\\ell$ values above what is expected from lensing. Accurately characterizing this excess will require very good predictions of the expected lensing signal. If all goes well, polarization experiments could eventually see strings with tensions as low as $G\\mu \\sim 10^{-9}$ \\cite{Seljak:2006hi}. Because strings can source B-modes in the aftermath of a hybrid inflationary epoch that cannot directly source B-modes, the class of inflationary models that B-mode experiments can test is broadened. Detecting neither string or primordial gravity wave sourced B-modes would strongly constrain many inflationary models. These constraints would be different from, and perhaps stronger than, the constraints derived from measurement of other cosmological parameters (e.g., the spectral index $n_s$).  In this paper, we outline the physics behind cosmic strings generation of B-mode polarization, placing a heavy emphasis on how much variability might be expected from various models of string networks. This gives us an idea of how much a successful detection of cosmic string B-modes might teach us about these networks. We also review some of the currently operating, soon to be launched, and future planned experiments to measure the CMB's B-mode polarization to give estimates for what range of string tensions they will be able to detect. ",
        "conclusions": "\\vspace{-0.1in}  Cosmic string networks can produce  B-mode polarization in the CMB at observable levels if the strings source between 1 and 10\\% of the TT-power in the CMB. This prediction holds across a wide variety of string network parameter values, as we have discussed above. Smooth ($\\alpha \\sim 1$) and/or slow-moving ($v_{rms} \\sim 0.1$) strings produce the strongest B-modes, because they generate the largest ratio of anisotropic to isotropic stress sourced by cosmic strings.  The first B-mode polarization experiments will bin many $\\ell$ multipoles together. We calculate that a cosmic string-sourced B-mode signal could mimic the B-mode power expected from inflationary gravity waves in the range between $2 < \\ell < 100$. Since low-$\\ell$ power from strings versus inflationary gravity waves scale approximately as $G\\mu  \\leftrightarrow 1.4 \\times 10^{-6} \\sqrt{r} \\, \\alpha \\, ( 0.65 /v_{rms} ) \\sqrt{\\xi/0.3}$, expectations that future experiments could probe $r \\sim 10^{-3}$ \\cite{Amarie:2005in} imply that cosmic strings with tensions as low as $5 \\times 10^{-8}$ may eventually be detectable via their low-$\\ell$ B-mode polarization. In the meantime, without fine resolution in $\\ell$, the only way to discriminate between these two sources would be to look at high-$\\ell$ ($100 < \\ell < 1500$), where the string sourced B-mode spectrum reaches its peak. In this range, gravitational lensing of E-mode polarization into B-mode polarization is expected to dominate, so a string contribution would appear as a systematic excess of power over what is expected (see Fig. \\ref{detection}).  Disentangling the string-sourced and primordial gravity wave signals will require either good spectral information, via narrow $\\ell$ bins, or accurate and reliable calculations of the spectrum and amplitude of E to B lensing signal. Some existing methods for calibrating lensing, however, rely on the assumption that all of the B-mode polarization present for $\\ell > 150$ is sourced by lensing \\cite{Seljak:2003pn}. In our model, strings generate the same B-mode power as is expected from the E to B lensing signal for $G\\mu \\simeq 1.8 \\times 10^{-7}$ (for standard network parameters and smooth strings); thus, strings with a tension in this range would lead to a doubling of the B-mode power in the ``high\" region.  If low-$\\ell$ B-modes or an excess of power in the high-$ell$ region are seen, we should look for cosmic strings with tensions $G\\mu\\gtrsim 10^{-8}$ in another observational probe, such as gravitational lensing \\cite{Mack:2007ae, Chernoff:2007pd} or small-angle anisotropies in the CMB temperature \\cite{Fraisse:2007nu, Amsel:2007ki}. Such an observation would also motivate a reworking of lensing reconstruction algorithms to account for the cosmic string contribution. Even in the absence of such analysis, detailed spectral information should potentially allow one to distinguish between the inflationary and stringy reionization peaks at low-$\\ell$.  In practice, we expect that this will give only ambiguous results in the near term. Accurately calibrating the E to B lensing spectrum and amplitude is more promising, since it allows direct access to the peak in string-generated B-mode power, which -- as noted before -- will be present as an apparent excess of power in this high multipole range.  The discovery of a cosmic string signal in the B-mode polarization of the CMB would teach us about the string network that generated it since the shape of the string-sourced spectrum is sensitive to network parameters. It would also help us to determine the energy scale at which inflation occurred. As we have demonstrated, the expected string-sourced B-mode spectrum is robust to theoretical uncertainty in the string network model. Thus, as these experiments turn on and collect polarization data, a cosmic string network will either be seen through its polarization spectrum or will be restricted to a considerably lower tension in the case of non-observation, with polarization-based limits on $G\\mu$ being comparable to those expected from the analysis of the data from PLANCK \\cite{Battye:2007si}.  \\vspace*{-0.35in}"
    },
    "0711/0711.2574_arXiv.txt": {
        "abstract": "The nearby dwarf spheroidal galaxy Draco with its high mass to light ratio is one of the most auspicious targets for indirect dark matter (DM) searches. Annihilation of hypothetical DM particles can result in high-energy $\\gamma$-rays,  e.g. from neutralino annihilation in the supersymmetric framework.\\\\ With the MAGIC telescope a search for a possible DM signal originating from Draco was performed during 2007. The analysis of the data  results in a flux upper limit ($2\\sigma$) of  $1.1~\\times~10^{-11}~\\mathrm{photons}~\\mathrm{cm}^{-2}~\\mathrm{sec}^{-1}$ for photon energies above $140\\,$GeV, assuming a point like source. Furthermore, a comparison with predictions from supersymmetric models is given. While our results do not constrain the mSUGRA phase parameter space, a very high flux enhancement can be ruled out. ",
        "introduction": "Astronomical observations provide strong evidence for the existence of a new type of non-luminous, non-baryonic matter, contributing to the total energy density of the universe about six times more than baryonic matter \\citep{Spergel2007}. This so-called Dark Matter (DM) makes its presence known through gravitational effects, and could be made of so-far undetected relic particles from the Big Bang. Weakly Interacting Massive Particles (WIMP) are candidates for DM, with the lightest supersymmetric particle (neutralino) being one of the most favored candidates in the list of possible WIMPs. Stable neutralinos are predicted in many supersymmetric (SUSY) extensions of the Standard Model of Particle Physics~\\citep{Jungman1995}. Since the neutralinos are Majorana particles, pairs of neutralinos can annihilate and produce Standard Model particles. Direct annihilations into $\\gamma\\gamma$ or Z$\\gamma$ produce a sharp line spectrum with a photon energy depending on the neutralino mass. Unfortunately, these processes are loop-suppressed and therefore very rare. Neutralinos can also annihilate to  pairs of $\\tau$ or quarks, leading in subsequent  processes to $\\pi^0$-decays, resulting in a continuous photon spectrum.\\\\ Draco is a dwarf spheroidal galaxy accompanying the Milky Way at a galactocentric distance of about $82\\,$kpc. It is characterized by a high mass to light ratio $M/L > 200$, implying a high DM concentration~\\citep{Mayer2007,Hooper2006} so complying with the trend generally deduced for low-luminosity galaxies (e.g. \\citet{Persic1996}). ",
        "conclusions": "We present the first search for $\\gamma$-rays from the direction of Draco using an Imaging Air Cherenkov Telescope (IACT). No signal was detected. The $2\\sigma$ u.l. on a steady $\\gamma$-ray emission above $140\\,\\mathrm{GeV}$ originating from Draco does not exceed $1.1\\times 10^{-11}\\,\\mathrm{photons}\\,\\mathrm{cm}^{-2}\\,\\mathrm{s}^{-1}$ if the underlying spectrum follows a power law with spectral index -1.5.\\\\ For the mSUGRA benchmark models defined in \\citep{Battaglia2003} and assuming a smooth DM density distribution for Draco as given in \\citep{Sanchez-Conde2007}, our flux upper limits are O($10^3$ - $10^9$) above the predicted values. It is therefore not possible to constrain the mSUGRA phase space by these results, but a very high flux enhancement can be excluded. \\\\ Even though an indirect DM detection by measuring  $\\gamma$-rays from neutralino annihilation within the halo of Draco seems for present IACTs out of reach, future satellite telescopes like GLAST with lower energy thresholds might be sensitive enough to reach the mSUGRA parameter space."
    },
    "0711/0711.3373_arXiv.txt": {
        "abstract": "We present a new set of theoretical evolutionary synthesis models, PopStar. This grid of Single Stellar Populations covers a wide range in both, age and metallicity. The models use the most recent evolutionary tracks together with the use of new NLTE atmosphere models for the hot stars (O, B, WR, post-AGB stars, planetary nebulae) that dominate the stellar cluster's ionizing spectra. The results of the models in VO format can be used through VOSpec. ",
        "introduction": "We have used the synthesis code by \\citet{gmb:1998}, updated by \\citet{mgv:2000} and newly revised now. The basic grid is composed by Single Stellar Populations (SSP) for five different IMFs: two based on Salpeter power law \\citep{sal:1955}, with mass range 0.85 - 120 M$_{\\odot}$ and 0.15 - 100 M$_{\\odot}$ respectively, and those of \\citet{fpp:1990}, \\citet{kro:2002} and \\citet{cha:2003} functions, with masses between 0.15 and 100 M$_{\\odot}$. These models do not include binaries either mass segregation. The isochrones are those from \\citet{bgs:1998} for 6 different metallicities: Z $=$ 0.0004, 0.001, 0.004, 0.008, 0.02 and 0.05. The age coverage is from $log{t}=$ 5.00 to 10.30, with a variable time resolution which is $\\Delta(log{t})=0.01$ in the youngest stellar ages. The WC and WN stars are identified in the isochrones according to their surface abundances.  \\begin{figure*} \\centering \\includegraphics[angle=90,scale=0.50]{fig1.eps} \\caption{PopStar SEDs for young clusters displayed through the VOSpec interface \\label{fig:single}} \\end{figure*}  The atmosphere models are from \\citet{lcb:1997} with an excellent coverage in effective temperature, gravity and metallicities, for stars with T$_{eff}$ $ \\leq 25 000$ K. For O, B and WR we have taken the NLTE blanketted models by \\citet{snc:2002} at 5 different metallicities. There are 110 models for O-B stars, calculated by \\citet{phl:2001}, with $25 000$ K $ <$ T$_{eff}$ $\\leq 51 500$ K and $2.95 \\leq \\log{g} \\leq 4.00$, and 120 for WR stars (60 WN and 60 WC), from \\citep{hm:1998}, with 30 000 K $\\leq T^{*}\\leq 120 000$ K and $1.3R_{\\odot}\\leq R^{*}\\leq 20.3 R_{\\odot}$ for WN, and with $ 40 000 K \\leq\\ T^{*} \\leq 140 000$ K and $ 0.8R_{\\odot}\\leq\\ R^{*}\\leq 9.3 R_{\\odot}$ for WC. T$^{*}$ and R$^{*}$ are the temperature and the radius at a Roseland optical depth of 10. To assign an appropriate model to each WR star, we use the relationships among opacity, mass loss and velocity wind: $d\\tau=-\\kappa(r)\\rho(r)dr$, where $\\kappa(r)=-0.2(1+X_{\\rm S})$ and $X_{\\rm S}$, the H surface abundance, is taken as 0.2 for WN and 0 for WC. The mass loss is: $dM/dt =4\\pi r^{2}\\rho(r)v(r)dr$ with $v(r)=v_{\\infty}(1-R_{S}/r)^{\\beta}$ taking $\\beta=1$.  Integrating these equations we find R$^{*}$ from the $R_{S}$ of isochrones selecting the closer atmosphere model to this R$^{*}$. Above 54ev emergent flux is determined by the wind density as a function of metallicity and in this set of models the hardness of ionizing radiation decreases, producing a lower flux of HeI continuum, for $Z > 0.4Z_{\\odot}$ and t $< 7$ Myr. HeII 4686 is very low and HeI 5786 / H$_\\beta$ decreases when metallicity increases due to line blanketing.  For post-AGB and Planetary Nebulae with T$_{eff}$ ranging from 50 000 K to 220 000 K we take the NLTE models by \\citet{rau:2003} with values of T$_{eff}$  between 50 000 K and 190 000 K and logg between 5.00 and 8.00. These models include all elements from H to Ni. For higher temperatures we use black bodies. ",
        "conclusions": "We have computed a new grid of SEDs with updated isochrones and atmosphere models in a wide range of age and metallicity. The use of NLTE blanketed models produce SEDs with less hard ionizing photons, able to explain the observed emission line spectrum in low excitation high metallicity H{\\sc ii} regions without changing the IMF parameters. Previous models \\citep{gbd:1995} could not explain these observations without eliminating the more massive stars artificially (with a flatter IMF, a lower value of the upper mass limit or using the standard IMF but considering the large clusters divided in small less massive sub-cluster with a low probability of forming very massive stars). H-R Diagrams are available as VO tables. The resulting SEDs are in VO format accessible through SSAP protocol and can be managed with VOSpec."
    },
    "0711/0711.1376_arXiv.txt": {
        "abstract": "We present a newly developed approach to solar coronal magnetic field extrapolation from vector magnetograms, based on the Principle of Minimum Dissipation Rate (MDR). The MDR system was derived from a variational problem that is more suitable for an open and externally driven system, like the solar corona. The resulting magnetic field equation is more general than force-free. Its solution can be expressed as the superposition of two linear (constant-$\\alpha$) force-free fields (LFFFs) with distinct $\\alpha$ parameters, and one potential field. Thus the original extrapolation problem is decomposed into three LFFF extrapolations, utilizing boundary data. The full MDR-based approach requires  two layers of vector magnetograph measurements on solar surface, while a slightly modified practical approach only requires one. We test both approaches against 3D MHD simulation data in a finite volume. Both yield  quantitatively good results. The errors in the magnetic energy estimate are within a few percents. In particular, the main features of relatively strong perpendicular current density structures, representative of the non-force freeness of the solution, are well recovered. ",
        "introduction": "Plasma relaxation processes are ubiquitous in astrophysical and laboratory plasmas. One well-known approach is to invoke the variational principle, minimizing the total magnetic field energy subject to the constraint of constant total magnetic helicity \\citep{taylor74}. Such an approach yields the magnetic field equation for a relaxed plasma state, in which the Lorentz force vanishes, i.e., the so-called linear force-free field (LFFF) with a constant $\\alpha$ parameter, \\begin{equation} \\nabla\\times\\mb{B}=\\alpha\\mb{B}. \\end{equation}\\label{lfff} Although the above equation~(\\ref{lfff}) has been widely applied in modeling magnetic field configurations in laboratory plasma devices and in solar atmosphere, it is found inadequate in certain situations \\citep[e.g.,][]{gary01, met95}. In particular, \\citet{ama00}, among others, showed by 3D numerical MHD simulation that in certain solar physics situation, after initial helicity drive, the final ``relaxed state is far from the constant-$\\alpha$ linear force-free field that would be predicted by Taylor's conjecture\" with helicity conserved. They suggested to derive alternative variational problem.  An alternative approach, also based on the variational principle, was recently developed for an open system with external drive, and applied to theoretical investigation of the solar arcade structures with flow \\citep{BJ04, Bha07}. It was later applied to develop a new approach to the extrapolation of solar coronal magnetic field in non-force free state \\citep{hu06, hu06sol, hu07aip}. The core of the approach is to construct a variational problem based on the Principle of Minimum Dissipation Rate (MDR), which states ``The steady state of an irreversible process is characterized by a minimum value of the rate of entropy production\" \\citep{pri47}. In most cases, the entropy production rate is equivalent to the energy dissipation rate. The basis is the set of generalized momentum balance equations, and the system, with flow, is always in broadly-defined force-balanced dynamic equilibrium, satisfying the MHD equations \\citep{BJ04}.  Prior works based on MDR \\citep{mon88, das98, das02} had successfully shown that it was able to yield a pressure-balanced configuration, supporting a finite plasma pressure gradient found in plasma confinement devices. Recently a proof of MDR was rigorously sought by 3D numerical simulations \\citep{shk07}. We refer the readers to the above referenced literatures for detailed description and justification of the theoretical basis of MDR.  In short, analogous to the principle of minimum energy, but for a more complex, open system like the solar corona, the MDR yields the following equations for the magnetic field $\\mb{B}$ and flow vorticity $\\vec\\omega=\\nabla\\times\\mb{v}$, $\\mb{v}$ being the plasma flow velocity \\citep{BJ04, Bha07, hu06sol}: \\begin{eqnarray} \\nabla\\times\\nabla\\times {\\bf B}+  a_1 \\nabla\\times {\\bf B}+ b_1{\\bf B}&=&\\nabla \\psi,  \\label{nffmag} \\\\ \\nabla\\times\\nabla\\times {\\vec \\omega}+ a_2 \\nabla\\times {\\vec \\omega} + b_2{\\vec \\omega}&=&\\nabla \\chi.    \\label{nfflow} \\end{eqnarray} Here,  coefficients $a$ and $b$ are constants, and involve the parameters of the system. The right-hand sides of both equations are arbitrary, undetermined functions.  In what follows, we first briefly provide a heuristic justification of the MDR in Section~\\ref{just}, based on a simple but standard analysis. We then present a brief but comprehensive description of the full MDR-based coronal magnetic field extrapolation approach in Section~\\ref{fullapp}, including a new test case study utilizing 3D MHD simulation data of a bright point region, highlighting the non-force free features of the solution. In Section~\\ref{praapp}, we develop a practical approach that requires only one single-layer vector magnetogram and show that it achieves the same satisfactory results as the full approach for the same test case. In the last Section~\\ref{conclu}, we conclude and discuss the significance and limitations of our approach. ",
        "conclusions": "In conclusions, we develop an approach to extrapolate the coronal magnetic field from vector magnetograms based on the Principle of Minimum Dissipation Rate (MDR). Analogous to, but yet opposed to the principle of minimum energy, the MDR yields a generally non-force free magnetic field through a different variational approach. The full MDR-based approach requires two layers of vector magnetograms, while a practical approach, representing a class of special solutions to the MDR system, requires only one layer, which is more amiable for practical applications to currently available data. A test case study using numerical simulation data shows that both  approaches recover the solution to a good degree of accuracy, as measured by a set of quantitative measures. The errors in energy estimate are both within a few percents. Moreover, the non-force free features of the solution, mainly the strong perpendicular current density concentrations, are well retained in the extrapolation results as well.  We hereby provide an alternative approach that is fast, and easy to implement, and allows one to extrapolate coronal magnetic field in a more general non-force free state with manageable effort in a solo work. For instance, all the calculations reported herein (128$\\times$128$\\times$63 grid) were performed on a single-processor 2.8GHz PC in IDL within a reasonable time frame. It potentially can be made much faster. It is apparent that the Eq.~(\\ref{nff3}) includes solutions to LFFF ($\\nabla\\times\\mb{B}=\\alpha\\mb{B}$), for $-\\alpha_1=\\alpha_3=\\alpha$. It is not explicit whether this is also the case for the nonlinear force-free field (NLFFF) when the parameter $\\alpha$ is allowed to vary, although several numerical experiments indicated that this might be so \\citep{hu06sol} sometimes.  We are fully aware of the limitations of the method, due to the fact that the currently employed LFFF solver originates from  an ill-posed problem. It is adversely affected by the increasing resolution of vector magetograms since $\\alpha_{max}$ is inversely proportional to $N$. This severely limits the search domain of $\\alpha_i$ that will yield unique physical solutions \\citep[e.g.,][]{ali81, gary89}. Recent progresses  in the algorithm for NLFFF extrapolations \\citep[e.g.,][]{son06}, in addition to the existing LFFF algorithm \\citep[e.g.,][]{abr96}, that modified the problem into a well-posed one, look promising and may be adapted.  It is worth noting that our approach is wholly applicable to extrapolating the plasma flow field, by simply replacing magnetic field $\\mb{B}$ with flow vorticity $\\vec{\\omega}$. However the boundary conditions are probably harder to obtain since it is the vorticity that is involved. A good start is to utilize the flow field data contained in the same 3D MHD simulation test case. It is not straightforward to obtain the plasma state since the flow field  has to be resolved as well, and to be consistent, the full set of MHD equations has to be considered \\citep[see, e.g.,][]{mon88}. One may also foresee that extension of this approach to the whole solar surface data is of great significance. Further theoretical investigation is underway."
    },
    "0711/0711.4264.txt": {
        "abstract": "{Motivated by the observed cosmic matter distribution, we present the following conjecture: due to the formation of voids and opaque structures, the average matter density on the path of the light from the well-observed objects changes from $\\Omega_M \\simeq 1$ in the homogeneous early universe to $\\Omega_M \\simeq 0$ in the clumpy late universe, so that the average expansion rate increases along our line of sight from EdS expansion $Ht \\simeq 2/3$ at high redshifts to free expansion $Ht \\simeq 1$ at low redshifts. To calculate the modified observable distance-redshift relations, we introduce a generalized Dyer-Roeder method that allows for two crucial physical properties of the universe: inhomogeneities in the expansion rate and the growth of the nonlinear structures. By treating the transition redshift to the void-dominated era as a free parameter, we find a phenomenological fit to the observations from the CMB anisotropy, the position of the baryon oscillation peak, the magnitude-redshift relations of type Ia supernovae, the local Hubble flow and the nucleosynthesis, resulting in a concordant model of the universe with 90\\% dark matter, 10\\% baryons, no dark energy, $15~{\\rm Gyr}$ as the age of the universe and a natural value for the transition redshift $z_0=0.35$. Unlike a large local void, the model respects the cosmological principle, further offering an explanation for the late onset of the perceived acceleration as a consequence of the forming nonlinear structures. Additional tests, such as quantitative predictions for angular deviations due to an anisotropic void distribution and a theoretical derivation of the model, can vindicate or falsify the interpretation that light propagation in voids is responsible for the perceived acceleration.}  \\preprint{HIP-2007-64/TH} ",
        "introduction": "The widely adopted framework of modern cosmology is the spatially homogeneous and isotropic Friedman-Robertson-Walker spacetime, with the growth of structure described as linear perturbations evolving on the smooth FRW background. Accordingly, the standard model of cosmology is built on the assumption that the effects of the evident nonlinear inhomogeneities on the detectable light average out\\footnote{Apart from gravitational lenses that are occasionally taken into account as additional corrections.} over cosmological distances. This assumption -- although perhaps in concordance with Newtonian intuition -- lacks a convincing demonstration within general relativity, not to mention a mathematically rigorous proof. In fact, the assumption was questioned already in the 60's by Zel'dovich \\cite{Zeldovich}, Feynman\\footnote{Apparently, R. P. Feynman was also one of the pioneers of this subject, as \\cite{Gunn1,Kantowski1} cite a colloquium by Feynman in 1964.}, Bertotti \\cite{Bertotti}, Gunn \\cite{Gunn1}, Kantowski \\cite{Kantowski1} and in the 70's by Dyer \\& Roeder \\cite{DyerRoeder1,DyerRoeder2}, who suggested that the clumpiness of matter in the real universe has consequences on the observable distance-redshift relations that the simplified FRW description fails to capture.  At the time of this pioneering research, the cosmological observations were too inaccurate to distinguish the predictions of the proposed inhomogeneous models from the simple and elegant FRW solutions. Consequently, the efforts towards a more thorough description of the effects of the nonlinear structures on the distance-redshift relations never became part of the mainstream research. Although the structure formation itself has since then developed into a major part of modern cosmology, the emphasis is usually given either to the general relativistic but perturbative description of structures on large scales, or to the Newtonian description of the small scale structure e.g.\\ in the context of galaxy formation.  At present, the unparalleled precision of the recent cosmological observations has raised the debate about the effect of the nonlinear structures on the distance-redshift relations once again topical. Indeed, the observations indicate the existence of inhomogeneities that seem to be on large enough scales to have cosmological significance but too strong to be described within the linear theory --- in particular the observed voids and the clustering of matter as filaments of opaque galaxies between the voids \\cite{Hoyle:2003hc,Gott:2003pf,Tikhonov:2007di,vonBendaBeckmann:2007wt,Einasto,Rudnick:2007kw}. The increase in the precision of the data naturally demands a corresponding improvement in the theoretical model; otherwise there is an increased risk of drawing incorrect conclusions from the observations.  In fact, the discovery that the standard FRW cosmology needs a fine-tuned dark energy component \\cite{Copeland:2006wr,Straumann:2006tv,Sahni:2006pa} in order to account for observations \\cite{Riess:2004nr,Riess:2006fw,Eisenstein:2005su,Spergel:2006hy}, suggests that the precision of the data may already have surpassed the precision of the employed theoretical model. This is also backed up by a mysterious coincidence, pointed out by Schwarz \\cite{Schwarz:2002ba} and Wetterich \\cite{Wetterich:2001kr} and elaborated by R\\\"as\\\"anen \\cite{Rasanen:2003fy,Rasanen:2006kp,Rasanen:2008it,Rasanen:2008be}: the perceived increase in the cosmic expansion, ascribed to dark energy in the standard cosmology, happens to take place during the same era when nonlinear structures start to form at cosmologically significant scales.  Whereas the fine-tuning of dark energy only indicates that something fundamental may be lacking from the standard picture, the coincidence problem suggests a causal connection between cosmic acceleration and structure formation. Perhaps the most natural explanation would be, that the standard FRW description breaks down when the nonlinearities become dominant in the late universe and the need for a tiny cosmological constant is a manifestation of this breakdown, not evidence for dark energy \\cite{Rasanen:2003fy}. In recent years, more and more authors have addressed this issue and substantially increased the general understanding of the subject, but crucial open questions still remain \\cite{Buchert:2007ik}. In this work, we address the issue by considering a phenomenological model that aims to: 1) clarify the physical interpretation for \\emph{how} the structure formation can actually mimic dark energy, and 2) provide a quantitative fit for the observations.  The paper is organized as follows. In Sects.\\ \\ref{noijannuoli}--\\ref{selectioneffects}, we review three physically independent mechanisms for how the cosmic inhomogeneities can induce accelerated expansion along our line of sight: global acceleration due to spatial variations in the expansion rate, faster local expansion rate due to a large local void, and biased light propagation through voids that expand faster than the average; possibilities to describe the total effect of the inhomogeneities have been discussed in Sect.\\ \\ref{TotalEffect}. After reviewing the original Dyer-Roeder method in Sect.\\ \\ref{originalDR}, and motivating the need to modify it, we introduce a generalized Dyer-Roeder method in Sect.\\ \\ref{dyerroeder} to allow for the total effect of the inhomogeneities on the observable distance-redshift relations. The distance-redshift relations are then calculated for an inhomogeneous model, based on the conjecture that the average matter density on the path of the light from the observed objects changes from $\\Omega_M \\simeq 1$ in the homogeneous early universe to $\\Omega_M \\simeq 0$ in the clumpy late universe. The physical foundations of the conjecture are summarized in Sect.\\ \\ref{physics}. In Sect.\\ \\ref{derivation}, we construct a simple description of the structure formation in terms of a function $\\alpha(z)$ that measures the deviation of the matter density along our line of sight from the FRW matter density. A comparison of the model with cosmological observations is performed in Sect.\\ \\ref{observations}, and a discussion of the results is given in Sect.\\ \\ref{results}. Finally, Sect.\\ \\ref{konkluusiot} contains our conclusions. ",
        "conclusions": "\\label{konkluusiot}  We have studied the effect of structure formation on the observable distance-redshift relations by generalizing the Dyer-Roeder method to allow for two important physical properties of the universe: inhomogeneities in the expansion rate and the growth of nonlinear structures. We noted that light propagation through voids can cause the expansion along our line of sight to increase from flat dust FRW-expansion at high redshifts to free expansion at low redshifts and thereby have a similar effect on the observed distance-redshift relations as dark energy. Conjecturing such a transition, we showed that the inhomogeneities offer a natural physical explanation for the cosmological observations without dark energy even with a single physically well-grounded phenomenological parameter, the transition redshift $z_0$ to the era when the matter density becomes negligible along our line of sight or when the detectable light starts to propagate in freely expanding space.  Indeed, applying the generalized Dyer-Roeder method in Sect.\\ \\ref{observations}, we found a simple inhomogeneous model that fits the observations from the CMB anisotropy, the position of the baryon acoustic oscillation peak inferred from the SDSS data, the magnitude-redshift relations of type Ia supernovae, the local Hubble flow and the Big Bang nucleosynthesis, plus is potentially consistent with the estimated age of the quasar APM $08279 + 5255$ at $z=3.91$. The main features of the resulting concordant model are: 90\\% dark matter, 10\\% baryons, $z_0=0.35$ as the transition redshift to the void-dominated era, the values $h=0.66$ and $\\bar{h}=0.44$ for the observable and FRW Hubble constants that yield $t_0 = 14.8~{\\rm Gyr}$ as the age of the universe, consistent with astronomical estimates, and no dark energy of any kind or $\\Lambda=0$. At large redshifts, this nonlinear CHDM model is similar to the linear CHDM model by Hunt and Sarkar \\cite{Hunt:2004vt,Hunt:2007dn}. A non-trivial point in the best fit model is that the distance-redshift relation is not a reproduction of the corresponding relation in the $\\Lambda$CDM model. The model is based on inhomogeneities of the observed kind so, unlike a large void in the local matter density, there is no violation of the cosmological principle: it is our special location in time, not in space, that is required for the acceleration. Therefore, a prediction is that \\emph{every} observer in the present-day universe should perceive the inhomogeneity induced acceleration.  Although homogeneous models can also provide a phenomenological fit to the cosmological data by assuming dark energy, they entail various issues that can be formulated into five questions: \\begin{enumerate} \\item What justifies the use of the FRW metric to model light propagation in the real universe? \\item Why would dark energy appear at a very low temperature $T \\simeq 4 ~ {\\rm K}$ if it has had the huge temperature range from $T_{{\\rm P}} \\simeq 10^{32}~ {\\rm K}$ to $T_0 \\simeq 2.7$ K available? An expectation for the probability of a quantum field to arise at a temperature $T$ is $P \\sim e^{-T_{{\\rm P}}/T}$. \\item Why would the effects of dark energy appear when large voids start to form? \\item Why would $\\Lambda$ be fine-tuned such that the universe undergoes nearly free expansion today, $H_0 t_0 =\\frac{ 2}{3 \\sqrt{\\Omega_\\Lambda}} {\\rm {arsinh} } \\sqrt{\\frac{\\Omega_\\Lambda}{1-\\Omega_\\Lambda}} \\simeq 1$, whereas a slightly bigger $\\Lambda$ would yield $H_0 t_0 >1$? This issue was pointed out by Kutschera and Dyrda in Ref.\\ \\cite{Kutschera:2006bh}. \\item If nonzero, why would the cosmological constant be so tiny, $\\Lambda \\sim 10^{-123}~G^{-1}$, whereas naive dimensional analysis suggests $\\Lambda \\sim G^{-1}$? \\end{enumerate} The structure formation offers a smart solution to the coincidence problem: if induced by the voids, the onset of the perceived acceleration naturally coincides with the formation of the voids. Furthermore, the observed $H_0 t_0 \\sim 1$ corresponds to the characteristic expansion rate of freely expanding voids, $t \\sim t_0/2$ is a natural time scale for gravity to form structures on the relevant length scales and $\\Lambda = 0$ is compatible with the data if void formation is responsible for the acceleration. In addition, there are various observed features in the universe that are anomalies for the standard FRW cosmology, such as the cold spot in the CMB, the low multipoles of the CMB, the CMB ellipticity and the scatter in the supernova data, but could be explained naturally by the anisotropy of the cosmic void distribution in the sky; a more exhaustive list of such features can be found in Ref.\\ \\cite{Wiltshire:2007zj}. As we have only discussed how to construct models of the observed inhomogeneous universe within relativistic classical gravity, the answer to the long-standing debate of \\emph{why} would $\\Lambda=0$ is naturally beyond the scope here.  The transition to the nonlinear era, roughly given by $z_0=0.35$, might seem to happen rather late, only $\\Delta t = 5.4 ~{\\rm Gyr}$ back in time, whereas the oldest stars have formed already $\\Delta t > 12 ~{\\rm Gyr}$ ago. However, there are in fact several natural reasons to expect such a seemingly late transition. Firstly, while traveling through structures that are still collapsing in earlier times, light is blueshifted which can partially counterbalance the excess redshift caused by the light propagation in voids. Later, however, when the structures have virialized also at larger scales, there are more and more structures which neither collapse nor expand. Those photons that are able to penetrate these virialized overdense regions, then do not blueshift anymore, diminishing this partial counterbalance at late times. Secondly, the large amount of weakly interacting dark matter in the universe prefers to form large halos which cluster into smaller clumps later than the luminous matter which nonetheless indicates the clumping to us. It also appears likely that the voids can grow large enough only in the late universe. Finally, when contrasting the acceleration from structure formation with dark energy driven acceleration, it is essential to note that the rate at which gravitational phenomena happen is largely independent of the temperature, whereas non-gravitational interactions happen faster at higher temperatures. Therefore, while it seems extremely difficult to explain why a quantum field would arise at $T \\simeq 4~{\\rm K}$ especially if it has had the huge temperature range from $T_{{\\rm{P}}} \\simeq 10^{32}~{\\rm K}$ to $T_0 \\simeq 2.7~{\\rm K}$ available, it is natural that the inhomogeneity induced acceleration would start as \"late\"\\ as $\\Delta t \\simeq 7 ~{\\rm Gyr}$ ago (corresponding to redshift $z \\simeq 0.52$ that can be read off e.g.\\ from Fig.\\ \\ref{figuuri}).  The analytic approach employed in this work necessitates approximations, so the precision of the numerical values for the best-fit parameters presented here is not expected to be as high as an ideal analysis could yield. Thus, it is difficult to say anything conclusive about the subtler features of the concordant nonlinear CHDM model, such as the bump in the primordial spectrum or the small hot dark matter component. Instead, given the physically well-grounded foundations of the proposed conjecture about light propagating mostly in voids, the observational evidence for the nonlinear structures as well as the large volume taken up by voids, the concordant fit to the cosmological data and the offered smart solutions to some of the perhaps most puzzling theoretical problems in the modern cosmology, it seems perfectly conceivable that the main physical implications -- most notably that dark energy may not exist -- could be correct.  An important future test is to study quantitatively the observed angular deviations in various data that are an anomaly for the homogeneous standard $\\Lambda$CDM model, but a predicted feature in the inhomogeneous models arising from the anisotropic void distribution. Another crucial test is to demonstrate the proposed conjecture within less approximative solutions of general relativity. Ideally, a thorough derivation of the conjecture from theory would involve a numerical computation employing the full Einstein equations starting with some initial perturbations in the early universe. More realistically, the feasible next step is to replace some of the analytic approximations with more careful numerical procedures; we leave this analysis to a future work, where e.g.\\ the transition redshift $z_0$ will be a derived, not phenomenological, parameter. Although potentially laborious, these tests should be realizable with present-day knowledge, so the proposed conjecture is certainly falsifiable."
    },
    "0711/0711.2719_arXiv.txt": {
        "abstract": "We use a modified outer gap model to study the multi-frequency phase-resolved spectra of the Crab pulsar. The emissions from both poles contribute to the light curve and the phase-resolved spectra. Using the synchrotron self-Compton mechanism and by considering the incomplete conversion of curvature photons into secondary pairs, the observed phase-averaged spectrum from $100$ eV - $10$ GeV can be explained very well. The predicted phase-resolved spectra can match the observed data reasonably well, too. We find that the emission from the north pole mainly contributes to Leading Wing 1. The emissions in the remaining phases are mainly dominated by the south pole. The widening of the azimuthal extension of the outer gap explains Trailing Wing 2. The complicated phase-resolved spectra for the phases between the two peaks, namely Trailing Wing 1, Bridge and Leading Wing 2, strongly suggest that there are at least two well-separated emission regions with multiple emission mechanisms, i.e.~synchrotron radiation, inverse Compton scattering and curvature radiation. Our best fit results indicate that there may exist some asymmetry between the south and the north poles. Our model predictions can be examined by GLAST. ",
        "introduction": "It is generally believed that the phase-resolved spectra can provide the most detailed information about the structure of the pulsar magnetosphere, acceleration mechanism, and, pair creation and radiation processes in the outer magnetosphere. \\citet{a31} and \\citet{DH96} firstly calculated the phase-resolved spectra of the Vela pulsar. Then, \\citet{a1} and \\citet{ZC02} calculated the phase-resolved spectra of the Crab pulsar in the $\\gamma$-ray and the X-ray regimes separately. It will be interesting to find a general scenario that can produce both the $\\gamma$-ray and the X-ray regimes in the phase-resolved spectra of the Crab pulsar. Our preliminary results have been reported in the 2006 COSPAR conference \\citep{COSPAR07}. We have continued to study the phase-resolved spectra of the Crab pulsar in more details. Here, we present our recent results with a different parametric fitting \\citep{myphdthesis}.  According to \\citeauthor{CHRa} (\\citeyear{CHRa,CHRb}; hereafter CHR), the outer gap starts at the null charge surface, ends at the light cylinder, is bounded below by the last closed field line and is bounded above by a layer of electric current which replenishes charges to the open field lines outside the gap to maintain a steady charge density, the Goldreich-Julian charge density \\citep{a27}, \\begin{equation} \\rho _{\\rm GJ}\\sim-\\frac{\\mbox{\\boldmath$B$}\\cdot\\mbox{\\boldmath{$\\Omega$}}}{2\\pi c}. \\end{equation} The charge depletion within the outer gap due to global flows of charged particles causes a large electric field along the magnetic field lines so that $\\mbox{\\boldmath$E$}\\cdot\\mbox{\\boldmath$B$}\\neq 0$ inside the gap. This gap thus acts as an accelerator to boost the charged particles to relativistic speeds. Through a cascade process, high energy $\\gamma$-ray photons and e$^\\pm$ pairs are produced. Recently, the classical outer gap model is being challenged by \\citet{HHS03}. By solving the set of Maxwell and Boltzmann equations, they find that a current at nearly Goldreich-Julian rate can shift the position of the inner boundary of the outer gap. Therefore, we adopt a modified version of the CHR outer gap model such that the inner boundary of the outer gap is shifted inwards.  The photon emission mechanism starts with curvature radiation of the accelerated charged particles in the gap. The emission direction is tangent to the local magnetic field lines. As a photon escapes, it may encounter a low energy photon. The low energy photon may be a thermal photon from the stellar surface or a magnetospheric soft synchrotron photon emitted by the secondary e$^\\pm$ pairs which are created by the curvature photons from the inner field lines. The primary curvature photons will then be converted into the secondary e$^\\pm$ pairs via photon-photon pair production, i.e.~$\\gamma+\\gamma\\rightarrow{\\rm e}^++{\\rm e}^-$ or $\\gamma+X\\rightarrow{\\rm e}^++{\\rm e}^-$ inside and outside the gap. As pointed out in \\citet{a1}, although pair production inside an outer gap is limited to a small region, pair production outside the outer gap can cover a much wider area because the synchrotron photons produced by the secondary pairs are more abundant than the thermal photons from the stellar surface. Although the secondary synchrotron photons cannot get into the outer gap due to the field line curvature, they can convert most of the primary curvature photons from the outer gap into secondary pairs. Since the synchrotron radiation is beamed to a small angle by the relativistic beaming effect \\citep{radiative}, synchrotron photons will also be seen as more or less tangent to the field lines if observed.  Besides curvature radiation and synchrotron radiation, inverse Compton scattering is another important radiation mechanism in the neutron star magnetosphere. It occurs when a fast moving electron or positron collides with a photon and net energy is transferred from the particle to the photon. In the far region of the magnetosphere, the relativistic particles collide with the soft synchrotron photons through the inverse Compton scattering process.  \\citet{CR92,CR94} and \\citet{a29} calculated the light curves by considering a single outer gap with the photon emission to be beamed to the outside alone in order to avoid the multiple-peak feature that does not occur in true observation data of the Crab pulsar. \\citet{a1} pointed out the lack of a reason to explain the ignorance of the incoming photons (the photons that are beamed towards the star). However, in Section~\\ref{emigeo}, we estimated that inward emission is much fainter than the outward emission (see also \\citeauthor{a1} \\citeyear{a1}). Furthermore, since the previous models by \\citet{a29} and \\citet{a1} considered emission beyond the null charge surface only, they suggest that the observer can measure the photons from one gap only, that is, a single-pole model is considered in the canonical model. However, if the gap extends below null charge surface, the photons originating from another gap become measurable by the observer as well. In this paper, therefore, although we follow \\citet{CR92,CR94} and \\citet{a29} to calculate the light curve, we consider both gaps, that is, a two-pole model is examined.  In this paper, we calculate the pulse profile and the phase-resolved spectra for the Crab pulsar with an outer gap accelerator model. Because the Crab pulsar is one of the brightest $\\gamma$-ray sources in the sky, the detailed observation for the pulse profile and the phase-resolved spectra have been obtained. These provide useful information to the study of the non-thermal processes in the pulsar magnetosphere. Since the Crab pulsar is still young and is believed to have a thin gap, calculating the contribution from one layer suffices. By fixing an inclination angle and an observer angle, the photon emission locations that produce the light curve are used to calculate the phase-resolved photon spectra which are compared with the observation data. For the Crab pulsar, we adopt $R=10$~km for the stellar radius and $B_p=3.8 \\times 10^{12}$~Gauss for the stellar magnetic field strength.  Another major feature in this paper is the relaxing of the assumption that all curvature radiation has been converted into e$^\\pm$ pairs as in the previous studies. In fact, by letting trace amount of the curvature radiation, which is emitted far away from the star, to survive and escape, we can obtain better fitting curves for the phase-resolved spectra of the Crab pulsar, especially in the phases of Trailing Wing 1, Bridge and Leading Wing 2.  In Section~\\ref{section:lightcurve}, we review briefly the emission geometry in the magnetosphere. Then we discuss the location of the inner boundary of the outer gap and produce the light curve for the Crab pulsar. In Section~\\ref{section:spectra}, we discuss the electric field component along the magnetic field lines and the three major emission mechanisms, namely synchrotron radiation, inverse Compton scattering and curvature radiation. We argue that although most of the primary curvature photons have been converted to secondary pairs to produce synchrotron photons, some curvature photons which are emitted far away from the star can survive the cascade process and escape. They contribute to the peak in the high energy regime at several GeV of the phases Trailing Wing 1, Bridge and Leading Wing 2. In the final section, we conclude our results and discuss briefly about the justification of our assumptions. ",
        "conclusions": "We have used the modified three dimensional outer magnetosphere gap model with the inner boundary of the outer gap being extended from the null charge surface to near the stellar surface. The exact location of the inner boundary does not affect the fitting results. The ``inwardly-extended\" part of the outer gap contributes to LW1 and TW2 of the light curve with a slight modification of P1 and P2. Such modified outer gap geometry also plays a vital role in explaining the optical polarization properties of the Crab pulsar \\citep{Ta06}. Together with the results of \\citet{Ta06} on the gap accelerator, this model explains the pulse profile, phase-resolved spectra and polarization of the Crab pulsar. Also, the outer gap model can explain the observed complex morphology change of light curve as a function of the photon energy \\citep{ChTa07}.  Four adjustable parameters are used to simulate the light curve: the inclination angle of the magnetic axis to the rotational axis $\\alpha$, $a_1$, the viewing angle also to the rotational axis $\\zeta$ and the emission cone pitch angle due to geometry of the emission location $\\sin\\varphi(R_L)$. As constrained by the phase separation of the double peaks, we choose the values $\\alpha=50^\\circ$ and $\\zeta=76^\\circ$. From radio observations, \\citet{Ra93} estimated $\\alpha\\approx 84^\\circ$ with $\\zeta$ unknown. By using the polarimetric observations at frequencies between $1.4$ and $8.4$ GHz, \\citet{MH99} calculated $\\alpha\\approx 56^\\circ$ and $\\zeta=117^\\circ$. Therefore, further observations is required in order to determine these two parameters more accurately.  For the photon spectra, our model fitting requires $f(R_L)\\approx 0.2$, which is larger than the theoretical estimation ($f(R_L)\\approx 0.12$) by a factor of $\\sim 2$. However, the theory \\citep{ZC97,a1} has assumed a vacuum dipole potential. It has been pointed out that in order to explain the observed radiation power on high energy region, current must flow in the accelerator and hence the potential at the gap will be reduced \\citep{hirotani06,hirotani07,TaShHiCh06}. Consequently, the gap size with current flow should be larger than the vacuum one.  Moreover, we have taken the stellar radius of a neutron star to be $10^6$ cm. Due to the lack of a direct method to determine the size of a neutron star \\citep{LaPr04}, and the equation of state inside a neutron star given by different theoretical models do not provide a unanimous value for the neutron star size, we can only determine the magnetic moment, i.e.~$B_pR^3$ , of the pulsar from the energy loss rate. However, although the dipole radiation comes from the spin down power of the star, it is not all of it. For example, the vacuum formula of the dipole radiation for the aligned rotator  is usually used to infer the dipole moment of the pulsar. On the other hand, MHD solution \\citep{Sp06} suggested that the effect of the current increases the spindown luminosity from the  the vacuum formula. Therefore, we may have overestimated the value of the magnetic moment and therefore the strength of the magnetic field. Since $E_{syn}\\propto B(r)$ hence $B_pR^3$, the synchrotron peak may have been overestimated. Similarly, the peak synchrotron flux $F_{syn}\\propto B(r)$ may also be overestimated. Since the power of synchrotron radiation and inverse Compton scattering can be compared by the ratio of the local magnetic energy density and the photon energy density, $\\frac{P_{syn}}{P_{ICS}}\\sim\\frac{U_B}{U_\\gamma}\\propto\\frac{B^2(r)/2\\pi}{E_{syn}(r)n_{sy% n}(E_{syn},r)}$ which is independent of $B_pR^3$ shows that $F_{ICS}$ may not be affected. On the other hand, since $F_{syn}$ may have been overestimated, $n_{syn}$ should be smaller. Hence less curvature photons should be absorbed.  The Crab pulsar is one of the most important pulsars because of its close distance and strong radiation signals in all frequencies. Although observation data of it have been gathered for about 40 years, the underlying physics involved in this pulsar are not completely understood. For example, the structure of the charged particle accelerator in the magnetosphere can be different in the polar cap model, the outer gap model and the newly proposed caustic model \\citep{DR03} which proposed a two pole model with the thin gap regions starting from each polar cap to the light cylinder. We hope that further observation by the more sensitive satellites in the future missions, e.g.~GLAST, can bring more crucial and critical information so that we can determine the mechanisms that are truly taking place. For example, as shown in Figure~\\ref{fig:spectra}, we predict that there is a clear component ($30$ MeV - $30$ GeV), which is mainly dominated by the survival curvature photons, in TW1, Bridge and LW2. Our model also suggests that the high energy turn over in TW1, Bridge and LW2 is $\\sim 10$ GeV, whereas the high energy turn over of LW1 is only $\\sim 100$ MeV (much lower than the other phases). This difference results from the consequence of the radiation in LW1 being emitted from a different pole in the region between the inner boundary and the null charge surface, where the local electric field is much weaker (cf.~Eq.~\\ref{eq:Er}). With an increase in sensitivity and a widening of the energy range than EGRET, GLAST may be able to confirm our predictions."
    },
    "0711/0711.0370_arXiv.txt": {
        "abstract": "{} {In view of the substantial uncertainties regarding the possible dynamics of the dark energy, we aim at constraining the expansion rate of the universe without reference to a specific Friedmann model and its parameters.} {We show that cosmological observables integrating over the cosmic expansion rate can be converted into a Volterra integral equation which is known to have a unique solution in terms of a Neumann series. Expanding observables such as the luminosity distances to type-Ia supernovae into a series of orthonormal functions, the integral equation can be solved and the cosmic expansion rate recovered within the limits allowed by the accuracy of the data.} {We demonstrate the performance of the method applying it to synthetic data sets of increasing complexity, and to the first-year SNLS data. In particular, we show that the method is capable of reproducing a hypothetical expansion function containing a sudden transition.} {} ",
        "introduction": "Astronomical measurements can constrain cosmology through two functions, the cosmic expansion rate and the growth rate of cosmic structures. If space-time is on average homogeneous and isotropic and topologically simply connected, it must be described by a Robertson-Walker metric characterised by a time-dependent scale factor $a(t)$. General relativity enters when the scale factor is to be related to the energy content of the universe. However, the geometry of space-time, in particular the distance measures, are determined already once the scale factor and its first derivative are specified by the expansion function $H(a):=\\dot a/a$.  The dynamics of structure growth requires more than that. Linear structure growth is commonly based on the Newtonian approximation to the continuity and Euler equations, together with Poisson's equation. They provide a valid foundation if Minkowskian space-time is a good local approximation on the scales of the structures considered. Then, the total inhomogeneous energy density drives its further evolution in time, and the relation of time to observables such as redshift again requires the expansion function $H(a)$.  We are writing this to emphasise that the expansion function is the central mathematical object underlying all cosmological constraints, augmented by the assumption of local Newtonian dynamics if structure growth in the late universe is to be included. This suggests that measurements of the expansion function itself, without any reference to Friedmann models, should be possible and much more fundamental than the common constraints of cosmological parameters entering the expansion function once the density contributions to the Friedmann models are specified.  This approach may become interesting for instance in the context of dark energy, for which our lack of understanding allows an only poorly constrained, vast range of possible models. Usually, dark-energy models are characterised by a small set of parameters and placed into the cosmic expansion rate through Friedmann's equation, substituting the conventional cosmological-constant term. While this is doubtlessly reasonable when testing specific dark-energy models, the question remains interesting what can be inferred on the cosmic expansion rate from observations without any reference to a specific model for the energy content of the universe and how it may affect cosmic dynamics.  The importance of a model-independent reconstruction of the cosmic expansion rate from luminosity distance data has been largely discussed in the literature. In \\cite{STA98.1} the relations between the observational data and the expansion rate are presented, and several different techniques have been developed since then to appropriately treat the data in order to perform such a reconstruction (see,~e.g.,~\\cite{HUT99.1, HUT00.1, TEG02.1, WAN05.1}). Recent reconstruction techniques with applications to data can be found in \\cite{DAL03.1, DAL04.1}, where luminosity-distance data from type-Ia supernovae are combined with angular-diameter distances inferred from radio galaxies, \\cite{FAY06.1}, where constraints from baryon acoustic oscillations are added to supernova data, and \\cite{SHA06.1, SHA07.1}. All these methods employ a smoothing procedure in redshift bins. Also principal component analysis has been used to reconstruct the dark energy equation of state as a function of redshift \\citep{HUT03.1, CRI05.1, SIM06.1}.  We are here proposing a method aiming at a direct determination of the cosmic expansion rate without reference to any specific model for the energy content of the universe. As proposed and used here, it only employs the minimal assumptions (1) that the universe is topologically simply connected, homogeneous and isotropic on average, and (2) that the expansion rate is reasonably smooth. We introduce the method in Sect.~2 and demonstrate its performance with synthetic and real type-Ia-supernova data in Sect.~3. We finish with the discussion in Sect.~4. ",
        "conclusions": "We are proposing a method to constrain the expansion function of the universe without assuming any specific model for Friedmann-type expansion. If the universe is isotropic, homogeneous and simply connected, it is described by a Robertson-Walker metric. Cosmological measurements can generally only constrain one of two functions of time, the expansion rate and the growth of structures. We have shown here how the expansion function can be observationally constrained without reference to specific assumptions on the time evolution of the terms in Friedmann's equation and their parameterisation in terms of density parameters. The issue may become important in search of constraints for a dynamical dark energy component, whose behaviour is so far only very poorly known. Since it is unclear how its energy density contribution to the cosmic fluid may change in time, any guessed parameterisation may be erroneous, hence a parameter-free recovery of the cosmic expansion rate may turn out advantageous.  We demonstrate our method here using the luminosity-distance measurements obtained from type-Ia supernovae as a model. Since the luminosity distance is a cosmological observable depending on space-time geometry only, the dynamics of structure growth does not enter yet. The method proceeds in two essential steps. First, the integral relation between the expansion function and the luminosity distance is transformed into a Volterra integral equation of the second kind. Under the relevant conditions, its solutions are known to exist and to be uniquely described in terms of a convergent Neumann series. In other words, the method is guaranteed to return \\textit{the} unique expansion rate of the universe within the accuracy limits allowed by the data.  The drawback of the transformation to a Volterra integral equation is that the derivative of the luminosity distance with respect to the scale factor is needed to start the Neumann series. Derivatives of data are notoriously noisy and should be avoided. We propose to expand the luminosity distance into an initially arbitrary orthonormal function set, fit its expansion coefficients to the data and then use the derivative of the series expansion instead of the derivative of the data. Suitable orthornormal function sets can be constructed by Gram-Schmidt orthonormalisation from any linearly independent function set. The only condition so far is that the number of coefficients required to fit the data should be minimal.  Once the orthonormal function set is specified, the Neumann series can be constructed beforehand for all its members. The measured coefficients of the series expansion directly translate to the solution for the expansion function. The convergence criterion for the Neumann series is determined by the data quality, as is the number of orthonormal modes in the series expansion of the data.  Applications to synthetic data samples of increasing complexity are very promising. In particular, we showed that an expansion function containing a sudden transition can be faithfully recovered by our method provided sufficient quality of the input data. We also apply our method to the first-year SNLS data and show the expansion function recovered, and notice how the inclusion of higher-redshift ($z>1$) objects could improve the performance of the method.  In future studies, we shall investigate how our method can be used to determine the single expansion function underlying all cosmological measurements used."
    },
    "0711/0711.4373_arXiv.txt": {
        "abstract": "We forecast the sensitivity with which the Murchison Widefield Array (MWA) can measure the 21 cm power spectrum of cosmic hydrogen, using radiative transfer simulations to model reionization and the 21 cm signal. The MWA is sensitive to roughly a decade in scale (wavenumbers of $k \\sim 0.1 - 1 h$ Mpc$^{-1}$), with foreground contamination precluding measurements on larger scales, and thermal detector noise limiting the small scale sensitivity. This amounts primarily to constraints on two numbers: the amplitude and slope of the 21 cm power spectrum on the scales probed.  We find, however, that the redshift evolution in these quantities can yield important information about reionization.  We examine a range of theoretical models, spanning plausible uncertainties in the nature of the ionizing sources, and the abundance of gas-rich mini-halos during reionization. Although the 21 cm power spectrum differs substantially among these models, a generic prediction is that the amplitude of the 21 cm power spectrum on MWA scales ($k \\sim 0.4 h$ Mpc$^{-1}$) peaks near the epoch when the intergalactic medium (IGM) is $\\approx 50\\%$ ionized.  Moreover, the slope of the 21 cm power spectrum on MWA scales flattens as the ionization fraction increases and the sizes of the HII regions grow. Considering detection sensitivity, we show that the optimal MWA antenna configuration for power spectrum measurements would pack all $500$ antenna tiles as close as possible in a compact core. Provided reionization occurs in the MWA observing band, this instrument is sensitive enough in its optimal configuration to measure redshift evolution in the slope and amplitude of the 21 cm power spectrum. Detecting the characteristic redshift evolution of our models will help confirm that observed 21 cm fluctuations originate from the IGM, and not from foregrounds, and will provide an indirect constraint on the evolution of the volume-filling factor of HII regions during reionization. After two years of observations under favorable conditions, the MWA can constrain the filling factor at an epoch when $\\avg{x_i} \\sim 0.5$ to within roughly $\\pm \\delta \\avg{x_i} \\sim 0.1$ at $2-\\sigma$ confidence. It can also constrain models for the ionizing sources and the abundance of mini-halos during reionization. ",
        "introduction": "\\label{sec:intro}  Detecting 21 cm emission from the high redshift IGM will provide fully three-dimensional information on the epoch of reionization (EoR). Futuristic experiments like the Square Kilometer Array (SKA) will have the sensitivity to produce maps of reionization as a function of redshift (Zaldarriaga et al. 2004, McQuinn et al. 2006). These maps, taken in various frequency bands and with sufficiently high angular resolution, will amount to a reionization `movie': they will depict the growth of HII regions around individual sources and their subsequent mergers with neighboring HII regions, and detail the completion of the reionization epoch, whereby the entire volume of the Universe becomes filled with ionized hydrogen.  While producing a reionization movie is perhaps the ultimate goal of 21 cm studies, first generation surveys like the MWA\\footnote{The array formerly known as the Mileura Widefield Array, is now known as the Murchison Widefield Array, owing to a change of site.}  will already provide valuable insights into the reionization process. These first generation surveys will lack the sensitivity required to produce detailed maps, but will allow for a statistical detection.  For example, the MWA is expected to measure the power spectrum of 21 cm brightness temperature fluctuations over roughly a decade in scale, in each of several redshift bins (McQuinn et al. 2006, Bowman et al. 2006).  In this paper, we use radiative transfer simulations (Sokasian et al. 2001, 2003, Zahn et al. 2007, McQuinn et al. 2007a, 2007b) to quantify how well power spectrum measurements with the MWA can constrain the reionization process. How well can we deduce the gross features of the reionization movie from a statistical detection?  In particular, we aim to check whether the MWA can constrain the volume filling factor of HII regions as a function of time.  Reliable estimates of the volume filling factor of HII regions -- the `ionization fraction' of the IGM -- will pinpoint the timing and duration of reionization.  This information, marking a key event in the formation of structure in our universe, will constrain models for the ionizing sources, and solidify existing measurements from cosmic microwave background (CMB) polarization (Page et al. 2007), quasar spectra (Fan et al. 2006, Lidz et al. 2006, 2007a, Becker et al. 2007, Bolton \\& Haehnelt 2007a, 2007b, Mesinger \\& Haiman 2007, Wyithe et al. 2005, Wyithe et al. 2007, Gallerani et al. 2007), Ly-$\\alpha$ emitter (LAE) surveys (Malhotra \\& Rhoads 2005, Kashikawa et al. 2006, Furlanetto et al. 2006c, Dijkstra et al. 2007, McQuinn et al. 2007b, Mesinger \\& Furlanetto 2007a), and gamma-ray burst (GRB) optical afterglow spectra (Totani et al. 2006, McQuinn et al. 2007c.), which currently provide tantalizing, yet controversial and subtle to interpret, clues.  The outline of this paper is as follows. In \\S \\ref{sec:rt_sig} we describe our radiative transfer simulations, and examine the simulated 21 cm power spectrum and its redshift evolution. In \\S \\ref{sec:pspec_noise} we forecast the sensitivity with which the MWA can measure the 21 cm power spectrum, paying attention to foreground contamination, and examining how the sensitivity depends on the, as yet unfinalized, configuration of the MWA's antenna tiles (\\S \\ref{sec:array_config}).  In \\S \\ref{sec:amp_slope} we show that the MWA sensitivity generally boils down to constraints on each of the slope and amplitude of the 21 cm power spectrum at $k \\sim 0.1 -1 h$ Mpc$^{-1}$.  We then quantify how well the MWA can constrain these two numbers in different redshift bins. In \\S \\ref{sec:fill_hii}, we show how constraints on the slope and amplitude of the 21 cm power spectrum in several redshift bins translate into constraints on the redshift evolution of the volume-filling factor of HII regions during reionization. In \\S \\ref{sec:conclusions} we summarize our main results and discuss future research directions.  Throughout we consider a $\\Lambda$CDM cosmology parameterized by: $n_s = 1$, $\\sigma_8 = 0.8$, $\\Omega_m = 0.27$, $\\Omega_\\Lambda = 0.73$, $\\Omega_b = 0.046$, and $h=0.7$, (all symbols have their usual meanings), consistent with the results from Spergel et al. (2007). ",
        "conclusions": "\\label{sec:conclusions}  In this paper, we considered the sensitivity of the MWA for 21 cm power spectrum measurements and examined the resulting insights into reionization.  Rather generically, the 21 cm power spectrum on MWA scales increases in amplitude with decreasing redshift, as HII regions grow and boost the level of large-scale fluctuations in the 21 cm signal, until the time at which $\\sim 50\\%$ of the volume of the IGM is ionized. Subsequently, later in the EoR, the amplitude of the 21 cm power spectrum drops as neutral hydrogen becomes scarce. In conjunction, the slope of the 21 cm power spectrum flattens as HII regions grow and the 21 cm power spectrum transitions, on the scales of interest, from simply tracing density fluctuations to tracing fluctuations in the ionization field.  These generic features imply that, although MWA measurements will be limited to a dynamic range of $\\sim$ a decade in scale, the experiment can significantly constrain reionization by comparing measurements in several redshift bins. In particular, the MWA may detect the power spectrum amplitude rise with decreasing redshift, before subsequently turning-over and dropping in amplitude, with the slope of the power spectrum flattening in conjunction.  This behavior is a signature of the IGM passing through a redshift where $\\sim 50\\%$ of its volume is filled with ionized bubbles.  It is unlikely that residual foreground contamination would share the characteristic redshift evolution in power spectrum amplitude and slope found in our models. Measuring this characteristic redshift evolution will hence help confirm that detected 21 cm fluctuations originate from the high redshift IGM, and cement the case for neutral material in the high redshift IGM.  We argued that the MWA is sensitive enough to detect this characteristic redshift evolution in the 21 cm power spectrum amplitude and slope, especially if it adopts a very compact configuration for its antenna tiles, and that reionization occurs at sufficiently moderate redshifts. A compact configuration optimizes power spectrum sensitivity because proposed configurations, with antennae distributed out to large $\\sim 1$ km distances, populate long baselines too sparsely to be useful.  Some antenna tiles are needed at long baselines for point source detection and antenna calibration, but we estimate that these requirements are not severe.  Another possible advantage of the super-core configuration relates to its well-behaved beam. A potential difficulty for 21 cm observations is that point sources far from the primary beam of an interferometer can enter through the beam's sidelobes, which are frequency dependent. This contamination will have structure in frequency space, and escape the usual foreground removal strategies (e.g. Oh \\& Mack 2003, Furlanetto et al. 2006a). This contamination will be significantly reduced by adopting a very compact antenna distribution, like the super-core configuration, which will have a very clean beam with minimal sidelobes.  Adopting a sufficiently compact antenna distribution, we believe that the MWA can move past a mere detection of 21 cm fluctuations, and constrain the volume filling factor of HII regions within two years of observing. In particular, measurements of the 21 cm power spectrum amplitude and slope at a redshift where $50\\%$ of the volume is ionized translate into $2-\\sigma$ allowed ionization fractions of roughly $\\avg{x_i} = 0.45-0.65$.  Moreover, we find that the MWA can rather easily distinguish between our fiducial model and each of our rare source and mini-halo models with $2,000$ hours of data, at least if the bulk of reionization occurs in the range of redshifts probed by the MWA.   In this paper, we focused on the 21 cm power spectrum, but investigating other statistical measures would be interesting.  One of the primary goals of reionization studies is to use observational measures to constrain the sizes and filling factor of HII regions during reionization. Clearly the ionization field, and the 21 cm signal, are highly non-Gaussian in the EoR, implying that there is, in principle, more information than contained in the 21 cm power spectrum alone.  We have argued that the redshift evolution of the 21 cm power spectrum contains relatively robust information regarding the filling factor of HII regions during reionization, but these constraints may not be optimal and are rather indirect.  On the other hand, it is unclear how beneficial higher order statistics will be in the low $S/N$ regime relevant for the MWA and other first generation surveys. We plan further investigation regarding the utility of various non-Gaussian statistical measures, considered as a function of instrumental sensitivity. Particularly interesting is that some low-$k$ bins become sample-variance dominated in the super-core configuration, implying that the MWA can actually {\\em image} large scale modes in this configuration.  Another useful endeavor would be to perform a mock MWA simulation, incorporating thermal noise, foreground models, and the MWA instrumental response and observing strategy. Our forecasts are inevitably simplified in ignoring these details, and likely optimistic in this regard. Mock simulations should quantify how much the sensitivity of the MWA is degraded when incorporating such real world details. Moreover, it would be interesting to examine how possible residuals from incomplete foreground cleaning might bias our constraints. Since the power spectrum variance increases strongly with wavenumber, the total $S/N$ at which the MWA can detect the 21 cm power spectrum depends strongly on the precise foreground cut. This is an important issue for further investigation.  Finally, in this paper we focused on a single parameter, the filling factor of ionized regions at different redshifts, and gave simple arguments for how one can constrain this quantity with the redshift evolution of the 21 cm power spectrum alone. A complementary approach for forecasting MWA parameter constraints would be to perform a multi-parameter Fisher matrix analysis.  We anticipate that the MWA will swiftly move beyond a mere detection of 21 cm emission from the high redshift IGM, and obtain valuable insights regarding the filling factor of HII regions at different stages of reionization."
    },
    "0711/0711.3881_arXiv.txt": {
        "abstract": "Dark matter (comprising a quarter of the Universe) is usually assumed to be due to \\emph{one and only one} weakly interacting particle which is neutral and absolutely stable.  We consider the possibility that there are several coexisting dark-matter particles, and explore in some detail the generic case where there are \\emph{two}. We discuss how the second dark-matter particle may relax the severe constraints on the parameter space of the Minimal Supersymmetric Standard Model (MSSM), as well as other verifiable predictions in both direct and indirect search experiments. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.4974_arXiv.txt": {
        "abstract": "{Study the surface properties of the transneptunian population of Solar-system bodies} { We investigate the surface characteristics of the large dwarf planet (136199)\\,Eris.  } { With the FORS1 instrument of the ESO VLT, we have obtained Bessell broadband $R$ linear polarimetry and broadband $V$ and $I$ photometry of Eris. We have modelled the observations in terms of the coherent-backscattering mechanism to constrain the surface properties of the object.} { Polarimetric observations of Eris show a small negative linear polarization without opposition surge in the phase angle range of $0.15-0.5\\degr$.  The photometric data allow us to suppose a brightness opposition peak at phase angles below $0.2-0.3\\degr$. The data obtained suggest possible similarity to the polarimetric and photometric phase curves of Pluto. The measured absolute magnitude and broadband colors of Eris are $H_V$=-1.15 mag, $V-R=0.41$ mag, and $V-I=0.75$ mag.} {The observational data with theoretical modelling are in agreement with the surface of Eris being covered by rather large inhomogeneous particles.} ",
        "introduction": "The discovery of the Pluto-sized transneptunian object 2003~UB$_{313}$ \\citep{Broetal05}, later named (136199)~Eris, was succeeded by an intense physical characterization phase and no less than a discussion on the definition of a planet.  After the resolution of IAU (24 August 2006) Eris is classified, together with Pluto, as a member of the new category of ``the dwarf planets''. Eris was discovered near its aphelion at 97\\,AU from the Sun and has an orbit with high eccentricity (0.44) and inclination ($44\\degr$). According to its dynamical characteristics, it belongs to the so-called ``detached objects'' with pericenters decoupled from Neptune \\citep{Glaetal07}.  The first size estimation of Eris was carried out by \\citet{Beretal06} based on thermal emission, resulting in a diameter of $3000 \\pm 400$\\,km and a geometric albedo of $0.60 \\pm 0.15$. Later on, \\citet{Broetal06}, using the Hubble Space Telescope, measured directly the angular size of Eris corresponding to the diameter of $2400 \\pm 100$\\,km, about 5\\,\\% larger than that of Pluto, and derived the albedo of $0.87 \\pm 0.05$ in the $V$ band. The discrepancy in the albedo is rather large but, undoubtedly, Eris is a high-albedo object like Pluto \\citep[see][]{Buretal03}. According to \\citet{BroSchal07}, the mass of Eris exceeds the mass of Pluto by a factor of 1.27.  The spectral properties of Eris are also found to be rather similar to those of Pluto. The near-infrared spectrum is dominated by methane ice absorptions \\citep{Broetal05}. However, the central wavelengths of the CH$_{4}$ bands do not show the shift in wavelength seen in the spectra of Pluto and Triton which is attributed to the fact that most (Triton) or some (Pluto) of the CH$_{4}$ is dissolved in solid N$_{2}$. The quality of Eris' spectrum does not allow the detection of N$_{2}$. \\citet{Licetal06} reported a visible spectrum of Eris in which the CH$_{4}$ ice bands show a small shift in the band at around 0.89\\,$\\mu$m and they suggested that the shift could indicate the presence of methane diluted in N$_{2}$. \\citet{Dumetal07}, using the new SINFONI instrument installed at the ESO VLT (Paranal Observatory, Chile), obtained a high-quality near-infrared spectrum and, on the basis of modelling, they obtained the best fit of the spectrum including N$_{2}$ on the surface of Eris. Due to the highly eccentric orbit, they expect a thin atmosphere close to perihelion, which freezes out onto the surface when the object is located at a larger heliocentric distance. They modelled the surface using two types of terrains of distinct composition: about $50\\,\\%$ covered with pure methane ice, while the other part is made of an intimate mixture of methane, nitrogen, and water and tholin ices. The best model fits are obtained when the icy grains are rather large, from sub-mm to a few tens of mm in size \\citep{Dumetal07}.  In the present paper, we report the first polarimetric observations of Eris made as part of the Large-Program observations at ESO VLT, aimed to set constraints on its surface properties.  \\begin{table*} \\caption{Polarimetry of (136199)\\,Eris in the Bessell $R$ band. $\\pq$ and $\\pu$ are the transformed Stokes parameters in a way that $\\pq$ represents the flux perpendicular to the plane Sun-Object-Earth (the scattering plane) minus the flux parallel to that plane, divided by the sum of the two fluxes, and $\\Theta$ is the position angle measured counterclockwise from the perpendicular to the scattering plane.} \\label{Tab_Pol} \\centering \\begin{tabular}{ccccrrr} \\hline \\hline Date                             & Time (UT)                        & Sky                              & Phase angle                      & \\multicolumn{1}{c}{$\\pq$}        & \\multicolumn{1}{c}{$\\pu$}        & \\multicolumn{1}{c}{$\\Theta$}     \\\\ \\multicolumn{1}{c}{(yyyy mm dd)} & \\multicolumn{1}{c}{(hh:mm)}& \\multicolumn{1}{c}{}        & \\multicolumn{1}{c}{(DEG)} & \\multicolumn{1}{c}{(\\%)} & \\multicolumn{1}{c}{(\\%)}         & \\multicolumn{1}{c}{(DEG)}        \\\\ \\hline 2006 10 19 & 03:27 & PHO & 0.15 & $-0.09 \\pm 0.07$ & $ 0.00 \\pm 0.07$ & $ 91 \\pm 23$ \\\\ 2006 11 12 & 02:21 & CLR & 0.30 & $-0.06 \\pm 0.07$ & $-0.12 \\pm 0.07$ & $122 \\pm 19$ \\\\ 2006 11 18 & 02:25 & THN & 0.35 & $-0.32 \\pm 0.12$ & $ 0.00 \\pm 0.10$ & $ 90 \\pm  9$ \\\\ 2006 12 14 & 01:53 & PHO & 0.51 & $-0.12 \\pm 0.07$ & $-0.05 \\pm 0.07$ & $100 \\pm 21$ \\\\ \\hline \\end{tabular} \\end{table*} \\begin{table} \\caption{Photometry of (136199)\\,Eris. The errors on $R$ magnitudes are estimated \\textit{a priori} (see text).}\\label{Tab_Phot} \\centering \\begin{tabular}{cccccr} \\hline \\hline Date                & Time (UT)  & Phase angle& Band       & Magnitude  &\\\\ \\multicolumn{1}{c}{ (yyyy mm dd)} & \\multicolumn{1}{c}{(hh:mm)} & \\multicolumn{1}{c}{(DEG)} & & & \\\\ \\hline 2006 10 19 & 03:27 & 0.15 &$R$&$18.34 \\pm 0.05$ \\\\ 2006 10 20 & 05:42 & 0.16 &$V$&$18.74 \\pm 0.02$ \\\\ 2006 10 20 & 05:45 & 0.16 &$I$&$18.00 \\pm 0.02$ \\\\ 2006 10 20 & 05:49 & 0.16 &$V$&$18.75 \\pm 0.02$ \\\\ 2006 10 20 & 05:52 & 0.16 &$I$&$17.99 \\pm 0.02$ \\\\ 2006 11 12 & 02:21 & 0.30 &$R$&$18.37 \\pm 0.05$ \\\\ 2006 12 14 & 01:53 & 0.51 &$R$&$18.40 \\pm 0.05$ \\\\ \\hline \\end{tabular} \\end{table} ",
        "conclusions": "\\label{Sect_Discussion} The present polarimetric observations of Eris constitute the first observations of a high-albedo transneptunian object at very small phase angles. A sharp surge of negative polarization expected for high-albedo objects within the mechanism of coherent backscattering was not found in the measured range of phase angles ($0.15-0.5\\degr$). It was also shown that, if a brightness opposition peak exists for Eris, it must be constrained to phase angles below $0.2-0.3\\degr$. The data obtained for Eris suggest possible similarity with the polarimetric and photometric phase curves of Pluto. Their modelling puts constraints on the surface properties suggesting that the regoliths of Eris and Pluto most probably consist of transparent inhomogeneous particles large compared to the wavelengths of visible light."
    },
    "0711/0711.1410_arXiv.txt": {
        "abstract": "We systematically perform the merger simulation of black hole-neutron star (BH-NS) binaries in full general relativity, focusing on the case that the NS is tidally disrupted.  We prepare BH-NS binaries in a quasicircular orbit as the initial condition in which the BH is modeled by a nonspinning moving puncture. For modeling the NS, we adopt the $\\Gamma$-law equation of state with $\\Gamma=2$ and the irrotational velocity field.  We change the BH mass in the range $M_{\\rm BH} \\approx 3.3$--$4.6M_{\\odot}$, while the rest mass of the NS is fixed to be $M_{*}=1.4 M_{\\odot}$ (i.e., the NS mass $M_{\\rm NS} \\approx 1.3M_{\\odot}$).  The radius of the corresponding spherical NS is set in the range $R_{\\rm NS} \\approx 12$--15 km (i.e., the compactness $GM_{\\rm NS}/R_{\\rm NS}c^2 \\approx 0.13$--0.16). We find for all the chosen initial conditions that the NS is tidally disrupted near the innermost stable circular orbit.  For the model of $R_{\\rm NS}=12$ km, more than 97 \\% of the rest mass is quickly swallowed into the BH and the resultant torus mass surrounding the BH is less than $0.04M_{\\odot}$. For the model of $R_{\\rm NS} \\approx 14.7$ km, by contrast, the torus mass is about $0.16M_{\\odot}$ for the BH mass $\\approx 4M_{\\odot}$.  The thermal energy of the material in the torus increases by the shock heating occurred in the collision between the spiral arms, resulting in the temperature $10^{10}$--$10^{11}$ K. Our results indicate that the merger between a low-mass BH and its companion NS may form a central engine of short gamma-ray bursts (SGRBs) of the total energy of order $10^{49}$ ergs if the compactness of the NS is fairly small $\\alt 0.145$. However, for the canonical values $M_{\\rm NS}=1.35M_{\\odot}$ and $R_{\\rm NS}=12$ km, the merger results in small torus mass, and hence, it can be a candidate only for the low-energy SGRBs of total energy of order $10^{48}$ ergs. We also present gravitational waveforms during the inspiral, tidal disruption of the NS, and subsequent evolution of the disrupted material. We find that the amplitude of gravitational waves quickly decreases after the onset of tidal disruption. Although the quasinormal mode is excited, its gravitational wave amplitude is much smaller than that of the late inspiral phase. This reflects in the fact that the spectrum amplitude sharply falls above a cut-off frequency which is determined by the tidal disruption process. We also find that the cut-off frequency is 1.25--1.4 times larger than the frequency predicted by the study for the sequence of the quasicircular orbits and this factor of the deviation depends on the compactness of the NS. ",
        "introduction": "The merger of black hole (BH)-neutron star (NS) binaries is one of the most promising sources of kilo-meter size laserinterferometric gravitational wave detectors such as LIGO and VIRGO.  Although such system has not been observed yet in our Galaxy in contrast to NS-NS binaries \\cite{Stairs}, statistical studies based on the stellar evolution synthesis suggest that the merger will happen 1--10\\% as frequently as the merger of the NS-NS binaries in the universe \\cite{BHNS}. Because the BH mass should be more than twice as large as the canonical NS mass $\\sim 1.35M_{\\odot}$, the typical amplitude of gravitational waves from the BH-NS binaries will be larger than that from the NS-NS binaries even if the distance is larger. This indicates that the detection rate of the BH-NS binaries by the gravitational wave detectors may be more than $10\\%$ as high as that for the NS-NS binaries, suggesting that the detection of such system will be achieved in the near future.  The final fate of the BH-NS binaries is divided into two cases depending primarily on the BH mass: (1) when the BH mass is small enough, the NS will be tidally disrupted before it is swallowed by the BH; (2) when the BH mass is large enough, the NS will be swallowed into the BH without tidal disruption.  The tidal disruption occurs when the tidal force of the BH is stronger than the self-gravity of the NS. Such condition is approximately written as \\beqn {M_{\\rm BH}R_{\\rm NS} \\over r^3} \\geq C^2 {M_{\\rm NS} \\over R_{\\rm NS}^2}, \\label{eq1.0} \\eeqn where $M_{\\rm NS}$ is the NS mass, $R_{\\rm NS}$ is the circumferential radius of the NS, $M_{\\rm BH}$ is the BH mass, $r$ is the orbital separation, and $C$ is a nondimensional coefficient. Using the orbital angular velocity $\\Omega$, Eq. (\\ref{eq1.0}) may be written as \\beqn \\Omega^2 {M_{\\rm BH} \\over M_0} \\geq C^2 {GM_{\\rm NS} \\over R_{\\rm NS}^3}, \\eeqn where $G$ is the gravitational constant and $M_0=M_{\\rm BH}+M_{\\rm NS}$. Actually, the latest high-precision numerical study for the quasicircular states of the BH-NS binaries shows that the tidal disruption occurs if the following condition is satisfied \\cite{TBFS,TBFS2}: \\beqn \\Omega \\geq C \\biggl({GM_{\\rm NS} \\over R_{\\rm NS}^3}\\biggr)^{1/2} \\biggl(1+{M_{\\rm NS} \\over M_{\\rm BH}}\\biggr)^{1/2}, \\eeqn where the value of $C$ is $\\approx 0.270$ for $\\Gamma=2$ polytropic equation of state (EOS).  We note that in \\cite{TBFS,TBFS2} the NS is assumed to be irrotational and the BH spin is set to be zero.  According to an approximate general relativistic study \\cite{ISM}, the value of $C$ depends weakly on the stiffness of the EOSs for $2 \\leq \\Gamma \\leq 3$, and hence, it would be close to 0.27 for the NS which likely has such stiff EOS.  Substituting the typical values considered in this paper, the tidal disruption is expected to set in at \\beqn && {G M_0 \\Omega \\over c^3}=0.0708 \\biggl({C \\over 0.270}\\biggr) \\biggl({M_{\\rm NS} \\over 1.30M_{\\odot}}\\biggr)^{3/2} \\nonumber \\\\ &&~~~~~~~~~~~~~\\times \\biggl({R_{\\rm NS} \\over 13.0~{\\rm km}}\\biggr)^{-3/2} \\biggl({q \\over 1/3}\\biggr)^{-1} \\biggl( \\frac{1+q}{4/3} \\biggr)^{3/2}, \\label{eq1.1} \\eeqn where $c$ is the speed of the light and $q$ denotes the mass ratio $q \\equiv M_{\\rm NS}/M_{\\rm BH}$. The corresponding frequency of gravitational waves is calculated from $f \\equiv \\Omega/\\pi$ as \\beqn && f_{\\rm tidal}=8.79 \\times 10^2 {\\rm Hz} \\biggl({C \\over 0.270}\\biggr) \\biggl({M_{\\rm NS} \\over 1.30M_{\\odot}}\\biggr)^{1/2} \\nonumber \\\\ &&~~~~~~~~~~~~~~~\\times \\biggl({R_{\\rm NS} \\over 13.0~{\\rm km}}\\biggr)^{-3/2} \\biggl({1+q \\over 4/3}\\biggr)^{1/2}. \\label{eq1.2} \\eeqn According to the third post-Newtonian (PN) study, the innermost stable circular orbit (ISCO) for $q \\approx 1/3$ is located at $GM_0\\Omega/c^3 \\approx 0.11$ \\cite{Luc}.  The tidal effect reduces this value slightly to be $GM_0\\Omega/c^3 \\sim 0.08$--0.09 \\cite{TBFS,TBFS2}. Adopting this value, we expect that the tidal disruption of the NSs of mass $1.3M_{\\odot}$ occurs for $q \\agt 1/4$ if $R_{\\rm NS}=13$ km, and for $q \\agt 1/3$ if $R_{\\rm NS}=11$ km.  This indicates that the tidal disruption occurs only for the binaries of low-mass BH and NS at an orbit very close to the ISCO.  The tidal disruption of NSs by a BH has been studied with great interest because of the following reasons: (i) Gravitational waves during the tidal disruption may bring the information about the NS radius because the orbital frequency at the onset of tidal disruption depends strongly on the compactness of the NS ($G M_{\\rm NS}/R_{\\rm NS} c^2$) \\cite{ISM,TBFS,TBFS2}. The NS mass will be determined by the data analysis for observed gravitational waves in the inspiral phase \\cite{CF}. If the NS radius could be determined from observed gravitational waves emitted during the tidal disruption, the resultant relation between the mass and the radius of the NSs may constrain the EOSs of the high density nuclear matter \\cite{lindblom,valli}. (ii) The tidally disrupted NSs may form a disk or torus of mass larger than $0.01M_{\\odot}$ around the BH if the tidal disruption occurs outside the ISCO. Systems consisting of a rotating BH surrounded by a massive, hot torus have been proposed as one of the likely sources for the central engine of gamma-ray bursts with a short duration \\cite{grb2} (hereafter SGRBs). Hence, the merger of a low-mass BH and its companion NS can be a candidate central engine. According to the observational results by the {\\it Swift} and {\\it HETE-2} satellites \\cite{Swift}, the total energy of the SGRBs is larger than $\\sim 10^{48}$ ergs, and typically $10^{49}$--$10^{50}$ ergs. The question is whether or not the mass and thermal energy of the torus are large enough for driving the SGRBs of such huge total energy.  The tidal disruption of NSs by a BH has been investigated in the Newtonian \\cite{Newton} and approximately general relativistic \\cite{FBST,FBSTR} simulations. However, the criterion of the tidal disruption and the evolution of the tidally disrupted NS material would depend strongly on general relativistic effects around the BH, and hence, a simulation in full general relativity is obviously required. In the previous paper \\cite{SU06}, we presented our first numerical results for fully general relativistic simulation performed by our new code which had been improved from the code used for the NS-NS merger \\cite{STU0,STU}: We handle an orbiting BH employing the moving puncture method, which has been recently developed (e.g., \\cite{BB2,BB4}).  We prepared a quasicircular state composed of a nonspinning moving puncture BH and a corotating NS as the initial condition. We computed the BH-NS quasicircular state by employing our new formalism based on the moving puncture framework.  From the merger simulation, we found that a torus with mass $\\sim 0.1M_{\\odot}$ is an outcome for $M_{\\rm NS}=1.3M_{\\odot}$ and $R_{\\rm NS}=13$--14 km, and for $M_{\\rm BH}=3$--$4M_{\\odot}$.  In the previous work, we assumed that the NSs are corotating around the center of mass of the system for simplicity. However, this is not very realistic velocity field for the NS in close compact binaries \\cite{KBC}. Recently, we have developed a new code for computing quasicircular states of a BH and an irrotational NS. For obtaining the irrotational velocity field of the NS, one has to solve the equation for the velocity potential \\cite{ST}. We employ the same method as that for computing the NS-NS binaries (e.g., \\cite{GGTMB,TG}), which has been used in our project for the NS-NS merger \\cite{STU0,STU}. (For solving the basic equations for the initial data, we use the spectral method library LORENE developed by the Meudon relativity group \\cite{LORENE}.)  In this paper, we perform the simulations employing these new quasicircular states as the initial conditions. In addition, we modify our simulation code for the Einstein evolution equation to the fourth-order scheme for more accurate computation. Changing the BH mass and NS radius, we systematically investigate the dependence of the torus mass formed after the merger on $M_{\\rm BH}$ and $R_{\\rm NS}$. We also compute gravitational waveforms during the tidal disruption and their spectrum.  The paper is organized as follows.  Section II summarizes the initial conditions chosen in this paper.  Section III briefly describes the formulation and numerical methods for the simulation. Section IV presents the numerical results of the simulation for the merger of BH-NS binaries. Section V is devoted to a summary.  In the following, the geometrical units of $c=G=1$ are used. ",
        "conclusions": "We have presented the numerical results of fully general relativistic simulation for the merger of BH-NS binaries, focusing on the case that the NS is tidally disrupted by a nonspinning low-mass BH of $M_{\\rm BH}\\approx 3.3$--$4.6M_{\\odot}$.  The $\\Gamma$-law EOS with $\\Gamma=2$ and irrotational velocity field are employed for modeling the NSs. To see the dependence of the tidal disruption event on the NS radius, we choose it in a wide range as $R_{\\rm NS}=12.0$--1.47 km whereas the NS mass is fixed to be $\\approx 1.3M_{\\odot}$. The resulting mass ratio $q \\equiv M_{\\rm NS}/M_{\\rm BH}$ is in the range between $\\approx 0.28$ and 0.4.  As predicted by the study for the quasicircular states \\cite{TBFS,TBFS2}, for all the chosen models, the NS is tidally disrupted at the orbits close to the ISCO.  The BH of the spin of 0.5--0.6 is formed and a large fraction of the material is quickly swallowed into the BH, whereas 2--12\\% of the material forms a hot and compact torus around the BH. The resultant mass and density of the torus depend strongly on the mass ratio $q$ and NS radius $R_{\\rm NS}$, in particular on $R_{\\rm NS}$. For $R_{\\rm NS}=12.0$ km, the torus mass is at most $0.05M_{\\odot}$ even for the large value of the mass ratio $q \\approx 0.39$. For $R_{\\rm NS}=14.7$ km, by contrast, the torus mass is $\\approx 0.16M_{\\odot}$ for $q \\approx 0.33$. Extrapolating the results in this paper, the torus mass would be smaller than $0.01M_{\\odot}$ for $R_{\\rm NS}=11$ km even for $q \\sim 0.4$.  The stiff nuclear EOSs predict the radius of $\\approx 11$--12 km for $M_{\\rm NS}=1.3$--$1.4M_{\\odot}$ \\cite{EOS}. This suggests that the torus mass is likely to be $\\ll 0.1M_{\\odot}$ even for the BH mass 3--$4M_{\\odot}$. This point should be more rigorously clarified for the future simulation employing the realistic nuclear EOSs.  For the optimistic cases in which $R_{\\rm NS} \\agt 13$ km or of $q \\agt 0.4$, the torus mass can be $\\agt 0.05M_{\\odot}$. As we found from the value of the $\\varep_{\\rm th}$, the resultant torus likely has high temperature as $10^{10}$--$10^{11}$ K.  This suggests that such outcomes are promising candidates for driving SGRBs. According to the latest simulations for the hot and compact torus around the BH by Setiawan et al. \\cite{GRBdisk1}, the total energy deposited by neutrino-antineutrino annihilation is $\\sim 10^{49} (M_{\\rm torus}/0.01M_{\\odot})$ ergs for $0.01 M_{\\odot} \\alt M_{\\rm torus} \\alt 0.1M_{\\odot}$ and $a=0.6$. Here, $M_{\\rm torus}$ is the initial torus mass for the simulation. Assuming that the conversion rate to the gamma-ray energy is 10\\% \\cite{Aloy}, the total energy of SGRBs is $\\sim 10^{48} (M_{\\rm torus}/0.01M_{\\odot})$ ergs according to their numerical results.  This indicates that if the NS radius is $\\agt 13$ km with $M_{\\rm BH}=3.3$--$4M_{\\odot}$, the SGRB of the total energy $\\sim 10^{49}$ ergs may be driven.  By contrast, for $R_{\\rm NS} \\alt 12$ km with $M_{\\rm BH} \\agt 3.3 M_{\\odot}$, the total energy is likely to be at most several $\\times 10^{48}$ ergs; i.e., only the weak SGRBs can be explained.  We note that the final spin parameter depends on the initial spin of the BH.  In the merger of the rapidly rotating BH and NS, the resuling spin of the BH will be close to unity. For such BHs, the total deposited energy will be enhanced by the spin effects \\cite{GRBdisk3}.  Gravitational waves are also computed. It is found that the amplitude of gravitational waves damps during the tidal disruption, and when the tidal disruption completes, the amplitude becomes much smaller than that at the last inspiral orbit.  Although we find that the waveforms in the final phase of the tidal disruption are characterized by the QNM ringing, its amplitude is much smaller than that of gravitational waves at the last inspiral phase, in particular, for the large value of $R_{\\rm NS}$.  The reason for this is that the NS is tidally disrupted before plunging into the BH, and hence, a significant excitation of the QNM by the coherently infalling material is not achieved in contrast to the BH-BH merger (e.g., \\cite{BB2,BB4,BHBH0,BHBH1}).  For the case that the NS is compact with $R_{\\rm NS}=12$ km, the tidal disruption occurs at an orbit very close to the ISCO and a large amount of the material falls fairly coherently into the BH, resulting in a relatively high amplitude of the QNM. However, the amplitude is still much smaller than that at the last inspiral orbit.  The Fourier spectrum of gravitational waves is also analyzed. Because the amplitude of gravitational waves quickly damps during the tidal disruption, the spectrum amplitude steeply decreases above a cut-off frequency $f_{\\rm cut}$. The noteworthy feature is that the cut-off frequency does not agree with the predicted frequency at which the tidal disruption sets in, $f_{\\rm tidal}$. The reason for this is that the NS in the compact binaries has a radial-approach velocity of order $\\sim 10\\%$ of the orbital velocity and is not immediately tidally disrupted at $f=f_{\\rm tidal}$. We have found that $f_{\\rm cut}/f_{\\rm tidal}$ is in the range between 1.25 and 1.4 in our results depending primarily on the NS radius $R_{\\rm NS}$; this value is larger for the smaller values of $R_{\\rm NS}$. These results imply that it is not straightforward to determine $f_{\\rm tidal}$ from gravitational wave observation. If the dependence of $f_{\\rm cut}/f_{\\rm tidal}$ on the NS radius and mass ratio is clarified in detail by the numerical simulation, $f_{\\rm tidal}$ may be inferred from $f_{\\rm cut}$. If such a relation is found, $f_{\\rm cut}$ will be useful for determining the properties of NSs such as the radius and density profile. For this purpose, a detailed simulation taking into account the nuclear EOSs is necessary. We plan to perform such simulation in the next step.  \\vskip 1mm"
    },
    "0711/0711.2908_arXiv.txt": {
        "abstract": "As suggested by some extensions of the Standard Model of particle physics, dark matter may be a super-weakly interacting lightest stable particle, while the next-to-lightest particle (NLP) is charged and meta-stable.  One could test such a possibility with neutrino telescopes, by detecting the charged NLPs produced in high-energy neutrino collisions with Earth matter.  We study the production of charged NLPs by both atmospheric and astrophysical neutrinos; only the latter, which is largely uncertain and has not been detected yet, was the focus of previous studies.  We compute the resulting fluxes of the charged NLPs, compare those of different origins, and analyze the dependence on the underlying particle physics setup.  We point out that even if the astrophysical neutrino flux is very small, atmospheric neutrinos, especially those from the prompt decay of charmed mesons, may provide a detectable flux of NLP pairs at neutrino telescopes such as IceCube.  We also comment on the flux of charged NLPs expected from proton--nucleon collisions, and show that, for theoretically motivated and phenomenologically viable models, it is typically sub-dominant and below detectable rates. ",
        "introduction": "\\label{sec:intro}  The fundamental nature of dark matter poses a profound challenge to contemporary theoretical particle physics. Observations constrain the neutrino---the only electrically and color neutral non-baryonic elementary particle within the Standard Model of particle physics---to have a negligible contribution to the overall dark matter budget~\\cite{Fukugita:2006rm,Spergel:2006hy}. Dark matter is regarded as one of the most compelling hints towards new physics beyond the Standard Model. The question of its elementary essence has triggered enormous theoretical and phenomenological efforts~\\cite{Bergstrom:2000pn}.  The existence of suitable dark matter particle candidates in several theoretically cogent extensions of the Standard Model, like low-energy supersymmetry~\\cite{Jungman:1995df} or extra-dimensional scenarios~\\cite{Cheng:2002ej,Hooper:2007qk}, focused a strong interest on a class of dark matter candidates known as weakly interacting massive particles (WIMPs). Similarly to other Standard Model particles, WIMPs would fall out of thermal equilibrium and freeze out in the early Universe, leaving a relic abundance compatible with the inferred amount of dark matter~\\cite{Wolfram:1978gp}. These WIMPs can be directly detected by experiments looking for the minuscule energy deposition caused by dark matter particles scattering nuclei~\\cite{Munoz:2003gx}. The pair annihilation of WIMPs into energetic gamma rays, neutrinos, and antimatter, is a second, yet indirect, handle on the presence and potential imprint of galactic particle dark matter~\\cite{Bergstrom:1998hd}.  The connection of the aforementioned scenarios to the electroweak scale, soon to be probed at the Large Hadron Collider (LHC), motivated the exploration of complementarity between collider searches for new physics and the question of the elementary nature of dark matter~\\cite{Drees:2000he,Baer:2003wx,Weiglein:2004hn,Masiero:2004ft,Baltz:2006fm}. In most cases, if dark matter is a WIMP, the anticipated experimental signature at LHC would be the production of strongly interacting massive particles which promptly decay to the lightest and stable WIMP, plus a number of energetic jets and leptons. The neutral dark matter particle would escape the detector unobserved, leading to large missing transverse energy as well~\\cite{howie}. Conclusive identification of escaping neutral particles at LHC with dark matter permeating our Galaxy and other cosmic structures would, however, require some evidence from the direct and/or indirect WIMP searches listed before~\\cite{Bergstrom:1998hd}.  WIMPs are indeed attractive dark matter candidates, but are not the only theoretically envisioned possibility.  The dark matter particle could exhibit even feebler interactions with ordinary Standard Model particles than a WIMP, making direct and indirect searches completely hopeless.  For instance, the supersymmetric gravitino~\\cite{Blumenthal:1982mv} or the Kaluza-Klein graviton of universal extra dimensions~\\cite{Feng:2003nr,Hooper:2007qk} are perfectly plausible `super-weakly interacting'~\\cite{Feng:2003xh} dark matter candidates (super-WIMPs).  If Nature chose an option like this, collider signatures of new physics, if any, would strongly depend upon the nature of the next-to-lightest supersymmetric particle (NLSP). Since a super-WIMP is also very weakly coupled to the other new-physics heavier states, the NLSP would likely be quasi-stable.  If the NLSP is neutral, the qualitative experimental landscape would look like that of a standard WIMP scenario.  However, we would lack the needed proof of a connection between the weakly interacting long-lived particles produced at colliders and galactic dark matter.  If the NLSP is instead charged (constituting a charged massive particle, or CHAMP), the LHC would potentially observe the extremely distinct signature of a `heavy muon': charged tracks and penetration of the outer muon sub-detector, possibly at very low relativistic beta.  CHAMPs are constrained by direct collider searches at LEP2~\\cite{lep2}, as well as at the Tevatron~\\cite{tevatron}.\\footnote{These and future searches at LHC are not trivial, but advanced work is being done to ensure that such potential signals would not be missed~\\cite{Fairbairn:2006gg}. Detection for masses $\\lesssim 1$~TeV is essentially guaranteed.}  If a CHAMP were stable on collider scales, it could still decay on cosmological scales, and thus impact precision astrophysical measurements~\\cite{chargeddecay}, including the chemical potential associated with the cosmic microwave background black-body spectrum~\\cite{Zhang:2007zzh}, the extragalactic gamma-ray background~\\cite{Cembranos:2007fj}, the reionization history of the universe~\\cite{Pierpaoli:2003rz}, the formation of small scale structures~\\cite{Sigurdson:2003vy,Profumo:2004qt} and the synthesis of light elements in the early Universe~\\cite{BBNLDP,Jedamzik:2007cp} (see, for implications of neutral particle, Ref.~\\cite{BBNLDP2}). Anomalies in the above mentioned quantities, however, could hardly be considered smoking gun evidence that the collider CHAMPs were indeed related to lighter, super-weakly interacting dark matter.  If CHAMPs were stable on collider timescales, but featured a short lifetime on cosmological timescales, say on the order of a year or less, CHAMPs produced in colliders might be trapped in large water tanks surrounding the detectors~\\cite{Feng:2004yi}.  The tanks would then be periodically drained to underground reservoirs where CHAMP decays might be observed in low-background conditions~\\cite{Feng:2004yi}.  While certainly not straightforward experimentally, such a technique might provide important information first on the actual meta-stability of the charged species, and, secondly on the nature of the super-weakly-interacting particle the CHAMP would decay into.  However, even if CHAMP decays were actually observed, this would still not suffice as conclusive evidence that the elusive particle CHAMPs decay into is indeed the dark matter constituent.  To our knowledge, beyond high-energy collider experiments, the only direct experimental handle on a super-weakly interacting dark matter particle featuring a heavier, meta-stable charged partner is CHAMP pair production via neutrino--nucleon collisions, followed by direct observation at neutrino telescopes.  This idea, originally proposed in Ref.~\\cite{Albuquerque:2003mi}, relies on the fact that the energy losses of CHAMPs in Earth are significantly smaller than those of muons, therefore CHAMP pairs (unlike muon pairs) can reach the detector even if they were produced far away.  This makes the relevant target volume for neutrino--nucleon interactions much larger. The CHAMP pairs can be efficiently separated from muon pairs, due to large track separations in the detector. The original proposal was subsequently followed up by related studies~\\cite{Ahlers:2006pf,Albuquerque:2006am,Reno:2005si,Huang:2006ie,Ahlers:2007js,Reno:2007kz}, which focused on the specific case of a gravitino lightest supersymmetric particle (LSP), and a stau NLSP playing the role of the CHAMP.  Among other aspects, these studies investigated in detail stau energy losses in Earth and in the detector, computed expected event rates for a few sample models and the relevant background, and addressed the possibility of discriminating single-stau from single-muon events.  Given the steady progress in the deployment of next-generation km${}^3$-size neutrino telescopes---particularly IceCube at the South Pole, already under construction and taking data---we consider it timely to address in detail a few points relevant for improving our understanding of the prospects for implementing the above outlined technique. In particular, as background rejection is not a substantial issue, the crucial point appears to be the evaluation of the CHAMP pair event rate at IceCube. To this end, we focus on the following four aspects: \\begin{itemize} \\item[(i)]  {\\em Incoming neutrino flux}: So far, all long-lived CHAMP analyses for neutrino telescopes have considered a relatively optimistic flux of astrophysical neutrinos---as large as the Waxman-Bahcall (WB) bound~\\cite{WBBound}---as the primary source for CHAMP pair production.  However, the WB bound is only a {\\it theoretical upper limit} on the flux of astrophysical neutrinos expected from optically thin sources; therefore, the absolute normalization as well as the spectral shape of the true astrophysical neutrino flux remain largely unknown (of course the biggest reason for this is that these neutrinos are as yet undetected!).  On the other hand, atmospheric neutrinos have been detected and their flux is rather accurately known (below $\\sim$100~TeV~\\cite{AtmExp}).  At larger energies ($\\gtrsim 100$~TeV), although there is no detection so far, one can rather reliably extrapolate the flux of conventional atmospheric neutrinos from lower energies.  In addition, one also expects a significant flux of so-called prompt-decay atmospheric neutrinos, which originate from the decay of short-lived charmed mesons and feature a harder spectral index than the conventional component.  While we again have only an upper limit on this prompt decay component, we know from particle physics that it necessarily exists, at some level, in the high-energy regime, and will be measured accurately by IceCube.  In any event, atmospheric conventional and prompt-decay neutrinos evidently contribute as well to CHAMP pair production in neutrino--nucleon collisions.  In this paper, we study the role of these standard, {\\it guaranteed} neutrino sources, and compare it to the contribution from astrophysical neutrino flux models including the WB bound.  \\item[(ii)]  {\\em Underlying particle physics model}: The event rate depends not only on the incoming neutrino flux, but also on the nature of the assumed particle physics model.  Here, we consider generic supersymmetric models featuring a gravitino LSP and stau NLSP, and study how the stau pair production cross section and event rates at neutrino telescopes depend on the given mass spectrum. Neutrino--nucleon interactions produce slepton--squark pair final states, the amplitudes mediated by supersymmetric fermion exchange (neutral and charged gauginos and higgsinos).  \\item[(iii)] {\\em Other CHAMP sources}: Unlike the high-energy neutrino flux, the flux of very energetic protons is accurately measured and well-known up to extremely high energies; proton--nucleon interactions feature a supersymmetric pair production cross section significantly larger than that of neutrino--nucleon processes. It therefore seems reasonable to quantitatively assess the flux of stau pairs produced in the interaction of incident primary protons with nuclei in the atmosphere. The trade-off is the enormous cross section for proton--nucleus scattering into Standard Model particles, which depletes the incoming proton flux, highly suppressing any stau pair production rate. Since quark--anti-quark processes can directly produce stau pairs, however, the kinematic threshold for stau pair production, as a function of the incoming primary particle energy, is lower for proton--nucleus than for neutrino--nucleus processes. Yet we find that, for reasonable and phenomenologically acceptable particle models, the expected stau pair event rates from proton--nucleus collisions are experimentally negligible at all energies, leaving neutrinos as the only relevant primary source particles for CHAMP pair production.  \\item[(iv)] {\\em Simplified analytic approach:} We present the computation of the flux of staus from neutrino--nucleon interactions from first principles, and argue that, to an acceptable degree of accuracy, the total number of expected staus can be computed as one simple integral of three factors. Specifically, we show that the quantities of physical relevance are (1) the incident flux of primary neutrinos, (2) the ratio of the cross sections of neutrino plus nucleon into supersymmetric particle pairs to the total neutrino--nucleon cross section, and (3) a geometric efficiency factor.  \\end{itemize}  Hereafter, we specifically use supersymmetric staus as charged meta-stable NLPs, but note that the following arguments are applicable to any other possible candidates of long-lived CHAMPs. The outline of the remainder of the paper is as follows.  We discuss the various components of the high energy neutrino flux and their connected uncertainties in \\Sref{sec:nuflux}.  We introduce the new physics scenarios and compute the stau pair production cross sections in \\Sref{sec:particle}. \\Sref{sec:rates} outlines the computation of the stau event rate at IceCube, including the above-mentioned simplified analytic treatment.  The stau rate dependence on the particle physics framework is addressed in \\Sref{sec:susy}, and we draw conclusions in \\Sref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  We reassessed the flux of meta-stable staus produced by neutrino--nucleon and nucleon--nucleon interactions that might be detectable at km$^3$ neutrino telescopes.  We derived the flux of staus from first principles, and showed that, under the approximation that Earth is opaque to very high energy neutrinos, the number of staus at the detector is given by a simple integral over the neutrino energy of the product of three factors: the incident neutrino flux, the ratio of the neutrino--nucleon cross section into supersymmetric particle pairs over the total neutrino--nucleon cross section, and a geometric efficiency factor.  We showed that this approximation reproduces an exact numerical computation within a factor 2, which in turn is much better than the level of our knowledge of the first two factors entering the stau flux computation---namely, the incident neutrino flux and the features of the supersymmetric particle setup.  We focused on each of the factors relevant to the computation of the final flux of staus.  We concentrated on four well-motivated supersymmetric benchmark models, and independently evaluated the relevant production cross sections.  We pointed out that previously-neglected atmospheric neutrinos from prompt charmed meson decays could give a potentially large stau flux, even in the absence of a (yet to be discovered) astrophysical high-energy neutrino flux. This will depend on the prompt atmospheric neutrino flux being at the upper end of the theoretically expected range. Nucleon--nucleon processes, even for the most optimistic benchmark models, would not contribute sizably to the final stau flux.  Finally, we numerically and analytically studied how the relevant cross sections depend on the supersymmetric model mass spectrum, and how the relative importance of primary neutrino sources depends on the mass scale of supersymmetric scalars.  In particular, we predict that if the signal discussed here is indeed detected, a very sizable fraction of it would originate from conventional and prompt-decay neutrinos."
    },
    "0711/0711.2887_arXiv.txt": {
        "abstract": "The hemispheric pattern of solar filaments is considered using newly-developed simulations of the real photospheric and 3D coronal magnetic fields over a 6-month period, on a global scale. The magnetic field direction in the simulation is compared directly with the chirality of observed filaments, at their observed locations. In our model the coronal field evolves through a continuous sequence of nonlinear force-free equilibria, in response to the changing photospheric boundary conditions and the emergence of new magnetic flux. In total 119 magnetic bipoles with properties matching observed active regions are inserted. These bipoles emerge twisted and inject magnetic helicity into the solar atmosphere. When we choose the sign of this active-region helicity to match that observed in each hemisphere, the model produces the correct chirality for up to 96\\% of filaments, including exceptions to the hemispheric pattern. If the emerging bipoles have zero helicity, or helicity of the opposite sign, then this percentage is much reduced. In addition, the simulation produces a higher proportion of filaments with the correct chirality after longer times. This indicates that a key element in the evolution of the coronal field is its long-term memory, and the build-up and transport of helicity from low to high latitudes over many months. It highlights the importance of continuous evolution of the coronal field, rather than independent extrapolations at different times. This has significant consequences for future modelling such as that related to the origin and development of coronal mass ejections. ",
        "introduction": "The recent work of \\inlinecite{mackay3} gives a convincing explanation for the hemispheric pattern of filament chirality on the Sun. A filament may be classified as having either dextral or sinistral chirality (handedness) depending on whether its axial magnetic field points right or left, as seen from the side of the filament with positive magnetic polarity in the photosphere. Observations show that quiescent filaments in the northern hemisphere tend to be dextral, while those in the southern hemisphere tend to be sinistral \\cite{rust1967,leroy1983,martin1994,pevtsov2003}. The latter paper found that $80\\%$ -- $85\\%$ of quiescent filaments follow this hemispheric pattern, which remains the same when the Sun's polar field reverses every 11 years.  Using the model of \\inlinecite{vanballegooijen2000} to simulate the coupled evolution of photospheric and coronal magnetic fields, \\inlinecite{mackay3} have shown that the dominant chirality may be produced for the dominant range of bipole tilt angles \\cite{wang1989} and dominant sign of helicity \\cite{pevtsov1995} observed in each hemisphere. The model suggests how the evolution of bipolar active regions, and their interaction as they are transported poleward, may explain both the sheared magnetic fields in filament channels and their hemispheric pattern. Importantly, the authors showed how exceptions to the hemispheric pattern could occur in a natural way, for bipoles with very large tilt angles or the minority sign of helicity for their hemisphere. However, as in the filament-formation simulations of \\inlinecite{litvinenko2005} and \\inlinecite{welsch2005}, \\inlinecite{mackay3} considered only a pair of magnetic bipoles, in a localised region of the solar corona. In this paper we test their conclusions in a direct comparison between global coronal simulations and observations of many real filaments on the Sun, over a 6-month period in 1999.  In the first paper (\\opencite{yeates2007a}, hereafter ``Paper 1''), we described how barb orientations were used to determine the chirality of 255 filaments observed in H$\\alpha$ images from Big Bear Solar Observatory (BBSO). In the present paper, the skew of the simulated magnetic field is compared with the chirality of these filaments, at their observed locations. This comparison requires the extension of the previous simulations to cover the global corona, and its evolution over a period of many months.  In Paper 1, we described the photospheric simulations, which provide a continuously evolving lower boundary condition for the 3D coronal field. By inserting new active regions into the simulation as they were observed, we showed how accuracy to the observed synoptic magnetograms could be maintained without ever resetting the simulation to the observed photospheric field. In particular large areas of positive and negative flux were well reproduced, as were the shape of the polarity inversion lines (PILs) at the majority of observed filament locations. This represents a significant improvement over previous simulations of filament formation in the real corona \\cite{mackay2}, which lost accuracy over time if the photospheric field was not reset to the observed magnetogram every 27 days. The solution they adopted was to construct a new potential field in the corona at this point and begin the evolution again, thus removing any long-term memory in the coronal field of previous interactions, as well as any helicity that had built up. Magnetic helicity is a quantity that measures the topological properties of field lines such as linking, twist, or kinking (\\opencite{berger1998}, \\citeyear{berger1999}). We believe that the gradual build-up of helicity, in the form of sheared non-potential field lines, is key to the hemispheric pattern of filaments. Our new simulations are able to maintain accuracy in the photospheric field over many months without resetting, enabling a continuous evolution of the coronal field and so no loss of memory.  Potential magnetic fields, which contain no electric currents, are the simplest way to model the magnetic field in the solar corona. Extrapolations using an upper ``source-surface'' where the field is assumed to be radial were first proposed in the 1960s \\cite{altschuler1969,schatten1969}, and have since been used by many authors to study the evolution of the coronal field based on photospheric motions (\\textit{e.g.} \\opencite{wang2002}; \\opencite{mackay5b}; \\opencite{lockwood2003}; \\opencite{schrijver2003}). However, in these models, the coronal field is extrapolated from the photospheric simulations at discrete intervals. It is independently computed at each time and contains no memory of its past evolution. The potential field model does not describe \\emph{how} the magnetic field of the Sun evolved into a particular configuration. It precludes the long-term build-up of magnetic helicity and shear which have been shown to be essential for filament formation \\cite{mackay2003}, and in particular to model the observed hemispheric pattern \\cite{mackay3}. Indeed potential fields represent the state of minimum magnetic energy, so cannot represent the sheared fields of active regions \\cite{riley2006, regnier2007} or filament channels \\cite{mackay1997}.  A more realistic model for such regions is a force-free equilibrium, where the Lorentz force $\\vec{j}\\times\\bb$ vanishes everywhere but the current density $\\vec{j}$ may be non-zero. This approximation is thought to be good for magnetic fields in the corona \\cite{priest1982}, where the magnetic pressure is generally much greater than the gas pressure, so that other forces may be neglected and the Lorentz force must balance itself. This condition may be written \\begin{equation} \\nabla\\times\\bb = \\alpha\\bb \\label{eqn:forcefree} \\end{equation} for some $\\alpha$. If $\\alpha$ is everywhere constant, this is called a linear force-free field, and Equation \\eqref{forcefree} is readily solved. The special case $\\alpha=0$ is the potential field discussed earlier, while constant non-zero values have been used successfully to model individual filaments \\cite{aulanier1998,aulanier2000}. However, such constant-$\\alpha$ fields are limited to localised regions of the corona, and in general $\\alpha$ can be a function of position, although it must be constant along any individual field line. The more realistic fields where $\\alpha$ is a function of position are called nonlinear force-free, and are considerably more difficult to compute, particularly with regard to the required boundary conditions. Various methods have been developed to compute nonlinear force-free fields based on a single surface vector-magnetogram (see \\opencite{schrijver2006} for a summary), and are being developed to compute the global coronal field based on a full-disk vector-magnetogram \\cite{wiegelmann2007}.  The technique we adopt here is described in Section \\ref{sec:model} and differs significantly from these extrapolation methods. It effectively produces a continuous sequence of nonlinear force-free equilibria as the 3D coronal field responds to flux emergence and photospheric motions, rather than a single-time extrapolation. In this paper we consider whether our model can produce the correct hemispheric pattern observed in our filament sample. Section \\ref{sec:skew} presents the simulation results; firstly we consider the overall percentage of simulated filaments following the hemispheric pattern (Section \\ref{sec:pattern}), and then we compare the chirality of individual filaments with the observations (Section \\ref{sec:indiv}). The performance of the simulations is discussed in Section \\ref{sec:discussion}. Conclusions are given in Section \\ref{sec:conclusion}, along with consequences for future modelling. ",
        "conclusions": "\\label{sec:conclusion}  We have developed new simulations of the large-scale magnetic field of the solar corona, over a period of many months. Based on the flux transport and magnetofrictional model of \\inlinecite{vanballegooijen2000}, the coronal field evolves in response to flux emergence and the changing photospheric boundary conditions. Using a reduced form of the MHD equations, the corona evolves in essence through a continuous sequence of force-free equilibria. In Paper 1 we described our technique for incorporating newly-emerging active regions based on observations, and showed how the large-scale photospheric field remains accurate over a 6-month period. In this paper we have demonstrated how our 3D model is able to reproduce an important observed property of filament magnetic fields, namely the hemispheric chirality pattern. By direct comparison with H$\\alpha$ observations of real filaments over a 6-month period in 1999, we have shown that the emergence, surface transport, and reconnection of large-scale twisted active-region magnetic fields are sufficient to produce the chirality pattern.  Our simulations represent a significant improvement over coronal field models where independent extrapolations are taken at different times, particularly those which assume a potential field in the corona. In particular, the newly emerged fields in our simulation are assumed to be twisted, injecting magnetic helicity into the solar atmosphere. The high electric conductivity of the coronal plasma causes magnetic helicity to be conserved \\cite{priest2000}, and gives the coronal field a certain degree of memory. This memory is retained because the simulation is a continuous evolution in which the magnetic field is never reduced to its lowest energy state. This allows helicity to be transported from low to high latitudes as the coronal field evolves in response to photospheric motions. We believe such helicity transport to play a key role in filament chirality; if high latitude quiescent filaments were formed \\emph{in situ} by differential rotation of a preexisting coronal field then their chirality would be opposite to the observed pattern.  For each filament in our sample of 109 with definite chirality (as inferred from the H$\\alpha$ observations--see Paper 1 for details), we have compared the local skew direction of the simulated coronal magnetic field with the observed filament chirality. We find the following main results: \\begin{enumerate} \\item{The simulation produces the correct skew for up to 96.9\\% of filaments, including exceptions to the hemispheric pattern. This is for simulations where the sign of newly-emerging bipole twist matches the observed hemispheric sign of active-region helicity \\cite{pevtsov1995}.} \\item{If the emerging regions have zero twist, or their sign of twist is opposite to the observed pattern, then the percentage of filaments with correct skew is significantly lower.} \\item{The simulation produces a higher proportion of filaments with correct chirality at later times.} \\end{enumerate}  The results suggest that the four ingredients of differential rotation, meridional circulation, surface diffusion, and emerging twisted active regions are sufficient to produce the observed filament chirality pattern. The results presented in this paper give a comparison of the chirality at a large number of filament locations, and we have not discussed the circumstances of individual locations in any detail. Previous papers using the simpler configuration of two interacting bipoles (\\opencite{mackay4}, \\citeyear{mackay6}) have considered in detail the basic physical interactions which occur in the model. Very similar features are found at multiple locations in the global simulations presented in this paper.  In addition to forming the necessary sheared magnetic fields, the non-potential simulations have allowed us to vary the twist of the emerging regions, and our results support the conclusions of \\inlinecite{mackay3} that this active-region helicity has an important influence on filament chirality. It is not however the only factor; even the simulation run with bipole twist opposite to the observed pattern produced the correct chirality for over 60\\% of filaments.  The third result, namely that the performance of the simulation improves with time, is a key conclusion of this study. A large number of wrongly-classified filaments were early in the simulation, before magnetic helicity could build up from the shearing and evolution away from the initial potential field, and from the emergence of new bipoles. This highlights the importance of continuous long-term evolution of the coronal field, rather than independent extrapolations at different times.  There are several improvements that we hope to make in future simulations. Firstly, in view of the results above, future simulations will have a longer ``ramp-up'' time from the initial potential field. Further improvements concern the observations of newly-emerging active regions, which are crucial to our model. In the present study, all newly-inserted regions were given the same magnitude of twist in any individual simulation run. However, using photospheric vector magnetograms, we hope in future to measure the twist of individual active regions \\cite{nandy2007} and incorporate this into the simulations. Another uncertainty in our model is the exact date of emergence of each active region, but this can not be seriously addressed without observations of the far-side of the Sun, although helioseismic imaging techniques \\cite{lindsey2000} offer some promise in this respect.  In summary, this study has demonstrated the importance of long-term evolution and build-up of non-potential magnetic structure in the solar corona. The transport of helicity from low to high latitudes over many months is a fundamental element in the coronal evolution. After the successful reproduction of filament chirality, we hope that the model will be useful for further study of the corona. An interesting feature to explore is the temporary loss of equilibrium and ejection of magnetic flux ropes, as reported by \\inlinecite{mackay4}; such behaviour also occurs in our global simulations. These eruptions may correspond to coronal mass ejections \\cite{zhang2005,hudson2006}, and an understanding of the magnetic topologies responsible offers the possibility of space weather forecasting on the time-scale of weeks or months.  \\begin{acks} Financial support for ARY and DHM was provided by the UK Science and Technology Facilities Council (STFC). The simulations were performed on the UKMHD parallel computer in St Andrews, funded jointly by SRIF/STFC. We would like to thank S.F. Martin for supplying the original filament data set, based on data from BBSO/NJIT, as well as travel support from the PROM group (NSF grant ATM-0519249). Synoptic magnetogram data from NSO/Kitt Peak was produced cooperatively by NSF/NOAO, NASA/GSFC, and NOAA/SEL and made publicly accessible on the World Wide Web. Finally we thank the anonymous referee for his/her detailed comments which have improved the paper. \\end{acks}"
    },
    "0711/0711.0882_arXiv.txt": {
        "abstract": "{Observations of the cosmic microwave background, light element abundances, large-scale distribution of galaxies, and distant supernovae are the primary tools for determining the cosmological parameters that define the global structure of the Universe.} {Here we illustrate how the combination of observations related to strong gravitational lensing and stellar dynamics in elliptical galaxies offers a simple and promising way to measure the cosmological matter and dark-energy density parameters.} {A gravitational lensing estimate of the mass enclosed inside the Einstein circle can be obtained by measuring the Einstein angle, once the critical density of the system is known. A model-dependent dynamical estimate of this mass can also be obtained by measuring the central velocity dispersion of the stellar component. By assuming the well-tested homologous $1/r^{2}$ (isothermal) profile for the total (luminous+dark) density distribution in elliptical galaxies acting as lenses, these two mass measurements can be properly compared. Thus, a relation between the Einstein angle and the central stellar velocity dispersion is derived, and the cosmological matter and the dark-energy density parameters can be estimated from this.} {We determined the accuracy of the cosmological parameter estimates by means of simulations that include realistic measurement uncertainties on the relevant quantities. Interestingly, the expected constraints on the cosmological parameter plane are complementary to those coming from other observational techniques. Then, we applied the method to the recent data sets of the \\emph{Sloan Lens ACS} (SLACS) and the \\emph{Lenses Structure and Dynamics} (LSD) Surveys, and showed that the concordance value between 0.7 and 0.8 for the dark-energy density parameter is included in our 99\\% confidence regions.} {The small number of lenses available to date prevents us from precisely determining the cosmological parameters, but it still proves the feasibility of the method. When applied to samples made of hundreds of lenses that are expected to become available from forthcoming deep and wide surveys, this technique will be an important alternative tool for measuring the geometry of the Universe.} ",
        "introduction": "The Universe appears to be dominated by dark-energy and dark matter. Although the physical nature of these dark components is still unknown, the standard cosmological $\\Lambda$\\emph{CDM} model with only a few parameters fits most of the current data well: precision measurements of the anisotropies in the cosmic microwave background (\\emph{CMB}; Bennett et al. \\cite{ben}; Spergel et al. \\cite{spe1}, \\cite{spe2}), the observed abundances of light elements (Burles et al. \\cite{bur}; Cyburt et al. \\cite{cyb}), the large-scale distribution of galaxies (\\emph{LSS}; Tegmark et al. \\cite{teg}; Cole et al. \\cite{col}), and the luminosity-distance relationship for distant type Ia supernovae (\\emph{SNIa}; Riess et al. \\cite{rie1}, \\cite{rie2}; Perlmutter et al. \\cite{per}). In this standard model, the Universe is homogeneous and isotropic on its largest scales, and its geometry appears to be flat ($\\Omega = \\Omega_{m} + \\Omega_{\\Lambda} \\approx 1$); the total mass-energy density is mainly in the form of dark-energy ($\\Omega_{\\Lambda} \\approx 0.7$) and matter ($\\Omega_{m} \\approx 0.3$), ordinary and dark. These values of the cosmological parameters imply a fairly recent transition from a decelerating to an accelerating universal expansion. The \\emph{Hubble Space Telescope} (\\emph{HST}) \\emph{Key Project} has measured the current expansion rate, the Hubble parameter $H_{0}= (72 \\pm 8)\\,\\mathrm{km\\,s^{-1}\\,Mpc^{-1}}$ (Freedman et al. \\cite{fre}). The estimates of different cosmological parameters from a single observational method are often correlated (hence ``degenerate'') and exhibit significant uncertainties. For instance, from the \\emph{WMAP} three year data alone, without a prior on the flatness of the Universe, the best-fit model is characterized by $\\Omega_{m} = 0.42$, $\\Omega_{\\Lambda} = 0.63$, $H_{0}=55\\,\\mathrm{km\\,s^{-1}\\,Mpc^{-1}}$ (Spergel et al. \\cite{spe2}), values that are quite different from the concordance values reported above. This suggests that precise measurements of the cosmological parameters can only be obtained by using complementary techniques. In fact, considerable efforts are still being made in order to secure accurate measurements of these parameters (in particular, see the scientific goals of the forthcoming \\emph{PLANCK} and \\emph{SNAP} missions).  The deflection of light due to gravitational lensing is sensitive to the total matter density of the structures in the Universe, independently of the nature or dynamical state of the deflecting mass. Therefore, strong and weak gravitational lensing provide valuable tools for measuring the distribution of mass. In particular, cosmic shear estimates of the amplitude of the weak lensing distorsions of distant sources over a wide range of angular scales (Bartelmann \\& Schneider \\cite{bar2}) are a very encouraging way to study the large-scale structure of the Universe, and therefore to probe the parameters that define the relevant cosmological model (Refregier \\cite{ref}). Further constraints on the geometry of the Universe can also be provided by the abundance of lenses or arcs observed in lens surveys (see Bartelmann \\& Weiss \\cite{bar1} for a numerical approach; Mitchell et al. \\cite{mit} for observational results from the CLASS Survey described by Myers et al. \\cite{mye} and Browne et al. \\cite{bro}) and by strong (Link \\& Pierce \\cite{lin}; Soucail et al. \\cite{sou}) or weak (Lombardi \\& Bertin \\cite{lom}; Jain \\& Taylor \\cite{jai}) lensing mass reconstructions in clusters of galaxies. Isolated strong lens galaxies offer another possibility to measure the cosmological parameters (Kochanek \\cite{koc3}, \\cite{koc4}; Myungshin et al. \\cite{myu}), including the Hubble parameter (Refsdal \\cite{refs}; Koopmans et al. \\cite{koo4}; M\\\"ortsell \\& Sunesson \\cite{mor}).  The various techniques that have been proposed are limited by several assumptions. For example, a measurement of the matter density parameter from cosmic shear is degenerate with that of the normalization of the amplitude of the power spectrum of matter perturbations ($\\sigma_{8}$); moreover, an extremely large number of high-quality galaxy images and some modeling on the growth of the structure in the Universe are required. For techniques based on gravitational lens statistics, the luminosity function, the relation between luminosity and velocity dispersion, and the density profile of the lens galaxies play an important role. The use of arcs statistics awaits more realistic simulations of clusters and observations, since different studies have led to contrasting results (e.g., Bartelmann et al. \\cite{bar3}; Dalal et al. \\cite{dal}). The estimates of the cosmological parameters from cluster mass reconstructions have also two distinct limitations: the presence of possible substructure in the region of multiple image formation (for strong lensing) and the need for a high density of background galaxies (for weak lensing).  In this paper we propose a new technique that, starting from strong gravitational lensing and stellar dynamics observations in elliptical galaxies, is able to probe the geometry of the Universe in a different and effective way. The paper is organized as follows. In Sect.~2, we describe our method to estimate the matter and dark-energy density parameters. Then, in Sect.~3, we determine through simulations the precision attainable in these measurements of the cosmological parameters. Good estimators for the quantities relevant to the problem are identified in Sect.~4. In Sect.~5, we apply our technique to the data collected and just published of two surveys of lens galaxies for which stellar dynamical measurements are available. Finally, in Sect.~6, we summarize and discuss the results obtained. ",
        "conclusions": "Lensing studies have been recognized to be very valuable in providing a testing ground for theories of formation and evolution of early-type galaxies (Rusin \\& Kochanek \\cite{rus2}; Koopmans et al. \\cite{koo3}). In these studies the structure of the lens galaxies can be investigated, once a specific cosmological model is adopted. From a complementary perspective, a measurement of the cosmological parameters can in principle be accomplished by assuming some empirically justified general properties for a sample of lenses. In this paper, we have shown that indeed, starting from the $1/r^{2}$ paradigm for the total density profile of elliptical galaxies, the cosmological parameters $\\Omega_{m}$ and $\\Omega_{\\Lambda}$ can be estimated. The adopted paradigm has been validated by stellar dynamics (e.g., Gerhard et al. \\cite{ger}), strong gravitational lensing (e.g., Rusin et al. \\cite{rus}), and dynamical studies based on several other tracers: globular clusters and planetary nebulae (e.g., Peng et al. \\cite{pen}),  X-ray halos (e.g., Humphrey et al. \\cite{hum}), and HI disks and rings (e.g., Franx et al. \\cite{fra2}).  In detail, the main results of this paper can be summarized as follows:  \\begin{itemize} \\item We have developed a new method to investigate the geometry of the Universe by combining measurements of strong gravitational lensing and stellar dynamics in a sample of elliptical galaxies. The basic idea is to compare lensing and dynamical mass estimates in order to find an observable cosmology-dependent relation. For this purpose, it is necessary to model the total (luminous and dark) density profile inside the Einstein radius of the ellipticals under investigation. Here, we have assumed a one-component $1/r^{2}$ density profile.  \\item We have performed a feasibility study of our technique by simulating measurements of the cosmological parameters. We have studied the precision of the method, which depends on the level of uncertainty with which the Einstein angle and the central stellar velocity dispersion are known. Hence, we have demonstrated that the current uncertainty with which these quantities can be measured is sufficiently small to obtain precise information about cosmology, once a statistically significant sample of lenses is available. In particular, the method has been shown to be best suited to measure $\\Omega_{\\Lambda}$.  \\item Then, we have studied the relation between the stellar kinematics and the velocity parameter of the one-component $1/r^{2}$ (isothermal) model for the total density distribution in ellipticals. For a sample of eight bright, round, and nearby early-type galaxies, we have shown that the stellar central velocity dispersion is a good tracer of the velocity dispersion characterizing the one-component model.  \\item We have applied the proposed method to the SLACS Survey data set. This sample is composed of fifteen massive early-type galaxies at intermediate redshift acting as lenses. In a flat cosmology this method leads to a value of $\\Omega_{\\Lambda}$ lower than 0.80 at 95\\% CL. By including five more distant lenses from the LSD Survey, at the same confidence level and in the same flatness hypothesis, $\\Omega_{\\Lambda}$ has been measured to be greater than 0.16. This latter sample ruled out, at greater than 99\\% CL, general cosmological models with values of $\\Omega_{\\Lambda}$ smaller than 0.4.  \\item Finally, we conclude that future surveys expected to identify a large number of lenses (one or two hundred) will allow this method to measure the values of the cosmological parameters with a precision comparable to that of other standard techniques. This will be a new and important test for the $\\Lambda$\\emph{CDM} concordance model. \\end{itemize}"
    },
    "0711/0711.2249_arXiv.txt": {
        "abstract": "{ Until now, only one planetary nebula (PN) has been known in the Fornax dwarf spheroidal galaxy. } { The discovery of a second PN candidate, associated with one of the 5 globular clusters in the Fornax dwarf, is reported. } { Spectra of the globular cluster H5, obtained with the UVES echelle spectrograph on the ESO Very Large Telescope, show \\oiii\\ line emission at a radial velocity consistent with membership of the Fornax dwarf.  A possible counterpart of the \\oiii\\ emission is identified in archival images from the Wide Field Planetary Camera 2 on board the Hubble Space Telescope.  The source of the emission is located about $1\\farcs5$ (less than one core radius) southwest of the centre of the cluster. } { The emission line source is identified as a likely PN, albeit with several peculiar properties. No H$\\beta$, \\ion{He}{i}, or \\ion{He}{ii} line emission is detected and the \\oiii /H$\\beta$ ratio is $>25$ (2$\\sigma$). The expansion velocity inferred from the \\oiii\\ 5007\\AA\\ line is $V_e \\approx 55$ km s$^{-1}$, which is large for a PN.  The diameter measured on the HST images is about $0\\farcs23$ or 0.15 pc at the distance of the Fornax dSph. } { This object doubles the number of known PNe in Fornax, and is only the 5th PN associated with an old GC for which direct imaging is available. It may be a member of the rare class of extremely H-deficient PNe, the second such case found in a GC. } ",
        "introduction": "Like many other dwarf spheroidal galaxies, the Fornax dwarf spheroidal (dSph) has a remarkably complex star formation history. Colour-magnitude diagrams show evidence of star formation until as recently as $\\sim10^8$ years ago (Stetson et al.\\ \\cite{stet98}; Buonanno et al. \\cite{buon99}). The metallicity distribution of the field stars peaks at [Fe/H]$\\sim-1.0$, but with tails extending up to nearly solar metallicity and down to [Fe/H]$\\approx-2.7$ (Helmi et al.\\ \\cite{hel06}). Battaglia et al.\\ (\\cite{bat06}) identified three more or less distinct stellar populations in the Fornax dSph, namely an old, metal-poor component, a dominant intermediate-age (2--8 Gyr) moderately metal-rich component, and a younger ($\\sim 1$ Gyr) metal-rich component.  Five globular clusters (GCs) are known in the Fornax dSph, catalogued by Hodge (\\cite{hod61}) and hereafter referred to as H1\\ldots H5. When normalised to the (low) luminosity of the galaxy, this leads to an extremely high GC specific frequency of $S_N\\approx 29$ (e.g.\\ van den Bergh \\cite{van98}). Historically, it is interesting to note that one of the GCs (NGC 1049 = H3) was already discovered by John Herschel nearly a century before the parent galaxy itself.  From high-dispersion spectroscopy of individual stars in clusters H1, H2, and H3, Letarte et al.\\ (\\cite{let06}) derived [Fe/H]=$-2.5\\pm0.1$, [Fe/H]=$-2.1\\pm0.1$, and [Fe/H]=$-2.4\\pm0.1$, respectively.  From low-resolution spectroscopy, Strader et al.\\ (\\cite{stra03}) found an average [Fe/H]$\\sim-1.8$ and old ages (comparable to Milky Way GCs), with H5 possibly a few Gyrs younger than the rest.  Thus, the GCs are most likely associated with the old, metal-poor component in Fornax.  The Fornax dSph is known to host a single planetary nebula (PN), discovered by Danziger et al.\\ (\\cite{dan78}).  In a recent study of this object, Kniazev et al.\\ (\\cite{knia07}) derived an oxygen abundance of 12+log(O/H) = $8.28\\pm0.02$ ([O/H] = $-0.55$ for the solar oxygen abundance of Grevesse \\& Sauval \\cite{gs98}). They also estimated an iron abundance of [Fe/H] = $-1.13\\pm0.18$. It seems likely that this PN is associated with the intermediate-age component, but as for most PNe, pinpointing the exact age (and hence the mass of the progenitor star) is difficult.  This \\emph{letter} reports the discovery of a second PN candidate in the Fornax dSph, apparently associated with one of the GCs (H5). This not only doubles the number of known PNe in Fornax, but also constitutes a significant addition to the small number of PNe currently known to be associated with GCs in general (four in Milky Way GCs; Jacoby et al.\\ \\cite{jac97}). The identification of PNe in stellar clusters holds special significance since it gives a better handle on the age and metallicity of the progenitor star (although binary star evolution may be an important path for PN formation in clusters). Furthermore, the Fornax dSph is close enough that spatially resolved imaging of the PN is within reach. ",
        "conclusions": "\\begin{figure} \\begin{minipage}{42mm} \\includegraphics[width=42mm]{larsen_f4a.ps} \\end{minipage} \\hspace{1mm} \\begin{minipage}{42mm} \\includegraphics[width=42mm]{larsen_f4b.ps} \\end{minipage} \\caption{Left: HST/WFPC2 F555W image of H5. Right: Ratio of F555W and F814W exposures (black=higher ratio of F555W/F814W). The circle (diameter = 1$\\arcsec$) is centred on the location of the emission line object derived from the spectra (Fig.~\\ref{fig:twodfig}). } \\label{fig:hstfig} \\end{figure}  The new PN candidate in H5 doubles the number of PNe in the Fornax dSph. That it has escaped detection so far, despite many studies of the Fornax GCs, is probably due to its relative faintness and location far enough from the centre ($\\sim1\\farcs5$) of H5 to be easily missed in an ordinary long-slit spectrum. As seen in Fig.~\\ref{fig:hstfig}, it is also inconspicuous even in deep (broad-band) HST images, although it should be easily identifiable in narrow-band WFPC2, ACS/HRC, or WFC3 images.  As already noted by Danziger et al.\\ (\\cite{dan78}), comparison with other Local Group galaxies suggests that about one PN should be expected in the Fornax dSph. More recently, Magrini et al.\\ (\\cite{mag03}) concluded from a census of PNe in the Local Group that one PN is expected per $10^{6.92} L_V/L_{V,\\odot}$. Assuming $M_V\\approx-13.1$ for Fornax (Demers et al.\\ \\cite{dem94}), 1.7 PNe are then expected. Thus, the discovery of a second candidate does not change the conclusion that the PN population of Fornax is consistent with what is observed in other Local Group galaxies. The Sagittarius dSph, which is somewhat brighter than Fornax, is now known to host four PNe that span a range in metallicity and thus, like those in Fornax, appear to trace different populations (Zijlstra et al.\\ \\cite{zij06}). However, none of these is associated with any of the $\\sim8$ known GCs in the Sagittarius dSph (van den Bergh \\cite{van07}). What is the likelihood of detecting a PN in one of the Fornax GCs specifically?  Clearly, one should be careful about drawing statistical conclusions based on a sample of one object.  A naive scaling of the four PNe known in the $\\sim150$ Milky Way GCs (Harris \\cite{har96}), predicts about 0.13 PNe in the 5 Fornax GCs, but Poisson statistics would still yield 1 PN in a Fornax-like GC system in about 11\\% of the cases.  The nebula in H5 is a rather peculiar object. The expansion velocity is high for a PN, although hydrodynamical simulations (Villaver et al.\\ \\cite{vil02}) reproduce the observed combination of radius and $V_e$ at late stages of the PN evolution for a 1 $M_\\odot$ progenitor. There is some other evidence linking high PN $V_e$ values to old stellar populations (Richer \\cite{rich06}). High $V_e$ values are also observed in bipolar PNe with intermediate-mass progenitors (Corradi \\& Schwartz \\cite{cs95}), but the nebula in H5 does not appear to belong to this class (Fig.~\\ref{fig:hstfig}). The non-detection of emission lines other than \\oiii\\ is even more atypical.  In this regard the nebula resembles the peculiar PN IRAS 18333-2357 in the metal-poor Galactic GC M22, where the absence of emission lines other than \\oiii\\ 4959/5007\\AA\\ and [\\ion{Ne}{iii}] 3869/3967\\AA\\ is thought to imply an extreme H depletion (Cohen \\& Gillett \\cite{cg89}). However, IRAS 18333-2357 is exceedingly faint: The \\oiii\\ flux is comparable to that of the nebula in H5, although M22 is a factor of 40 closer! Extreme H depletion has also been found in knots within some Galactic PNe, e.g.\\ Abell 30 (Wesson et al.\\ \\cite{wes03}). A possible scenario is that this is the result of a late thermal pulse, occurring after most of the H-rich envelope had already been ejected (Iben et al.\\ \\cite{ib83}). Depending on when exactly this happened, an outer low-surface brightness PN might also be present, as in Sakurai's object (Duerbeck \\& Benetti \\cite{db96}).  Further observations of the nebula in H5 might help shed more light on its nature.  HST narrow-band imaging might be rewarding, and a comparison of the morphology with that of PNe in Galactic GCs and in the Sagittarius dSph would be interesting, e.g.\\ in relation to differences in the motion with respect to an ambient interstellar/intergalactic medium. The velocity dispersion of the Fornax GC system is only 6 km s$^{-1}$ (Cohen \\cite{coh83}), but Fornax may still be interacting with intergalactic gas in the Local Group (although this point remains uncertain; Piatek et al.\\ \\cite{pia07}). It would also be desirable to search for emission lines other than those from \\oiii , such as [\\ion{N}{ii}], [\\ion{Ne}{iii}], ionised He, and (of course) the hydrogen Balmer lines."
    },
    "0711/0711.2631_arXiv.txt": {
        "abstract": "We present results on spectral index distributions, number counts, redshift distribution and other general statistical properties of extragalactic point sources in the NEWPS$_{5\\sigma}$ sample \\citep{can07}. The flux calibrations at all the WMAP channels have been reassessed both by comparison with ground based observations and through estimates of the effective beam areas. The two methods yield consistent statistical correction factors. A search of the NASA Extragalactic Database (NED) has yielded optical identifications for $\\sim 89\\%$ of sources in the complete sub-sample of 252 sources with $S/N\\geq5$ and $S\\geq 1.1$ Jy at 23 GHz; 5 sources turned out to be Galactic and were removed. The NED also yielded redshifts for $\\simeq 92\\%$ of the extragalactic sources at $|b|>10^\\circ$. Their distribution was compared with model predictions; the agreement is generally good but a possible discrepancy is noted. Using the 5 GHz fluxes from the GB6 or PMN surveys, we find that $\\sim 76\\%$ of the 191 extragalactic sources with $S_{23\\rm GHz}> 1.3\\,$Jy can be classified as flat-spectrum sources between 5 and 23 GHz. A spectral steepening is observed at higher frequencies: only 59\\% of our sources are still flat-spectrum sources between 23 and 61 GHz and the average spectral indexes steepen from $\\langle\\alpha_{5}^{23}\\rangle= 0.01\\pm 0.03$ to $\\langle\\alpha_{41}^{61}\\rangle= 0.37\\pm 0.03$. We think, however, that the difference may be due to a selection effect. The source number counts have a close to Euclidean slope and are in good agreement with the predictions of the cosmological evolution model by \\citet{dz05}. The observed spectral index distributions were exploited to get model-independent extrapolations of counts to higher frequencies. The risks of such operations are discussed and reasons of discrepancies with other recent estimates are clarified. ",
        "introduction": "\\label{sec:intro}  \\begin{table*} \\centering \\begin{tabular}{ccccc|cc} WMAP band \\& &  Average   &  Symmetrized area &  Self-calibration correction & &&Calibrators correction\\\\ central freq. (GHz) & effective area (sr)                    & (sr) & average value + error & &&   average value + error    \\\\ \\hline\\hline% K (23)  & $(2.595 \\pm 0.014) \\times 10^{-4} $ & $2.46 \\times 10^{-4} $ & $1.050 \\pm 0.006$ &&& 0.99 $\\pm$ 0.06\\\\ Ka (33) & $(1.569 \\pm 0.012) \\times 10^{-4} $ & $1.44 \\times 10^{-4} $ & $1.086 \\pm 0.008$ &&& 1.12 $\\pm$ 0.07\\\\ Q  (41) & $(1.016 \\pm 0.007)\\times 10^{-4}  $ & $8.94 \\times 10^{-5} $ & $1.136 \\pm 0.008$ &&& 1.15 $\\pm$ 0.08\\\\ V  (61) & $(4.81 \\pm 0.07)\\times 10^{-5}    $ & $4.19 \\times 10^{-5} $ & $1.15 \\pm 0.02$ &&& 1.32 $\\pm$ 0.09\\\\ W  (94) & $(3.36 \\pm 0.17) \\times 10^{-5}   $ & $2.07 \\times 10^{-5} $ & $(1.62 \\pm 0.08)  $ &&&(1.41 $\\pm$ 0.10)\\\\% \\hline\\hline% \\end{tabular}  \\caption{Results of the self-calibration process (effective area/symmetrized area; see \\S~\\ref{sec:auto_cal}) and of the source calibration by comparison with ground based measurements  (ground based fluxes/NEWPS fluxes or average ratio of the corrected flux to the uncorrected one; see \\S~\\ref{sec:sou_cal}). The correction factors for the W band are in parenthesis because they are very uncertain (see text).} \\label{tb:cal_fac} \\end{table*}  The statistical properties of extragalactic sources  above $\\sim 10$ GHz are still largely unknown. This is due to the great difficulty of carrying out extensive surveys at high radio frequencies, because of the very small fields of view of ground-based radio telescopes and of the need of relatively long on-source integrations. Up to the end of the last century, the highest radio frequency at which large-area sky surveys had been performed was 5 GHz (GB6, \\citet{gre96}; PMN, \\citet{gri94,gri95}; \\citet{wri94,wri96}). Recently, the situation has considerably improved thanks to the Pilot ATCA survey that covered the declination region between $-60^\\circ$ and $-70^\\circ$ at 18 GHz down to 100 mJy \\citep{ricci04,sadler06}, to the 9C survey covering $520\\,\\hbox{deg}^2$ at 15.2 GHz to a flux limit of 25 mJy but going down to 10 mJy on small areas \\citep{waldram03,taylor01}, and to data from the VSA in the declination range $-5^\\circ <\\delta<60^\\circ$ at 33 GHz \\citep{cleary05}. The on-going Australia Telescope 20 GHz (AT20G) Survey \\citep{ekers}, once completed, will cover the whole Southern hemisphere down to 50 mJy.  Moreover, the Wilkinson Microwave Anisotropy Probe (WMAP) mission has produced the first all-sky surveys of extragalactic sources at 23, 33, 41, 61 and 94 GHz \\citep{ben03b,hin07}. >From the analysis of the first three years survey data the WMAP team have obtained a catalogue of 323 extragalactic point sources (EPS; \\citet{hin07}), substantially enlarging the first-year one that included 208 EPS detected above a flux limit of $\\sim 0.8$-1 Jy \\citep{ben03b}, with an estimated completeness limit of $\\sim 1.2$ Jy. The average estimated spectral index for detected sources is $\\simeq 0.0$, thus confirming that the source class which dominates the counts at bright fluxes is essentially made of flat-spectrum QSOs and BL Lacs, with minor contributions coming from steep- or inverted-spectrum sources \\citep{tof98,dz99,tof99,ben03b}.  In order to take full advantage of all the information content of the WMAP 3-year survey data, \\citet[hereafter LC07]{can07} have used the MHW2 filter \\citep{jgn06} to obtain estimates of (or upper limits on) the flux densities at the WMAP frequencies of a complete all-sky sample of 2491 sources at $|b|>5^\\circ$, brighter than 500 mJy at 5 GHz, or at 1.4 or 0.84 GHz in regions not covered by 5 GHz surveys but covered by either the NVSS \\citep{con98} or the SUMSS (Mauch et al. 2003). This work yielded $5\\sigma$ detections of 368 extragalactic sources, including 98 sources not present in the WMAP 3-year catalogue. The results were organized in a catalogue dubbed NEWPS$_{5\\sigma}$ (New Extragalactic WMAP Point Source) Catalogue.  As pointed out by LC07, the new flux estimates showed small systematic differences, increasing with frequency, with those obtained by the WMAP team. These differences were attributed to different approximations of the beam shapes. In fact, the beams are complex and asymmetrical. Taking also into account that sources are observed with different beam orientations, it is clear that the effective beam areas and, therefore, the flux calibrations, are uncertain.  To check the calibration of LC07 fluxes we have looked for ground based measurements of NEWPS$_{5\\sigma}$ sources at frequencies close to the WMAP ones, finding small but appreciable systematic differences. This prompted us to investigate in more detail the calibration problem. The adopted methods and the results are presented in \\S\\,\\ref{sec:calib}. In \\S\\,\\ref{sec:red_distrib_llf} we present the optical identifications and the redshift distribution of our sources, based on literature data. After having recalibrated the NEWPS$_{5\\sigma}$ flux densities we have investigated the spectral index distributions in different frequency intervals (\\S\\,\\ref{sec:pop_spec}) and used them to obtain model independent estimates of the mm-wave source counts, beyond the flux density intervals where they are directly determined (\\S\\,\\ref{sec:counts}). Finally, in \\S\\,\\ref{sec:Conclusions}, we summarize our main conclusions.  Throughout this paper we have used the Bayes Corrected version of the NEWPS Catalogue to take into account the well-known Eddington bias \\citep[see $\\S$ \\ref{sec:edd_bias} or LC07 for more details]{edding}\\footnote{The Bayes-corrected version of the NEWPS Catalogue can be found at http://max.ifca.unican.es/caniego/NEWPS/}. The spectral index, $\\alpha$, is defined as $\\alpha\\equiv -\\log(S_2/S_1)/\\log(\\nu_2/\\nu_1)$, where $S_1$ and $S_2$ are the fluxes at the frequencies $\\nu_1$ and $\\nu_2$, respectively (i.e., $S_\\nu \\propto \\nu^{-\\alpha}$). ",
        "conclusions": "\\label{sec:Conclusions}  We have presented and discussed a variety of statistical properties that characterize the extragalactic sources in the NEWPS Catalogue (LC07). The flux calibration has been investigated with 2 different methods and correction factors have been derived, except for the highest frequency WMAP channel, for which the available information is insufficient.  A search of the NED database has yielded optical identifications for 89\\% of the 252 sources in the complete NEWPS sample with  $S_{23\\rm GHz}\\geq 1.1\\,$Jy; 5 sources turned out to be Galactic, and were dropped. At high Galactic latitudes ($|b|>10^\\circ$) we obtain a sub-sample of 207 sources (92\\%) with spectroscopic redshifts (again from the NED), for a total number of extragalactic sources of 225 . The corresponding redshift distribution is in generally good agreement with the predictions of the model by \\citet{dz05}, except possibly around $\\log(z)\\simeq -1$ where the model predicts a significant contribution from FR II sources, that is not observed.  The distribution of spectral indices between 23 and 61 GHz, $\\alpha_{23}^{61}$, was obtained with the Kaplan-Meier estimator, taking into account the upper limits on 61 GHz fluxes. It is shifted towards steeper values compared with the distribution of $\\alpha_{5}^{23}$, obtained using the GB6 or PMN 5 GHz fluxes. Although the steepening can be just a selection effect, hints of a real high-frequency steepening are noted.  The source number counts obtained from our sample have close to Euclidean slope and are in good agreement with the prediction of the cosmological evolution model by \\citet{dz05}.  We have shown that the use of our spectral index distribution to extrapolate the 5 GHz counts to high frequencies leads to inconsistent results below $S_{23\\rm GHz}\\simeq 1\\,$Jy because of a substantial change in the mixture of flat- and steep-spectrum sources. The change, clearly visible comparing our distribution of spectral indices with that obtained by \\citet{waldram07} for a much fainter sample at the nearby frequency of 15 GHz, accounts for the difference between our high frequency counts and extrapolations from 9C counts. This is an example of the risks inherent in extrapolations with incomplete information of source properties.  Extrapolations to higher frequencies of the 23 GHz counts using the estimated distribution of $\\alpha_{23}^{61}$ tend to under-predict the observed counts, especially at 61 GHz, implying that the Kaplan-Meyer estimator probably overpopulates the steep portion of the spectral index distribution. The observed counts are recovered if the distribution is decreased by a factor of 0.2 for $\\alpha_{23}^{61}>0.5$."
    },
    "0711/0711.1542_arXiv.txt": {
        "abstract": "A variety of observations now indicate that intergalactic helium was fully ionized by $z \\sim 3$.  The most recent measurements of the high-redshift quasar luminosity function imply that these sources had produced at least $\\sim 2.5$ ionizing photons per helium atom by that time, consistent with a picture in which the known quasar population drives \\ion{He}{2} reionization.  Here we describe the distribution of ionized and neutral helium gas during this era.  Because the sources were rare and bright (with the photon budget dominated by quasars with luminosities $L \\ga L_\\star$), random fluctuations in the quasar population determined the morphology of ionized gas when the global ionized fraction $\\bxion$ was small, with the typical radius $R_c$ of a \\ion{He}{3} bubble $\\sim 15$--$20$ comoving Mpc.  Only when $\\bxion \\ga 0.5$ did the large-scale clustering of the quasars drive the characteristic size of ionized regions above this value.  Still later, when $\\bxion \\ga 0.75$, most ionizing photons were consumed by dense, recombining systems before they reached the edge of their source's ionized surroundings, halting the bubble growth when $R_c \\sim 35$--$40 \\Mpc$.  These phases are qualitatively similar to those in hydrogen reionization, but the rarity of the sources makes the early stochastic phase much more important.  Moreover, the well-known characteristics of the $z=3$ intergalactic medium allow a much more robust description of the late phase in which recombinations dominate. ",
        "introduction": "\\label{intro}  For most of the Universe's history, the intergalactic medium (IGM) evolved rather slowly and smoothly.  But there were two major exceptions:  the reionization of hydrogen (at $z \\ga 6$) and of helium (at $z \\sim 3$).  Recently, hydrogen reionization has received a great deal of attention in both the observational and theoretical communities (see, e.g., reviews by \\citealt{barkana01, ciardi05-rev, fan06-review, furl06-review}).  Helium reionization has received less attention, especially from theorists, but has actually been more accessible observationally and has important ramifications for quasar populations, galaxy formation, and the structure of the IGM itself.  Because \\ion{He}{2} has an ionization potential of 54.4 eV, the relatively soft photons produced by (known populations of) hot stars do not ionize it to any large degree (they can, on the other hand, singly ionize helium along with hydrogen).  As a result, helium remained singly ionized until the quasar population built up in sufficient numbers:  quasars have hard spectra, with luminosity densities $L_\\nu \\propto \\nu^{-1.6}$ \\citep{telfer02} and can produce enough photons to complete the reionization process by $z \\sim 3$ \\citep{sokasian02, wyithe03, gleser05}.  Observations do indeed point to helium reionization at $z \\sim 3$.  The strongest evidence comes from far-ultraviolet spectra of the \\ion{He}{2} \\lya forest along the lines of sight to several bright quasars at $z \\sim 3$:  the observed wavelength of such an absorption system is $\\lambda = 304(1+z)$ \\AA.  First suggested as a probe of the low-density IGM \\citep{miralda93, giroux95, miralda96, croft97}, the \\ion{He}{2} forest has also proven to be a powerful probe of reionization and of the background radiation field \\citep{fardal98}.  For our purposes, the most important point is that the apparent \\ion{He}{2} optical depth decreases rapidly at $z \\approx 2.9$, with a spread of $\\Delta z \\approx 0.1$ along different lines of sight \\citep{jakobsen94, davidsen96, anderson99, heap00, smette02, zheng04, shull04, reimers04, reimers05, fechner06}.  This apparent transition is analogous to the rapidly-increasing \\ion{H}{1} optical depth observed at $z \\sim 6$ and usually attributed to reionization \\citep{fan01, fan06}, though see \\citet{becker07} for a different interpretation.  Although appropriate lines of sight are rare (they require a bright quasar with no intervening Lyman-limit absorbers that attenuate the far-UV flux), recent large-scale surveys such as SDSS have opened up many more targets (e.g., \\citealt{zheng04-sdss}).  We may thus be able to observe the helium reionization era in detail.  Particularly interesting are the large opacity fluctuations observed along several lines of sight \\citep{anderson99, heap00, smette02, reimers05}.  At least one of these coincides with a nearby quasar \\citep{dobrzycki91, jakobsen03}.  This illustrates a crucial advantage of studying helium reionization over hydrogen:  we understand the $z \\sim 3$ Universe much better than the $z \\sim 6$ Universe, in terms of the ionizing sources, the IGM, and other galaxies.  We can thus provide much sharper tests of reionization models.  There are several other indirect lines of evidence for helium reionization at $z \\sim 3$, although in each case they are controversial.  Helium reionization should at least double the IGM temperature.  \\citet{schaye00} detected a sudden increase in the IGM temperature at $z \\sim 3.3$ by examining the Doppler widths of \\ion{H}{1} \\lya forest lines (see also \\citealt{schaye99, theuns02-reion}); at about the same time, the equation of state of the IGM also appears to become nearly isothermal, another indication of recent reionization \\citep{schaye00, ricotti00}.  However, temperature measurements via the \\lya forest flux power spectrum show no evidence for any sudden change, although the errors are rather large \\citep{zald01, viel04, mcdonald06}.  Recent models of the helium reionization process predict sudden jumps in the temperature as measured by the \\citet{schaye00} method, although with somewhat smaller magnitudes, along with smoother evolution in the mean temperature, which may help to resolve the controversy \\citep{gleser05, furl07-igmtemp}.  Such a temperature jump should also decrease the recombination rate of hydrogen, thereby decreasing the \\ion{H}{1} opacity \\citep{theuns02-sdss}.  \\citet{bernardi03} detected such a jump at $z \\sim 3.2$ in a large sample of SDSS spectra; however, \\citet{mcdonald06} found no such feature with similar data and precision.  Most recently, \\citet{fauch07} studied a sample of high-resolution \\lya forest spectra and found a feature nearly identical to that in \\citet{bernardi03}.  Finally, one would also expect the (average) metagalactic ionizing background to harden as helium is reionized, because the IGM would become transparent to high-energy photons.  This should manifest itself in the \\ion{He}{2}/\\ion{H}{1} ratio (as possibly observed along one line of sight; \\citealt{heap00}) and also in optical spectra.  In particular, the ionization potentials of \\ion{Si}{4} and \\ion{C}{4} straddle that of \\ion{He}{2}, so their ratio should evolve during helium reionization.  \\citet{songaila98, songaila05} found a break in their ratio at $z \\sim 3$; modeling of the ionizing background from optically thin and optically thick metal line systems also shows a significant hardening at $z \\sim 3$ \\citep{vladilo03, agafonova05, agafonova07}.  However, other data of comparable quality show no evidence for rapid evolution \\citep{kim02, aguirre04}.  This approach is made more difficult by the large fluctuations in the \\ion{He}{2}/\\ion{H}{1} ratio even after helium reionization is complete \\citep{shull04}.  Despite this wealth of (often controversial) data, helium reionization has received relatively little theoretical attention.  There have been a few attempts to estimate the evolution of the globally-averaged \\ion{He}{3} fraction with redshift, given models for the ionizing sources as input (e.g., \\citealt{wyithe03}).  There has also been one numerical simulation of helium reionization \\citep{sokasian02}, which confirmed that quasars provided enough photons to complete the process at $z \\sim 3$--$4$ and that the evolving opacity of the \\ion{He}{2} \\lya forest was consistent with the tail end of reionization.  Most recently, \\citet{gleser05} used a semi-analytic Monte Carlo model to examine both the optical depth and temperature evolution in the context of helium reionization.  However, there have been no efforts to understand the fundamental question of how the ionized and neutral gas are organized in the IGM (the ``morphology\" of reionization), which provides the overall paradigm in which observations must be interpreted.  On the other hand, in the past several years hydrogen reionization has received an enormous amount of theoretical attention, including both analytic models and simulations (see \\citealt{furl06-review} for a recent summary).  In particular, we now appreciate the importance of source clustering for reionization, which makes the process inhomogeneous on extremely large scales ($\\ga 10 \\Mpc$) and substantially affects the interpretation of many observables (\\citealt{barkana04, furl04-bub}, hereafter FZH04).  The purpose of this paper is to apply the simplest of these models to the helium reionization epoch as a first step toward understanding the large-scale inhomogeneity of the process.  In particular, we will describe three distinct stages in the evolution of the morphology, and we will emphasize the similarities and differences in the two eras.  We base our models on FZH04 and \\citet[hereafter FO05]{furl05-rec}.  This paper is organized as follows.  We begin in \\S \\ref{history} by considering simple reionization histories.  We then describe stochastic and clustering-driven models for the bubble sizes in \\S \\ref{stochastic} and \\S \\ref{bub}, respectively.  We next consider the role of inhomogeneous recombinations in \\S \\ref{recomb}.  Finally, we conclude in \\S \\ref{disc}.  In our numerical calculations, we assume a cosmology with $\\Omega_m=0.26$, $\\Omega_\\Lambda=0.74$, $\\Omega_b=0.044$, $H=100 h \\hunits$ (with $h=0.74$), $n=0.95$, and $\\sigma_8=0.8$, consistent with the most recent measurements \\citep{spergel06}.  Unless otherwise specified, we use comoving units for all distances. ",
        "conclusions": "\\label{disc}  We have examined several aspects of helium reionization using simple analytic models.  We first showed that the \\citet{hopkins07} luminosity function, together with standard template quasar spectra, produces $\\sim 2.5$ ionizing photons per helium atom at $z=3$, nicely consistent with observational hints that the IGM truly becomes fully ionized at that time.  Indeed, even with an average clumping factor $\\bar{C}=3$, the known quasar population can ionize helium by that time.  Interestingly, the quasar emissivity increases rapidly at $z \\ga 3$, with $\\sim 75\\%$ of the ionizing photons produced after $z=4$.  This is somewhat faster than naive models based on the collapse fraction predict, implying that supermassive black holes in quite massive halos dominate the photon budget.  In our calculations, we have used a quasar luminosity function with a shallow faint-end slope at $z \\ga 3$ \\citep{hunt04, cristiani04, hopkins07}.  This implies that most ionizing photons are produced by bright quasars ($L \\ga L_\\star$), which are of course relatively rare.  As a result, the initial phases of reionization are dominated by rare, large bubbles (with characteristic size $R_{\\rm char} \\sim 15 \\Mpc$)  whose overlap is random.  We described this regime with a ``stochastic\" model where sources are distributed according to Poisson statistics.  In this model, the bubble size distribution follows a power law (roughly), so there is no true characteristic scale -- although the power law is so steep that, in practice, nearly all bubbles contain one or at most a few sources anyway.  Thus, in the early stages of reionization, we expect bubbles to have a characteristic size corresponding to a single or at most a few sources.  However, more recent data shows a somewhat steeper faint-end slope \\citep{fontanot07, bongiorno07, siana08}.  This leads to a broader distribution of luminosities, making the stochastic phase somewhat less important.  Large-scale clustering of the ionizing sources will eventually cause the bubbles to grow beyond this point.  We use the model of FZH04 to describe this phase, in which large-scale overdensities are ionized first.  This imprints a well-defined characteristic scale, which first exceeds the size of the stochastic bubbles when $\\bxion \\sim 0.5$ and would exceed $100 \\Mpc$ by the end of reionization if allowed to grow unchecked.  The scale is nearly independent of redshift at a fixed $\\bxion$, but it does depend on the clustering of the quasar hosts, with more massive host galaxies driving larger bubbles. In our calculations, we used a simple halo-based model to represent quasar hosts.  This is reasonable if the observed relations between black hole mass and halo properties \\citep{magorrian98, gebhardt00, ferrarese00} hold at high redshift as well.  But many detailed models of the quasar luminosity function rely on mergers to trigger quasar phases.  This would change the distribution of quasar luminosities over halos and so change some of the predictions of our model (see also \\citealt{cohn07}).  We intend to examine these effects with the semi-numeric approach of \\citet{mesinger07}.  However, even during this phase nearly all quasars exist in isolation:  provided that their duty cycle $t_{\\rm QSO} H(z) \\ll 1$, very few bubbles will contain more than one \\emph{active} source.  Thus the ionizing background around each source is easy to compute (modulo the nearly uniform background of high-energy photons with large mean free paths).  Together with a model for the IGM, this allows us to estimate the mean free path of helium-ionizing photons as they propagate through the ionized bubbles and specifically to compute how far they can travel before being consumed by dense, optically thick absorbers.  We find that, at the ionization edge, $R_{\\rm max} \\sim 35$--$40 \\Mpc$ when averaged over the entire quasar population.  Once bubbles grow to this size, only a fraction of the photons would ionize new helium atoms.  As described in FO05, the effective clumping factor of the IGM will increase as more and more photons get consumed by these dense systems, so reionization will slow down.  Eventually, the ionized gas distribution will map onto the more uniform background characteristic of the post-reionization era (although even then large fluctuations will persist; \\citealt{shull04}).  Our model does not include recombinations within bubbles that contain zero active quasars over a period of time, except as an overall scaling factor to the the ionizing efficiency (through $\\bar{N}_{\\rm rec}$).  The helium recombination time is short enough, and the IGM clumpiness large enough, that a significant fraction of helium atoms (mostly in high density regions) can recombine, although $\\ga 90\\%$ of the volume will remain highly ionized.  We examine these ``fossil\" bubbles more closely in \\citet{furl08-fossil}.  Compared to hydrogen reionization, the helium era passes through the same qualitative phases: stochastic, isolated bubbles followed by a rapid increase in the bubble size when clustering becomes significant and finally a ``saturation\" when the mean free path of an ionizing photon becomes smaller than the bubble size.  However, quasars are much more rare than star-forming galaxies at $z \\ga 6$ (at least in most models of the high-redshift universe), so the early stochastic phase is much more important; for hydrogen, it ends by $\\bxion \\sim 0.1$.  There is thus a much shorter window in which density-driven clustering dominates the morphology, when $R_{\\rm stoch} < R < R_{\\rm max}$.  Because there is typically no more than one quasar per bubble, the scale $R_{\\rm max}$ is much easier to estimate than for hydrogen.  FO05 estimated it during hydrogen reionization by balancing the total recombination rate within the bubble against the emissivity of the ionizing sources.  This depends on the recombination history and uncertain IGM density distribution, and it assumes that, within each bubble, the ionization always begins at low densities.  As a result, $R_{\\rm max}$ is quite sensitive to many of the input parameters (varying by about a factor of two between case-A and case-B recombination, for example).  For helium reionization, $R_{\\rm max}$ can be much more robustly determined because it depends only on ionization balance around a source population (whose luminosity function is quite well-known, at least at $z \\la 4$) in an IGM whose density structure is extremely well-determined by the \\ion{H}{1} \\lya forest (although the spread in luminosities also makes it more difficult to incorporate directly into the simple analytic model).  The helium reionization era will thus offer a much more stringent test of the transition to the recombination-dominated phase than the hydrogen era.  There is, however, a possibility that hydrogen and helium reionization are more similar than it may appear.  We showed that observations at least allow the possibility that $z \\approx 3$ galaxies produce nearly as many helium-ionizing photons as quasars.  Their importance for reionization depends entirely on their escape fraction, which is difficult to estimate.  If they do contribute, we would expect a slower evolution of $\\bxion$ with redshift, a softer ionizing background, many more small ionized bubbles, and a more uniform, smaller mean free path (and hence $R_{\\rm max}$) for the helium-ionizing photons.  We have focused exclusively on analytic models here, but of course a full picture of helium reionization requires numerical simulations that capture both the stochastic and large-scale clustering elements of our model, as well as the complex geometry of the IGM and source distribution.  The models described here can be extended to ``hybrid\" semi-numeric schemes \\citep{zahn07-comp, mesinger07} and also full radiative transfer simulations (e.g., \\citealt{sokasian03}).  However, some elements -- principally recombinations -- will require either small-scale hydrodynamic simulations or analytic prescriptions similar to the ones used here to model accurately.  Our analytic models thus provide a basic underpinning for understanding and improving more complex numerical realizations of reionization.  Finally, we must acknowledge that observing the morphology of helium bubbles directly will likely be rather difficult.  During hydrogen reionization, the two most promising techniques are to use the 21 cm transition of neutral hydrogen to map the bubble distribution (e.g., \\citealt{furl06-review}) and to infer the distribution of neutral gas from \\lya line-selected galaxies.  Neither is available here, because $^4$He lacks a hyperfine transition and galaxies are already so faint in the far-UV.  Instead we must turn to more subtle approaches.  One possibility is through \\ion{He}{2} \\lya forest absorption spectra of distant quasars:  large bubbles may manifest themselves as transmission spikes separated by long ``dark gaps\" of saturated absorption.  Another is to search for large-scale coherent fluctuations of some proxy for the ionized helium abundance, such as in the thermal properties of the IGM \\citep{furl07-igmtemp}, absorption in the \\ion{H}{1} \\lya forest, or the abundance ratios of metal ions.  We intend to explore all of these issues in future work."
    },
    "0711/0711.4367_arXiv.txt": {
        "abstract": "We report precise Doppler measurements of \\npllet\\ stars, obtained at Lick Observatory as part of our search for planets orbiting intermediate--mass subgiants. Periodic variations in the radial velocities of both stars reveal the presence of substellar orbital companions. These two stars are notably massive with stellar masses of \\mstarC~\\msun\\ and \\mstarA~\\msun, respectively, indicating that they are former A--type dwarfs that have evolved off of the main sequence and are now K--type subgiants. The planet orbiting \\starC\\ has a minimum mass \\msini~$=$~\\msiniC~\\mjup, eccentricity $e = \\eC$ and a \\pC~day period, corresponding to a semimajor axis $a = \\arelC$~AU. The planet around \\starA\\ has a minimum mass \\msini~$=$~\\msiniA~\\mjup\\ and a \\pA~day orbit, corresponding to a semimajor axis $a = \\arelA$~AU. The eccentricity of \\starA\\,b is consistent with circular ($e = \\eA \\pm \\eeA$), adding to the rare class of known exoplanets in long--period, circular orbits similar to the Solar System gas giants. Like all of the planets previously discovered around evolved A stars, \\starC b and \\starA b orbit beyond 0.8~AU. ",
        "introduction": "Most of what is known about planets outside of our Solar System comes from Doppler surveys of Sun--like stars, with spectral types ranging from K0V to F8V and masses between 0.8 and 1.2~\\msun\\ \\citep{valenti05, butler06, takeda06}. However, recent results from planet searches around low--mass M dwarfs \\citep[][]{bonfils05b,endl06,butler06b} and evolved, intermediate--mass stars \\citep{reffert06, johnson06b, sato07, nied07, lovis07} have expanded into the frontiers at either end of the stellar mass range. These surveys have begun to reveal important relationships between stellar mass and the properties of exoplanets. For example, planets around ``retired'' (evolved) A--type stars reside preferentially in wide orbits, with semimajor axes $a \\gtrsim 0.8$~AU \\citep{johnson07} and appear to have larger minimum masses than planets around Sun--like stars \\citep{lovis07}. Also, the occurrence rate of Jovian planets (\\msini$\\geq 0.5$~\\mjup, $a \\leq 2.5$~AU) increases with stellar mass, rising from $< 2$\\% around M dwarfs, to approximately 9\\% around F and A stars \\citep{johnson07b}.  The effects of stellar mass on other characteristics of exoplanets, such as orbital eccentricity and multiplicity, will become evident as the sample of planets around intermediate--mass stars grows. Here, we report the detection of two Jovian planets orbiting stars with masses of \\mstarA~\\msun\\ and \\mstarC~\\msun. These planet detections come from our sample of intermediate--mass subgiants that we have been monitoring at Lick and Keck Observatories \\citep{johnson06b}. We describe our observations in \\S~\\ref{observations}. In \\S~\\ref{stellar} we present our data, and describe the characteristics of the host stars and the orbits of their planets. We conclude in \\S~\\ref{summary} with a summary of our results and a discussion of the properties of planets orbiting evolved A--type stars. ",
        "conclusions": "\\label{summary}  We report the detection of \\npllet\\ Jovian planets orbiting intermediate--mass subgiants. These detections come from our sample of evolved stars that we are monitoring at Lick and Keck Observatories. \\starA\\ and \\starC\\ have masses significantly larger than solar, with $M_* = \\mstarA$~\\msun\\ and  \\mstarC~\\msun, respectively. Examination of the stars' theoretical mass tracks reveals that these current--day subgiants began life on the main sequence as A--type dwarfs (Figure~\\ref{sg_hr}).  The relatively low--amplitude Doppler variations induced by \\starC b and \\starA b would not have been detectable if the stars were not in their current evolved state. Younger A-- and F--type dwarfs have large rotational velocities (\\vsini~$\\gtrsim 50$~\\ks) and excessive pulsation--induced velocity jitter that can mask the reflex velocity signal caused by a planet--mass companion \\citep{galland05}. On the other hand, older stars ascending the red giant branch exhibit stochastic velocity variations in excess of 20~\\ms \\citep{hekker06}, which would make the  planet orbiting \\starC\\ particularly difficult to detect since it induces a velocity amplitude of only \\kC~\\ms. Subgiants, with their low rotation rates (\\vsini~$ \\lesssim 10$~\\ks) and low jitter ($< 10$~\\ms), occupy an observational ``sweet spot'' in the H--R diagram allowing our Doppler survey to probe to relatively low planetary masses in a wide range of orbital configurations.  The growing sample of planets around intermediate--mass stars is beginning to reveal important relationships between stellar mass and the properties of exoplanets. Of the 15 planets detected around evolved A--type stars ($M_* > 1.5$~\\msun), none have been found orbiting closer than $\\sim 0.8$~AU, with the majority orbiting at or beyond 1~AU \\citep{johnson07}. \\starC b and \\starA b are no exception to this trend, with semimajor axes of \\arelC~AU and \\arelA~AU, respectively. As noted by \\citet{johnson07}, this cannot be due to a decrease in detection sensitivity since, for a given planet and stellar mass, the amplitude of the induced Doppler wobble scales as $a^{-1/2}$. Thus, the observed semimajor axis distribution of planets around A stars is significantly different from that of planets around lower--mass stars.  \\citet{johnson07} considered the possibility that the lack of close--in planets around K giants and clump giants may be attributable to engulfment by the expanding atmospheres of the central stars. However, stars crossing the subgiant branch do not undergo significant expansion, with radii smaller than $\\sim 5$~\\rsun\\ even at the base of the red giant branch. The fact that none of the 7 planets detected around massive subgiants ($M_* > 1.5$~\\msun) has $a < 0.9$~AU strongly suggests that the lack of close--in planets around A stars is due to the effects of stellar mass on planet formation and migration, rather than post--main--sequence engulfment.  Stellar mass also plays an important role in the likelihood that a star harbors a detectable giant planet. By measuring the fraction of stars with planets in three widely spaced stellar mass bins, \\citet{johnson07b} found that the occurrence rate of planets with \\msini~$>0.8$~\\mjup\\ and $a < 2.5$~AU is a rising function of stellar mass. This analysis reveals that A stars appear to be planet--enriched by a factor of 4.5  compared to low--mass M dwarfs, with a measured giant planet occurrence rate of 9\\% for $1.3 < M_*/M_\\odot \\leq 1.9$~\\msun. The planet fraction for higher--mass stars was based in part on three strong planet candidates from our survey of subgiants, one of which is announced here (\\starA b; \\starC b orbits beyond the 2.5 AU cut--off). This result indicates that stellar mass is a strong tracer of planeticity, which has important implications for the target selection of future planet searches. Just as stellar metallicity is exploited in the search for short--period planets around Sun--like stars \\citep{fischer05b, fischer05a, dasilva06}, stellar mass should be an important consideration in the selection of targets for future photometric, astrometric, and high--contrast direct imaging surveys.  The rising trend in planet occurrence toward higher stellar masses also informs models of planet formation. Several theoretical studies of the effects of stellar mass on planet formation have predicted that M stars should harbor fewer Jovian planets than Sun--like dwarfs \\citep{laughlin04, ida05b}. More recently \\citet{kennedy07} studied planet formation at higher masses by accounting for the evolving disk mid--plane temperature due to stellar irradiation and viscous evolution. They predict that the fraction of stars with Jupiters should rise up to a peak of $\\approx 20$\\% near $M_* = 3$~\\msun.  Testing this prediction requires searching for planets around stars with masses greater than 2~\\msun. Unfortunately, the mass range of suitable subgiants is limited to $M* \\lesssim 2.2$~\\msun, primarily due to our absolute magnitude criterion of $M_V > 2.0$ (see \\S~\\ref{observations}). Doppler surveys of massive K giants with $M_* > 2.5$~\\msun\\ provide perhaps the most favorable avenue for exploring the occurrence rate of planets around more massive A stars \\citep{frink01, sato05b, nied07}. In particular, \\citet{lovis07} are searching for planets around K giants in open clusters, which allows them to take advantage of the uniform ages of the cluster members to derive accurate stellar masses. The preliminary results from their survey indicate that stars more massive than $\\sim2$~\\msun\\ have an enhanced abundance of super--Jupiters and brown dwarfs compared to lower--mass stars.  To enlarge the statistical ensemble of planets detected around intermediate--mass stars, we have expanded our planet search by adding 300 additional subgiants at Lick and Keck Observatories. If our current 9\\% detection rate holds, our expanded planet search is expected to yield an additional 20-30 planets over the next 3 years. Together with the detections from Doppler surveys of K giants and clump giants, the increased number of planets orbiting intermediate--mass stars will add significantly to our understanding of the effects of stellar mass on planet formation and planetary system architecture."
    },
    "0711/0711.1297_arXiv.txt": {
        "abstract": "Dark matter caustics are small scale, high density structures believed to exist in galaxies like ours. If the dark matter consists of Weakly Interacting Massive Particles, these caustics may be detected by means of the gamma rays produced by dark matter particle annihilation. We discuss particle annihilation in outer and inner caustics and provide sky maps of the expected gamma ray distribution.  \\vspace{0.34cm} \\noindent \\pacs{98.80 Cq} ",
        "introduction": "}  A consequence of the presence of cold dark matter in galactic halos is the formation of caustics. Caustics are regions of infinite density in the limit where the dark matter particles have zero velocity dispersion. The leading dark matter candidates, namely axions and Weakly Interacting Massive Particles(WIMPs), have very small primordial velocity dispersions ($\\sim 0.03$ cm/s for a 100 GeV WIMP, $\\sim 10^{-6}$ cm/s for a 10 $\\mu$eV axion \\cite{ps_crs}) and may be expected to form caustics with large density enhancements. If the dark matter is made up of WIMPs, the particles can annihilate producing gamma rays. It may be possible to detect this gamma ray flux by means of future experiments such as GLAST \\cite{glast}. Caustics may possibly be revealed by this technique.  Caustics form because cold dark matter particles occupy a thin, three dimensional hypersurface in phase space \\cite{ps_crs,caustic_papers}. There are two kinds of caustics in galactic halos: outer and inner \\cite{ps_crs}.  The outer caustics are topological spheres surrounding galaxies. The inner caustics have a more complicated geometry which depends on the details of the dark matter angular momentum distribution \\cite{ns2}. When the angular momentum distribution is dominated by a net rotational term, the inner caustics have the appearance of rings \\cite{ps_crs,ns2}. Observational evidence for caustic rings was found by analyzing the rotation curves of external galaxies \\cite{ks} and the rotation curve of the Milky Way \\cite{iras}. A triangular feature seen in the Infrared Astronomical Satellite (IRAS) map of the Milky Way is also suggestive of the presence of a nearby inner caustic ring \\cite{iras}. Recently \\cite{os}, possible evidence was found for a caustic ring in the galaxy cluster Cl0024+1654.  In this article, we discuss dark matter annihilation in caustic spheres and caustic rings and present simulated gamma ray sky maps.  The existence of dark matter caustics is sometimes disputed on the basis of large scale cosmological simulations. However, current simulations use particles with mass of order $\\sim 10^6 M_\\odot$. Such massive particles have spurious collisions which destroy fine details in the dark matter phase space hypersurface. Thus the absence of caustics in most large scale simulations is probably due to the finite resolution of these simulations and the large particle masses involved. Note that discrete flows and caustics are observed in $N$-body simulations when special care is taken.  The simulations of Stiff and Widrow \\cite{Sti}, which increase the resolution in the relevant regions of phase space, do show discrete flows.  More recently, Vogelsberger et al. \\cite{Voge} employed a new technique to study small scale structures such as caustics.  Further progress in numerical techniques may enable the study of caustic properties in galactic halos using $N$-body simulations.  In cold dark matter cosmology, galactic halos form through a process of ``hierarchical clustering\".  Small halos form first.  Larger halos, such as that of the Milky Way, are the outcome of past mergers, the accretion of smaller halos of various sizes, and the accretion of unclustered dark matter. This general scenario must surely be correct without, however, telling us anything precise.  The critical issue for the existence of a caustic is the velocity dispersion of the flow that would form it.  To be specific, consider the flow of dark matter that the Milky Way halo is accreting {\\it today}.  If the dark matter in this flow has not clustered on any small scales yet, its velocity dispersion has the primordial value mentioned earlier, of order 0.03 cm/s for a 100 GeV WIMP.  If it has clustered already, its effective velocity dispersion is the velocity dispersion of the structure it has become part of. The largest structures for which we have evidence of recent accretion are small sattelites, such as the Sagittarius dwarf, with velocity dispersion of order 10 km/s.  So one may conservatively assume that the flow of WIMP dark matter accreting onto the Galaxy today has variable velocity dispersion between 0.03 cm/s and 10 km/s.  Dark matter caustics are generically surfaces on one side of which the density diverges, in the limit of zero velocity dispersion, as \\begin{equation} d(x) = {A \\over \\sqrt{x}} \\label{foldden} \\end{equation} where $x$ is the distance to the surface and $A$ is a constant called the ``fold coefficient\".  For finite velocity dispersion $\\delta v$, the dark matter caustics in a galactic halo are smoothed over a distance \\begin{equation} \\delta x \\sim {\\delta v \\over v} R \\label{cutoff} \\end{equation} where $v$ is the flow velocity (of order 300 km/s) and $R$ is the overall size of the halo (of order 300 kpc).  The velocity dispersion cuts off the divergence of the density at the caustic, i.e. the ${1 \\over \\sqrt{x}}$ behaviour stops for $x < \\delta x$. For the above range of velocity dispersion, $\\delta x$ ranges from $10^{15}$ cm to 10 kpc.  Since outer caustics have size of order hundreds of kpc, they are not erased by the expected velocity dispersion.  Inner caustics which have size of order 10 kpc are smoothed out only when the flow is entirely composed of sattelites with velocity dispersion of order 10 km/s.  As was mentioned, evidence  for caustic rings of dark matter in the Milky Way was found in the existence and distribution of sharp rises in the Galactic rotation curve and the presence of a triangular feature in the IRAS map of the Galactic plane \\cite{iras}.  This evidence implies an upper limit on the velocity dispersion of order 50 m/s, or an upper limit on the smoothing scale $\\delta x$ of order $10^{-4}~R \\sim 30$ pc.  The WIMP annihilation rate per unit surface of caustic is proportional to the integral \\begin{equation} \\int dx~d(x)^2 \\simeq \\int_{\\delta x}^R {A^2 \\over x} = A^2 \\ln\\left({R\\over \\delta x}\\right)~~~\\ . \\label{logenh} \\end{equation} This equation shows that caustics enhance the annihilation signal only by a logarithmic factor \\cite{detecting} which varies from $\\sim 3$ to 20, depending on the velocity dispersion.  In this article, we limit ourselves to the simple cases of a spherically symmetric outer caustic and an axially symmetric inner caustic ring.  Our symmetry assumptions are made solely for the purpose of computational ease.  In previous work \\cite{ns2}, we verified that caustics are stable under perturbations, in particular perturbations that break any assumed symmetry. Moreover, the symmetry assumptions are consistent with the observational evidence mentioned earlier.  Previous work on the possibility of detecting dark matter caustics by indirect means include \\cite{detecting,pieri,an,mss}. Caustics and cold flows are also relevant to WIMP and axion direct detection experiments \\cite{direct}.  The number of photons received by a detector per unit area, per unit time, per unit solid angle and per unit energy due to WIMP annihilation is given by the expression\\cite{pieri,gondolo} \\begin{equation} \\frac{d\\Phi}{d\\Omega dE_\\gamma }(E_\\gamma,l,b) = \\frac{dS}{dE_\\gamma}(E_\\gamma) \\times \\frac{E(l,b)}{4\\pi}. \\label{wimp_formula} \\end{equation} The term $dS/dE_\\gamma$ depends only on the particle physics: \\begin{equation} \\frac{dS}{dE_\\gamma}(E_\\gamma) = \\frac{ <\\sigma_a v>}{2 m^2_\\chi} \\sum_i b_i \\; \\frac{dN_{\\gamma,i}}{dE_\\gamma}(E_\\gamma).  \\end{equation} $<\\sigma_a v>$ is the WIMP annihilation cross section times the relative velocity averaged over the momentum distribution of the WIMPs, $m_\\chi$ is the WIMP mass, $dN_{\\gamma,i}/dE_\\gamma$ is the number of photons produced per annihilation per unit energy $E_\\gamma$ in the annihilation channel $i$, and $b_i$ is the branching fraction of channel $i$.  The term $E(l,b)$ is called the emission measure and depends solely on the dark matter distribution. \\begin{equation} E(l,b) = \\int_{los} dx \\; d^2(x) \\end{equation}  where $x(l,b)$ is the distance measured along the line of sight($los$) in the direction of the galactic co-ordinates $(l,b)$, and $d$ is the dark matter density. Since all detectors have a finite resolution, the emission measure is often expressed in terms of $<E>$, which represents an average over the resolution of the detector. In this article, we limit ourselves to computing $E(l,b)$.  As an example, let us compute the emission measure due to a smooth isotropic halo with the density profile \\begin{equation} d(r) = d(r_0) \\, \\frac{ r_0^2 + a^2 }{ r^2 + a^2 }. \\label{smooth} \\end{equation} Consider an observer at a distance $r_0$ from the center of the halo, and a line of sight making an angle $\\psi$ with the line joining the observer with the center (we choose $\\psi = 0$ to be the direction of the halo center). $a$ is the core radius and keeps the density finite. We choose $d = 0.3$ GeV/cm$^3$ when $r = r_0 = 8.5$ kpc.  The emission measure $E$ is plotted as a function of angle $\\psi$ in Fig. \\ref{figure1}, for the cases $a = 1,2$ and 4 kpc. ",
        "conclusions": "}  We have discussed dark matter annihilation in outer and inner caustics and provided simulated skymaps of the emission measure. In Section \\ref{section2}, we considered spherical outer caustics. When the observer is located outside the sphere, the emission measure $E$ is largest when the line of sight is tangent to the sphere, falling off rapidly to zero thereafter (Figs \\ref{figure2},\\ref{figure3}). When the observer is within the sphere, the emission measure only changes gradually with the angle $\\psi$ (Figs. \\ref{figure4},\\ref{figure5}).  In Section \\ref{section3}, we considered circular, symmetric tricusp ring caustics. When the observer is outside the caustic ring, the emission measure is largest when the line of sight is tangent to the ring. Looking tangentially, one may hope to identify the tricusp cross section of the ring[Fig. \\ref{figure6}(b)]. When the observer is within the tricusp, the emission measure is largest when the lines of sight pass close to the cusps, resulting in the bright ring seen in Fig. \\ref{figure7}. We find that for inner caustics, $E$ can be significantly larger than that produced by a smooth halo, especially in directions away from the halo center."
    },
    "0711/0711.2682_arXiv.txt": {
        "abstract": "It is well known that the mere location of a Beat Cepheid model in a Period Ratio \\vs Period diagram (Petersen diagram) puts constraints on its metallicity $Z$.  But these bounds are sensitive to the mixture of elements that are lumped into the parameter $Z$.  In this short paper we update the previous results that were based on the Grevesse-Noels solar mixture to the recent, revised \\cite{ags05} (AGS) solar mixture.\\\\ We also examine the effect of the envelope depth on the accuracy of the computed pulsation periods.  We find that for low period Cepheids with high $Z$ the customary approximation of envelope pulsation breaks down.  It is necessary to compute stellar models that extend to the center and to include burning and composition inhomogeneities in the modeling.  Fortunately, however, most Beat Cepheids that have been observed so far seem to avoid that regime. ",
        "introduction": "\\label{sec:intro}  A recent paper by \\cite{BS07} (Paper~I, hereafter) described the results of a modeling survey of Beat Cepheids that pulsate in the fundamental (F) mode and the first overtone (O1) simultaneously.  In order to allow a convenient comparison with the observations, Paper~I presented the results in a $P_{10}$ vs. $P_0$ diagram, a so called Petersen diagram, as a function of metallicity $Z$, where $P_0$ and $P_1$ are the fundamental and first overtone periods, respectively, and $P_{10} = P_1/P_0$.  The two periods can easily be extracted from the observed light curves.  From the metallicity dependent Petersen diagram one can set bounds on the metallicity $Z$ of the star.  This way of determining the metallicity is independent of and complementary to all other stellar and galactic metallicity probes.  \\cite{BS07} found that the Petersen diagrams are rather insensitive to stellar rotation rates and to the galactic helium enrichment function $Y = Y(Z)$ that they used.  However, the results exhibited sensitivity to the chemical mixture that is lumped into the metallicity parameter $Z$, mainly through their effect on the opacities.  The survey of paper I used the OPAL GN93 opacities (\\cite{IR96} for the solar mixture of \\cite{GN93}, merged with the \\cite{AF94} opacities at the lower temperatures (log $T<$3.95).  In the last couple of years \\cite{ags05} have dramatically revised the solar chemical composition of metals.  The most important change is a decrease of elements with atomic numbers $Z<11$, but dominated by O, in favor of elements with $Z>11$ of which Fe has the most important effect on the opacities.  In view of the sensitivity of the Petersen diagrams to the chemical mixture we have thought it useful to update our survey with OPAL opacities that reflect these new abundances.  In this study, as in Paper I, we use a solar elemental mixture even in low metallicity Cepheids out of expediency, because there is no observationally derived information available that would let one improve much upon such an approximation at this time.  In a linear study such as this the Beat Cepheids are defined as being simultaneously unstable in the fundamental (F) and the first overtone (O1) modes.  However, it is well known from nonlinear Cepheid studies (\\eg \\cite{kbsc02,kbby98}) that not all stars in this 'linear Beat Cepheid instability strip' will actually undergo stable Beat pulsations (a large fraction will be either F {\\sl or} O1 Cepheids).  This computationally convenient, linear definition therefore considerably overestimates the 'Beat Cepheid instability strip' (broader by at least a factor of 3 -- 4 in \\Log \\Teff\\ in a theoretical HR diagram \\citep{sb08}.  This in turn increases considerably the separation of the upper and lower boundaries in the Petersen diagrams with a concomitant increase in the uncertainty of the stellar metallicities.  The linear properties of stars do not allow one to infer the modal selection which is intrinsically nonlinear \\citep{stellingwerf74,bk86}.  Therefore, for a tighter, and correct, definition and study of Beat Cepheids it is necessary to resort to a full amplitude hydrodynamic survey of models which requires a huge computing effort.  Such a survey will be presented in a separate paper (\\cite{sb08}).  \\begin{figure} \\vskip 30pt \\epsscale{1.1} \\ifthenelse{\\boolean{color}} {\\plotone{f1_color.eps}}{\\plotone{f1.eps}} \\caption{\\small Helium galactic enrichment function $Y(Z)$ that is used in the model computations. } \\label{figyz} \\end{figure}  \\begin{figure*} \\epsscale{0.9} \\ifthenelse{\\boolean{color}} {\\plotone{f2a_color.eps}}{\\plotone{f2a.eps}} \\vskip 10pt \\ifthenelse{\\boolean{color}} {\\plotone{f2b_color.eps}}{\\plotone{f2b.eps}}  \\caption{\\small $P_1/P_0$ \\vs \\Log $P_0$ \\th ($P_0$ on top axis) where $P_0$ and $P_1$ are the fundamental and first overtone periods.  The lines delimit the ranges for which both F and O1 are linearly unstable: Lower boundaries in the left, and upper boundaries in the right panels.  The metallicity increases downward in the figure from $Z$ = 0.001 (top line) to 0.016 in steps of 0.001, and then from in steps of 0.002 to 0.020 (bottom line).  (In the left panels the lowest $Z$ values are not shown to avoid overcrowding).  The long-dashed line refers to the specific AGS composition of the Sun, $Z_\\odot=0.0122$ and $Y_\\odot =0.2486$. For reference we have superposed the location of most known Beat Cepheids.}  \\label{figags} \\end{figure*}  In \\S2 we display the computed Petersen diagrams for the AGS solar elemental composition.  In \\S3 we discuss the adequacy of envelope models for Cepheids, and find that it is necessary to compute full stellar models for the short period, low metallicity Cepheids.  We conclude in \\S4. ",
        "conclusions": "We have computed the Beat Cepheid models that make up the composite Petersen diagrams as a function of metallicity $Z$ with the \\cite{ags05} solar mixture.  This change causes a noticeable reduction of $\\sim$ 25 - 30 \\% in the average $\\langle Z\\rangle$ of the four galaxies in which a sizable number of Beat Cepheids are known.  It considerably improves the agreement between the average Galactic metallicity and the Sun's metallicity over what was obtained with the GN93 abundances.  While the elemental mixture of the low $Z$ stars is most likely not quite solar, this is the best approximation we can use in the absence of detailed determinations of chemical compositions from high resolutoin spectra of such stars.  Because we needed to construct Beat Cepheid models down to periods as short as 0.5\\thinspace d as have been found in the SMC by \\cite{mbb07}, we have been led to reexamine the accuracy of commonly used envelope models for Cepheid stars. We find that because of the high accuracy demands of Petersen diagrams such envelope models become inadequate for short period, high $Z$ Cepheids.  Good fortune has it though that the higher $Z$ Beat Cepheids in the Galaxy and in M33 have longer periods, and that the Beat Cepheids with smaller periods are observed in the lower metallicity Magellanic Clouds.  Thus the vast majority of the known Beat Cepheids avoid to some extent the troublesome region of uncertainty of envelope models."
    },
    "0711/0711.0152_arXiv.txt": {
        "abstract": "The energy--zenith angular event distribution in a neutrino telescope provides a unique tool to determine at the same time the neutrino-nucleon cross section at extreme kinematical regions, and the high energy neutrino flux. By using a simple parametrization for fluxes and cross sections, we present a sensitivity analysis for the case of a km$^3$ neutrino telescope. In particular, we consider the specific case of an under-water Mediterranean telescope placed at the NEMO site, although most of our results also apply to an under-ice detector such as IceCube. We determine the sensitivity to departures from standard values of the cross sections above 1 PeV which can be probed independently from an a-priori knowledge of the normalization and energy dependence of the flux. We also stress that the capability to tag downgoing neutrino showers in the PeV range against the cosmic ray induced background of penetrating muons appears to be a crucial requirement to derive meaningful constraints on the cross section. ",
        "introduction": "High energy neutrino astronomy is one of the most promising research lines in astroparticle physics. Neutrinos are in fact one of the main components of the cosmic radiation in the high energy regime, and although their fluxes are uncertain and depend on the specific production process, their detection would provide valuable information concerning the sources and the acceleration mechanism in extreme astrophysical environments. Similarly to photons and unlike cosmic rays, they keep directional information which can be used to perform astronomy. Differently from gamma rays, they are emitted only in hadronic processes and travel unimpeded through cosmological distances well above the TeV energy. Moreover, due to the peculiar charge assignment of neutrinos in the electroweak standard model, they may be the particles most sensitive to exotic physics. An intriguing possibility is that a measurement of their cross section on ordinary matter in extreme and still unexplored kinematical regions might unveil new physics (see e.g. \\cite{Anchordoqui:2005is}).  From the experimental point of view, after the first generation of telescopes has proved the viability of the Cerenkov detection technique under deep water (\\verb\"Baikal\" \\cite{Balkanov:1999up}) and ice (\\verb\"AMANDA\" \\cite{Ahrens:2002gq}) by detecting atmospheric neutrinos, one is probably on the verge of the first detections at the \\verb\"IceCube\" \\cite{Ahrens:2003ix} telescope, being completed at the South Pole, and possibly at the smaller \\verb\"ANTARES\" \\cite{antares} telescope under construction in the Mediterranean. Additionally, \\verb\"ANTARES\" as well as \\verb\"NESTOR\" \\cite{nestor} and \\verb\"NEMO\" \\cite{nemo} are involved in R\\&D projects aimed at the construction of a km$^3$ Neutrino Telescope (NT) in the deep water of the Mediterranean sea, coordinated in the European network \\verb\"KM3NeT\" \\cite{km3net}.  In the present paper we address in a more quantitative way the question of how well a km$^3$ NT performs in the simultaneous determination of the neutrino flux and the neutrino-nucleon cross section by using as observable the energy--zenith angular event distribution.  Note that, unless the spectrum is known a priori, for a NT surrounded by an isotropic medium the cross section and the flux would be completely degenerate. However, under the sole assumption that the incoming spectrum is isotropic, the different opacity of the Earth in different directions --- which becomes relevant above $\\sim$100 TeV --- allows one to disentangle the two observables, as it has been discussed before e.g. in \\cite{Hooper:2002yq,Hussain:2006wg,Anchordoqui:2005pn}. In other words, a NT is only sensitive to the product of cross section and flux, but the incoming isotropic flux is shielded (via a non-trivial function of the cross section) in a direction-dependent way, which makes possible to break the degeneracy by means of directional information. Here we study this property using a parametric and phenomenological approach, rather than illustrating it with a few discrete cases from models proposed in the literature.  In particular, we consider the specific case of an under-water Mediterranean telescope, although most of our results also apply to an under-ice detector such as IceCube. Since the three different proposed sites for the under-water km$^3$ telescopes have shown event rate differences of the order of 20\\% \\cite{Cuoco}, for the sake of brevity we report the results of our analysis for the \\verb\"NEMO\" site only, which shows intermediate performances.  This paper is structured as follows. In Sec. \\ref{parameteriz} we introduce the parameterization used for the flux and cross section. In Sec.~\\ref{formalism} we describe the formalism used, while our results are reported in Sec. \\ref{results}. Finally, we conclude in Sec. \\ref{conclusions}. ",
        "conclusions": "In this paper we have performed an analysis of the capability of a km$^3$ NT to disentangle the high energy diffuse neutrino flux and the  neutrino-nucleon cross section in the PeV energy region for the incoming neutrinos. As well known, a separation is possible by exploiting the energy--zenith angular event distribution. Differently from previous treatments, we have used a simple phenomenological parametrization for fluxes and cross sections, presenting a sensitivity analysis in a given parameter space, rather than for a few benchmark cases. Our forecast is thus independent from strong assumptions on the normalization or energy dependence of the flux. The results we obtain are summarized in Fig.s~\\ref {fig2}, \\ref{contour1cc}, \\ref{contour2cc}, \\ref{contour1nc}, \\ref{contour2nc}, and refer both to the cases where CC and NC change proportionally to each other, or the case where only NC events are affected by new physics.  One important conclusion is that the capability to tag downgoing neutrino showers in the PeV energy range against the cosmic ray induced background of penetrating muons appears to be a crucial requirement to derive meaningful constraints on the cross section. In this case energy {\\it and} zenith angle information on diffuse neutrino flux event rate in a NT (upgoing vs. downgoing) would greatly improve the diagnostic power for disentangling cross section from flux.  Of course, our analysis presents several approximations. In particular, while we took into account the geometry of the site (for illustration we considered the \\verb\"NEMO\" site in the Mediterranean), we did not consider specific information on the detector geometry, efficiency, etc. Also, we assumed that in muon and tau tracks the energy loss is deterministic,  rather than stochastic. A more accurate account of the telescope and of the energy losses is beyond the scope of this work. Yet, we expect that it would confirm our qualitative findings, while possibly worsening the accuracy with  which parameters can be realistically reconstructed. Since our forecasts are likely more optimistic than a realistic treatment would show, our conclusion on the importance of downgoing event tagging is reinforced. In general, the background due to single energetic muons generated from cosmic rays and surviving  deep underground drops faster with energy than the expected neutrino signal, and well above the PeV region the contamination should be negligible. However, detailed simulation are required to settle down this issue quantitatively;  an active veto may help pushing down in energy the region where the background can be disentangled from the signal on an event-by-event basis.  For the IceCube telescope, in the PeV energy range downgoing muons from cosmic ray showers should be vetoed thanks to the IceTop array. For the Mediterranean km$^3$ neutrino telescopes no final design is available, yet. Given the importance of the present topic, we believe this is an important issue to be addressed in the design phase of these telescopes.  Even in the case where no downgoing tagging for neutrinos will be available for a Mediterranean NT, still its upgoing events at the PeV would provide important constraints on the flux and cross section if a combined analysis of its data with IceCube ones is performed. At very least, it should allow to {\\it test} the hypothesis that the diffuse neutrino fluxes at the PeV from the two hemispheres are consistent with each other, thus validating the hypothesis of an overall isotropic flux."
    },
    "0711/0711.2966_arXiv.txt": {
        "abstract": "{Measurements of standard candles and measurements of standard rulers  give an inconsistent picture of the history of the universe.  This discrepancy can be explained if photon number is not conserved as computations of the luminosity distance must be modified.  I show that photon number is not conserved when photons mix with chameleons in the presence of a magnetic field. The strong magnetic fields in a supernova mean that the probability of a photon converting into a chameleon in the interior of the supernova is high, this results in a large flux of chameleons at  the surface of the supernova.  Chameleons and photons also mix as a result of the intergalactic magnetic field.   These two effects combined cause the image of the supernova to be brightened resulting in a model which fits both observations of standard candles and observations of standard rulers.}  \\bibliographystyle{unsrt} ",
        "introduction": "Observations of standard candles  and of standard rulers allow us to learn how the universe has evolved in recent times.  Both sets of observations favour a universe containing a large component of dark energy but the cosmologies preferred by the two sets of data are not consistent; in fact at higher redshifts they appear to diverge.  This indicates that there is still something missing in our understanding of the accelerated expansion of the universe.  In this paper I will show how the discrepancy between the two sets of observations can be explained if photon number is not conserved because photons are allowed to mix with chameleon particles.  In metric theories of gravity where photons travel on unique null geodesics there is a well-defined notion of distance as long as photon number is conserved. The luminosity distance, $d_L(z)$, to an object at redshift $z$  and the angular diameter distance, $d_A(z)$, to the same object are related by the reciprocity relation \\cite{Etherington} \\begin{equation} \\label{duality} d_L(z)=d_A(z)(1+z)^2 \\end{equation} If photon number is not conserved then this relation no longer holds and a discrepancy in the distance measures should be observed.  There are a number of current proposals which do not conserve photon number; the light from distant objects could be scattered by dust or free electrons in the intergalactic medium, photons may decay,  or the photon could mix with another light state such as the axion.  It had been thought that loss of photons in the intergalactic medium could explain the observed dimming of type 1a supernova without the need for dark energy  \\cite{Csaki:2001yk,Song:2005af,Kaloper:2007gq}. Photon number non-conservation changes the luminosity distance to an object but does not affect distance measurements from standard rulers.  Therefore, if loss of photons is the explanation for the dimming of supernova, angular diameter distance measurements should be consistent with an $\\Omega_{\\Lambda}=0$ universe. This is not the case; observations of standard rulers imply a universe with a significant dark energy component.  \\begin{figure} \\begin{center} \\includegraphics[width=14cm,height=8cm]{notchameleon.eps} \\caption{Distance modulus [$m-M=5\\log(d_L)-25=5\\log(d_A(1+z)^2)-25$] compared to a $\\Lambda$CDM universe plotted against redshift, showing the best fits to observations of standard candles and standard rulers.} \\label{fig:DeltaM} \\end{center} \\end{figure}  When luminosity distances and angular diameter distances  are compared there is a noticeable disagreement at $z \\gtrsim0.5$ \\cite{Bassett:2003vu,Bassett:2004}.  This is shown in Figure \\ref{fig:DeltaM} where distance modulus (compared to a $\\Lambda$CDM universe) is plotted against redshift for measurements coming from standard rulers and standard candles.  The standard candles used are type 1a supernovae \\cite{Riess:2004nr,Tonry:2003zg,Barris:2003dq,Knop:2003iy} and the standard rulers used are FRIIb radio galaxies \\cite{Daly:2003iy,Daly:2002kn}, compact radio sources \\cite{Gurvitis,Jackson:2006ib,Jackson:2003jw} and X-ray clusters \\cite{Allen:2002eu}. Considering the two sets of data separately the best fit to the angular diameter distance data is a universe with $\\Omega_m=0.22$, $\\Omega_{\\Lambda}=0.79$ but the best fit to the luminosity distance data is a very closed universe with $\\Omega_m=0.46$, $\\Omega_{\\Lambda}=0.98$, such a universe is already ruled out by measurements of the CMB.  If the explanation for this apparent violation of the reciprocity relation is that photon number is not conserved, then computed luminosity distances need to be modified.  Photon number non-conservation does not affect the angular diameter distance to an object so the cosmology which best fits observations of standard rulers is the correct one.  If this is the case Figure \\ref{fig:DeltaM} implies that non-conservation of photon number actually brightens the image of a supernova.  This is in contrast to all previously studied mechanisms for non-conservation of photon number which all predict that photons should be lost in the intergalactic medium, and thus that the supernova is dimmed \\cite{Bassett:2003vu,Bassett:2004}. In this paper I show that the images of supernova can be brightened if photons are allowed to mix with chameleons.  The chameleon is a scalar particle which arises in certain models of scalar-tensor gravity \\cite{Khoury:2003rn,Khoury:2003aq}.  In the Einstein frame the action is \\begin{eqnarray} S&=&\\int d^4x\\sqrt{-g}\\left(\\frac{M_P^2R}{2}-\\frac{1}{2}(\\partial \\phi)^2-V(\\phi)\\right)\\nonumber\\\\ & &-\\int d^4x \\mathcal{L}_m(\\psi_m^{(i)},g_{\\mu, \\nu}^{(i)}) \\end{eqnarray} where $\\psi_m^{(i)}$ are matter fields which couple to the metric $g^{\\mu \\nu}_{(i)}=e^{-2\\phi/M_i}g^{\\mu \\nu}$.  Assuming that the chameleon couples to all forms of matter in the same way so that $M_i=M$ $\\forall i$ the current bound on $M$ is $10^6\\mbox{ GeV}<M$ \\cite{Mota:2006fz}. I assume a runaway potential of the form \\begin{equation} \\label{pot} V(\\phi)=\\Lambda^4e^{\\Lambda^n/\\phi^n} \\end{equation} but the chameleon feels an effective potential \\begin{equation} V_{eff}(\\phi)=V(\\phi)+\\rho e^{\\phi/M} \\end{equation} which depends on the local matter density $\\rho$.  The chameleon sits in the minimum of the effective potential with a mass which depends on the local matter density.  In high density regions such as on earth the chameleon becomes very massive and so evades current experimental bounds, but in low density regions it can become almost massless.  The way in which the scalar field can change its mass is known as  the `chameleon mechanism' and this allows the scalar field to evade all current searches for fifth-force effects or violations of the equivalence principle if  $\\Lambda \\sim 10^{-3}\\mbox{ eV}$. In addition it has been shown that the chameleon is compatible with recent searches for axion like particles at PVLAS and CAST \\cite{Brax:2007hi,Brax:2007ak}.  As the universe evolves the local energy density falls and the chameleon changes its properties as the universe cools.  In particular at early times the chameleon behaves like dust but as the universe cools below  temperatures $T \\approx 10\\mbox{ MeV}$ the chameleon  begins to behave like a fluid with equation of state $\\omega=-1$ \\cite{Brax:2004qh,Brax:2004px}. In this way the chameleon is a natural explanation for dark energy.  The paper is organised as follows.  Section \\ref{homogeneous} shows how photons mix with chameleon particles and gives the probability of a photon converting into a chameleon in a homogeneous magnetic field.  Section \\ref{varying} discusses the probability of conversion as the particles pass through a magnetic field made up of many randomly oriented domains and section \\ref{supernova} explains how photons can be converted into chameleons in the interior of a supernova, resulting in a flux of chameleons at the surface.  Then Section \\ref{sec:brightening} describes how this can account for the brightening of the image of the supernova. In Section \\ref{CMB} I show that chameleon-photon mixing is in agreement with observations of the CMB.   I conclude in section \\ref{conc}. ",
        "conclusions": "\\label{conc} The chameleon model provides an explanation for the observed accelerated expansion of the universe.  Observations of standard candles and standard rulers do not give a consistent picture of this acceleration which implies that the reciprocity relation (\\ref{duality}) does not hold.  This can be explained if the chameleon couples to photons in the presence of a magnetic field.  The strong magnetic field inside a supernova means that chameleons and photons mix in the interior of the supernova and so there is a flux of chameleon particles at the surface of the supernova.  This is in contrast to axion-photon coupling which does not allow for a flux of axions at the surface of the supernova. Photons and chameleons also mix as they travel to earth through the intergalactic magnetic field. This requires $M\\lesssim10^{10}\\mbox{ GeV}$ to ensure the effect of the mixing is achromatic for optical photons.  A chameleon model with a coupling of this strength should be detectable in future experiments looking for chameleonic afterglow \\cite{Gies:2007su,Ahlers:2007st} or Casimir forces.  The overall effect of mixing between photons and chameleons is that observers on earth see a brightened image of the supernova. The brightening of supernovae eases the tension between observations of standard candles and standard rulers.  Future observations of high redshift supernova will significantly improve the constraints on this model."
    },
    "0711/0711.4895_arXiv.txt": {
        "abstract": "We consider direct dark matter detection rates and investigate the difference between a standard Maxwell-Boltzmann velocity distribution and a \"realistic\" distribution like the ones extracted from numerical $N$-body simulations. Sizable differences are observed when such results are compared to the standard Maxwell-Boltzmann distribution. For a light target both the total rate and the annual modulation are reduced by $\\sim25$\\%. For a heavy target the total rate is virtually unchanged, whereas the annual modulation is modified by up to 50\\%, depending on the WIMP mass and detector energy threshold. We also consider the effect of a possible velocity anisotropy, and the effect is found to be largest for a light target For the realistic velocity distribution the anisotropy may reduce the annual modulation, in contrast to the Maxwell-Boltzmann case. ",
        "introduction": "The universe is observed to contain large amounts of dark matter \\cite{spergel,WMAP06}, and its contribution to the total energy density is estimated to be $\\sim 25\\%$. This non-baryonic dark matter component, responsible for the growth of cosmological perturbations through gravitational instability, has still not been detected directly. Even though there exists firm indirect evidence from the halos of dark matter in galaxies and clusters of galaxies it is essential to detect matter directly.   The possibility of direct detection, however, depends on the nature of the dark matter constituents. Supersymmetry naturally provides candidates for these constituents \\cite{goodwit,KVprd,ellrosz,ref2}. In the most favored scenario of supersymmetry, the lightest supersymmetric particle (LSP) can be described as a Majorana fermion, a linear combination of the neutral components of the gauginos and higgsinos.  Since the LSPs (or WIMPs) are expected to be very massive ($m_{\\text{WIMP}} \\gta 30$ GeV) and extremely non-relativistic with average kinetic energy $ \\langle T \\rangle \\simeq$ 50 keV $\\left(m_{\\text{WIMP}}/ 100\\, {\\rm GeV} \\right)$, a WIMP interaction with a nucleus in an underground detector is not likely to produce excitation.  As a result, WIMPs can be directly detected mainly via the recoil of a nucleus ($A,Z$) in elastic scattering. The event rate for such a process can be computed from the following ingredients: \\begin{enumerate} \\item An effective Lagrangian at the elementary particle (quark) level obtained in the framework of the prevailing particle theory. For supersymmetry this is achieved as described in refs. \\cite{ref2,JDV96}, for example. \\item A well defined procedure for transforming the amplitude obtained using the previous effective Lagrangian from the quark to the nucleon level, i.e. a quark model for the nucleon. This step in SUSY models is non-trivial, since the obtained results depend crucially on the content of the nucleon in quarks other than u and d. \\item Knowledge of the relevant nuclear matrix elements \\cite{Ress,DIVA00}, obtained with reliable many-body nuclear wave functions. Fortunately, in the case of the scalar coupling, which is viewed as the most important, the situation is a bit simpler, as only the nuclear form factor is needed. \\item Knowledge of the WIMP density in our vicinity and its velocity distribution. Since the essential input here comes from the rotation curves, dark matter candidates other than the LSP are also characterized by similar parameters. \\end{enumerate}  In the past various velocity distributions have been considered for the dark matter gas in our galaxy. The most popular one is the isothermal Maxwell-Boltzmann (M-B) velocity distribution with $\\langle\\upsilon ^2\\rangle=3v^2_d\\simeq 3\\upsilon_0^2/2$, where $v^2_d =\\langle v^2_x \\rangle=\\langle v^2_y \\rangle=\\langle v^2_z \\rangle$ and $\\upsilon_0$ is the velocity of the sun around the galaxy, i.e. $\\upsilon_0\\simeq 220$ km/s.  Extensions of the M-B distribution have also been considered, in particular those that are axially symmetric with enhanced dispersion in the galactocentric direction \\cite{Druk,Verg00,evans00,fornen}. The need for such a velocity anisotropy is numerically well founded \\cite{HANSEN05}.  It has also been shown that the velocity anisotropy is, in principle, measurable in a direction sensitive experiment \\cite{Host}, and it has possibly been observed to be non-zero in clusters of galaxies~\\cite{rocco}.  Non-isothermal models have also been considered. Among these one should mention dark matter orbiting the Sun \\cite{Krauss}, or dark matter which is part of the Sagittarius tidal stream \\cite{GREEN02}. The velocity distribution has also been obtained in ``adiabatic'' models employing the Eddington method \\cite{EDDIN,UK01,VEROW06,Belli:2002yt}. In such an approach, given the density of matter, one can obtain a distribution that depends both on the velocity and the gravitational potential. Evaluating this distribution in a given point in space, e.g. in our vicinity, yields the velocity distribution at that point in a self-consistent manner. Unfortunately this approach is applicable only if the density of matter is spherically symmetric and the distribution depends only on energy.  Also variants of the M-B distribution resulting from a coupling of dark matter to dark energy \\cite{TETRVER06} have been considered.  In the present work we will consider a Tsallis type velocity distribution \\cite{TSALLIS88}. This has been found to be a good description of the velocity distributions in numerical simulations with realistic dark matter density and anisotropy profiles \\cite{HANSEN06}. We compare the direct dark matter detection rates obtained in this way with the results for an axially symmetric Maxwell-Boltzmann velocity distribution. ",
        "conclusions": "The event rate for direct WIMP detection depends on the nucleon cross section (which in turn depends on the assumed particle model), the WIMP mass, the structure of the target, and the WIMP velocity distribution. In the present study we focused on the last aspect by considering a realistic velocity distribution similar to the ones extracted from numerical $N$-body simulations, here parametrized by the Tsallis shape. Generally speaking the calculations using the M-B distribution and the present distribution yield similar results. This is particularly true of the time independent part of the event rate. Some differences show up, however, in the prediction of the time-dependent event rate, described by the modulation amplitude $h$ (the difference between the maximum and the minimum event rates being $2|h|$).  For intermediate targets at zero threshold the modulation amplitude obtained in the present calculation is small for light WIMPs, $|h|<2 \\% $, but it can reach values of $|h|\\approx5 \\%$ for heavy WIMPs. For light WIMPs in the presence of an energy threshold of a few keV the obtained values of $|h|$ may become larger, but this occurs at the expense of the number of counts. The predicted effect on the event rates of the asymmetry parameter of the velocity distribution is relatively small.  It will be interesting to explore the possible consequences of the realistic velocity distribution considered in this work on the event rates of experiments which measure the direction of the recoiling nucleus \\cite{DRIFT,VERFAE06,spoo07}."
    },
    "0711/0711.1749_arXiv.txt": {
        "abstract": "{Stars form from dense molecular cores, and the mass function of these cores (the CMF) is often found to be similar to the form of the stellar initial mass function (IMF).  This suggests that the form of the IMF is the result of the form of the CMF.  However, most stars are thought to form in binary and multiple systems, therefore the  relationship between the IMF and the CMF cannot be trivial.  We test two star formation scenarios - one in which all stars form as binary or triple systems, and one in which low-mass stars form in a predominantly single mode.  We show that from a log-normal CMF, similar to those observed, and expected on theoretical grounds, the model in which all stars form as multiples gives a better fit to the IMF. ",
        "introduction": "The origin of the stellar initial mass function (IMF) is one of the outstanding unsolved problems in astrophysics.  As stars form in dense molecular cores (see e.g. Ward-Thompson et al. 1994; Kirk et al. 2005; Ward-Thompson et al. 2007), it might well be expected that the IMF is related to the mass function of those cores (the CMF).  This idea is supported by observations of prestellar cores, which show that their mass functions are often similar to the IMF of  Galactic field stars (Motte et al. 1998; Testi \\& Sargent 1998; Johnstone et al. 2000; Johnstone et al. 2001; Motte et al.  2001; Johnstone \\& Bally 2006; Alves et al. 2007; Young et al. 2006;  Nutter \\& Ward-Thompson 2007; Simpson et al. 2007).  Further support is given by the observation that Taurus may have both an unusual CMF (Onishi et al. 2002) and an unusual IMF (Luhman 2004; see also Goodwin et al. 2004c), although Kroupa et al. (2003) show that the IMF in Taurus may be compatible with the field IMF.  However, the relationship between the CMF and the IMF cannot be simple, as many, if not the vast majority, of stars form in binaries or higher-order multiple systems (see Goodwin \\& Kroupa 2005; Duch\\^ene et al. 2007 and Goodwin et al. 2007 and references therein; see also Clark et al. 2007).  Observations suggest that the binary frequency amongst young stars is higher than in the field (see Goodwin et al. 2007 and references therein) implying that binaries are destroyed by dynamical interactions in clusters (see Kroupa 1995a,b).  However, Lada (2006) has argued that most M-dwarfs  form as single stars, since the {\\em field} M-dwarf binary fraction is  relatively low and there is no need to invoke dynamical destruction of low-mass binaries to form these (single) stars.  The opposing view is argued by Goodwin \\& Kroupa (2005) and  Goodwin \\& Whitworth (2007).  If stars (or at least relatively high-mass stars) usually form in small-$N$ multiples then there cannot be a trivial one-to-one relationship between the IMF to the CMF.  Firstly, the mass of a  core is distributed between a number of stars. Secondly, some  stars are expected to be ejected at an early age from small-$N$ multiples (e.g. Reipurth \\& Clarke 2001; Goodwin et al. 2007 and references therein; see also Section~3).  Thirdly, many binary systems are expected to be destroyed in  clusters (Kroupa 1995a,b; Kroupa et al. 2003; Goodwin \\&  Whitworth 2007; also see Goodwin et al. 2007 and references  therein).  Thus the CMF should relate most closely to the initial {\\em system} mass function which, in turn, is modified by dynamical effects to produce a mixture of single and multiple systems.  In this paper we examine the relationship between the IMF and the CMF, in particular we use the new results for the CMF in Orion from Nutter \\& Ward-Thompson (2007).  In Section 2 we review observations of the CMF, in Section 3 we present our general method, and in Section 4 we compare the IMFs we produce with the observations. ",
        "conclusions": "We have examined the relationship between the core mass function (CMF) and the stellar initial mass function (IMF).  We use the Orion CMF from Nutter \\& Ward-Thompson (2007) as a `standard' which we fit using a log-normal distribution.  We note that this CMF is not dissimilar to the stellar (Kroupa 2002) IMF shifted upwards in mass by a factor of $8$ (see also Alves 2007).  We randomly sample cores from the CMF and assumed that each core produces a certain number of stars with a random distribution of masses between the components.  The canonical IMF is reproduced very well by a scenario in which every low-mass cores fragment into binaries, and high-mass cores fragment into a multiple system with a ratio of binaries-to-triples of $3\\!:\\!1$ (see e.g. Goodwin \\& Kroupa 2005) and a star formation efficiency (SFE) of $\\sim\\! 30$\\%.  Dynamical disruption (Kroupa 1995a,b; Goodwin \\& Whitworth 2007) of systems then evolves the initial binary fraction of unity into the field population.  We find that a scenario in which low-mass stars preferentially form single systems (e.g. Lada 2006) cannot reproduce the observed IMF from a log-normal CMF.  Firstly, the slope of the high-mass IMF is too steep.  Secondly, and most seriously, this model cannot reproduce the correct numbers of brown dwarfs to high-mass stars.  The best-fit to the canonical IMF is found when the SFE is only $\\sim\\! 15$\\%.  Such a low SFE is required, as the only way in which brown dwarfs may be produced in significant numbers is through the formation of a single brown dwarf from a core.  Higher SFEs are required to produce sufficient high-mass stars, however such SFEs significantly under-produce brown dwarfs and low-mass stars.  A lingering question is the value of the star formation efficiency that must be applied to fit the IMF.  The best-fit value of $\\epsilon$ in the fully multiple model suggest that only $\\sim\\! 30$\\% of the mass in a core ends-up in the stars which that core forms (a similar value for the SFE is found by Alves et al. 2007).  This seems a very low value and may suggest that the determinations of the absolute core masses are wrong.  Another possibility is that feedback from jets is far more efficient than previously thought and manages to disperse most of the gas initially in the core.  A final possibility is that we are not observing `typical' cores which produce the IMF and that the observed CMFs will produce somewhat top-heavy IMFs (cf. Taurus, Goodwin et al. 2004c).  We conclude that a model in which {\\em all} stars and brown dwarfs form in multiple systems from a log-normal core mass distribution provides a very good fit the observed IMF."
    },
    "0711/0711.3266_arXiv.txt": {
        "abstract": "High resolution imaging observation of a sunspot umbra was done with Hinode Solar Optical Telescope (SOT). Filtergrams in wavelengths of blue and green continuum were taken during three consecutive days. The umbra consisted of a dark core region, several diffuse components and numerous umbral dots. We derived basic properties of umbral dots (UDs), especially their temperatures, lifetimes, proper motions, spatial distribution and morphological evolution. Brightness of UDs  is confirmed to depend on the brightness of their surrounding  background. Several UDs show fission and fusion. Thanks to the stable condition of space observation, we could first follow the temporal behavior of these events. The derived properties of internal structure of the umbra are discussed in viewpoint of magnetoconvection in a strong magnetic field. ",
        "introduction": "Umbral fine structure in sunspots has been studied by many authors. Recent reviews were given in Thomas and Weiss (2004) and in references cited therein. The study of umbral fine features is very essential for our understanding of the magnetoconvection in a strong magnetic field atmosphere of celestial bodies. As the spatial size of umbral fine structure, such as umbral dots (UDs), was very fine, it was very hard to get the basic characteristics of them. Especially it was very difficult to follow the temporal evolution of the fine features from the ground-based telescopes, due to the influences of variable atmospheric seeing conditions.  Solar Optical Telescope (SOT) on board Hinode successfully launched on September 23, 2006 was designed to observe the solar fine structure with a 50 cm mirror from space (Kosugi et al. 2007; Tsuneta et al. 2007; Suematsu et al. 2007; Ichimoto et al. 2007; Shimizu et al. 2007). The resolving power in flight condition was confirmed to have nearly the theoretical one of 0.2 arcsec. With the Hinode/SOT, we observed the temporal evolution of umbral fine structures during the period of March 2-4, 2007. Spatial distribution of umbral structure, its temporal evolution, lifetimes, proper motions, temperatures were studied during the three days period. Besides the basic characteristics stated above, we could follow temporal evolution of fission and fusion events of umbral dots. In the following sections, we describe the details of observation and analysis procedures in Section 2, give our results in section 3, and finally discuss and summarize our results in section 4. ",
        "conclusions": "Weiss et al. (2002), from their simulation of three-dimensional non-linear magnetoconvection in a strongly stratified compressible layer, gave two important suggestions on the internal structure of sunspots. First one was that UDs are manifestation of small scale convective cells, whose sizes are reduced strongly compared to normal granules by the strong magnetic field. Second one was that diffuse background components, which are larger in size than UDs, correspond to clumps of vigorously convecting plumes, from which magnetic flux is expelled. According to their suggestion, umbral area is separated into (a) regions of strong fields and small-scale convection and (b) regions of weak fields and large-scale vigorous convection.  Spectroscopic observations on the magnetohydrodynamic (MHD) nature of UDs have been done by several researchers. Wiehr and Degenhardt(1993) and  Socas-Navarro et al. (2004) got the results that the magnetic field strength is weaker and that small or virtually null upflows exist in UDs, while Lites et al (1991) detected no indication of reduced field strength in UDs. It seems that no comprehensive observational view of MHD behavior of UDs is obtained at present. Especially the differences of MHD behavior among the types of UUD/LUD/PUD are unknown.  As was suggested in previous works and confirmed in our present observation, PUDs have different characteristics from UUDs/LUDs in view points of their birth places, and proper motions. As UUDs/LUDs are immobile and hotter than their surroundings, they probably correspond to small-scale magnetoconvection in a strong magnetic field, as was suggested by Weiss et al. (2002). On the other hand, PUDs show systematic proper motions from penumbrae to umbrae and seem to be natural extension of so called penumbral grains. We think that the interchange instability model at penumbrae by Schlichenmaier et al. (1998) well explains PUDs. However, one important point remains unexplained. UUDs/LUDs and PUDs have a common size of around 200-300km. Why do UUDs/LUDs of magnetic convective origin in umbrae and PUDs of penumbral origin have nearly the same geometrical sizes?  What are physical factors which determine the UD size of 200-300km? Are there a common mechanism which controls the sizes of fine structures both in umbrae and penumbrae?  Diffuse components are observed to be hotter than the dark core of umbrae. So they probably will correspond to plumes of vigorous convection with weak magnetic flux, as was suggested by Weiss et al. (2002). As LUDs observed in these diffuse components are observed to have much higher temperature than UUDs, we suspect that convection is stronger in these diffuse components.  However, according to Weiss et al. (2002), small-scale convective plumes, i.e. UDs, are not expected to appear in large-scale vigorous plumes of convection, i.e. diffuse components. This expectation is contrary to our observed result. In our three days observation, it was observed that diffuse components and the dark core kept their identities. Garcia de la Rosa (1987), from the analysis of temporal evolution of many sunspots, suggested that sunspots consist of a cluster of several large fragments and that these fragments keep life-long identities from the birth to the decay of sunspots. In the fragments, UUDs will be formed as was suggested by Weiss et al. (2002). At the interfaces of fragments, convection from deeper layers is expected to intrude the interfaces, resulting in rather bright diffuse components.  In both the vigorous convective plumes model by Weiss et al.(2002) and the fragment model by Garcia de la Rosa (1987), the occurrence of LUDs in diffuse components remains unexplained. It may be related to the question why UDs have common geometrical sizes irrespective of their types. Similar mechanism such as the interchange instabilities proposed by Schlichenmaier et al. (1998) may work to form UDs in addition to magnetoconvetion.  Discussion and suggestions stated above, including the conjecture of the absence of UDs in the dark core (Lites et al., 1991), are to be studied by further observation and analysis. In our next paper, we will report our analysis of Spectropolarimeter data obtained on the same sunspot by Hinode/SOT. Temporal evolution of vector magnetic field and Doppler velocities in and around the UDs will give us more conclusive views of magnetoconvection in sunspots.  The authors acknowledge Drs. M. Sobotka and T. Yokoyama whose comments are valuable for the improvement of the paper. The authors are partially supported by a Grant-in-Aid for the 21st Century COE 'Center for Diversity and Universality in Physics' from the Ministry of Education, Culture, Sports, Science and Technology (MEXT) of Japan, and also partially supported by the Grant-in-Aid for 'Creative Scientific Research The Basic Study of Space Weather Prediction' (17GS0208, Head Investigator: K. Shibata) from the Ministry of Education, Science, Sports, Technology, and Culture of Japan. Hinode is a Japanese mission developed and launched by ISAS/JAXA, with NAOJ as domestic partner and NASA and STFC (UK) as international partners. It is operated by these agencies in co-operation with ESA and NSC (Norway)."
    },
    "0711/0711.0053_arXiv.txt": {
        "abstract": "A subset of neutrino-induced upward through-going muons in the Super-Kamiokande detector consists of high energy muons which lose energy through radiative processes such as  bremsstrahlung, $e^{+} e^{-}$ pair production and photonuclear interactions.  These ``upward showering muons'' comprise an event sample whose mean parent neutrino energy is approximately 1 TeV. We show that the zenith angle distribution of upward showering muons is consistent with negligible distortion due to neutrino oscillations, as expected of such a high-energy neutrino sample. We present astronomical searches using these high energy events, such as those from WIMP annihilations in the Sun, Earth and Galactic Center, some suspected point sources, as well as searches for diffuse flux from the interstellar medium. ",
        "introduction": "In order to select neutrino events with the highest energies, we consider muon neutrino interactions in the rock around the detector, because the effective target volume is very much increased~\\cite{Grillo}.  To separate neutrino-induced muons from cosmic ray muons, we select only upward-going muons, since the background from downward going cosmic ray muons overwhelms any neutrino-induced muons from above. Muons penetrating the detector have energies of at least several GeV and point along the neutrino direction within a few degrees, allowing astrophysical studies.  Neutrinos originating from cosmic point sources are expected to have harder energy spectra than the background of atmospheric neutrinos.  Some of these high energy neutrino-induced muons undergo radiative energy loss. We identify these muons  as ``showering muons''. High energy muons are correlated with high energy neutrinos, which allow us to statistically select an extremely high energy parent neutrino sample from the Super-Kamiokande data. Our aim is to extract this sample for physics and astronomy studies.  The Super-Kamiokande (Super-K) experiment~\\cite{sknim} has previously analyzed two topologically distinct categories of upward muons caused by neutrino interactions in the rock below the detector: muons which exit the detector (called ``through-going'') and those which stop inside the detector (called ``stopping'')~\\cite{skstop}. The parent neutrino energy of upward stopping and through-going muons for atmospheric neutrinos is peaked at $\\simeq$ 10 GeV and $\\simeq$ 100 GeV respectively.  In 1646 days of data, Super-Kamiokande detected 1856 upward through-going and 458 upward stopping muons.  Oscillation results using only upward through-going muons and stopping muons have been reported in Refs.~\\cite{skstop} and ~\\cite{upthru}. Both upward muon samples have been combined along with the contained events to do the most precise atmospheric neutrino oscillation studies~\\cite{fullpaper,3flavor}. Moreover, astrophysical searches for annihilation signatures of WIMPs~\\cite{skwimp}, point and diffuse sources~\\cite{sknuastro}, and neutrinos from GRBs~\\cite{skgrb} have been done with upward muons.  In this paper we bifurcate the upward through-going muon sample into upward showering and upward non-showering events.  This classification is done on the basis of the physical mechanism for muon energy loss above 1 TeV. The showering sample is then analyzed to see what these highest-energy events might reveal. ",
        "conclusions": "From the sample of upward through-going muons in 1646 days of data we have isolated a sample of 318 showering muon events which lose energy through radiative processes. The mean parent neutrino energy of these events is $\\simeq$ 1 TeV  which make this subset the highest energy neutrinos seen in Super-K. The fraction of these events is approximately one fifth of upward through-going muon sample.  At these high energies, the neutrino oscillation probability (for oscillation parameters obtained from Ref.~\\cite{fullpaper}) is negligible. We have verified that the zenith angle distribution of upward showering muons is consistent with null oscillations. This shows that the highest neutrino energy dataset (with mean parent neutrino energy of $\\simeq$ 1 TeV, path-lengths of order 10000 km, and $\\Delta m^2 \\sim 0.0025$~eV$^2$) is consistent with null oscillation. We also performed various  searches for extra-terrestrial neutrino sources with this subset, such as from WIMP annihilations, suspected known and unknown point sources, as well as diffuse sources. We do not see any evidence for astrophysical neutrinos in the upward showering muon dataset."
    },
    "0711/0711.0862_arXiv.txt": {
        "abstract": "{The Deep Extragalactic VLBI-Optical Survey (DEVOS) aims at constructing a large sample of compact radio sources up to two orders of magnitude fainter than those studied in other Very Long Baseline Interferometry (VLBI) surveys. Optical identification of the objects is ensured by selecting them from the Sloan Digital Sky Survey (SDSS) list.} {While continuing to build up the DEVOS data base, we investigated how the VLBI detection rate could be enhanced by refining the initial selection criteria introduced in the first paper of this series.} {We observed 26 sources in two adjacent, slightly overlapping $2\\degr$-radius fields with the European VLBI Network (EVN) at 5~GHz frequency on 2 March 2007. The phase-reference calibrator quasars were J1616+3621 and J1623+3909. The objects selected were unresolved both in the Faint Images of the Radio Sky at Twenty-centimeters (FIRST) survey catalogue and the SDSS Data Release 4.} {We present images of milli-arcsecond (mas) scale radio structures and accurate coordinates of 24 extragalactic sources. Most of them have never been imaged with VLBI. Twenty-two compact radio sources (85\\% of our initial sample) are considered as VLBI detections of the corresponding optical quasars in SDSS. We found an efficient way to identify quasars as potential VLBI targets with mas-scale compact radio stucture at $>1$~mJy level, based only on the FIRST and SDSS catalogue data by applying simple selection criteria.} {} ",
        "introduction": "Distant extragalactic radio sources detected with Very Long Baseline Interferometry (VLBI) are typically active galactic nuclei (AGNs): the rest-frame brightness temperature of a 1-mJy source radiating from a region subtending a few milli-arcseconds (mas) at 5~GHz is at least $10^6 - 10^7$~K, often much higher. The aim of the Deep Extragalactic VLBI-Optical Survey (DEVOS; Mosoni et al. \\cite{moso06}, Paper~I hereafter) is to build up a large sample of optically identified AGNs imaged with VLBI at mas angular resolution. Many of the potential astrophysical and cosmological utilisations of such a sample require the optical identification and spectroscopic (or photometric) redshift of the sources. These data can primarily be obtained from large optical sky surveys. In particular, the latest release\\footnote{Data Release 6 (DR6), \\tt{http://www.sdss.org/dr6}} of the Sloan Digital Sky Survey (SDSS; Adelman-McCarthy et al. \\cite{adel07}) contains imaging data for over 20\\,000, and spectroscopic data for more than 3000 unresolved radio sources from the Faint Images of the Radio Sky at Twenty-centimeters (FIRST) survey catalogue (White et al. \\cite{whit97})\\footnote{\\tt http://sundog.stsci.edu} with integrated flux density of at least 20~mJy. These sources could be targeted with phase-referenced radio interferometric observations by selecting nearby bright and compact calibrator sources from e.g. the VLBA Calibrator Survey (VCS)\\footnote{\\tt http://www.vlba.nrao.edu/astro/calib}. To indicate the potential of a FIRST- and SDSS-based VLBI imaging survey, we note that optical counterparts of at least 30\\% of all FIRST sources are expected to be found in SDSS (Ivezi\\'c et al. \\cite{ivez02}). A recent sensitive single-baseline Arecibo--Effelsberg VLBI survey of $\\sim1000$ FIRST sources indicated that about one third of {\\em all} sources had mJy-level compact stucture, irrespective of the total flux density (Porcas et al. \\cite{porc04}).  The technical feasibility of DEVOS has recently been demonstrated in a pilot experiment (Paper~I). However, at present reaching the ambitious goal of increasing the total number of VLBI-imaged extragalactic sources to $\\sim10^4$ is very demanding in terms of observing resources. This led us to somewhat modify the sample definition criteria used in Paper I in order to increase the throughput of the survey. We studied two adjacent celestial fields, each of $2\\degr$ radius, to search for mas-scale structures of radio-emitting quasars detected with the SDSS. Here we describe the target selection (Sect.~\\ref{sect2}), the phase-referenced VLBI observations (Sect.~\\ref{sect3}) and the data reduction procedure (Sect.~\\ref{sect4}). The results are presented in Sect.~\\ref{sect5}. The VLBI images of the target sources are available in the online version of this paper. Implications of our results are discussed in Sect.~\\ref{sect6}. ",
        "conclusions": "\\label{sect6}  Phase-referenced 5-GHz VLBI images of 24 radio-loud AGNs at mJy flux density level or higher are presented here. To our knowledge, 20 of these sources have never been imaged with VLBI before. Five of our target sources (J161446.9+374607, J161748.4+380141, J162229.3+400643, J162330.5+355933, J162422.0+392440) are nearly as bright and compact at 5~GHz as our phase-reference calibrator sources. It is interesting to note that two of them (J161446.9+374607 and J162422.0+392440) have so far escaped from being selected for any of the VLBI surveys. This indicates that the radio sky is still not fully explored even for relatively strong ($\\sim80$~mJy) mas-scale compact extragalactic sources.  We detected 22 out of 26 target SDSS quasars, 85\\% of our initial sample. (Two other VLBI detections are probably chance positional coincidences with SDSS quasars.) Combining this with the DEVOS pilot project results (Paper~I) and counting also the phase-reference calibrator sources, the VLBI detection rate is nearly 90\\% for the radio quasars that are known to be compact in both optical (SDSS) and in the FIRST survey, and have integrated flux density $S_{\\rm{1.4}}>20$~mJy in FIRST. Based on currently available optical and low-resolution radio sky surveys, we found a highly efficient way to pre-select extragalactic sources that are detectable with VLBI at 5~GHz. On the other hand, for any given object found in the SDSS and FIRST lists, the selection method we used is also an excellent tool for assessing whether it can be potentially successfully imaged with VLBI.  The SDSS DR6 contains nearly 9000 quasars that satisfy our selection criteria and thus can be targeted with VLBI. Extrapolating our detection rate for this entire sample, the number of potential VLBI detections can be as high as 7500. Such a big VLBI-optical sample --~partly with spectroscopic redshifts available~-- would eventually be invaluable for a multitude of astrophysical and cosmological studies.  Astrometric projects like global or local densifications of the ICRF, and linking the radio reference frame with the future optical reference frame would also benefit from the DEVOS data. In practice, the data base is to be built up field by field gradually, concentrating on celestial areas that are otherwise targeted by different projects, e.g. deep spectroscopy, or deep imaging at any waveband. The immediate science return of these observations could be enhanced this way."
    },
    "0711/0711.2056_arXiv.txt": {
        "abstract": "% A spectacular, transient mass-loss event in the hydrogen-deficient central star of the planetary nebula Longmore~4 was discovered in 1992 by Werner et al. During that event, the star temporarily changed from its normal PG~1159 spectrum to that of an emission-line [WCE] star. I have been monitoring the spectrum of Lo~4 since 2003.  Out of 31 spectra, two of them reveal mass-loss outbursts similar to the one seen in 1992, showing that the phenomenon recurs. I speculate on possible mechanisms for these unique  outbursts, but emphasize that we still have no fully convincing explanation. ",
        "introduction": "Longmore~4 (Lo~4) is a low-surface-brightness planetary nebula (PN), discovered by Longmore (1977) in the course of the ESO-SRC southern-hemisphere sky survey. The first spectroscopic observations of its $V=16.6$ central star by M\\'endez et~al.\\ (1985) led to a PG~1159$-$035 spectral classification.  The PG~1159-type hydrogen-deficient planetary-nebula nuclei (PNNi) and white dwarfs have spectra with conspicuous C~$\\scriptstyle\\rm IV$ and He~$\\scriptstyle\\rm II$ absorption, and no Balmer absorption. O~$\\scriptstyle\\rm VI$ 3811-3834~\\AA\\ is often seen in emission. In various spectral-classification schemes, stars of this type have also  been assigned to types ``O~$\\scriptstyle\\rm VI$'' or ``O(C).'' Werner et al.\\ (2007) have recently revised the atmospheric parameters for Lo~4 to $T_{\\rm eff}=170,000$~K, $\\log g=6$.  The spectrum of Lo~4 is quite similar to that of K~1-16, the first known pulsating PNN (Grauer \\& Bond 1984) and a member of the GW~Vir class of non-radial pulsators. The close similarity was borne out by the discovery of pulsations in Lo~4 by Bond \\& Meakes (1990). Lo~4 was, at the time, only the second known pulsating PNN, having a strong periodicity near 31~min (1850~s), along with at least 8 other pulsation modes ranging from 831 to 2325~s. ",
        "conclusions": ""
    },
    "0711/0711.2783_arXiv.txt": {
        "abstract": "The first images of 6.7-GHz methanol masers in the massive star-forming regions DR21(OH) and DR21(OH)N are presented. By measuring the shapes, radial velocities and polarization properties of these masers it is possible to map out the structure, kinematics and magnetic fields in the molecular gas that surrounds newly-formed massive stars. The intrinsic angular resolution of the observations was 43 mas ($\\sim$~100 AU at the distance of DR21), but structures far smaller than this were revealed by employing a non-standard mapping technique. By plotting the positions of the Gaussian-fitted maser emission centroids in each velocity channel, the internal velocity gradients of the masers were investigated at very high spectral and spatial resolution. This technique was used in an attempt to identify the physical structure (e.g. disc, outflow, shock) associated with the methanol masers. Two distinct star-forming centres were identified. In DR21(OH) the masers had a linear morphology, and the individual maser spots each displayed an internal velocity gradient in the same direction as the large-scale structure. They were detected at the same position as the OH 1.7-GHz ground-state masers, close to the centre of an outflow traced by CO and class I methanol masers. The shape and velocity gradients of the masers suggests that they probably delineate a shock. In DR21(OH)N the methanol masers trace an arc with a double-peaked profile and a complex velocity gradient. This velocity gradient closely resembles that of a Keplerian disc. The masers in the arc are 4.5$\\%$ linearly polarized, with a polarization angle that indicates that the magnetic field direction is roughly perpendicular to the large-scale magnetic field in the region (indicated by lower angular resolution measurements of the CO and dust polarization). The origin and nature of these maser structures is considered within the context of what is already known about the region. The suitability of channel-by-channel centroid mapping is discussed as an improved and viable means to maximise the information gained from the data. ",
        "introduction": "The DR21 star-forming complex is a large and highly active region, containing many sites of recent and ongoing star-formation as evidenced by observations of atomic and molecular gas and dust at a wide range of frequencies. It lies in the Cygnus X region of our Galaxy at a kinematical distance estimated to be between 2 and 3 kpc (Dickel \\& Wendker 1978; Campbell et al. 1982; Piepenbrink \\& Wendker 1988; Odenwald \\& Schwartz 1993). DR21 contains many sub-regions denoted by positions of Infra-Red Sources (IRSs), Far Infra-Red sources (FIRs) or Extremely Reddened Objects (EROs), which are interpreted as sites of ongoing star-formation (e.g. Kumar et al. 2007; Jackob et al. 2007). This paper will describe the results of Multi-Element Radio-Linked Interferometer Network (MERLIN) observations of class-II methanol masers at 6.7-GHz in two regions of active massive star-formation within DR21.  The aim of this work is to understand the role of outflows, circumstellar discs and magnetic fields in the formation of massive stars. This can be achieved by studying the 6.7-GHz methanol masers at high spatial and spectral resolution. It is possible to effectively enhance the resolving power of MERLIN by measuring the centroid positions of the maser emission in each spectral channel of the maser line. This, coupled with linear and/or circular polarization information on the masers allows us to form a 3-dimensional $+$ magnetic field picture of the physical region immediately surrounding emerging massive stars.  This work is part of an on-going study of several regions of massive star-formation that produce bright OH and methanol masers. Data from phase-referenced observations at 1.7, 4.7, 6.0 and 6.7-GHz are also available to the authors. This sample covers all the OH and methanol maser lines that can be observed within the working frequency range of MERLIN.  Parts of this survey have already been published (Harvey-Smith \\& Cohen 2005; Nammahachak et al. 2006; Green et al. 2007) but the mapping of several regions is on-going. Most of the observations in the multi-frequency study were correlated in full polarization mode, and it is therefore possible to measure the properties of magnetic fields in these regions, as well as the Faraday rotation (e.g. Vlemmings, Harvey-Smith \\& Cohen 2006). The goal of this survey is to map all the maser lines listed above in full polarizations, thus forming an in-depth picture of several regions of massive star-formation from which to draw specific conclusions about the relationship between masers, magnetic fields and stellar accretion. ",
        "conclusions": "By observing the DR21 star-forming complex at 6.7-GHz, two new sites of methanol maser emission in DR21(OH) and DR21(OH)N have been found. Two dimensional Gaussian components were fit to the emission in each spectral channel containing a signal greater than 5$\\sigma_{RMS}$. The positions of the emission centroids in each channel were then mapped, in an attempt to determine the internal velocity gradients of the maser spots. This technique added a great deal of information about the structures traced by the methanol masers, which is vital to our aim of linking methanol masers to physical features such as discs around protostars and outflows of circumstellar material. The maser feature in DR21(OH)N, component 1, is a strong candidate to be a circumstellar disc. It has a double-peaked spectral profile, a complex velocity gradient that resembles a Keplerian rotation curve and a magnetic field gradient across its major axis. It also coincides with a massive protostar that has been previously identified by studies of infra-red and dust emission. Further study of the existing MERLIN 1.7-GHz and 6.0-GHz OH maser data in full polarizations and high-resolution European VLBI Network data of methanol at 6.7-GHz will hopefully add to our understanding of the disc-like component 1 in particular. Numerical analysis of the feature in terms of a Keplerian rotation model will reveal whether or not this feature is likely to be a disc around a massive protostar."
    },
    "0711/0711.2326_arXiv.txt": {
        "abstract": "MAGIC, as well as HESS and VERITAS, is a \\v{C}erenkov Telescope unveiling $\\gamma$-ray sources above 60 GeV at vertical within noisy (hadronic) airshowering sky. These telescopes while facing the horizons may reveal rarest blazing UHECR as well as far fluorescence tails of downward PeV-EeV hadronic airshowers. Few of these inclined airshowers blazing on axis are spread by the geomagnetic field into twin spots. These twin flashes and their morphology may tag the UHECR origination site. There is a rich window of such reflecting \\v{C}erenkov lights visible by Telescopes on top of Mountains as MAGIC  (and partially VERITAS): the reflections from the nearby ground (possibly enhanced by rain or snow, ice white cover), from the Sea and from the cloudy sky; in particular, these cloudy sheets may lay above or below the observer. MAGIC looking downward to the clouds or the snow, may well reveal blazing Moliere disks diffusing \\v{C}erenkov spots (few events per night). Because of geomagnetic forces and splitting of the inclined air-shower, one should reveal for the first time (at tens PeV or above) \\v{C}erenkov airshowers whose flashes are skimming the MAGIC nearby Sea and opened into twin spots. Their morphology may tag the UHECR origination, its consequent cross-section and  composition. Magic telescopes looking upward into cloudy sky may observe very rare up-going UHE Tau, originated by UHE PeVs neutrinos skimming earth, air-showering into sky, reflecting into clouds. In particular Glashow resonant antineutrinos electron hitting into Earth electrons may lead to gauged boson $W^-$ , whose decay (inside the Earth) may produce a $\\tau+\\bar{\\nu}_{\\tau}$  \\cite{Fargion1999}, which later escape and decay in air is producing \\v{C}erenkov lights; these flashes may blaze into the clouds above MAGIC as upward dot spots. The Magic energy threshold for such UHE Neutrinos showers rises to PeV values.  EeV UHE tau neutrinos by guaranteed  GZK UHECR secondaries \\cite{Greisen:1966jv, za66}, via the muon-tau flavor mixing,  may skim the Earth, produce UHE tau particles whose escape in air and decay in flight may blaze Magic or reflect \\v{C}erenkov light at opposite far  cloudy sky edges. Any collaboration (MAGIC,HESS,VERITAS) will be  willing to dedicate cloudy time to UHECR physics since nothing else can be done in that time. These  reflections may open a totally new UHECR spectroscopy while unveiling a rare, but loud  Tau Neutrino astronomy\\cite{Fargion2007, Fargion1}. ",
        "introduction": "erenkov lights at the edges.} The common downward vertical airshowers, mostly of hadronic nature, are growing as a tree whose wide roots are spread in a conical shape, because at low altitudes vertical shower are randomized into a disk. The resultant secondaries shine \\v{C}erenkov light on a disk's size which is about twice a Moliere radius (a hundred meter size). At high altitudes, inclined shower (zenith angle $\\theta\\geqslant70^{\\circ}$) at lowest air density are split by geomagnetic Lorentz force into two twin charged beams. At their center, energetic  muon pair bundles define  the shower axis. These  fan-like shape are made by twin  pair lobes.  These lobes are shining \\v{C}erenkov lights in different beams and regions; such an exemplar UHECR inclined event has been foreseen in last years,\\cite{Fargion3, Fargion4,Fargion5}, recorded by the AUGER detector \\cite{Auger07} and discussed recently \\cite{Fargion1}: both a \\v{C}erenkov flash (seen at Coihueco) and a fluorescence flare (spot by Los Leones) were detected while a muon pair bundle hit a dozen of tanks in the heart of the AUGER array (see Fig.\\ref{fig2} and Fig.\\ref{fig3} \\cite{Fargion1}). Let us remark that only one side (of the two) of an off-axis blazing \\v{C}erenkov flash (the beam made by bent negative electrons) have been revealed by Coihueco: because of the wide fan-like shower size, \\v{C}erenkov lights are mostly seen in one main lobe. Only rarely \\v{C}erenkov jets are observed in axis being split far away into their twin lobes; the higher energy muon bundles survive, being harder than electron ones, reaching the ground almost undeflected along the main shower axis. Such a downward twin spot discover would offer a new UHECR spectroscopy at the horizon (inclined showers) \\cite{Fargion2}. \\v{C}erenkov earliest twin lights define, indeed, the UHECR original direction, its most probable cross-section and its consequent primary composition. \\v{C}erenkov light's intensity and color (spectra), zenith angle, traversed column depth, angular separation, time structure and muons component delay and multiplicity, defines primary energy and composition. Moreover, such a \\emph{pedagogical}  in-axis event, where the lobes split, even if rare could be enhanced by time correlated muon light rings (by secondary muon \\v{C}erenkov imprint \\cite{Fargion2, Fargion3, Fargion4, Fargion5}). The MAGIC-II telescope array \\cite{ANGELIS}, as well as VERITAS and HESS, should discover such inclined events if looking horizontally \\cite{zas05}. The hope is to prepare at the same time the calibration for rarest but exciting upgoing PeV-EeV neutrino induced airshowers \\cite{Fargion1999, Fargion2004}. In the present article we reflect on their Cherenkov reflections summarized by the last four figures. Let us remind that at highest energies (EeV), Fluorescence and \\v{C}erenkov hybrid events at AUGER (estimated about a dozen per year \\cite{Fargion1}) may offer an empirical calibration (see Fig.\\ref{fig2} and Fig.\\ref{fig3}, \\cite{Fargion1}). Hybrid events by only \\v{C}erenkov twin spots may be extended to lower energies, as tens PeV, at more frequent rate. In analogy, MAGIC may study twin split events: their double spot signature is larger at larger zenith angles though this separation is observed at more and more far edges. Therefore, the angular size between the two lobes first  vanishes at vertical than increases at larger zenith angle to a maximum angle. The \\v{C}erenkov lobes' angular size appears to shrink to smaller aperture at higher and higher zenith angles, because of their larger distances at the horizons. The twin signature should be resolved better by MAGIC telescope than by AUGER because better resolution. AUGER, on the contrary, may enhance and  tag the twin lobes  by adding tanks around the four fluorescence-\\v{C}erenkov telescopes at Coihueco, Los Leones, Los Morados and Norte: their trigger correlated muon signal may, in fact, trigger the reading of \\v{C}erenkov flashes at the edges, revealing  their hidden twin split signature. ",
        "conclusions": "Airshowers \\v{C}erenkov lights may be reflected while skimming the Earth, the Sea, the icy or rainy grounds.The MAGIC Telescopes (as well as the VERITAS ones) being on top of a mountain, may soon reveal these reflections whose role is also to calibrate the downward UHECR rate. In principle, MAGIC telescope observing large zenith angles, may also search for upgoing tau airshower probably related to GRB-SGRs or BL Lac brightening. Also Glashow resonant neutrinos may rise upgoing and blazing airshowers at horizons. Because of the atmosphere dimming of far hadronic airshowers at Earth atmosphere edges, the nearer \\cite{Fargion1999}, younger \\cite{Bertou2002}neutrino induced ones are emerging better,  over the far hadronic noise. In particular the Earth skimming Hortaus (horizontal $\\tau$s) \\cite{Fargion2004} while decaying and showering may, at the same time, blaze at the front and-or  flash on the opposite side to the (eventual) clouds or by fluorescence in air. The location where these reflection might hit is not a point anti-correlated to the source but a wide  arc parallel to the horizon line. All the hadronic inclined air-showers may and must show the geomagnetic splitting whose signature can be a test of UHECR  rate and composition from the knee up to GZK energy edges. Just where AUGER most recent and surprising results on UHECR  call for a confirmation \\cite{Yamamoto2007}. In larger AUGER sky, similar higher energy UHE tau events (possibly originated by EeV GZK \\cite{Greisen:1966jv, za66} tau neutrinos),  events whose spectra and local anisotropy has been recently revealed, must shine in next few years\\cite{Fargion2007}. \\vspace{-7pt}"
    },
    "0711/0711.4927_arXiv.txt": {
        "abstract": "The double pulsar system \\psr{0737}A/B is a highly relativistic double neutron star (DNS) binary, with a $2.4$-hour orbital period.  The low mass of the second-formed NS, as well the low system eccentricity and proper motion, point to a different evolutionary scenario compared to other known DNS systems.  We describe analysis of the pulse profile shape over 6 years of observations, and present the resulting constraints on the system geometry.  We find the recycled pulsar in this system, PSR J0737$-$3039A, to have a low misalignment between its spin and orbital angular momentum axes, with a $68.3\\%$ upper limit of $6.1\\degrees$, assuming emission from both magnetic poles. This tight constraint supports the idea that the supernova that formed the second pulsar was relatively symmetric, possibly involving electron-capture onto an O-Ne-Mg core. ",
        "introduction": "The \\psr{0737}A/B system \\citep{bdp+03,lbk+04} has provided pulsar astronomers and neutron-star theorists with an abundance of astrophysical phenomena to study in more detail, and with more precision, than ever before. Among these is the formation and evolution of binary pulsars, and in particular how this system, and perhaps others like it, have formed and evolved to arrive at their current configurations.  We have performed analysis confirming that perhaps the channel of binary evolution undergone by this system is somewhat different than that of other double neutron star (DNS) systems for which orbital parameters and component masses have been measured.  The remainder of this paper will focus on describing the direct observational constraints on this evolution that we have gained through study of this system. ",
        "conclusions": ""
    },
    "0711/0711.3393_arXiv.txt": {
        "abstract": "We investigate when the energy that pins a superfluid vortex to the lattice of nuclei in the inner crust of neutron stars can be approximated by the energy that binds the vortex to a single nucleus. Indeed, although the pinning energy is the quantity relevant to the theory of pulsar glitches, so far full quantum calculations have been possible only for the binding energy. Physically, the presence of nearby nuclei  can be neglected if the lattice is  dilute, namely with nuclei sufficiently distant from each other. We find that the dilute limit is reached only for quite large Wigner-Seitz cells, with radii $R_{_{\\rm WS}} \\gtrsim 55 $ fm; these are found only in the outermost low-density regions of the inner crust. We conclude that  present quantum calculations do not correspond to the pinning energies in almost the entire inner crust and thus their results are not predictive for the theory of glitches. ",
        "introduction": "Pulsar glitches are sudden spin-ups in the otherwise steadily decreasing frequency of rotating magnetized neutron stars. According to the vortex-model, glitches may represent direct evidence for the existence of a macroscopic superfluid inside such stars \\cite{AI}. This scenario involves the inner crust of the star, namely the density range $\\rho_d \\leq \\rho \\leq 0.6 \\rho_o $, where $\\rho_d = 4 \\times 10^{11}$ g/cm$^{3}$ is the neutron drip density and $\\rho_o = 2.8 \\times 10^{14}$ g/cm$^{3}$ is the nuclear saturation density. Here, the quantized vortices that form in the superfluid of unbound neutrons to carry its angular momentum can attach themselves to the lattice of neutron-rich nuclei co-existing with the neutrons in the inner crust, thus freezing part of the superfluid angular momentum. As a consequence of the star's spin-down, hydrodynamical forces then develop (Magnus force) which tend to detach the vortices from the lattice in order to let the angular momentum free to decrease. If a large number of vortices can be unpinned simultaneously and deliver their angular momentum to the star surface, its corresponding sudden spin-up is observed as a glitch. The process naturally repeats itself so that part of the superfluid angular momentum is released in discrete time steps, thus explaining the approximate periodicity of glitches.  A crucial  microscopic  input of the model is the structure of the nuclear lattice  at the different densities found in the inner crust. Calculations predict a bcc structure for the coulomb lattice, namely each nucleus is at the center of a cubic cell of side $\\ell_p = 2 R_{_{\\rm WS}}$ with nuclei at each vertex. The bcc lattice can also be seen as a superposition of layers of cubic cells; each layer, of thickness $\\ell_p$, is  delimited by two planes of nuclei (those at the vertices of the cells) and contains in its middle a third plane of nuclei  (those at the center of the cells). In Table~1 we list the results from the classical paper by Negele and Vautherin \\cite{NV} for five zones along the inner crust. The parameters which characterize  the lattice structure are the baryon mass density ($\\rho_{_{\\rm B}}$), the number density ($n_{_{\\rm G}}$) and corresponding Fermi momentum ($k_{_{\\rm F,G}}$) of the gas of unbound neutrons, the radius of the Wigner-Seitz (WS) cell ($R_{_{\\rm WS}}$) and the radius of the nucleus ($R_{_{\\rm N}}$) at the center of the cell. \\begin{table} \\tbl{Properties of five zones along the inner crust, as obtained in Ref.~\\refcite{NV}.} {\\begin{tabular}{@{}cccccc@{}} \\toprule Zone & $\\rho_{_{\\rm B}}$ & $n_{_{\\rm G}}$ & $k_{_{\\rm F,G}}$ & $R_{_{\\rm WS}}$ & $R_{_{\\rm N}}$ \\\\ & (g/cm$^{3}$) & (fm$^{-3}$) & (fm$^{-1}$) & (fm) & (fm) \\\\ \\colrule 1 & $1.5\\times 10^{12}$ & $4.8\\times 10^{-4}$ & $0.24$ & $44.0$  & $6.0$  \\\\ 2& $9.6\\times 10^{12}$ & $4.7\\times 10^{-3}$ & $0.52$ & $35.5$  & $6.7$  \\\\ 3& $3.4\\times 10^{13}$ & $1.8\\times 10^{-2}$ & $0.81$ & $27.0$  & $7.3$  \\\\ 4 & $7.8\\times 10^{13}$ & $4.4\\times 10^{-2}$ & $1.09$ & $19.4$  & $6.7$  \\\\ 5 & $1.3\\times 10^{14}$ & $7.4\\times 10^{-2}$ & $1.30$ & $13.8$  & $5.2$  \\\\ \\botrule \\end{tabular}} \\end{table}  The other crucial microscopic input is the vortex-lattice interaction potential as a function of the vortex position in the lattice, from which the pinning force could be derived. Due to the complicated spatial geometry of the general problem, however, a simpler approach has been followed so far which calculates only the \\emph{pinning energy}. This is defined as the difference in energy between two vortex-lattice configurations which are relevant to the pinning mechanism and possess a reasonable degree of symmetry. More precisely, one starts by taking the vortex axis oriented along a symmetry direction of the lattice and then considering only the interaction of the vortex with the nuclei lying in a plane perpendicular to the vortex axis (the middle plane of a layer). This allows to determine the pinning energy 'per site' (i.e. per layer). The pinning energy 'per unit length', which is the one relevant to the vortex-model for glitches, can then be obtained if one knows  the number of pinning sites (layers) per unit length. Statistical estimates of this quantity allow to evaluate pinning also for random vortex-lattice mutual directions.  In Figure~1~(left) the black disks represent the nuclei in the given plane, each positioned at the center of a cube. The nuclei at the vertices of the cubes are shown with white disks; they lie in the planes that delimit the layer, at a distance $R_{_{\\rm WS}}$ above and below the original plane. The vortex core position is schematically represented by the dashed line and the vortex axis, perpendicular to the plane, is shown by a star. \\def\\figsubcap#1{\\par\\noindent\\centering\\footnotesize(#1)} \\begin{figure}[t] \\begin{center} \\parbox{2.4in}{\\epsfig{figure=pizzochero1.eps,width=2.4in}} \\hspace*{6pt} \\parbox{2.4in}{\\epsfig{figure=pizzochero2.eps,width=2.4in}} \\caption{ Left: vortex positions in a bcc lattice.  \\ Right: configurations for the pinning energy. } \\end{center} \\end{figure} The nuclear pinning (NP) configuration corresponds to the vortex pinned on a nucleus, e.g. in position (1).  The interstitial pinning (IP) configuration corresponds to the vortex away from and equidistant between two nearby nuclei in the plane, e.g. in position (2). Notice that in the NP configuration the vortex also pins to the nuclei in the layers above and below, which are aligned along its axis with spacing $2 R_{_{\\rm WS}}$. Conversely, in the IP configuration the vortex does not pin to any nucleus in the lattice. Position (4) is not allowed as an interstitial pinning site; actually, it  is physically equivalent to nuclear pinning, since the vortex will pin to the nuclei in the layers above and below; in other words, positions (1) and (4) are equivalent once we consider the pinning energy per unit length.  The pinning energy is the energy necessary to move the vortex by a distance $ R_{_{\\rm WS}}$, from the NP to the IP configuration. The \\emph{pinning energy per site} is then defined as $E_{\\rm pin} = E_{_{\\rm NP}}-E_{_{\\rm IP}}$, i.e. the difference in energy between the two configurations. Physically, the vortex interacts significantly only with first neighbors; therefore, it is enough to restrict the attention to a system of two adjacent cells, as shown in Figure~1~(right), since contributions from cells further away will cancel out. If $E_{\\rm pin} < 0$ the vortex is attracted by the nucleus ('nuclear pinning') and the average pinning force per site is obtained as $F_{\\rm pin} = |E_{\\rm pin}|/R_{_{\\rm WS}}$. If $E_{\\rm pin} > 0$ the vortex is repelled by the nucleus ('interstitial pinning'); is such a case the pinning force is several orders of magnitude smaller than $F_{\\rm pin}$ and practically negligible. This because the vortex can be moved around the lattice completely avoiding the nuclei, e.g. from (2) to (3), which cost almost no energy \\cite{LE}.  If one considers a \\emph{single} nucleus instead of the actual lattice of nuclei, a simpler quantity can be introduced: the \\emph{binding energy} of the vortex-nucleus system is the energy necessary to detach the vortex from the nucleus and take it to infinity, where the background matter is uniform (unaffected by the presence of the nucleus or the vortex). It can be defined as the difference in energy between two configurations, one with the vortex bound to the nucleus (i.e. with vortex-nucleus separation $d=0$) and the other with the vortex far away and unbound from the nucleus (i.e. with vortex-nucleus separation $d \\rightarrow \\infty$). The short range nature of nuclear forces and the rapidly decreasing kinetic potential ($\\propto 1/r^{2} $) associated to the vortex flow, however, imply that the vortex and the nucleus interact significantly only when they are relatively close to each other. Therefore, to a good approximation the binding energy is reached not for $d \\rightarrow \\infty$, but already for some finite separation $d = D_{\\rm conv}$. The convergence distance $D_{\\rm conv}$ corresponds to a configuration presenting a zone between the nucleus and the vortex where neutron matter has practically reached uniformity; taking $d > D_{\\rm conv}$ only adds more uniform matter in between, whose contributions cancel out in the binding energy.  In Figure~2~(left) the vortex and the nucleus are shown at a distance $d$ apart. Such a system can be described by two cylinders of radius $R = \\frac{1}{2} d$ and height $h$, which are also natural quantization boxes to evaluate the energy of this geometry. Indeed, the vortex-nucleus bound (B) and unbound (U) configurations are those represented in Figure~2~(right) and the binding energy is defined as $E_{\\rm bind} = E_{_{\\rm B}}-E_{_{\\rm U}}$. Of course, for this to make sense the result must converge; as explained, this will happen only when the separation is $ d \\geq D_{\\rm conv}$ and thence when the cylinders have radius larger than  $R_{\\rm conv} = \\frac{1}{2} D_{\\rm conv}$. The physical interpretation  of the convergence radius is the same as before: when $R = R_{\\rm conv}$, matter has practically reached uniformity at the surface of both cylinders of Figure~2. Taking larger boxes would only add more uniform matter at their boundaries, whose energy contributions cancel out in the binding energy. In this respect, we notice that the choice of cylindrical boxes, although natural, is by no means unique; for example one could as well take two cubic boxes of side $\\ell = d$. As long as $\\ell \\geq 2 R_{\\rm conv}$, energy contributions from the additional uniform matter will again cancel out in the final binding energy.  \\begin{figure}[b] \\begin{center} \\parbox{2.4in}{\\epsfig{figure=pizzochero3.eps,width=2.4in}} \\hspace*{6pt} \\parbox{2.4in}{\\epsfig{figure=pizzochero4.eps,width=2.4in}} \\caption{Left: vortex-nucleus system with separation $d$ and cylindrical quantization boxes of radius $R = \\frac{1}{2} d$. \\ Right: configurations for the binding energy. } \\end{center} \\end{figure}  The complicated geometry of the pinning configurations, with its interplay of different symmetries (spherical for the nucleus, axial for the vortex, periodic for the lattice), makes the calculation of pinning energies a very difficult problem at the quantum level. Conversely, binding energies are simpler to calculate since the convergence requirement allows to split both the U and B configurations into  independent axially symmetric problems (the pairs of cylinders). The question is under which conditions pinning and binding energies can be expected to be equivalent. ",
        "conclusions": ""
    },
    "0711/0711.4099_arXiv.txt": {
        "abstract": "In order to understand the origin of clustered anisotropies detected in {\\it Spitzer} images between 3.6 and 8\\,$\\mu$m, we stack the {\\it Spitzer} IRAC/Great Observatories Origins Deep Survey (GOODS) images at pixel locations corresponding to faint, $z_{AB}\\sim27$ mag, optical sources with no obvious IR counterparts. We obtain a strong detection of the sources with a stacked median flux at 3.6$\\mu$m of 130$\\pm$5 nJy above the background. The wealth of multi-wavelength imaging data in GOODS enables a similar stacking analysis to be undertaken at various wavelengths between the ultraviolet and near-infrared bands. We obtain strong stacked detections of these optically faint sources over the entire wavelength range which places constraints on the average properties of these sources.  We find that the flux spectrum of the median, stacked source is consistent with a $L\\lesssim0.03~L_{*,{\\rm UV}}$ galaxy with a 90\\% confidence interval for the redshift of $1.9-2.7$. These sources produce a 3.6 $\\mu$m absolute background intensity between 0.1 and 0.35 nW m$^{-2}$ sr$^{-1}$ and the clustered IR light could account for $\\sim30-50$\\% of fluctuation power in the IR background at 4 arcminute angular scales. Although the exact redshift distribution of these sources is unknown, these galaxies appear to contain $5-20$\\% of the co-moving stellar mass density at $z\\sim2.5$. ",
        "introduction": "The intensity of the cosmic near-infrared background (IRB) is a measure of the total light emitted by stars and galaxies in the universe which is not thermally reprocessed by dust. The absolute background has been estimated with the Diffuse Infrared Background Experiment (DIRBE; Hauser \\& Dwek 2001) and the Infra-Red Telescope in Space (IRTS; Matsumoto et al. 2005) which results in values that are about a factor of two higher than the intensity obtained by integrating the light from individually detected galaxies. This is most likely attributable to the large uncertainties associated with the removal of foreground zodiacal light from the DIRBE observations (Dwek et al. 2005). This hypothesis has been strengthened by the TeV spectrum of gamma-ray blazars which indicate that the total IRB intensity is smaller than the DIRBE/IRTS estimates (Aharonian et al. 2005) and that $\\sim$ 90\\% of the IRB light probably arises in known galaxy populations. Yet, attempts have been made to explain the difference between the measured and resolved IRB intensity with sources at the epoch of reionization (e.g., Santos, Bromm \\& Kamionkowski 2002; Salvaterra \\& Ferrara 2003; Cooray \\& Yoshida 2004; Fernandez \\& Komatsu 2006). If first-light galaxies are to explain   the missing intensity completely, then an extreme scenario is needed with the conversion of at least 5\\% of all baryons to stars (Madau \\& Silk 2004).  Instead of the absolute intensity, recent works have concentrated on spatial   fluctuations of the IRB (Cooray et al. 2004; Kashlinsky et al. 2004). While an interpretation of unresolved fluctuations is subject to extremely uncertain astrophysical modeling of underlying faint populations, fluctuations in deep {\\it Spitzer}  images have  been fully attributed to first-light galaxies containing Pop~III stars during reionization (Kashlinsky et al. 2005). An alternative study shows that a reasonable fraction ($>50$\\%) of unresolved fluctuations is arising from faint, unresolved sources at lower redshifts (Cooray et al. 2007). Thus, while two independent studies find a  similar clustering amplitude for IRB fluctuations (Kashlinsky et al. 2004; Cooray et al. 2007), they differ in the interpretation related to the contribution from Pop~III stars due to differences in the background light ascribed to faint foreground galaxies unresolved by {\\it Spitzer}.  The suggestion for a low redshift origin in Cooray et al. (2007)  comes from predictions using a halo model (Cooray \\& Sheth 2002) for the 3.6 $\\mu$m population matched to their luminosity functions  (Babbedge et al. 2006) and clustering power spectra (Sullivan et  al. 2007). In Cooray et al. (2007) this low-redshift population was identified as the faint optical galaxies that are resolved in the {\\it Hubble} Advanced Camera for Surveys (ACS; Giavalisco et al. 2004) images of the Great Observatories Origins Deep Survey (GOODS; Dickinson et al. 2003) fields, but not detected in  deep {\\it Spitzer} Infrared Array Camera (IRAC)   images  of the same fields. With models, it was estimated  that this faint galaxy population  has a 3.6 $\\mu$m  intensity between 0.1 and 0.8 nW m$^{-2}$ sr$^{-1}$. The analysis in Cooray et al. (2007) did not exclude a high redshift contribution to the IRB, with an  upper  limit on the 3.6 $\\mu$m absolute background intensity from $z>8$ sources of  0.6   nW m$^{-2}$ sr$^{-1}$. In comparison, if all unresolved IR fluctuations are generated by faint $z>6.5$ sources which are below the detection threshold of {\\it Spitzer}, then the IRB intensity for such sources is $\\gtrsim$1 nW m$^{-2}$ sr$^{-1}$ at 3.6 $\\mu$m (Kashlinsky et al. 2007c).  In this {\\it Letter}, we further study the physical origin of fluctuations in the IR background. Since faint ACS optical sources are undetected in deep {\\it Spitzer} IRAC images, we stack their pixel locations in IRAC to establish the average IR flux. We also utilize the wealth of multiwavelength imaging data in the GOODS fields to obtain a stacked broadband spectral energy distribution of these faint optical galaxies between UV and near-IR bands. Leaving the redshift as a free parameter, we fit the average optical to IR flux spectrum with galaxy population synthesis models and find that the stacked flux spectrum is best fit by a $\\sim0.03$~L$_{*, {\\rm UV}}$ galaxy at $1.9<z<2.7$, similar to a scaled-down population of Lyman-break galaxies. We also measure the expected IRB fluctuation clustering spectrum from these sources.  We summarize the stacking analysis in the next section and discuss results on the average optical to IR flux spectrum in \\S~3. We discuss clustering of these faint optical sources in \\S~4. In \\S~5 we discuss the implications of our measurements. The magnitudes quoted throughout this paper are AB magnitudes.   \\begin{figure}[!t] \\centerline{\\psfig{file=Fig1.ps,width=3.3in,angle=0}    } \\caption{ The average-stacked 3.6 $\\mu$m image from GOODS-N with pixel locations of all 6160 ACS sources unmatched and unmasked in IRAC. The median stacked flux is  $(130 \\pm 5)$ nJy (or $=26.1 \\pm 0.1$ magnitude) and is detected with a statistical signal-to-noise ratio higher than 25. } \\label{fig:stack} \\end{figure}  \\begin{figure}[!t] \\centerline{\\psfig{file=Fig2.ps,width=3.9in,angle=0}    } \\caption{ The stacked ultraviolet to IR flux spectral energy distribution of faint galaxies detected by ACS but not detected individually by {\\it Spitzer}/IRAC in GOODS-S (black line and data points). The SED is consistent with the sub-L$_{*,{\\rm UV}}$ galaxy population at $z\\sim2.5$. The parameters of the stacked galaxy for GOODS-S are as shown in the legend. The convolution of the fit SED with the different passbands is shown with empty circles. The stacking in GOODS-N results in a best fit redshift of $z\\sim1.9$ and the SED shown as the dotted line (See text for details). } \\label{fig:spec} \\end{figure}  \\begin{deluxetable}{ccc} \\tablecaption{Results from Stacking Analysis\\label{table:data}} \\tablehead{ \\colhead{} & \\colhead{GOODS-N} & \\colhead{GOODS-S} \\\\ } \\startdata $\\Delta \\Omega$/arcmin$^2$ &   169  & 171 \\\\ N$_{\\it ACS}$ &  6160 & 5441 \\\\ U & 28.3 $\\pm$ 0.1 & 28.4 $\\pm$ 0.1 \\\\ B & 28.0 $\\pm$ 0.1 & 27.8 $\\pm$ 0.1 \\\\ V & 27.5 $\\pm$ 0.1 & 27.4 $\\pm$ 0.1 \\\\ i & 27.2 $\\pm$ 0.1 & 27.1 $\\pm$ 0.1 \\\\ z & 27.1 $\\pm$ 0.1 & 27.0 $\\pm$ 0.1 \\\\ HK$\\arcmin$ & 27.3 $\\pm$ 0.3 & ... \\\\ J & ... & 26.9 $\\pm$ 0.15 \\\\ H & ...\t& 26.7 $\\pm$ 0.15 \\\\ Ks & ... & 26.6 $\\pm$ 0.15 \\\\ 3.6 $\\mu$m  & 26.1 $\\pm$ 0.1 &  26.0 $\\pm$ 0.1 \\\\ 4.5 $\\mu$m & 26.2 $\\pm$ 0.1 & 26.2 $\\pm$ 0.1 \\\\ 5.8 $\\mu$m & 26.2 $\\pm$ 0.1 & 26.0 $\\pm$ 0.1 \\\\ 8.0 $\\mu$m & 26.3 $\\pm$ 0.15 & 26.0 $\\pm$ 0.15 \\\\ \\enddata \\tablecomments{N$_{\\it ACS}$ is the total number of faint ACS sources over an effective area $\\Delta \\Omega$ that were included in the stack. Uncertainties in the stacked flux (tabulated in AB magnitudes) are mostly dominated by calibration systematics.} \\end{deluxetable} ",
        "conclusions": "The Kashlinsky et al. (2007c) interpretation for IR fluctuations involves $z>6.5$ galaxies with an IR background intensity of $\\gtrsim$ 1 nW m$^{-2}$ sr$^{-1}$ at 3.6 $\\mu$m. This is a large intensity given that all resolved sources, so far, lead to an IR background intensity between $\\sim$ 6 nW m$^{-2}$ sr$^{-1}$ (Sullivan et al. 2007; Fazio et al. 2004) and $\\sim$ 10 nW m$^{-2}$ sr$^{-1}$ (Levenson \\& Wright 2008). To avoid individual  detections of a large number of $z>6.5$ galaxies, such a high background intensity must then be hidden in a source population with a large surface density but with individual faint fluxes of around 10 nJy for each source. Given the large surface density of the   population, pixel to pixel intensity variations  are a small fraction of the total intensity produced by those sources.  The alternative explanation is that a reasonable fraction of fluctuations arises from a population of sources that have a low surface density but are just below the {\\it Spitzer}/IRAC detection threshold. The results from the stacking analysis presented here suggest that faint optical galaxies, with an average source flux of around 130 nJy at 3.6 $\\mu$m, contribute between $\\sim$30$-$50\\% of the power in the fluctuations on the largest angular scales that we are able to measure.   The average IR absolute intensity produced by these sources is in the range 0.12-0.35 nW m$^{-2}$ sr$^{-1}$ at 3.6 $\\mu$m, in agreement with the Kashlinsky et al. (2005) estimate that the contribution from faint extragalactic sources is $\\sim$0.15 nW m$^{-2}$ sr$^{-1}$. However, they claim that such faint objects cannot account for the strong clustering signal in unresolved IR light. We have demonstrated in this paper that a significant fraction of the fluctuations do arise from these faint galaxies. These sub-L$_{*}$ galaxies contribute up to 20\\% of the stellar mass density at $z \\sim 2.5$ and can be described as a low luminosity version of well-studied Lyman-break galaxies.  While these sources do not fully explain all IR fluctuations measured with deep IRAC images, in Cooray et al. (2007), we placed a conservative upper limit that any contribution from $z>6.5$ sources must have an absolute intensity less than about 0.6 nW m$^{-2}$ sr$^{-1}$.  Beyond {\\it Spitzer} IRAC, studies have also been conducted at lower near-IR wavelengths with NICMOS. Though limited by the small field-of-view of NICMOS, they show that   IR fluctuations at 1.6 and 1.25 $\\mu$m are more consistent with a $z<8$ origin  than a high redshift interpretation (Thompson et al. 2007a, 2007b). Beyond IR fluctuations, it will be interesting to further study the nature of the faint, dwarf galaxy population at $z=2$ to 3 to understand the contribution it makes towards bridging the systematic offset between the co-moving stellar mass density and the integrated star-formation rate density (Hopkins \\& Beacom 2006).  {\\it Acknowledgments:} We thank Mark Dickinson for helpful suggestions. We also acknowledge the contributions of various members of the GOODS team who are associated with key aspects of data processing and handling. This research was funded by NASA grant NNX07AG43G, NSF CAREER grant AST-0645427, award number 1310310 from {\\it Spitzer} for Archival Research, and program number HST-AR-11241.01-A by NASA through a grant from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Incorporated, under NASA contract NAS5-26555."
    },
    "0711/0711.0997_arXiv.txt": {
        "abstract": "We analyse Keck spectra of 24 candidate globular clusters (GCs) associated with the spiral galaxy NGC~2683. We identify 19 bona fide GCs based on their recession velocities, of which 15 were suitable for stellar population analysis.  Age and metallicity determinations reveal old ages in 14 out of 15 GCs.  These old GCs exhibit age and metallicity distributions similar to that of the Milky Way GC system. One GC in NGC~2683 was found to exhibit an age of $\\sim$3~Gyr. The age, metallicity and $\\alpha$-element abundance of this centrally located GC are remarkably similar to the values found for the galactic centre itself, providing further evidence for a recent star formation event in NGC~2683. ",
        "introduction": "\\label{intro}  Despite their historical importance in understanding the formation processes of our own Galaxy (Eggen, Lynden-Bell \\& Sandage 1962; Searle \\& Zinn 1978; Mackey \\& Gilmore 2004; Forbes, Strader \\& Brodie 2004), detailed studies of the stellar populations of globular cluster (GC) systems in spiral galaxies beyond the Local Group are somewhat limited.  It is important that this be rectified, not only to inform formation models of spiral galaxies, but also to constrain formation models of other morphological types. For example, Ashman \\& Zepf (1992) proposed that the GC systems of elliptical galaxies represent the merged systems of spiral galaxies plus the addition of newly formed red (metal-rich) GCs. Bedregal et al. (2006) have argued that the GC systems of S0s are consistent with faded spirals. A better understanding of GC systems in spirals, with a range of types and luminosities, is needed to test these ideas.  Imaging studies of GC systems exist for about a dozen spirals (e.g. Kissler-Patig et al. 1999; Larsen, Forbes \\& Brodie 2001; Goudfrooij et al. 2003). When sufficient numbers of GCs are present, they reveal a bimodal colour distribution (similar to those seen in elliptical galaxies) with the red (metal-rich) subpopulation associated with the galaxy bulge component (see Forbes, Brodie \\& Larsen 2001). Spectroscopic studies of GCs in spirals beyond the Milky Way and M31 (Burstein et al. 1984; Beasley et al. 2004) are even more limited. Schroder et al. (2002) investigate the stellar population properties of 16 individual GCs in M81. Similar, or smaller, numbers have been investigated in M104 (Larsen et al. 2002), NGC~253 and NGC~300 (Olsen et al. 2004), and M33 (Chandar et al. 2006). These studies generally find old ages with a wide range of metallicities for the GCs. Some GC systems reveal bulk rotation, while others do not, but small numbers and the lack of edge-on systems make such analyses uncertain.   \\begin{table*} \\begin{centering} \\begin{tabular}[b]{|c|c|c|c|c|c|c|c|c|} \\hline Galaxy    & Hubble       &Distance &M$_V$      &Bulge r$_{\\rm{eff}}$ & Disc Scale &      Mass HI       &  N$_{GC}$        &  S$_N$   \\\\ &    Type      & (Mpc)   &(mag)      & (kpc)          &  (kpc)     &  (10$^9$M$_{\\odot}$) &                &          \\\\ \\hline Milky Way & S(B)bc$^1$ &  --     &-20.9$^1$    & 2.5$^3$        &  5.0$^5$        &  4.0$^5$        & 160$\\pm$20$^1$ & 0.70     \\\\ M31       & Sb$^1$     & 0.78    &-21.2$^1$    & 2.4$^3$        &  6.4$^5$        &  3.0$^5$        & 400$\\pm$55$^1$ & 1.32     \\\\ NGC~2683  & Sb$^2$     & 7.2     &-20.3        & 2.5$^4$        &  1.7$^2$        &  0.6$^2$        & 120$\\pm$40$^6$   & 0.90      \\\\ \\hline \\end{tabular} \\caption{Comparison of galaxy parameters for NGC~2683, the Milky Way and M31. The values for M$_V$, bulge effective-radius and disc scale-length for NGC~2683 have been adjusted, where necessary, to a distance of 7.2~Mpc. The specific frequency (S$_N$) is calculated from the number of GCs in the galaxy (N$_{GC}$) and the galaxy luminosity (M$_V$) by S$_N$=N$_{GC}$.10$^{0.4(M_V+15)}$. References: 1. Courteau \\& van den Bergh (1999) and references therein; 2. Broeils \\& van Woerden (1994); 3. van den Bergh (1999); 4. Kent (1985); 5. Gilmore, King \\& van der Kruit (1989); 6. Rhode et al. (2007). } \\label{gal_params} \\end{centering} \\end{table*}  Based on the HST Advanced Camera for Surveys (ACS) imaging study of Forde et al. (2007), we have obtained Keck telescope spectra of GC candidates in the nearby, edge-on Sb spiral NGC~2683. A variety of distance estimates exist in the literature for NGC~2683. Here we adopt the surface brightness fluctuation distance modulus of Tonry et al. (2001) modified by the correction found by Jensen et al. (2003). This is gives m--M = 29.28 $\\pm$ 0.36 or 7.2 $\\pm$ 1.3 Mpc which lies near the midpoint of the literature estimates. With a luminosity of M$_V$ = --20.31~mag it has a lower luminosity (by a factor of 2) than the Milky Way or M31.  We note that, although possessing a Hubble type and bulge size similar to the Milky Way and M31, NGC~2683 exhibits a disc size and HI gas mass that are significantly smaller.  Rhode et al. (2007) show that the extent of the GC system of NGC~2683 is also rather small, with the projected density of the system falling to background levels within $\\sim$8~kpc. This can be compared to the Milky Way GC system in which a fraction of the GC system lies outside $\\sim$30~kpc.  Some properties of NGC~2683 are compared to those of the Milky Way and M31 in Table \\ref{gal_params}.    In Section \\ref{obs} we present our observations and data reduction methods.  The measurement and analysis of recession velocities and Lick indices are given in Section \\ref{spec_anal}. The results of our chemical and kinematic analysis of the sample is outlined in Section \\ref{results}.  Our conclusions are presented in in Section \\ref{disc}. ",
        "conclusions": "\\label{disc} We have analysed the recession velocities and stellar populations of a small sample of GCs in the spiral galaxy NGC~2683 and compared the results with the Galactic GC system.  Our stellar population analysis identified one GC, located near the centre of NGC~2683, with the relatively young age of 3.3$\\pm$0.5~Gyr. The age, metallicity and [E/Fe] of this young GC appear remarkably similar to the values found by Proctor \\& Sansom (2002) for the centre of NGC~2683 itself. This result therefore suggests the possibility that this GC and the recent burst in the central regions were formed at the same time, from the same gas supply, and provide further evidence for a star-formation event in NGC~2683 about $\\sim$3~Gyr ago.  The stellar population parameters of the 14 \\emph{old} globular clusters in our sample show many similarities to the old globular clusters of the Milky Way. The metallicity distribution spans a similar range to those found in studies of the Milky Way and other spiral galaxy systems, i.e from $\\sim$--2.5 to 0.0~dex. The data for NGC~2683 GCs are also consistent with the trends in [Fe/H] with radius observed in red and blue Galactic GC subpopulations.\\\\"
    },
    "0711/0711.2676_arXiv.txt": {
        "abstract": "Two Algol-type eclipsing binary systems (EW Lyr and IV Cas) have been investigated for period changes. Our study was primarily focused on the light-time effect with alternative explanation by magnetic activity cycles. In~the case of EW Lyr we have found third body in the orbit with period about 78 years, amplitude $A=0.052$ days and orbital eccentricity $e=0.57$. For IV~Cas the long period is 58 years, amplitude $A=0.034$ days and zero eccentricity. With these results we are also able to calculate mass functions and minimal masses of these components. ",
        "introduction": "} Eclipsing binaries are known as ideal cosmic laboratories for studying the properties of stars. One method is the period analysis of these systems. In some cases there is a possibility of detection of other bodies in these systems using this method. But the $O\\!-\\!C$ diagrams have to show periodic oscillations with well-defined shapes, sufficiently large amplitudes and adequately long periods. Very short periods have only small amplitudes and the long periods are not well covered with data yet. All these conditions and the accuracy of our data set have to be considered for the limitations of our proposed hypothesis of the light-time effect. \\index{$O-C$ diagram}  The light-time effect (hereafter LITE), caused by the orbital motion of the eclipsing pair around the barycenter of the multiple system, was   explained by Irwin (1959) and its necessary criteria have been listed by Friebos-Conde \\& Hertzeg (1973) and also by Mayer (1990). Presence of a third body in the system is possible only if the times of minimum, also the secondary ones, behave in agreement with a theoretical LITE curve, the resultant mass function has reasonable value and corresponding variations in radial velocities are measured. During the last decade also a confirmation by astrometry seems to be plausible. But regrettably in both investigated systems (EW Lyr and IV Cas) no photometric or spectroscopic study was performed, so these LITE systems are still only hypothetic, because of no other confirmation of the LITE. But there is still a possible explanation of this variations in minima times by another phenomenon, such as magnetic activity. We discuss this possibility later. \\index{light-time effect} \\index{EW Lyr} \\index{IV Cas}  In both cases we have calculated new light elements of the eclipsing pair and also the parameters of the third body orbit. Their $O-C$ diagrams are shown in the Figs.~1 and 2. The curves represent computed light-time effect caused by the third bodies with parameters given in Table~2 below. The individual primary and secondary minima are denoted by dots and circles, respectively. Larger symbols correspond to the photoelectric or CCD measurements which were given higher weights in our computations.  For the computation we used the least-squares method. All data were found in published papers. Table 2 presents our results, where $M_i$ are masses of components, $p_3$ computed period of the unresolved body, $A$ semiamplitude of LITE, $e$ eccentricity, $\\omega$ length of periastron, $f(M_3)$ computed mass function and $M_{3,{\\rm min}}$ minimum mass (for $i$ = 90$^\\circ$) of the third body, respectively. ",
        "conclusions": "We have derived new LITE parameters for two eclipsing binaries by means of the $O-C$ diagram analysis. We have suggested the light-time effect for explanation of the period variation. We have also mentioned a probability of presence of magnetic activity cycles, which could modulate the period. In case of secondary component of IV Cas this hypothesis should also be in stage, but we have not enough data to make a final decision. So the consequence is, that for the confirmation of presence of the LITE in these systems, we need further detailed photometric, spectroscopic or astrometric study of these binaries."
    },
    "0711/0711.0673_arXiv.txt": {
        "abstract": "Recently, Rindler and Ishak have argued that the bending of light is, in principle, changed by the presence of a cosmological constant since one must consider not only  the null geodesic equation, but also the process of measurement. I agree with the fact that both must be considered. Here, on the basis of the mathematically exact solution to the classical bending problem, and independent of the cosmological constant, I clarify the approximate argument found in the vast majority of texts (new and old) for the measured value of the bending of light for a single source and show that the result is in part due to the almost perfect cancelation of two terms, one of which is seldom considered. When one considers two sources, this cancelation is of no consequence, and, for example, if the sources are opposite with the same associated apsidal distance, the approximate argument gives the rigorously correct answer (up to numerical evaluation), an answer which is unaffected by the presence of a cosmological constant. ",
        "introduction": "Recently, Rindler and Ishak  \\cite{rindler} have corrected a long-standing error in the literature concerning the cosmological constant ($\\Lambda$) and the bending of light for a single source, an error perpetuated by the author \\cite{lake}. Here I reexamine the classic subject of the bending of light (with a concentration on a solar mass deflector - not gravitational lensing in general) and arrive at some further results which should be of wide - spread interest. The background geometry is the spherical vacuum given by \\begin{equation} \\label{vacuum} ds^2 = \\frac{d\\textsf{r}^2}{f(\\textsf{r})} + \\textsf{r}^2(d\\theta^2 + \\sin^2 \\theta  d\\phi^2)-f(\\textsf{r})dt^2, \\label{line} \\end{equation} where \\begin{equation} \\label{f} f(\\textsf{r}) = 1-\\frac{2 m}{\\textsf{r}}-\\frac{\\Lambda \\textsf{r}^2}{3} \\label{factor} \\end{equation} and we restrict our analysis to $\\theta=\\pi/2$.  \\bigskip ",
        "conclusions": "The standard argument for explaining the measured value of the deflection of light is to set the emitter and observer at $u=0$ so that the total deflection for a single source for grazing incidence with the Sun is given by $(2\\phi(0)-\\pi)|_{u_{_{\\bigodot}}}$, which, when approximated by first order corrections to the null geodesic equation, gives $4u_{_{\\bigodot}}$ $('')$. For wide-angle deflections a similar argument leads to (\\ref{shward}). On the basis of the mathematically exact solution to the classical bending problem I have shown that these are excellent approximations. However, as regards \\textit{explaining} the measured deflection of light, these approximate arguments fall short in the sense that the bending prior to and after an apse are fundamentally different. The post apsidal bending involves the process of measurement and is dominated by the cancelation of two terms, one of which is not considered in the usual approximate argument. Whereas these more involved arguments may not be numerically important, conceptually they are.  \\bigskip  Aside form these details, the present discussion differs from others in the sense that analytic approximations to the exact solution (\\ref{solution}) is not central to the discussion, as is the usual case, but rather (\\ref{solution}) is viewed as an elementary function in the sense that it can be trivially approximated numerically to arbitrary accuracy using modern computational platforms. However, approximations have been included, relegated to the Appendices, as these serve to amplify and ``explain\" the results obtained."
    },
    "0711/0711.0359_arXiv.txt": {
        "abstract": "We develop a simple method of dark energy reconstruction using a geometrical form of the luminosity-distance relation. In this method the FRW dynamical system with dark energy is reconstructed instead of the equation of state parameter. We give several examples which illustrate the usefulness of our method in fitting the redshift transition from the decelerating to accelerating phase as the value of the Hubble function at the transition. ",
        "introduction": "Recent observations of distant SNIa and other astronomical observations of WMAP, SDSS, Chandra X-ray, etc. indicate that current Universe is experiencing an accelerated expansion. There are different cosmological models which explain this phenomenon by postulating the hypothesis that the Universe is filled by fluid which violates the strong energy condition (so called dark energy) or by postulating that the Friedmann equations do not describe well the evolution of the Universe. Due to the lack of a satisfactory cosmological model (while the simplest $\\Lambda$CDM model seems to be the best one we have still some problems with it) the model-independent reconstruction of the expansion history of the Universe and the nature of dark energy seems to be important. These reconstruction method can be divided into parametric and non-parametric approach \\cite{Sahni:2006}. Here we consider the reconstruction of dynamics in terms of single potential function using the relation between the luminosity distance and the Hubble function. This reconstruction allow us to find the moment of transition from decelerating to accelerating phase of expansion as well as to probe the evolution of the dark energy. ",
        "conclusions": "Different cosmological models can be observationally differentiate in terms of the potential function of the dynamical system. We reconstructed the potential function of the FRW dynamical system directly from SNIa observations using the relation between the luminosity distance and the Hubble function and the parametric method of smoothing the function $d_L(z)$. As one concludes the shape of the reconstructed potential is similar to the shape of the $\\Lambda$CDM model potential for small redshift ($z<1$). We compute the moment of transition using the reconstructed potential as well as the potential of the $\\Lambda$CDM model. We also consider the cosmological model with a single, homogeneous and minimally coupled to gravity scalar field $\\phi$ with the potential $U(\\phi)$. We derive the relation between $w_{\\phi}(a)$/$\\rho_{\\phi}(a)$/$U(a)$ function and the potential function of the FRW dynamical system $V(a)$ which allow us to reconstruct the evolution of such functions without making assumptions on the form of the $w_{\\phi}$ and $\\rho_{\\phi}$. We compare evolution of such quantities with the evolution predicted by the $\\Lambda$CDM model."
    },
    "0711/0711.2512_arXiv.txt": {
        "abstract": "{  We prove that inflation is forbidden in the most well understood class of semi-realistic type IIA string compactifications: Calabi-Yau compactifications with only standard NS-NS 3-form flux, R-R fluxes, D6-branes and O6-planes at large volume and small string coupling. With these ingredients, the first slow-roll parameter satisfies $\\epsilon\\ge\\frac{27}{13}$ whenever $V>0$, ruling out both inflation (including brane/anti-brane inflation) and de Sitter vacua in this limit. Our proof is based on the dependence of the 4-dimensional potential on the volume and dilaton moduli in the presence of fluxes and branes. We also describe broader classes of IIA models which may include cosmologies with inflation and/or de Sitter vacua. The inclusion of extra ingredients, such as NS 5-branes and geometric or non-geometric NS-NS fluxes, evades the assumptions used in deriving the no-go theorem. We focus on NS 5-branes and outline how such ingredients may prove fruitful for cosmology, but we do not provide an explicit model.  We contrast the results of our IIA analysis with the rather different situation in IIB.  } ",
        "introduction": "Our desire to understand the large-scale properties of our Universe is one of the motivations for studying fundamental microphysical theories such as string theory. Indeed, observing the early Universe may be our most promising path toward confronting string theory with data. The leading paradigm for explaining the large-scale isotropy, homogeneity, and flatness of the Universe, as well as its $\\mathcal{O}(10^{-5})$ seed fluctuations, is cosmological inflation \\cite{ Guth,Linde,AlbrechtSteinhardt82,Linde83}. Specifically, by assuming that there exist one or more scalar fields that undergo slow rolling in the early Universe in a potential energy function of just the right shape, one can explain these large-scale properties and predict the numerical values of as many as eight cosmological parameters, many of which have now been accurately measured \\cite{8pars,WMAP06,SDSS06}. The leading candidate for a fundamental microphysical theory is string theory, and so we would like to know how generically string theory can accommodate such potential energy functions.  It is rather well known that the conditions for inflation do not arise easily in string theory, or in other words, that a {\\em generic} point in field space may not be expected to satisfy the slow-roll conditions \\cite{Hellerman, Fischler, Brustein}.  In part this is because of issues like the $\\eta$ problem (essentially, that the potential varies too quickly), which also complicate attempts to build inflationary models in quantum field theory and supergravity theories. There are, however, three reasons to suspect a priori that string theory can accommodate inflation: Firstly, the potential energy is typically a function of hundreds of fields, which means that there is a  large field space to explore.  Secondly, there are an exponentially large number of infinite families of potential energy functions, parameterized by typically hundreds of discrete fluxes in the compact space \\cite{fluxrevs1,fluxrevs2,fluxrevs3,fluxrevs4,fluxrevs5}. Thirdly, there are at least many millions of different topologies, such as Calabi-Yau manifolds, that in general give rise to qualitatively different physical theories in 4 dimensions.  It is reasonable to suspect that occasionally in this vast space of possibilities, the conditions for inflation are satisfied.  For this reason, the past few years have seen intense investigation into the possibility of inflation driven by closed string moduli \\cite{modulione,modulitwo, modulithree,modulifour,modulifive}, axions \\cite{axionone,axiontwo,axionthree,axionfour,axionfive}, or brane positions \\cite{braneone,braneshiu,branetwo,branethree,branefour,branefive,branesix,braneseven} in the extra dimensions. In the most intensely studied case of IIB string compactifications on Calabi-Yau orientifolds, the conclusion at this point is that one can probably build working models, at the cost of fine-tuning the relevant potentials. Some of these models could even have interesting observable signatures \\cite{strings,NG}. Reviews of this general subject appear in \\cite{reviews1,reviews2,reviews3,reviews4,reviews5}.  One feature of the existing constructions is that they are implicit, relying at some point on either non-compact models of regions of the compactification space, or on the ability to perform tunes which (though seemingly possible based on detailed theoretical considerations) are not performed explicitly. One would ideally like to build simpler models, where all of the calculations are performed in a completely explicit and reliable way. Commonly, compactifications suffer from unstabilized moduli in the low energy description, or the calculations stabilizing moduli  apply in a regime that, while apparently numerically controlled, is not under parametric control.\\footnote{This in particular applies in any class of stabilized compactifications where the number of choices of fluxes, branes, etc., while perhaps very large, is ${\\it finite}$. In such models, there is perforce a limit on how small $g_s$ can be, namely the smallest $g_s$ obtained in the finite list.  Of course if the finite number is sufficiently large, the smallest attainable coupling may be quite small, so this may not be a serious limitation.}  In some limits of string theory, however,  we now have explicit examples of stabilized models with parametric control of the moduli potential.  The best understood case occurs in massive IIA string theory; namely IIA string theory with R-R 0-form flux, compactified on a Calabi-Yau orientifold.  The 10-dimensional massive IIA supergravity action was suggested in \\cite{Romans}.  The compactification of this theory on a Calabi-Yau orientifold was performed in a 4-dimensional supergravity formalism in \\cite{Grimm} and the stabilization was obtained in \\cite{DGKT}.\\footnote{With additional ingredients (geometric flux) stabilization was achieved in \\cite{vz}.} The 10-dimensional description of these compactifications was further studied in \\cite{Acharya}. Since these models carry at most ${\\cal N}=1$ supersymmetry in 4 dimensions, and gauge groups and chiral matter can be incorporated in this context, we consider them to be at least semi-realistic.  An investigation of cosmology in these IIA compactifications was initiated in Ref.~\\cite{Hertzberg} by considering some specific simple examples and showing that inflation could not occur in these examples. In the present work we extend that study.  We use a simple scaling analysis of various terms which appear in the low-energy 4-dimensional potential to rule out inflation and de Sitter vacua at large volume and weak coupling in all IIA Calabi-Yau models with conventional fluxes, D-branes, and O-planes. This means that inflation imposes the constraint that our Universe is not in this portion of the landscape. We emphasize that the derivation of this no-go result is only valid in the large volume limit and can be evaded by various other structures including NS 5-branes as well as geometric and non-geometric NS-NS fluxes, indicating that IIA compactifications containing these ingredients may be a good place to look for string models with inflation and/or de Sitter vacua.  Indeed, as we were completing this work we received a copy of \\cite{Eva}, in which more explicit IIA de Sitter models are constructed by using geometric NS-NS fluxes, 5-branes, and various other ingredients.  The structure of this article is as follows:  In Section \\ref{SR} we summarize the IIA supergravity theory in 10 dimensions and outline the dimensional reduction to 4 dimensions.  We explain the key step in the analysis of the paper, which involves considering 2-dimensional slices in the full moduli space parameterized by the volume and dilaton moduli of the compactification.  The behavior of the four (space-time) dimensional potential on these 2-dimensional slices of moduli space allows us to place a lower bound on the slow-roll parameter $\\epsilon$.  In Section \\ref{Nogo} we compute the scaling of the various terms appearing in the 4-dimensional potential energy function $V$ in terms of the two model-independent moduli. We prove that both inflation and de Sitter vacua are forbidden at large volume and weak string coupling when standard fluxes, D6-branes, and O6-planes are included; the slow-roll parameter is bounded below in this case by $\\epsilon\\ge\\frac{27}{13}$ whenever $V>0$.  In Section \\ref{Evade} we describe some additional ingredients such as NS 5-branes, geometric fluxes and non-geometric NS-NS fluxes which can be included in type IIA and which lead to terms in the 4-dimensional potential with scaling properties allowing us to evade the no-go theorem.  In Section \\ref{sec:IIB} we discuss the type IIB theory.  We show how the structure of the IIB theory differs from the IIA theory from the point of view taken in this paper and discuss the connection of our IIA results with previous work on inflation in IIB models.  We discuss our results in Section \\ref{Discussion}. More details regarding the kinetic energy and potential energy are provided in Appendix \\ref{App}. ",
        "conclusions": "\\label{Discussion}  In this paper we have demonstrated that a large class of flux compactifications of type IIA string theory cannot give rise to inflation in the regime of moduli space where we have parametric control of the potential.  This result applies to large-volume, weak coupling compactifications on arbitrary Calabi-Yau spaces with NS-NS 3-form flux, general R-R fluxes, D6-branes and O6-planes.  These ingredients are arguably the most well understood in IIA compactifications.  The no-go theorem of Section \\ref{Nogo} applies in particular to the $T^6/\\mathbb{Z}_3^2$ orientifold model of Ref.~\\cite{DGKT} and the $T^6/\\mathbb{Z}_4$ orientifold model of Ref.~\\cite{IW}, and explains the numerical results of Ref.~\\cite{Hertzberg} which suggested that inflation is impossible in these models.  The no-go theorem we have derived here, however, applies to all other Calabi-Yau compactifications of this general type as well.  So the simplest part of the IIA flux compactification landscape does not inflate.  This implies the following constraint: the portions of the landscape that are possibly relevant to phenomenology will necessarily involve interplay of more diverse ingredients, as has also been found in the IIB theory.  We emphasize that while our derivation has only involved two moduli (the volume $\\volu$ and the dilaton $\\dil$) we are not assuming that either of those moduli necessarily play the role of the inflaton.  Instead, inflation by {\\em any} modulus (or brane/anti-brane) is always spoiled due to the fast-roll of $\\volu$ and/or $\\dil$. This follows because a necessary condition for slow-roll inflation is that the potential be flat in {\\em every} direction in field space, as quantified by $\\epsilon$. In fact because the first slow-roll parameter is so large $\\epsilon\\geq \\frac{27}{13}$, there can never be many e-foldings, even ruling out so-called {\\em fast-roll inflation} \\cite{Fast}. We point out that this result is slightly non-trivial because it requires analyzing {\\em both} $\\volu$ and $\\dil$, as in eq.~(\\ref{operator}), and cannot be proven by focussing on only one of them.   A simple corollary of our result is that no parametrically controlled de Sitter vacua exist in such models.  We emphasize, though, that proving the non-existence of inflation is a much stronger statement than proving the non-existence of de Sitter vacua.  In particular we can imagine a priori a scenario where, although the vacuum is anti-de Sitter (or Minkowski) or there is no vacuum at all, inflation is still realized somewhere in a region where $V>0$. We find, however, that this does not occur.  It is intriguing that the proof of the non-existence of inflation is so closely related to that of the non-existence of de Sitter.  This {\\em suggests} that there may be a close connection between building de Sitter vacua and realizing inflation.  Our result can be interpreted as giving a necessary condition for inflation in IIA models.  To have inflation, some additional structure must be added which gives rise to a potential term in 4 dimensions with scaling such that $-\\rho \\partial V/\\partial \\rho -3 \\tau \\partial V/\\partial \\tau = \\alpha V$ with a coefficient $\\alpha < 9$ for a positive contribution and $\\alpha > 9$ for a negative contribution. We described various ingredients which give rise to such terms; compactifications including these ingredients may be promising places to look for inflationary models. Some of these ingredients take us outside the range of string compactifications which are understood from a perturbative/supergravity point of view.  Among the possibilities which evade the assumptions made in deriving the no-go theorem are other NS-NS fluxes, such as geometric and non-geometric fluxes.  It is currently difficult to construct models with such generic fluxes in a regime that is under control, but progress in this direction has been made in \\cite{Eva}, where IIA de Sitter vacua are found using a specific set of geometrical fluxes and other ingredients.  Another promising direction which we have indicated here (also incorporated in \\cite{Eva}) is to include NS 5-branes and anti-NS 5-branes on 2-cycles in the Calabi-Yau.  More work is needed to find explicit models where these branes are stabilized in a regime allowing de Sitter vacua and inflation, but this does not seem to be impossible or ruled out by any obvious considerations.  In addition to the mechanisms we have discussed, there are probably other structures (e.g., D-terms \\cite{Wrase2,Wrase3}) which violate the conditions of the no-go theorem.  Including any of these ingredients does not guarantee that inflation will be realized.  It may be the case that a slightly more general no-go theorem for inflation exists with certain combinations of additional ingredients. This may follow from studying other moduli, since any field can ruin inflation by fast-rolling. It may also be that with more work an elegant realization of inflation and  de Sitter vacua can be found in type IIA string theory.  This deserves further investigation."
    },
    "0711/0711.0384_arXiv.txt": {
        "abstract": "{} {We estimate stellar masses of galaxies in the high redshift universe with the intention of determining the influence of newly available Spitzer/IRAC infrared data on the analysis. Based on the results, we probe the mass assembly history of the universe.} {We use the \\emph{GOODS--MUSIC} catalog, which provides multiband photometry from the U--filter to the 8~$\\mu m$ Spitzer band for almost $15\\,000$ galaxies with either spectroscopic (for $\\approx 7~\\%$ of the sample) or photometric redshifts, and apply a standard model fitting technique to estimate stellar masses. We than repeat our calculations with fixed photometric redshifts excluding Spitzer photometry and directly compare the outcomes to look for systematic deviations. Finally we use our results to compute stellar mass functions and mass densities up to redshift $z = 5$.} {We find that stellar masses tend to be overestimated on average if further constraining Spitzer data are not included into the analysis. Whilst this trend is small up to intermediate redshifts $z \\la 2.5$ and falls within the typical error in mass, the deviation increases strongly for higher redshifts and reaches a maximum of a factor of three at redshift $z \\approx 3.5$. Thus, up to intermediate redshifts, results for stellar mass density are in good agreement with values taken from literature calculated without additional Spitzer photometry. At higher redshifts, however, we find a systematic trend towards lower mass densities if Spitzer/IRAC data are included.} {} ",
        "introduction": "Whilst the assembly of stellar mass over cosmic history became a matter of interest decades ago, substantial research on this problem has become feasible only in the past few years. Extensive surveys had to be carried out, and techniques to infer stellar masses of galaxies from multicolor photometry had to be developed, as large complete spectroscopic samples at cosmic distances are not yet available. In general, these methods rest upon fitting a grid of stellar population models to the data to compute stellar masses by multiplying the mass--to--light ratio of the best matching template with its absolute luminosity. This procedure has been used in numerous publications to calculate stellar mass functions or stellar mass densities up to intermediate redshifts \\citep[e.g.][]{2000ApJ...536L..77B, 2001ApJ...562L.111D, 2003ApJS..149..289B, 2006A&A...453..869B}. With the increasing availability of deeper surveys, the analysis has been extended to higher redshifts, e.g. up to $z = 2$ or 3 \\citep{2003ApJ...587...25D, 2003ApJ...594L...9F, 2003ApJ...599..847R, 2004A&A...424...23F}, or even to $z = 5$ \\citep{2005ApJ...619L.131D}. However, the derived results are typically based on observations from the U to the K--band and may be affected by extrapolation errors as the rest--frame wavelength coverage of the observed objects is shifted to the blue, and a good observational constraint to a galaxy's optical to near--infrared rest--frame luminosity is necessary to estimate its stellar mass reliably. Results for stellar masses may therefore be impaired by systematic uncertainties in the high redshift regime if observations are not extended to longer wavelengths.  Since the Spitzer space telescope has become operational, it has been possible to complement observations up to the K filter with high--quality infrared data in the adjacent wavelength range \\citep{2004ApJS..154....1W}. To benefit from this improvement we used the Great Observatories Origins Deep Survey, which provides deep publicly available observations from numerous facilities in different wavelength regimes \\citep{dickinson}. Based on these data we probe the influence of Spitzer photometry on the estimated stellar masses and examine the consequences of the results on inferences about the mass assembly history of the universe. In this work we address the effects of Spitzer data to the mass calculation process only, i.e.\\ we do not study whether additional systematic deviations of photometric redshifts emerge that, in general, can influence the result as well. Furthermore we adopt a specific set of template models for the calculations.  The paper is organized as follows. In Sect.~2 we give a short overview of the dataset, and discuss the procedure adopted to estimate stellar masses in Sect.~3. We focus on the influence of Spitzer data on the derived masses in Sect.~4 and use the outcome to calculate stellar mass functions (Sect.~5) and mass densities (Sect.~6). Finally, we summarise our results and draw our conclusions in Sect.~7.  Throughout the paper we assume $\\Omega_{\\mathrm{M}} = 0.3$, $\\Omega_{\\Lambda} = 0.7$ and $H_{0} = 70~\\mathrm{km\\,s^{-1}\\,Mpc^{-1}}$. Magnitudes are given in the AB system. ",
        "conclusions": "In this work we estimated the influence of newly available infrared data longward of the K--filter on stellar mass estimates. To do so we used the \\emph{GOODS--MUSIC} catalog published by \\citet{2006A&A...449..951G}, which combines photometric data in 10 filters from 0.35 to 2.3~$\\mu m$ with observations from the IRAC instrument of the Spitzer space telescope at 3.6, 4.5, 5.8 and 8.0~$\\mu m$. The catalog consists of 14\\,847 objects within an area of 143.2~$\\mathrm{arcmin^2}$ detected either in the z or the $\\mathrm{K_S}$--band. We computed stellar masses of this sample by fitting stellar population synthesis models \\citep{bruzual-2003-344} to the data and multiplying the k--corrected absolute $\\mathrm{K_S}$--band luminosity with the $M/L_\\mathrm{K}$--ratio of the best fitting model SED. To probe the influence of the IRAC data on the analysis we repeated the computation of stellar masses without Spitzer photometry, keeping the photometric redshifts fixed, and compared the outcome directly.  We found stellar masses to be overestimated on average, if further constraining infrared data from Spitzer were not included in the calculation. Whilst this trend is almost independent of mass itself, a closer analysis reveals a strong dependency on redshift.  While up to $z \\approx 2$ the systematic deviation in mass is only moderate ($|\\log \\, M_\\mathrm{U-4}/M_\\mathrm{U-K}| < 0.2~\\mathrm{dex}$) and comparable to the intrinsic uncertainty of the method adopted to estimate stellar masses, it increases strongly for higher redshifts and reaches a maximum of a factor of three at $z \\approx 3.5$. The reason for this systematic shift can be traced back to insufficient constraints on the slope of the spectra redward of the 4000~\\AA\\ break at high redshifts. It turns out that, on average, models that are too old, with excessively high absolute infrared luminosities and $M/L$--ratios were assigned to the data if Spitzer photometry is not included. Thus, a shift to lower stellar masses is likely to be correlated with a decreasing age of the best fitting model SED. The inclusion of one additional data--point longward of $3~\\mu m$ can already reduce the remaining error in mass significantly.  In the next step we used our results to calculate stellar mass functions in different redshift intervals utilizing the $V/V_\\mathrm{max}$--formalism of \\citet{1968ApJ...151..393S} to correct our sample partly for incompleteness. To assign errors we performed extensive Monte Carlo simulations where we considered uncertainties in the underlying $M/L_\\mathrm{K}$--ratios, redshifts and a general error in photometry. Afterwards, the data points were fitted via three free parameters using the analytical expression suggested by \\citet{1976ApJ...203..297S}. We found a pronounced general decrease in numberdensity of all objects with redshift. Beyond that, it is hard to say whether there is also a mass dependent evolution. Although the change in computed Schechter parameter may support this position, the effect can be caused by large--scale structure as well.  Finally, we computed stellar mass densities as a function of redshift. We summed the mass of our objects within each redshift interval at the high mass end and integrated the Schechter function derived on the basis of our stellar mass functions to complete the result for lower masses. To estimate errors we again used Monte Carlo simulations, but we pointed out that further effects such as biases in stellar population synthesis models may be dominating sources of uncertainty. By comparing the outcome to the local value of \\citet{2001MNRAS.326..255C} we found at least 42~\\% of the stellar mass density to be already in place at $\\langle z \\rangle = 1$. This value decreases to 23~\\% at $\\langle z \\rangle = 2$ and about 7~\\% at $\\langle z \\rangle = 3.5$. Therefore, up to intermediate redshifts our results are in good agreement with values taken from literature derived without additional Spitzer/IRAC data. However, at high redshifts a systematic deviation to lower densities is present as one would expect from the effect of Spitzer photometry on the calculation."
    },
    "0711/0711.2348_arXiv.txt": {
        "abstract": "The characteristics relating to the lateral and longitudinal development of EAS in the energy region of $10^{15} - 10^{19}$~eV have been analyzed in the framework of the QGSJET model and of mass composition of primary cosmic rays. It is found that at $E_{0} \\ge 5 \\times 10^{15}$~eV the mean mass composition of primary cosmic rays begins to vary as indicated by a rise of $\\left<\\ln A\\right>$ with increasing energy. The maximum value of $\\left<\\ln A\\right>$ is observed at $E_{0} \\sim (5 - 50) \\times 10^{16}$~eV. It is confirmed by data of many compact EAS arrays and does not contradict an anomalous diffusion model of cosmic ray propagation in our Galaxy. In the superhigh energy region $(\\ge 10^{18}$~eV) the value $\\left<\\ln A\\right>$ begins to decrease, i.e. the mass composition becomes lighter and consists of protons and nuclei of He and C. It does not contradict our earlier estimations for the mass composition and points to a growing role of the metagalactic component of cosmic rays in the superhigh energy region. ",
        "introduction": "An investigation of superhigh energy cosmic rays in the 1950s - 1960s and 1970s - 1990s has led to the discovery of irregularities in the energy spectrum at the energy $\\sim 3 \\times 10^{15}$~eV and $\\sim 10^{19}$~eV~\\cite{1, 2}. The problem of origin of breaks in the energy spectrum requires reliable knowledge about the mass composition of primary cosmic rays (PCR) from high energies $\\sim 10^{12} - 10^{14}$~eV up to ultra-high energies $10^{18} - 10^{21}$~eV.  The study of PCR mass composition at $E_{0} \\ge 10^{15}$~eV is possible only with the arrays registering extensive air showers (EAS) and by indirect way only using the EAS characteristics maximum sensitive to the mass composition. For these purposes, at the Yakutsk complex array the characteristics of both longitudinal and radial development of EAS are taken. The longitudinal development of showers is reconstructed measuring the EAS Cherenkov light, and the radial development measuring the structural functions of electron and muons with $E_{\\text{thr}} \\ge 1$~GeV.  In this paper PCR mass composition data obtained at the Yakutsk EAS array in recent year are summarized. ",
        "conclusions": "Fig.~\\ref{fig01} and Fig.~\\ref{fig02} present results on the PCR mass composition of the Yakutsk EAS array. Data were obtained in the framework of the QGSJET model and two-component mass composition (proton-iron nucleus). A few characteristics corresponding to the radial and longitudinal development of EAS~\\cite{3, 4, 5, 6} are used in this analysis. In each case, the value $\\left<\\ln A\\right>$ was determined by an interpolation method according to~(\\ref{eq2}). From Fig.~\\ref{fig01}, Fig.~\\ref{fig02} it follows that the mass composition changes reaching the heavier in the energy region of $(2 - 5) \\times 10^{17}$~eV and then, beginning from $3 \\times 10^{18}$~eV it becomes lighter.  \\begin{figure} \\centering \\includegraphics[width=0.45\\textwidth, clip]{fig01} \\caption{Mass composition of CR obtained by different characteristics of EAS in the region of ultrahigh energies.} \\label{fig01} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=0.45\\textwidth, clip]{fig02} \\caption{Mass composition of cosmic rays at superhigh and ultrahigh energies. The curve is a calculation by Lagutin et al (2001) and (2007) according to the anomalous diffusion model for cosmic ray propogation.} \\label{fig02} \\end{figure}  The curves are calculations according to the anomalous diffusion model of cosmic ray propagation in the Galaxy (dashed line in Fig.~\\ref{fig01})~\\cite{7} and the two-source model (solid line in Fig.~\\ref{fig02})~\\cite{8}. The calculations have been carried out separately for each of following groups of nuclei: p, He, CNO, N-Si, Fe. The summary values of $\\left<\\ln A\\right>$ are shown in Fig.~\\ref{fig01} and Fig.~\\ref{fig02} and for the first case in the Table~\\ref{t01}. Data in the Table~\\ref{t01} point to the change of mass composition. In the energy region of $5 \\times 10^{15} - 5 \\times 10^{18}$~eV the mass composition is heavier than at $E_{0} \\simeq 10^{19}$~eV. It contradicts the experimental data (Fig.~\\ref{fig02}). The calculations~\\cite{8} according to a scenario for the particle generation in the sources of two different types, assuming the analogous model for the CR propagation, showed the presence of a fine structure in the dependence of $\\left<\\ln A\\right>$ on energy (see Fig.~\\ref{fig02}). There are also the sharp peaks in our experimental data. In this connection, the investigation of the PCR mass composition with using a great numbers of EAS characteristics and statistical data as possible is of particular interest.  It should be noted that in the region after a ``knee'' the estimations of the CR mass composition obtained at the compact arrays agree well with each other. The same cannot be said of the energy region of $\\sim 10^{18}$~eV (see Fig.~\\ref{fig02}), where HiRes data point to the faster enriching of the primary radiation by light nuclei and protons in comparison with the Yakutsk array data, which point to the gradual change of the mass composition from heavier to light (proton and He nuclei, on the whole) at $E_{0} \\sim 10^{19}$~eV. In both cases, the data point to the existence of tendency to the ``protonization'' of PCRs at $E_{0} >3 \\times 10^{18}$~eV.  \\begin{table*} \\centering \\caption{Calculation results for $\\left<\\ln A\\right>$.} \\label{t01} \\renewcommand{\\arraystretch}{1.3} \\begin{tabular}{ccccccccc} \\hline & A & H, \\% &  He, \\% & CNO, \\% & Ni-Si, \\% & Fe, \\% & $\\left<\\ln A\\right>$ & $\\left<\\ln A\\right>$ \\\\ $E_{0}$, eV &  & & & & & & & experiment \\\\ \\hline $1 \\times 10^{15}$ & & $22$ & $21$ & $20$ & $17$ & $19$ & $2.14$ & $1.55 \\pm 0.25$ \\\\ $1 \\times 10^{16}$ & & $19$ & $19$ & $20$ & $18$ & $23$ & $2.32$ & $2.15 \\pm 0.24$ \\\\ $1 \\times 10^{17}$ & & $17$ & $17$ & $20$ & $19$ & $26$ & $2.43$ & $3.16 \\pm 0.23$ \\\\ $1 \\times 10^{18}$ & & $18$ & $16$ & $19$ & $19$ & $27$ & $2.44$ & $2.61 \\pm 0.21$ \\\\ $1 \\times 10^{19}$ & & $55$ & $17$ & $9 $ & $ 8$ & $11$ & $1.18$ & $1.66 \\pm 0.18$ \\\\ $1 \\times 10^{20}$ & & $68$ & $18$ & $ 6$ & $ 4$ & $ 3$ & $0.69$ &-\\\\ \\hline \\end{tabular} \\end{table*}"
    },
    "0711/0711.0498_arXiv.txt": {
        "abstract": "The efficacy of fast/slow MHD mode conversion in the surface layers of sunspots has been demonstrated over recent years using a number of modelling techniques, including ray theory, perturbation theory, differential eigensystem analysis, and direct numerical simulation. These show that significant energy may be transferred between the fast and slow modes in the neighbourhood of the equipartition layer where the Alfv\\'en and sound speeds coincide. However, most of the models so far have been two dimensional. In three dimensions the Alfv\\'en wave may couple to the magneto\\-acoustic waves with important implications for energy loss from helioseismic modes and for oscillations in the atmosphere above the spot. In this paper, we carry out a numerical ``scattering experiment'', placing an acoustic driver 4 Mm below the solar surface and monitoring the acoustic and Alfv\\'enic wave energy flux high in an isothermal atmosphere placed above it. These calculations indeed show that energy conversion to upward travelling Alfv\\'en waves can be substantial, in many cases exceeding loss to slow (acoustic) waves. Typically, at penumbral magnetic field strengths, the strongest Alfv\\'en fluxes are produced when the field is inclined $30^\\circ$--\\,$40^\\circ$ from the vertical, with the vertical plane of wave propagation offset from the vertical plane containing field lines by some $60^\\circ$--\\,$80^\\circ$. ",
        "introduction": "\\inlinecite{sc06} (hereafter SC) and \\inlinecite{cally07} recently presented a generalization of magneto\\-hydro\\-dynamic (MHD) ray theory that describes mode transmission and conversion between fast and slow magnetoacoustic waves in the context of solar active regions. A uniform, inclined magnetic field was imposed on a plane parallel atmosphere, and magnetoacoustic rays \\emph{in the vertical plane containing the field lines} were modelled using the general ray transmission/conversion formalism of \\inlinecite{tkb03}. The ray results were fully confirmed by numerical solution of the wave equations. By remaining in the plane defined by the magnetic field and gravity though, only the magneto\\-acoustic modes are involved. The Alfv\\'en wave, the plasma displacement of which is perpendicular to this plane, is entirely decoupled. The purpose of this paper is to extend the wave analysis to the full three dimensions (3D), thereby bringing the Alfv\\'en wave into play. This occurs because the Alfv\\'en wave displacement vector is rotated out of the horizontal plane, and therefore interacts with the atmosphere's stratification. The situation differs from the coronal analysis of \\inlinecite{melIII} and \\inlinecite{melIV}, where field-line curvature and twist are responsible for tying the magnetoacoustic and Alfv\\'en waves together.  Importantly, we integrate from interior to high atmosphere, so as to explore the coupling of helioseismic waves ($p$-modes) with atmospheric oscillations, especially in the vicinity of the Alfv\\'en/sound speed equipartition level $z_{eq}$ (typically shallow subsurface in sunspots) where fast/slow conversion occurs.  Several important lessons may be drawn from our analysis. \\begin{enumerate} \\item The 2D result (SC) that the atmospheric slow wave flux has a pronounced maximum at a particular value of magnetic field inclination from the vertical $\\theta$, (typically in the range $20^\\circ$--$30^\\circ$ depending on spherical harmonic degree $\\ell$), survives into 3D, at least out till $|\\phi|\\lesssim90^\\circ$, where $\\phi$ is the angle between the vertical plane of wave propagation and the vertical plane in which the field lines lie. \\item There is a distinct maximum in upward Alfv\\'enic energy flux in the atmosphere at around $\\theta=40^\\circ$, $\\phi=60^\\circ$, with these numbers only weakly dependent on model parameters. The Alfv\\'en wave can only be produced by mode conversion in this model. \\item Conversion of magnetoacoustic waves to Alfv\\'en waves happens over an extended height range, typically some hundreds of kilometres, as distinct from fast/slow conversion which is normally localized sharply near $z_{eq}$. \\end{enumerate}  In the interests of brevity, the reader is referred to SC for an extensive discussion of the solar context and significance to local helio\\-seism\\-ology of MHD waves in strong magnetic fields near the surface, and to several earlier papers exploring MHD mode conversion using a variety of techniques (\\opencite{sb92}; \\opencite{cb93}; \\opencite{cbz}; \\opencite{bc97}; \\opencite{cb97};\\opencite{cally00}; \\opencite{cc03}, \\citeyear{cc05}).  A 3D analysis from the ray perspective is deferred until Paper II. ",
        "conclusions": "\\label{conc}  The greater breadth of the acoustic maxima in the $\\phi$ direction compared to the $\\theta$ direction exhibited in Figures \\ref{lc5} and \\ref{lcOthers} may be broadly explained using the concept of \\emph{attack angle} discussed in SC, \\emph{i.e.}, the angle $\\alpha$ between the wavevector $\\k$ and the magnetic field $\\B_0$. Ignoring the acoustic cutoff and {\\bv} frequencies, generalized ray theory predicts an acoustic transmission coefficient of \\begin{equation} T=\\exp\\left[-\\pi\\,|\\k|\\,h \\sec\\psi\\sin^2\\alpha\\right]_{a=c},     \\label{T} \\end{equation} where $\\psi$ is the angle the wavevector makes to the vertical at the equipartition level $z_{eq}$, and $h=[d(a^2/c^2)/dz]_{a=c}^{-1}$ is the thickness of the $a\\approx c$ layer. In 3D, $\\sin^2\\alpha=1-(\\sin\\alpha\\cos\\phi\\sin\\psi+\\cos\\theta\\cos\\psi)^2$. If we crudely set $|\\k|=\\omega/c$, then for the case of Figure \\ref{lc5}, $|\\k|=3.7$ $\\rm Mm^{-1}$ and $h=0.29$ Mm, \\emph{i.e.}, $kh\\approx1.1$. Figure \\ref{Tcont} presents contours of $T$ in $\\phi$\\,--\\,$\\theta$ space for this case, with $\\psi=22^\\circ$, and clearly displays a greater breadth in $\\phi$ than in $\\theta$. If the acoustic cutoff effect is reintroduced, the corresponding flux contours must be squeezed upwards at the bottom to produce a ``blank zone'' below $\\theta=10.7^\\circ$ (as in Figure \\ref{lc5}) and an even more pronounced anisotropy. This is very much in qualitative agreement with the acoustic results of Figures \\ref{lc5} and \\ref{lcOthers}.  It is clear from the few examples presented in Section \\ref{num_mod} that 3D magneto\\-acoustic coupling to the Alfv\\'en wave is very strong, rivalling and in some cases surpassing the maximum fast-to-slow transmission. Once the angle ($\\phi$) between the vertical plane of the wave and the vertical plane of the magnetic field differs, the coupling turns on. It is typically strongest at $\\phi=50^\\circ$--\\,$70^\\circ$, and for $\\theta=30^\\circ$--\\,$40^\\circ$. This is a significant result, with important implications for our understanding of $p$-mode lifetimes, and atmospheric oscillations above sunspots.  The assumption of uniform magnetic field is made for simplicity and in the spirit of a simple scattering experiment. It is probably quite reasonable for the localized fast\\,-\\,slow mode conversion. However, the greater breadth of the magneto\\-acoustic\\,-\\,Alfv\\'en coupling region suggests that field-line curvature may be relevant to that process. Curvature can certainly provide another Alfv\\'en coupling mechanism in its own right.  \\begin{figure} \\centerline{\\includegraphics[width=.6\\textwidth]{T}} \\caption{Transmission coefficient $T$ in the $\\phi$--$\\theta$ plane as given by Equation (\\ref{T}) with $|\\k|h=1.1$ and $\\psi=22^\\circ$. The contour levels are 0.9, 0.8, 0.7, \\ldots.} \\label{Tcont} \\end{figure}  The upward (Alfv\\'en) and downward (slow plus Alfv\\'en) magnetic losses illustrated in Figure \\ref{pf} indicate a range of relative efficiencies. At fixed $\\theta$ and $\\phi$ though, it seems that upward loss through the atmosphere is more important at higher frequency. And of course, upward loss is typically of greater relative importance in that $\\theta$\\,--\\,$\\phi$ region where $F_{mag}$ is maximal.  The model presented here differs from the earlier \\inlinecite{cc05} analysis of the interaction of helioseismic $p$-modes with 3D inclined magnetic field in that (\\emph{i}) it consists of a realistic (albeit modified) solar interior model rather than a polytrope; (\\emph{ii}) it contains an overlying atmosphere, rather than a vacuum, which allows escape of waves upward as well as downward; and (\\emph{iii}) it includes a driver instead of solving for (undriven) eigenmodes. Since our focus is on coupling to atmospheric oscillations, point (\\emph{ii}) is especially important. Nevertheless, it is interesting that that study also found substantial (downward propagating) Alfv\\'en wave coupling.  Only the relative magnitudes of the acoustic and magnetic fluxes were quantified in Section \\ref{num_mod}. In practical terms though, how big are these fluxes, and how quickly can they deplete the underlying $p$-modes which drive them? Recall that the fluxes were normalized by setting $\\mathcal{E}=\\int_{z_b}^0 E_{ac}\\,dz$ to 1, with $z_b=-4$ Mm in all cases. Now, $F/\\mathcal{E}$ has the dimensions of inverse time, and broadly represents the inverse time scale over which the $p$-mode energy in $z_b<z<0$ depletes, assuming it were not being continually replenished by the driver at $z_b$. The maximum fluxes seen in Figures \\ref{lc5} and \\ref{lcOthers} are of order three in these normalized terms, where time is being measured in kiloseconds. Consequently, the minimum decay time is of order $\\tau\\sim\\frac{1}{3}$ ks $\\approx300$ seconds only, \\emph{i.e.}, not much more than one period! Of course, this is somewhat of an underestimate of $\\tau$, because: \\begin{enumerate} \\item it relies on a particular combination of $\\theta$ and $\\phi$; only waves travelling in a particular direction and meeting field of a particular inclination will be so strongly affected; \\item we have not integrated the acoustic energy over the entire acoustic cavity $z_1<z\\lesssim0$, so the available energy store is larger than indicated by $\\mathcal{E}$; \\item there is also a substantial store of \\emph{kinetic} energy in the acoustic cavity on which to draw. \\end{enumerate} Nevertheless, the surface magnetic field and its attendant mode conversion (to slow or Alfv\\'en waves) is clearly significant for $p$-modes in active regions.  Figure \\ref{pf} helps us understand the nature of the mode conversion. In all cases, we see a broad conversion region starting a little below $z_{eq}$, but extending well above it. This is quite different from the fast\\,-\\,slow transmission\\,-\\,conversion process, which is strongly localized around the equipartition level. The reason for this may be discerned using a WKB description, which will be addressed in the forthcoming Paper II."
    },
    "0711/0711.3980_arXiv.txt": {
        "abstract": "{Three eclipsing binary systems with astrometric orbit have been studied. For a detailed analysis two circular-orbit binaries (VW~Cep and HT~Vir) and one binary with an eccentric orbit ($\\zeta$~Phe) have been chosen. Merging together astrometry and the analysis of the times of minima, one is able to describe the orbit of such a system completely. The $O-C$ diagrams and the astrometric orbits of the third bodies were analysed simultaneously for these three systems by the least-squares method. The introduced algorithm is useful and powerful, but also time consuming, due to many parameters which one is trying to derive. The new orbits for the third bodies in these systems were found with periods $30$, $221$, and $261$~yr, and eccentricities $0.63$, $0.37$, and $0.64$ for VW~Cep, $\\zeta$~Phe, and HT~Vir, respectively. Also an independent approach to compute the distances to these systems was used. The use of this algorithm to VW~Cep gave the distance $d=(27.90~\\pm~0.29)~\\mathrm{pc}$, which is in excellent agreement with the previous \\emph{Hipparcos} result.} ",
        "introduction": "During the last century, many observations of close binary systems were collected, especially for eclipsing binaries. Most observations of these stars were made photometrically but in some cases also the spectroscopy was obtained. In a few cases, also astrometric observations were carried out and these systems, which were studied by several different techniques, are the most interesting, because one is able to find more relevant parameters of them. Especially, if it is possible to analyse the different measurements together (it means from different observational techniques in one least-squares fit) one is able to find the complete set of parameters describing the orbit of such a binary, which are not in contradiction. Such approach is potentially very powerful, especially in upcoming astrometric and photometric space missions.  There are a few well-known eclipsing binaries with astrometric measurements for which the light-time effect (hereafter LITE) was considered, or expected. However, the eclipsing binary (hereafter EB) nature and the astrometric orbit were usually studied separately.  For example, there are several published light-curve solutions for V505~Sgr with the third light included (see e.g. L\\'azaro et al. (2006), {\\.I}bano{\\v g}lu et al. (2000), etc), and an analysis of its apparent orbital period changes interpreted as the LITE due to the third body (e.g. Rovithis-Livaniou et al. 1991). The only paper which compares the astrometry and a period analysis of $O-C$ deviations from the constant orbital period was published by Mayer (1997). Despite existing astrometric measurements, there were no attempts to combine these two methods together. The results from different approaches were just compared to each other.  Another systems are for instance$\\,$ QS~Aql,$\\,$ 44i~Boo,$\\,$ QZ~Car, SZ~Cam, GT~Mus, or V2388~Oph. The coverage of the astrometric orbit is very poor for some of them. In the system QS~Aql the LITE could be determined precisely, but the astrometric orbit is covered only very poorly, see e.g. Mayer (2004). The opposite case is V2388~Oph, where the astrometric-orbit parameters were computed very precisely, but there are only a few minimum-time measurements (see e.g. Yakut et al. 2004). For SZ~Cam only a few usable astrometric observations were obtained, but the LITE is well-defined and also the third light in the light-curve solution was detected (Lorenz et al. 1998). QZ~Car is a more complicated, probably quadruple system consisting of an eclipsing and a non-eclipsing binary, but there are also only a few usable astrometric measurements. Also GT~Mus is a quadruple system, consisting of an eclipsing and RS~CVn component. In many other cases, only measurements of the times of minima are available, without astrometry. For some others, astrometry without photometry, is only available. Other systems where the astrometric observations were obtained and the LITE is observable or expected are listed in Mayer (2004).  The only paper on combining these two different approaches into one joint solution is that by Ribas et al. (2002), where a similar method as described in this paper was applied to the system R~CMa, but where only a small arc of the astrometric orbit was available. Besides the astrometry and LITE also the proper motion on the long orbit was analyzed. On the other hand one has to note, that in Ribas et al. (2002) the complete astrometric parameters (with proper motions, parallax, etc.) were used, while in this paper only a relative astrometry of the distant body relative to the eclipsing pair was analysed. ",
        "conclusions": "Although the number of systems where the astrometric orbit of the third component has been measured together with the presence of LITE is growing steadily, in the most cases only very limited coverage, both in astrometry and times of minima is available. Especially due to this reason the combined analysis of these systems is still difficult, the results are not very convincing and the resultant parameters are affected by relatively large errors.  During the last decade a few papers combining the approach of simultaneous solution of radial velocities, spectral analysis, astrometry, \\emph{Hipparcos} measurements or LITE were published. Besides the systems mentioned in the introduction (V505~Sgr, QS~Aql, 44~Boo, QZ~Car, SZ~Cam, GT~Mus, and V2388~Oph) there were also the analyses of V1061~Cyg (combining the light curve analysis, radial velocity analysis, light-time effect and \\emph{Hipparcos} measurements, see Torres et al. 2006), the papers where the radial velocity measurements and astrometry were combined (see Muterspaugh et al. (2006) for the solution of LITE system V819~Her, or Gudehus (2001) for $\\mu$~Cas), the paper on HIP 50796 combining the radial velocities with the \\emph{Hipparcos} abscissa data (Torres 2006), or the paper on $\\delta$~Lib comparing the results from the period analysis, light-curve analysis, spectral analysis, radio emission and astrometry, respectively; see Budding et al. (2005).  Three eclipsing binaries were studied in this paper. In the case of VW~Cep, where the orbits in both methods have relatively best coverage, the distant body satisfies the limits for the luminosities, and also the systemic velocity variations coincide with our hypothesis. New results are comparable with the previous ones. An additional fourth body was introduced to describe the long-term variation in times of minima. The system is probably more complicated than we assumed (chromospheric activity cycles, stellar spots and flares), and we decided to explain only the most pronounced effects in the $O-C$ diagram. Using this combined approach it is possible to derive the parallax of VW~Cep more precisely than in previous papers, resulting in $\\pi=(35.85~\\pm$ $\\pm~0.37)~\\mathrm{mas}$. The system $\\zeta$~Phe displays an apsidal motion together with LITE and this explanation fits the $O-C$ residuals quite well. The quality of the astrometry is worse, but the main result about the masses and the spectral type of the third body is in an excellent agreement with the previous photometric analysis. The last system HT~Vir is the case where the new resultant parameters of the distant body are about 2 times larger than one would expect. This could be due to the fact that only a few times of minima were observed in the linear part of the $O-C$ diagram, new minima are needed in the next decades.  The final result is that the method itself is potentially very powerful but it is also very sensitive to the quality of the input data, especially if the method is used for determining the distances of these binaries. It can only be applied successfully in those cases where the astrometric orbit and the LITE in the $O-C$ diagram are well defined by existing observations."
    },
    "0711/0711.4791_arXiv.txt": {
        "abstract": "It has been suggested that a potentially large fraction of supernovae could be accompanied by relativistic outflows that stall below the stellar surface. In this letter we point out that internal shocks that are believed to accelerate protons to very high energies in these flows will also accelerate secondary mesons and muons. As a result the neutrino spectrum from meson and muon decay is expected to be much harder compared to previous estimates, extending as a single power law up to $\\sim$$10^3$ TeV. This greatly improves the detection prospects. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.4933_arXiv.txt": {
        "abstract": "In our recent paper~\\citep{YKW07} we described a fast cubic (bispectrum) estimator of the amplitude of primordial non-Gaussianity of local type, $f_{NL}$, from a combined analysis of the Cosmic Microwave Background (CMB) temperature and E-polarization observations. In this paper we generalize the estimator to deal with a partial sky coverage as well as inhomogeneous noise. Our generalized estimator is still computationally efficient, scaling as $O(N^{3/2}_{pix})$ compared to the $O(N^{5/2}_{pix})$ scaling of the brute force bispectrum calculation for sky maps with $N_{pix}$ pixels. Upcoming CMB experiments are expected to yield high-sensitivity temperature and E-polarization data. Our generalized estimator will allow us to optimally utilize the combined CMB temperature and E-polarization information from these realistic experiments, and to constrain primordial non-Gaussianity. ",
        "introduction": "\\label{introduction}  Non-Gaussianity from the simplest inflation models, that are based on a slowly        rolling        scalar        field,        is        very small~\\citep{Salopek_Bond90,Salopek_Bond91,Falk_et93,Gangui_et94,Acquaviva02,Maldacena03}; however,  a  very large  class  of  more  general models, e.g., models with multiple scalar fields, features in inflation potential, non-adiabatic fluctuations,    non-canonical   kinetic   terms,    deviations   from the Bunch-Davies vacuum, among others, predict  substantially   higher level of primordial non-Gaussianity \\cite[for a review and references therein]{BKMR_04}.     Primordial non-Gaussianity can be described in terms of the 3-point correlation function of Bardeen's curvature perturbations, $\\Phi(k)$, in Fourier space: \\begin{eqnarray} \\langle \\Phi(\\mathbf{k_1})(\\mathbf{k_2})(\\mathbf{k_3})\\rangle = (2\\pi)^3\\delta^3(\\mathbf{k_1} + \\mathbf{k_2} + \\mathbf{k_3})F(k_1, k_2, k_3). \\end{eqnarray} Depending on the shape of the 3-point function, i.e., $F(k_1, k_2, k_3)$, non-Gaussianity can be broadly classified into two classes~\\citep{Babich_etal_04}. First, the local, ``squeezed,'' non-Gaussianity where $F(k_1, k_2, k_3)$ is large for the configurations in which $k_1 \\ll k_2, k_3$. Second, the non-local, ``equilateral,'' non-Gaussianity where $F(k_1, k_2, k_3)$ is large for the configuration when $k_1 \\sim k_2 \\sim k_3$.  The local form arises from a non-linear relation between inflaton and curvature perturbations~\\citep{Salopek_Bond90,Salopek_Bond91,Gangui_et94}, curvaton models~\\citep{Lyth_Ungarelli_Wands_03}, or the New Ekpyrotic models~\\citep{Koyama_etal07,Buchbinder_etal07}. The equilateral form arises from non-canonical kinetic terms such as the Dirac-Born-Infeld (DBI) action~\\citep{Alishahiha_etal04}, the ghost condensation~\\citep{Arkani_et_04}, or any other single-field models in which the scalar field acquires a low speed of sound ~\\citep{Chen_etal07,Cheung_Creminelli_etal07}. While we focus on the local form in this paper, it is straightforward to repeat our analysis for the equilateral form.  The local form of non-Gaussianity may be parametrized in real space as~\\citep{Gangui_et94,verde00,KS2001}: \\begin{equation} \\label{eqn:phiNG} \\Phi(\\mathbf{r}) = \\Phi_L(\\mathbf{r}) + f_{NL} \\left( \\Phi_L^2(\\mathbf{r}) - \\langle \\Phi_L^2(\\mathbf{r}) \\rangle \\right) \\end{equation} where $f_{NL}$ characterizes the amplitude of primordial non-Gaussianity. Different inflationary  models predict  different amounts  of  $f_{NL}$, starting from $O(1)$ to $f_{NL}\\sim 100$, beyond which values have been excluded by the Cosmic Microwave Background (CMB) bispectrum of WMAP temperature data, $-36 < f_{NL} <100$, at the $2\\sigma$ level \\citep{nong_wmap,creminelli_wmap2,wmap_2nd_spergel}.  So far all the constraints on primordial non-Gaussianity use only temperature information of the CMB. By also  having the E-polarization  information together with CMB temperature information, one  can  improve the sensitivity to the primordial  fluctuations \\citep{BZ04,YW05,YKW07}. Although the  experiments have already started characterizing    E-polarization    anisotropies    \\citep{dasi_pol_02, wmap_1st_pol, wmap_2nd_pol,boom_ee},  the   errors   are  large   in comparison to temperature anisotropy. The upcoming experiments such as Planck satellite will characterize E-polarization anisotropy to a higher accuracy. It is very timely to develop the tools which can optimally utilize the combined CMB temperature and E-polarization information to constrain models of the early universe.  Throughout this paper we use the standard Lambda CDM cosmology with the following cosmological parameters: $\\Omega_b = 0.042$, $\\Omega_{cdm} = 0.239$, $\\Omega_L = 0.719$, $h = 0.73$, $n = 1$, and $\\tau = 0.09$. For all of our simulations we used HEALPix maps with $N_{pix}\\approx 3\\times 10^6$ pixels.  \\subsection{Generalized Bispectrum Estimator of Primordial Non-Gaussianity} \\label{estimator_generalized} In our recent paper~\\citep{YKW07} we  described  a   fast  cubic (bispectrum)  estimator  of $f_{NL}$,  using  a  combined analysis  of the temperature and E-polarization  observations.  The estimator was optimal for homogeneous noise, where optimality was defined  by saturation of the Fisher matrix bound.  In this paper we generalize our previous estimator of $f_{NL}$ to deal more optimally with a partial sky coverage and the inhomogeneous noise. The generalization is done in an analogous way to how~\\cite{Creminelli_wmap1} generalized the temperature only estimator developed by \\citet{KSW05}; however, the final result of \\cite{Creminelli_wmap1} (their Eq.~(30)) is off by a factor of two, which results in the error in $f_{NL}$ that is much larger than the Fisher matrix prediction, as we shall show below.  The fast bispectrum estimator of $f_{NL}$ from the combined CMB temperature and E-polarization data can be written as $\\hat{f}_{NL} = \\frac{\\hat{S}_{prim}}{N}$, where \\citep{YKW07} \\begin{eqnarray} \\label{s_prim} \\hat{S}_{prim}=\\frac{1}{f_{sky}}\\int     r^2dr     \\int     d^2\\hat{n}\\; B^2(\\hat{n},r) A (\\hat{n},r), \\end{eqnarray}  \\begin{eqnarray} N=  \\sum_{i j  k  p q r}\\sum_{2 \\le \\ell_1\\le  \\ell_2\\le  \\ell_3} \\frac{1}{\\Delta_{\\ell_1 \\ell_2 \\ell_3}} B^{p q  r,prim}_{\\ell_1  \\ell_2 \\ell_3}(C^{-1})^{ip}_{\\ell_1}(C^{-1})^{j         q}_{\\ell_2}(C^{-1})^{k r}_{\\ell_3} B^{i j k,prim}_{\\ell_1 \\ell_2 \\ell_3 }, \\end{eqnarray}  \\begin{eqnarray} \\label{B} B(\\hat{n},r)\\equiv           \\sum_{ip}\\sum_{lm}(C^{-1})^{ip}a^{i}_{\\ell m}\\beta^p_{\\ell}(r)Y_{\\ell m}(\\hat{n}), \\end{eqnarray} \\begin{eqnarray} \\label{A} A(\\hat{n},r)\\equiv           \\sum_{ip}\\sum_{lm}(C^{-1})^{ip}a^{i}_{\\ell m}\\alpha^p_{\\ell}(r)Y_{\\ell m}(\\hat{n}), \\end{eqnarray}  \\begin{eqnarray}\\beta^i_{\\ell}(r)=\\frac  {2b^i_{\\ell}}{\\pi} \\int k^2  dk\\; P_{\\phi}(k)g^{i}_{\\ell}(k)\\, j_{\\ell}(kr), \\end{eqnarray}  \\begin{eqnarray}\\alpha^i_{\\ell}(r)=\\frac  {2b^i_{\\ell}}{\\pi} \\int k^2  dk\\; g^{i}_{\\ell}(k)\\, j_{\\ell}(kr), \\end{eqnarray} and $f_{sky}$ is a  fraction of the sky observed. Indices $i, j, k, p, q$ and $r$ can  either be  $T$ or  $E$. Here, $\\Delta_{\\ell_1  \\ell_2 \\ell_3}$ is 1 when  $\\ell_1 \\neq \\ell_2 \\neq \\ell_3$, 6 when $\\ell_1 =  \\ell_2 = \\ell_3$, and 2 otherwise, $B^{p q r,prim}_{\\ell_1 \\ell_2  \\ell_3}$ is the theoretical bispectrum for $f_{NL}=1$~\\citep{YKW07}, $P_{\\Phi}(k)$ is  the  power  spectrum of  the  primordial  curvature perturbations, and $g^i_{\\ell}(r)$ is the radiation transfer function of adiabatic perturbations.  It has been shown that the above mentioned estimator is optimal for the full sky coverage and homogeneous noise \\citep{YKW07}. To be able to deal with the realistic data, the estimator has to be able to deal with the inhomogeneous noise and foreground masks.  The estimator can be generalized to deal with a partial sky coverage and the inhomogeneous noise by adding a linear term to $\\hat{S}_{prim}$: $\\hat{S}_{prim}\\rightarrow \\hat{S}_{prim}+\\hat{S}_{prim}^{linear}$. For the temperature only case, this has been done in~\\cite{Creminelli_wmap1}. Following the same argument, we find that the linear term for the combined analysis of CMB temperature and polarization data is given by \\begin{eqnarray} \\label{slin} \\hat{S}^{linear}_{prim}=\\frac{-1}{f_{sky}}\\int  r^2dr  \\int d^2\\hat{n} \\left   \\{  2 B(\\hat{n},r)\\, \\langle A_{sim}(\\hat{n},r)B_{sim}(\\hat{n},r)\\rangle_{MC}  +     A(\\hat{n},r)\\, \\langle B^2_{sim}(\\hat{n},r) \\rangle_{MC} \\right \\}, \\end{eqnarray} where  $A_{sim}(\\hat{n},r)$ and $B_{sim}(\\hat{n},r)$ are the $A$ and $B$ maps generated from Monte Carlo simulations that contain signal and noise,  and $\\langle ..\\rangle$ denotes the average over the Monte Carlo simulations.  The generalized estimator is given by \\begin{equation} \\hat{f}_{NL} = \\frac{\\hat{S}_{prim}+\\hat{S}^{linear}_{prim}}{N}, \\end{equation} which is the main result of this paper. Note that $\\langle\\hat{S}^{linear}_{prim}\\rangle_{MC}=-\\langle\\hat{S}_{prim}\\rangle_{MC}$, and this relation also holds for the equilateral shape. Therefore, it is straightforward to find the generalized estimator for the equilateral shape: first, find the cubic estimator of the equilateral shape, $\\hat{S}_{equilateral}$, and take the Monte Carlo average, $\\langle\\hat{S}_{equilateral}\\rangle_{MC}$. Let us suppose that $\\hat{S}_{equilateral}$ contains terms in the form of $ABC$, where $A$, $B$, and $C$ are some filtered maps. Use the Wick's theorem to re-write the average of a cubic product as $\\langle ABC\\rangle_{MC}= \\langle A\\rangle_{MC}\\langle BC\\rangle_{MC} +\\langle B\\rangle_{MC}\\langle AC\\rangle_{MC} +\\langle C\\rangle_{MC}\\langle AB\\rangle_{MC}$. Finally, remove the MC average from single maps, and replace maps in the product with the simulated maps: $\\langle A\\rangle_{MC}\\langle BC\\rangle_{MC} +\\langle B\\rangle_{MC}\\langle AC\\rangle_{MC} +\\langle C\\rangle_{MC}\\langle AB\\rangle_{MC} \\rightarrow A\\langle B_{sim}C_{sim}\\rangle_{MC} +B\\langle A_{sim}C_{sim}\\rangle_{MC} +C\\langle A_{sim}B_{sim}\\rangle_{MC}$. This operation gives the correct expression for the linear term, both for the local form and the equilateral form.  One can find the estimator of $f_{NL}$ from the temperature data only by setting $i=j=k=p=q=r\\equiv T$. We have compared our formula in the temperature-only limit with  the original formula derived by~\\citet{Creminelli_wmap1} (their Eq.~(30)), and found a discrepancy. To see the discrepancy, let us re-write the estimator as: $\\hat{f}_{NL} = \\frac{\\hat{S}_{prim}+ x \\hat{S}^{linear}_{prim}}{N}$. Our formula gives $x=1$, while Eq.~(30) of \\citet{Creminelli_wmap1} gives $x=0.5$. \\footnote[1]{Equation~(30) in~\\citet{Creminelli_wmap1} is off by a factor  of 2. Since we used  $\\sum_{\\ell_1 \\ell_2 \\ell_3}  = 6\\sum_{\\ell_1 \\le \\ell_2  \\le  \\ell_3} \\frac{1}{\\Delta_{\\ell_1 \\ell_2 \\ell_3}}$, to compare our normalization factor $x$ with~\\citet{Creminelli_wmap1} one needs to divide the normalization in~\\citet{Creminelli_wmap1} by a factor of $6$.}  To make sure that our normalization gives the minimum variance estimator, we have done Monte Carlo simulations with varying $x$. We find that $x=1$ minimizes the variance (as shown in Fig.~\\ref{normalization}). We conclude that the analysis given in \\citet{Creminelli_wmap1} resulted in the larger-than-expected uncertainty in $f_{NL}$ because of this error in their normalization of the linear term.   \\begin{figure*}[t] \\begin{center} \\includegraphics[height=8cm,angle=0]{f1.eps} \\caption{Testing normalization of the linear term in the estimator of $f_{NL}$. The symbols show the standard deviation of $f_{NL}$ derived from the the Monte Carlo simulations using the estimator for a given normalization, $x$. The horizontal line shows the Fisher matrix prediction. Our formula gives $x=1$, while \\citet{Creminelli_wmap1} give $x=0.5$ (their Eq.~(30)). We have used simulated polarized Gaussian CMB maps with the Planck inhomogeneous noise as well as the WMAP Kp0 and P06 mask for temperature and polarization, respectively.} \\label{normalization} \\end{center} \\end{figure*}   The main  contribution to the linear term  comes from  the inhomogeneous noise and sky cut. For the temperature only case, most of the contribution to the linear term comes from the inhomogeneous noise, and the partial sky coverage does not contribute much to the linear term. This is because the sky-cut induces a monopole contribution outside the mask. In the analysis one subtracts the monopole from outside the mask before measuring $\\hat{S}_{prim}$, which makes the linear contribution from the mask small \\citep{Creminelli_wmap1}. For a combined analysis of the temperature and polarization maps, however, the linear term does  get a significant contribution from a partial sky coverage (see the right panel of Fig.~\\ref{fnl_stdev}). Subtraction of the monopole outside of the mask is of no help for polarization, as the monopole does not exist in the polarization maps by definition. (The lowest relevant multipole for polarization is $l=2$.)   \\begin{figure*}[t] \\begin{center} \\includegraphics[height=5.5cm,angle=0]{f2a.eps} \\includegraphics[height=5.5cm,angle=0]{f2b.eps} \\caption{Optimality of the generalized estimator. The solid  lines  show  the  Fisher matrix prediction for the standard deviation of $f_{NL}$, the triangles show the standard   deviation derived from the  Monte    Carlo simulations using the estimator without the linear term, and the stars show the  standard   deviation  derived from the  Monte    Carlo simulations using the generalized estimator (i.e., with the linear term). {\\bf Left panel:} The uncertainty vs the maximum multipole that is used in the analysis, $\\ell_{max}$. The simulations contain the Gaussian CMB signal, inhomogeneous noise (which simulates the Planck satellite), WMAP Kp0 and P06 masks. {\\bf Right panel:} The uncertainty vs a fraction of the sky observed, $f_{sky}$, for $\\ell_{max}=500$. The simulations include the Gaussian CMB signal, and flat sky-cut (which is azimuthally symmetric in the Galactic coordinates), while they do not include instrumental noise. This  figure therefore shows that the sky cut contributes significantly to the linear term of polarization. } \\label{fnl_stdev} \\end{center} \\end{figure*}   The estimator is still computationally efficient, taking only $N^{3/2}_{pix}$ (times the $r$ sampling, which is of order 100) operations in comparison  to the full bispectrum calculation which takes $N^{5/2}_{pix}$ operations. Here  $N_{pix}$ refers to the total number of  pixels.  For Planck,  $N_{pix}  \\sim  5 \\times  10^7$,  and  so the  full bispectrum analysis is not feasible while our analysis is. ",
        "conclusions": "Upcoming CMB experiments will provide a wealth of information about the CMB polarization anisotropies together with temperature anisotropies. The combined information from the CMB temperature and polarization data improves the sensitivity to primordial non-Gaussianity~\\citep{BZ04,YW05,YKW07}. The promise of learning about the early universe by constraining the amplitude of primordial non-Gaussianity is now well established. In this paper we have generalized the bispectrum estimator of non-Gaussianity described in~\\cite{YKW07}, to deal with the inhomogeneous nature of noise and incomplete sky coverage.  The generalization from \\cite{YKW07} enables us to increase optimality of the estimator significantly, without compromising the computational efficiency of the estimator: the estimator is still computationally efficient, scaling as $O(N_{pix}^{3/2})$ compared to the $O(N_{pix}^{5/2})$ scaling of the full bispectrum~\\citep{BZ04} calculation for sky maps with $N_{pix}$ pixels. For the Planck satellite, this translates into a speed-up by factors of millions, reducing the required computing time from thousands of years to just hours and thus making $f_{NL}$ estimation feasible. The speed of our estimator allows us to study its statistical properties using Monte Carlo simulations.  We have used Gaussian and non-Gaussian simulations to characterize the estimator. We have shown that the generalized fast estimator is able to deal with the partial sky coverage very well and in fact the variance of $f_{NL}$ saturates the Fisher matrix bound. In the presence of both the realistic noise and galactic mask, we find that the generalized estimator greatly reduces the variance in comparison to the~\\cite{YKW07} estimator of non-Gaussianity using combined CMB temperature and polarization data.  Since the estimator is able to deal with the partial sky coverage very effectively, the estimator can also be used to constrain primordial non-Gaussianity using the data from ground and balloon based CMB experiments which observe only a small fraction of the sky. The estimator also solves the problem~\\citep{YKW07} of non-trivial polarization mode coupling due to foreground masks. Earlier this issue was dealt with by removing the most contaminated $\\ell$ modes from the analysis (usually $\\ell <30 $).  The naive approach of using galactic masks to deal with the polarization contamination is to be refined. Both temperature and polarization foregrounds are expected to produce non-Gaussian signals. Some sources of non-primordial non-Gaussianity are CMB lensing, point sources, and the Sunyaev Zel'dovich effect. Understanding the non-Gaussianity from the polarization foreground sources and refining the estimator to be able to deal with it will be the subject of our future work."
    },
    "0711/0711.4334_arXiv.txt": {
        "abstract": "The hydrogen ionization and dissociation front around an ultraviolet radiation source should merge when the ratio of ionizing photon flux to gas density is sufficiently low and the spectrum is sufficiently hard. This regime is particularly relevant to the molecular knots that are commonly found in evolved planetary nebulae, such as the Helix Nebula, where traditional models of photodissociation regions have proved unable to explain the high observed luminosity in H$_2$ lines. In this paper we present results for the structure and steady-state dynamics of such advection-dominated merged fronts, calculated using the Cloudy plasma/molecular physics code. We find that the principal destruction processes for H$_2$ are photoionization by extreme ultraviolet radiation and charge exchange reactions with protons, both of which form H$_2{}^+$, which rapidly combines with free electrons to undergo dissociative recombination. Advection moves the dissociation front to lower column densities than in the static case, which vastly increases the heating in the partially molecular gas due to photoionization of He$^0$, H$_2$, and H$^0$. This causes a significant fraction of the incident bolometric flux to be re-radiated as thermally excited infrared \\Hmol{} lines, with the lower excitation pure rotational lines arising in $1000$~K gas and higher excitation \\Hmol{} lines arising in $2000$~K gas, as is required to explain the H$_2$ spectrum of the Helix cometary knots. ",
        "introduction": "\\label{sec:introduction}  The ultraviolet photons from hot stars, such as main-sequence OB stars or the central stars of planetary nebula (CSPN) will dissociate and ionize the surrounding circumstellar and interstellar gas. In the classical picture \\citep[e.g.,][]{1997ARA&A..35..179H}, the extreme ultraviolet (EUV) photons with energies $h\\nu > 13.6~\\mathrm{eV}$ photoionize the hydrogen gas, forming an \\hii{} region bounded by a relatively sharp ionization front (IF), while the far ultraviolet (FUV) photons with energies $6~\\mathrm{eV} < h\\nu < 13.6~\\mathrm{eV}$ penetrate the IF to form a neutral photodissociation region (PDR) that surrounds the \\hii{} region. The dissociation of \\Hmol{} in such a PDR is dominated by a two-step radiative process \\citep{1967ApJ...149L..29S, 2000A&AS..141..297A}, in which absorption of an FUV photon leaves the \\Hmol{} molecule in an electronically excited state, from which it has a certain probability ($\\simeq 10$\\%) of decaying to the vibrational continuum of the ground electronic state.  However, as shown by \\citet{1996ApJ...458..222B}, a classical extended PDR cannot exist if the FUV flux is sufficiently weak compared with the EUV flux at the IF, rather the IF and dissociation front (DF) should merge.  \\citeauthor{1996ApJ...458..222B} considered the case of \\hii{} regions around OB stars and found that this regime was most relevant for cases in which the dust optical depth through the ionized gas is low, which corresponds to a low ionization parameter (the ionization parameter is a dimensionless number equal to the ratio of the number density of ionizing photons to the number density of hydrogen nuclei). To date, no models have been calculated of the structure and emission properties of such fronts, which are also expected to show strong deviations from static chemical and ionization equilibrium.  In this paper, we calculate in detail the internal dynamics and chemistry of merged ionization/dissociation fronts (IF/DFs), concentrating on the particular case of compact photoevaporating molecular knots in evolved planetary nebulae (PNe), such as are seen in the Helix nebula (\\citealp{1999ApJ...522..387Y, 2005AJ....130.1784M, 2006ApJ...652..426H}; \\citealp[OHF07]{2007AJ....133.2343O}). The stellar spectrum from the hot central star of a PN is harder than that from an OB star, leading to an EUV luminosity that is much greater than the FUV luminosity. In addition, the ionization parameter of the knots is much lower than is typically seen in \\hii{} regions, resulting in very little attenuation of the EUV flux by recombinations in the ionized gas. Both these factors place the knots firmly in the merged IF/DF regime. The most comprehensive existing theoretical study of PDRs in planetary nebulae is that of \\citet{1998A&A...337..517N}, who present detailed modeling of the time-dependent evolution of an expanding circumstellar shell as the luminosity and spectrum of the CSPN evolves, finding that soft X-rays can be important in exciting the molecular gas at late times. However, \\citetalias{2007AJ....133.2343O} showed that this is not the case in the Helix since it is only in the EUV band that the CSPN luminosity is sufficient to excite the knot PDRs. \\citeauthor{1998A&A...337..517N} used an analytic treatment of the \\hii{} region, which is assumed to absorb all the EUV radiation, and were thus unable to model the case of a merged IF/DF. ",
        "conclusions": "\\label{sec:discussion}  Observations of the molecular hydrogen spectrum of the knots in the Helix Nebula (\\citealp{1998ApJ...495L..23C}; \\citetalias{2006ApJ...652..426H}; \\citetalias{2007AJ....133.2343O}) are indicated as error bars on Figure~\\ref{fig:excit}. It can be seen that the two advective models are in broad agreement with the observations, whereas the static model is not. Model~C06 best matches the distance of the spectroscopically observed knots from the ionizing star, and indeed shows a better agreement than Model~A06, which corresponds to the closer-in knots. The observed nebular flux in the 1--0 $S(1)$ line and in the sum of the 0--0 $S(1)$ to $S(7)$ lines are $\\simeq 1\\%$ and $\\simeq 4\\%$, respectively, of the bolometric stellar flux \\citepalias{2007AJ....133.2343O}. From Figure~\\ref{fig:lumfrac}, this can be satisfied by our models with $\\lambda_\\mathrm{ad} > 0.3/f$, where $f$ is the area covering fraction of knots. A strong prediction of our model is that higher excitation lines from the $v \\ge 1$ levels should show a higher population temperature of $\\simeq 2000$~K. A recent study of an inner knot \\citep{2007arXiv0709.3065M} finds a population temperature of about 1800~K for these levels, in agreement with the prediction of our relevant model (A06).  Ro-vibrationally warm \\Hmol{} has been detected in other PNe \\citep[e.g.,][]{2006AJ....131.1515L, 2007ApJ...659.1291M} and has frequently been interpreted as evidence for shocks \\citep{1988ApJ...324..501Z}. However, EUV-dominated advective PDRs may be a promising alternative in these cases too.  \\vspace*{-\\baselineskip}"
    },
    "0711/0711.2042_arXiv.txt": {
        "abstract": "% One of the major challenges providing large databases like the WFCAM Science Archive (WSA) is to minimize ingest times for pixel/image metadata and catalogue data. In this article we describe how the pipeline processed data are ingested into the database as the first stage in building a release database which will be succeeded by advanced processing (source merging, seaming, detection quality flagging etc.). To accomplish the ingestion procedure as fast as possible we use a mixed Python/C++ environment and run the required tasks in a simple parallel modus operandi where the data are split into daily chunks and then processed on different computers. The created data files can be ingested into the database immediately as they are available. This flexible way of handling the data allows the most usage of the available CPUs as the comparison with sequential processing shows. ",
        "introduction": "The WFCAM Science Archive (\\htmladdnormallinkfoot{WSA}{http://surveys.roe.ac.uk/wsa/}; Hambly et al. 2007, Collins et al. 2006) holds the image and catalogue data products generated by the Wide Field Camera (WFCAM) on UKIRT (United Kingdom Infrared Telescope). The data comprise pipeline processed multi-extension FITS files (multiframes) containing pixel/image and catalogue data for four detectors at one pointing. The latter contains all detections of stacked multiframes. The data are pipeline processed at the Cambridge Astronomical Survey Unit (CASU) and transferred to Edinburgh where the Wide Field Astronomy Unit (WFAU) processes it for ingestion into the database. Since the release database contains advanced products which can take a lot of CPU time to produce, it is preferable to carry out the ingest procedure as fast as possible. Another aspect is that the pixel/image and catalogue data need to be ingested completely before further processing is done. Another constraint is the uniqueness of multiframes and detections. Each multiframe has to have a unique identifier across the whole database and each detection must be unique for a given survey. ",
        "conclusions": ""
    },
    "0711/0711.3758_arXiv.txt": {
        "abstract": "{\\\\ \\emph{Aims.} We seek to probe the Galactic bulge IMF starting from microlensing observations.\\\\ \\emph{Methods.} We analyse the recent results of the microlensing campaigns carried out towards the Galactic bulge presented by the  EROS, MACHO and OGLE collaborations. In particular, we study the duration distribution of the events. We assume a power law initial mass function, $\\xi(\\mu)\\propto \\mu^{-\\alpha}$, and we study the slope $\\alpha$ both in the brown dwarf and in the main sequence ranges. Moreover, we compare the observed and expected optical depth profiles. \\\\ \\emph{Results.} The values of the mass function slopes are strongly driven by the observed timescales of the microlensing events. The analysis of the MACHO data set gives, for the main sequence stars, $\\alpha=1.7 \\pm 0.5$, compatible with the result we obtain with the EROS and OGLE data sets, and a similar,  though less constrained slope for brown dwarfs. The lack of short duration events in both  EROS and OGLE data sets, on the other hand, only allows the determination of an \\emph{upper} limit in this range of masses, making the overall result less robust. The optical depth analysis gives a very good agreement between the observed and the expected values, and we show that the available data do not allow one to discriminate between different bulge models. ",
        "introduction": "\\label{sec:intro}   Gravitational microlensing is an established tool for the study and the characterisation of faint compact objects located between the observer and the source stars. It was originally proposed as a tool for the detection of dark matter in the form of MACHOs \\citep{pacz86}. Searches towards the Magellanic Clouds by the MACHO \\citep{macho00} and the EROS groups \\citep{eros07} have placed strong constraints on the possible contribution of a MACHO population to the dark matter halo (for a discussion see, e.g., \\citealt{mancini04,novati06}). A few results have also been obtained with observational campaigns towards M31 by the POINT-AGAPE \\citep{novati05} and the MEGA \\citep{mega06} collaborations. On the other hand, the Galactic bulge soon proved to be an almost as interesting target, if not more \\citep{pacz91,kiraga94} and indeed, by now, the number of observed microlensing events along this line of sight is by two orders of magnitude larger than those observed towards the Magellanic Clouds and M31. In this case, any contribution from a dark matter MACHO population is expected to be extremely small compared to that of either bulge or disc stars \\citep{griest91b}, so that these studies in principle allow us to constrain the inner Galactic structure. In particular, microlensing observations in this direction have been very important for the assessment of the Galactic triaxial, bar-like, structure \\citep{pacz94,zhao95,zhao96a,zhao96b,bissantz97,gyuk99}.  Recently, the MACHO \\citep{popowski05}, OGLE  \\citep{sumi06} and EROS \\citep{hamadache06} collaborations presented the results out of their several-year campaigns towards the Galactic bulge. A remarkable result is the agreement among the different collaborations for the \\emph{optical depth}, in accord with theoretical expectations \\citep{evans02,bissantz02,hangould03}.  While the determination of the optical depth allows the study of the density distribution of the lenses, a more detailed analysis of the shape of the microlensing lightcurves carries much information on the parameters of the lenses. Of particular relevance is the possibility of studying the \\emph{mass spectrum} of the lenses. This approach  is based on the relationship between the observed event duration, the Einstein time, and the mass of the lens $t_\\mathrm{E} \\propto \\sqrt{\\mu_l}$. Even if the exact analytical formula shows a dependence also on other unknown physical parameters such as the distances of both lens and source, the relative velocity between them as well as the configuration of the particular lens event, a few conclusions are made possible by the rather large set of observed events at our disposal together with a few reasonable assumptions on the space and velocity distributions of both lenses and sources. A key step to reach the aforementioned agreement of the optical depth between theory and experiments has been the acknowledgment of the severe \\emph{blending} problem resulting in the choice of restricting the sample of source stars to the \\emph{red clump giant} subset \\citep{popowski01}. In turn, this is essential in the framework of a mass spectrum analysis because of the bias introduced in the evaluation of the Einstein time for blended events.  The determination of the mass function using the results of microlensing searches has been addressed by several authors. \\cite{hangould96} consider a sample of MACHO and OGLE events. Through a likelihood analysis they determine the slope of a power law mass function to be 2.1 in the mass range $(0.04-10)~\\mathrm{M}_\\odot$. \\cite{jetzer94b} and \\cite{grenacher99} use the mass moments method to place constraints on the lens masses. In particular, starting from a sample of 41 MACHO events and assuming a Salpeter profile in the mass range $(1-10)~\\mathrm{M}_\\odot$, \\cite{grenacher99} constrain the mass function minimum mass and slope below $1~\\mathrm{M}_\\odot$, finding $0.012~\\mathrm{M}_\\odot$ and $2.0$ respectively. Overall, therefore, there is an agreement in attributing a rather large fraction of events to the brown dwarf lens population. On the other hand, \\cite{peale98} finds no compelling evidence for such a contribution, and evaluates the slope for a power law mass function in the mass range $(0.08-2)~\\mathrm{M}_\\odot$ to be in the range $2.2-2.5$. All of these analyses, we recall, used the complete sample of detected events, not restricted to those with red clump giant sources. \\cite{bissantz04}, considering only red clump giant sources, find a good agreement with the MACHO observed timescale using a mass function with a large contribution from the brown dwarf population (with a power law slope 2.35 in the mass range $(0.04-0.35)~\\mathrm{M}_\\odot$). \\cite{woodmao05} extend the \\cite{zoccali00} slope 1.3 down to a minimum mass of 0.$03~\\mathrm{M}_\\odot$ compared with the OGLE observed timescale.  In the present paper our aim is to make use of the most recent observational results towards the Galactic bulge to study the mass spectrum of the \\emph{bulge} lens population. The structure of the paper is as follows. Sect.~\\ref{sec:models} is devoted to the description of the models we use. In Sect.~\\ref{sec:anaml} we point out a few particular features of the usual microlensing quantities upon which we base our method of analysis. In Sect.~\\ref{sec:res} we present and discuss our main results. In Sect.~\\ref{sec:end} we conclude. ",
        "conclusions": "\\label{sec:end}  We have considered the sample of microlensing events observed towards the Galactic bulge with red clump giant sources reported by the  the MACHO \\citep{popowski05}, OGLE  \\citep{sumi06} and EROS \\citep{hamadache06} collaborations to place constraints on the bulge mass function. In particular, through a likelihood analysis, we have studied the slopes, $\\alpha_\\mathrm{BD},\\alpha_\\mathrm{MS}$, of a power law mass function in the brown dwarf $(0.01-0.08)~\\mathrm{M}_\\odot$ and main sequence $(0.08-1.)~\\mathrm{M}_\\odot$ mass ranges.  For our fiducial model, comparing to the CGR sample of 42 MACHO events, we obtain $\\alpha_\\mathrm{MS}=1.7\\pm 0.5$. This result compares well to that obtained in the $(0.15-1)~\\mathrm{M}_\\odot$ range by \\cite{zoccali00}, $\\alpha\\sim 1.3$. The slope in the brown dwarf range turns out to be less well constrained, $\\alpha_\\mathrm{BD}=1.6\\pm 1$. Overall our maximum likelihood results indicate a rather significant contribution of low mass lenses, with $\\sim~30\\%$ of the events to be attributed to brown dwarfs. The last result is in agreement with previous analyses \\citep{hangould96,grenacher99}, while our value for the main sequence slope, somewhat smaller, may be explained because we are using a more suitable sample of red-clump-source events.  The analyses of the EROS and OGLE data sets give us somewhat different results. We derive a smaller value for the slope in the main sequence range, although compatible with the MACHO data set result, but we obtain only an upper limit for the slope of the brown dwarf population. This behaviour finds its explanation in the different observed timescales. In particular, very short timescale events ($t_\\mathrm{E}<5~\\mathrm{d}$), only observed in the MACHO data set, are essential to constrain the brown dwarf mass function. The lack of short timescale events has been noticed in the EROS analysis \\citep{hamadache06}. In all of the experiments, the detection efficiency of short durations events is extremely low, rendering the analysis in the brown dwarf regime difficult and, therefore, making the result less robust.  More reliable constraints on the mass function may come from a better understanding of the bulge model, but especially, as already stressed , by improving the statistics of observed short duration events.  Furthermore, we have carried out an analysis on the optical depth. The agreement with the expected values is recognized \\citep{hangould03}. Here we have considered the profile of the expected optical depth as compared to the observed one, finding a good agreement for both the models we have considered, the model E2 of \\cite{stanek97} and the model G2 of \\cite{dwek95}. To further constrain the bulge profile it would be useful to extend microlensing searches to cover a larger area in the sky plane."
    },
    "0711/0711.3791_arXiv.txt": {
        "abstract": "Dark matter (DM) halos formed in CDM cosmologies seem to be characterized by  a  power law phase- space density profile. The density of the DM halos is   often fitted by the   NFW profile but a better fit is provided by the   Sersic fitting formula. These relations are empirically derived from cosmological simulations of structure formation but have not yet been explained on a first principle basis. Here we solve   the Jeans equation under the assumption of a spherical DM halo in  dynamical equilibrium, that obeys a power law phase space density and either the NFW-like or the Sersic density profile. We then calculate the velocity anisotropy, $\\beta(r)$, analytically. Our main result is that for the NFW-like  profile the $\\beta ~ - ~ \\gamma$ relation is not a linear one (where $\\gamma$ is the logarithmic derivative of the density $\\rho[r]$). The shape of $\\beta(r)$ depends mostly on the ratio of the gravitational to kinetic energy within the NFW scale radius $R_{\\rm s}$. For the Sersic profile   a linear  $\\beta ~ - ~ \\gamma$ relation is recovered, and in particular for  the Sersic index of $n\\approx 6.0$ case the linear fit of Hansen \\& Moore is reproduced. Our main result is that the phase-space density power law, the Sersic density form and the linear $\\beta ~ - ~ \\gamma$ dependence constitute a consistent set of relations which obey the spherical Jeans equation and as such provide the framework for the dynamical modeling of DM halos. ",
        "introduction": "\\label{sec:intro}  In the standard cosmological model of structure formation the luminous matter is embedded in extended dark matter (DM) halos. The large scale structure emerges out of the primordial perturbation field \\via\\ gravitational instability. The model further assumes the DM  to be made of weakly interacting particles and its dynamics to be collisionless and therefore dissipationless (\\citeauthor{pad93} \\citeyear{pad93} for a review).  The problem of the dynamics of DM halos can be formulated as the classical gravitational $N$-body problem, subject to the assumption of cosmological initial and boundary conditions. As such the problem can be very easily formulated, yet it defies any rigorous analytical treatment. The spherical top-hat model provides the main analytical tool for shading light on the problem, but its scope of validity is rather limited \\citep{gg72}. The model can be extended to accommodate shell crossing  \\citep{gunn77} and cosmological initial conditions \\citep{hs85} resulting in the secondary infall model.  In the absence of a rigorous analytical theory the study of the evolution DM halos relies heavily on numerical $N$-body simulations. The advent of CPU power and improved numerical algorithm have led to a general consensus about the basic properties of DM halo, such as the spherically-averaged density profile, the spin and shape of halos. One of the pillars of the phenomenology of DM halos is the so-called NFW density profile (\\citeauthor{nfw96}, \\citeyear{nfw97}), \\begin{equation} \\label{eq:NFW} \\rho_{\\rm NFW}(r) = { 4 \\rho_{\\rm s} R_{\\rm s}^3\\over r (R_{\\rm s} + r)^2}. \\end{equation} It has been argued the NFW two-parameters fitting formula provides a good approximation for DM halos found in a wide range of mass scales and cosmological models, and hence can be considered universal. Subsequent numerical simulations have basically confirmed the functional form of the NFW profile but some controversy concerning the asymptotic slopes of the profile at small and large radii persists (\\eg\\  \\citeauthor{moore98} \\citeyear{moore98},  \\citeauthor{kly01} \\citeyear{kly01}, \\citeauthor{jing00} \\citeyear{jing00}). Here we refer to these generalizations of the NFW fitting formula as the gNFW profile.  A second basic phenomenological finding is the so-called phase-space density (PSD) profile \\citep{tay01}. These authors defined the PSD profile by \\begin{equation} \\label{eq:phase-space} Q(r)=\\frac{\\rho}{\\sigma_{\\rm r}^3}, \\end{equation} where $\\sigma_{\\rm r}$ is the radial velocity dispersion. The basic finding of \\citet{tay01} is that the PSD profile follows a power law of the form \\begin{equation} \\label{eq:Q(r)} Q(r) \\propto r^{-\\alpha}. \\end{equation} \\citet{tay01} found $\\alpha\\approx1.875$, and more recent studies find $\\alpha=1.92 \\pm 0.01$ \\citep{DM05} and $1.94 \\pm 0.01$ \\citep{HRSH07}.  Again, the power law behavior of $Q(r)$  is found over a wide range of mass scales and different cosmological models, suggesting a universal property of DM halos.  It has been recently suggested that a universal density slope- velocity anisotropy relation exists for relaxed DM halos \\citep{HM06}. These author found a linear relation between the velocity anisotropy parameter $\\beta$ and the density slope $\\gamma$ ($<0$), \\begin{equation} \\label{eq:linear} \\beta = - 0.2(\\gamma + 0.8), \\end{equation} where \\begin{equation} \\label{eq:beta} \\beta(r)=1-\\frac{\\sigma{_\\theta^2}+\\sigma{_\\phi^2}}{2\\sigma{_{\\rm r}^2}} \\end{equation} and \\begin{equation} \\label{eq:gamma} \\gamma={ r\\over \\rho} { d\\  \\rho \\over  d\\  r}, \\end{equation} where $\\sigma_\\theta$ and $\\sigma_\\phi$  are the dispersions of the two transversal velocity components. Again, the linear relation has been found for a variety of halos obtained in various simulations and setups, suggesting a universal nature of the linear relation.  The phenomenological gNFW density profile, the PSD power law and the linear $\\beta-\\gamma$ relation are supposed to provide us clues about DM halos. The universal nature of these relations suggests that they hold over a broad range of scales and models, at least within the parameter space of the  Cold Dark Matter (CDM)-like cosmogonies. The equilibrium structure of collisionless self-gravitating systems obey the Jeans equation (\\eg\\ \\citeauthor{BT} \\citeyear{BT}), which relates the velocity second moments and the density field. Under the assumption of a spherical symmetry, the Jeans equation relates $\\rho(r)$, $\\sigma_{\\rm r}(r)$ and $\\beta(r)$, and can be rewritten so as  to relate $\\rho(r)$, $Q(r)$ and $\\beta(r)$. Lacking a fundamental theory that can predict even one of the above relations we are motivated to study the internal consistency of the three conditions and to find out whether one of these can be found to depend on the other two. Given the large scatter around the linear $\\beta-\\gamma$ relation (\\citeauthor{HM06} \\citeyear{HM06}, \\citeauthor{HS06} \\citeyear{HS06}) we suspect that the $\\beta -  \\gamma$ relation is the 'weakest' and less certain among the three relations. We therefore assume that the density of spherical halos follows an exact gNFW  profile and that the PSD profile is a power law and solve the Jeans equation to obtain the $\\beta$ profile. This is to be compared with the linear $\\beta - \\gamma$ relation.   It has been recently suggested that the so-called  Sersic, also known as the    \\cite{ein65}, profile  provides a better fit to the density profile \\citep{mer05,prad06,gao07}.   Given the very different functional form of the Einaso profile it will be considered here as an alternative to the NFW-like family of profiles and the associated  $\\beta(r)$ will be calculated. This will enable us to check the sensitivity of $\\beta(r)$ to the assumed density fitting formula.   The structure of the paper is as follows. The Jeans equation is solved in \\S \\ref{sec:Jeans} and the solutions of the $\\beta$ profile are given in \\S \\ref{sec:res} and discussed in \\S \\ref{sec:disc}. ",
        "conclusions": "\\label{sec:disc}  The present work has been primarily motivated by the quest for understanding the linear $\\beta - \\gamma$ dependence  \\citep{HM06} and its relationship to the two so-called  universal relations, the PSD power law  \\citep{tay01} and the gNFW and the Sersic density profiles, that characterize the structure of DM halos. Here we have focused on whether the linear $\\beta - \\gamma$ relation is implied by the other two relations, and is it consistent with them.   We find that for no choice of parameters  an NFW-like density profile can yield a velocity anisotropy,  $\\beta(r)$, that is even in a rough agreement with the linear  $\\beta - \\gamma$. This stands in a sharp contrast with the Sersic fitting formula for the DM halos density profile. For a  Sersic profile index of   $n \\approx 6.0$ one recovers quite faithfully the linear $\\beta - \\gamma$ relation. For that model $\\beta(r)$ is a monotonically increasing function of the radius. For the nominal model we find $\\beta(x =0.1)=0.12$ and $\\beta(x=10)=0.33$. A by-product of the solution of the Jeans equation is that the PSD of the total velocity dispersion follows a power law, namely $Q_{\\rm tot}(r) \\propto r^{-\\alpha_{\\rm tot}}$, with $\\alpha_{\\rm tot}=1.84$, for both the gNFW and the Sersic density profiles. Yet, in the  fractional deviation in NFW-like case is of the order of twenty percents and in the Sersic case the deviation is of the order of one percent.   It has been realized in recent years that the Sersic profile provides a better fit to the density profile of DM halos  than the NFw  model \\citep{mer05,mer06,prad06,gao07}. The present work substantiates and strengthens that fact and it strongly suggests that the the Sersic model should be used for the dynamical modeling of DM halos.   The main result of the papers is that the three pillars of the DM halos  phenomenology, namely the PSD power law, the Sersic density profile and the linear $\\beta-\\gamma$ relation, constitute a consistent set of relations that obey the Jeans equation. These relations provide  a theoretical framework for a consistent dynamical modeling of DM halos.  A very different motivation for the calculation of the $\\beta$ profile has been to provide a practical tool for modeling the mass distribution in clusters of galaxies from kinematic data. A powerful way of modeling  the clusters of galaxies is based on taking moments of the velocity distribution of clusters' galaxies and fitting them to the solutions of the Jeans equation, under the assumptions of the NFW density profile and $\\beta= const.$ \\citep{LO06,WL07}. The presently calculated $\\beta$ profile certainly  provides a better approximation to the actual profile than the assumption of a constant value. We suggest that future analysis of the mass distribution of clusters will be based on the calculated profile of $\\beta$.  We conclude the paper with a final note of caution. The analysis presented here applies strictly to spherical DM halos in virial equilibrium. Inspection of the formation of DM halos in cosmological simulations of CDM-like cosmologies reveals an ongoing process where halos grow continuously by a slow accretion and by major mergers, alternating between phases of dynamical equilibrium and violent off-equilibrium. Therefore, one expects the DM halos to show some deviations from a strict virial equilibrium and this might limit the validity of the solutions of the Jeans equation. DM halos are not isolated island structures in an otherwise unperturbed Friedmann universe. Halos are experiencing an ongoing smooth, and occasionally no so smooth, accretion and the boundary between the halo and outer universe is not easily defined. The stationary Jeans equation, on the other hand  is applicable to   isolated systems. It follows that even in stationary, seemingly relaxed systems  the Jeans equation may not be strictly obeyed. A further complication arises from the deviation from sphericity of the DM halos. Cosmological simulations give rise to oblate and prolate ellipsoidal halos. This again might introduce a further deviation of the $\\beta$ profile of the DM halos from the solutions of the Jeans equation. These issues should be further investigated by means of N-body simulations."
    },
    "0711/0711.1107_arXiv.txt": {
        "abstract": "We present results of a 3 year monitoring campaign of the Seyfert~1 galaxy Markarian~509, using X-ray data from the {\\em Rossi X-ray Timing Explorer} and optical data taken by the SMARTS consortium.  Both light curves show significant variations, and are strongly correlated with the optical flux leading the X-ray flux by 15 days.  The X-ray power spectrum shows a steep high-frequency slope of $-2.0$, breaking to a slope of $-1.0$ at at timescale of 34 days.  The lag from optical to X-ray emission is most likely caused by variations in the accretion disk propagating inward. ",
        "introduction": "Variability has long been recognized as one of the defining characteristics of active galactic nuclei (AGN).  This variability is not limited to one wavelength regime, but is spread across the entire electromagnetic spectrum.  Produced in the innermost regions of the nucleus are X-ray and optical/UV continuum photons. Current models involve an accretion disk emitting thermal ($10^5$~K) photons in the optical and UV range \\citep{Malkan1983,Shields1978}, which are then inverse Compton scattered to X-ray energies by a corona of hot electrons above the accretion disk \\citep{Haardt1991,Haardt1993}.  By monitoring the variability at optical and X-ray wavelengths, we can learn information about the geometry and exact mechanism responsible for the emission. The traditional view is that because the X-rays are produced closer to the nucleus, variations should be the largest at that wavelength, with optical variations lagging behind.  On the other hand, the variations may originate in the optical seed photons, which then cause the X-ray emission to lag behind the optical.  Previous attempts at correlating X-ray and optical variability in Seyfert galaxies have been inconclusive.  On long timescales, \\citet{Uttley2003} detected a strong correlation in NGC~5548 with no lag, while \\citet{Chiang2000} found a weaker correlation on short timescales.  The same was true for NGC~4051, with \\citet{Peterson2000} finding a correlation on long but not short timescales. However this is not the case with all Seyferts, as \\citet{Maoz2002} found no correlation in NGC~3516 after 5 years of monitoring.  \\citet{Shemmer2003} and \\citet{Arevalo2005} both found correlations on short timescales in NGC~4051 and MCG--6-30-15, respectively, while \\citet{Papadakis2000} found no correlation on short timescales in Akn~564.  For a more comprehensive review of past efforts, see \\citet{Maoz2002} and references therein.  Mkn~509 is a nearby ($z=0.034$) Seyfert~1 galaxy.  First detected in the X-rays by {\\em Ariel V} \\citep{Cooke1978}, Mkn~509 was later found to have a soft X-ray excess by \\citet{Singh1989}.  Simultaneous observations with {\\em GINGA} and {\\em ROSAT} again showed a soft excess, with some flattening at higher energies due to reflection \\citep{Pounds1994}.  Later observations with {\\em ASCA} showed the presence of an Fe~K$\\alpha$ line at 6.4~keV, and the presence of a warm absorber rather than a soft excess \\citep{Reynolds1997}. More recent observations with {\\em XMM-Newton} show the origin of the soft X-ray excess lies mostly in thermal emission from the inner parts of the accretion disk \\citep{Pounds2001}.  We present here results from a 3.5 year X-ray and optical monitoring campaign on Mkn~509.  In \\S2, we show light curves and discuss data analysis methods for all of our observations.  In \\S3, we calculate the X-ray power density spectrum (PDS), using Monte Carlo methods.  In \\S4, we present results of cross-correlation between the X-ray and optical light curves and discuss the significance of the results.  Finally, in \\S5 we compare our result to previous efforts, and discuss possible implications for physical models of accretion. ",
        "conclusions": "We have shown that over long timescales, optical variations lead the X-rays by 15~days for Mkn~509.  Initially, this lag would appear to be far too large for reprocessing models where thermal optical/UV photons from the disk are inverse Compton scattered to X-ray energies \\citep{Haardt1991}.  However, such a lag is not unprecedented.  \\citet{Shemmer2003} found an optical to X-ray lag of 2.1~days in NGC~4051.  Such a lag is most likely related to the viscous or thermal timescale.  In that case, as variations in the accretion flow propagate inwards through the disk, they pass first through the optical emitting region, and then later through the X-ray emitting region closer to the black hole.  The additional high-frequency variability seen in the X-rays could come from the decreased light-travel time closer to the black hole, or from an additional emission component closer to the central engine.  If the lag is due to changes in the accretion flow, wouldn't we expect to find such a lag in all Seyfert galaxies?  Probably not, since the mass of the black hole in Mkn~509 is $1.43 \\times 10^{8} \\msun$ \\citep{Peterson2004}, roughly an order of magnitude greater than previously studied objects.  Many of the relevant timescales (orbital, thermal, viscous) increase linearly with mass, so we would expect to see lags of 1-2 days for other, less massive Seyferts.  Previous observing campaigns have involved weekly monitoring for a period of years, or more intensive monitoring for a period of a few days.  In those cases, it is distinctly possible that the data would appear to be correlated, but with zero lag if the monitoring is on a weekly basis.  Conversely, the data would appear to be uncorrelated if the lag time is greater than the monitoring period of just a few days.  By observing a more massive galaxy on a frequent basis for several years, we have been able to measure a lag between the optical and X-ray emission."
    },
    "0711/0711.0536_arXiv.txt": {
        "abstract": "We show high resolution spectra of the eclipsing brown dwarf binary 2MASS~J05352184$-$0546085 taken at the two opposite radial velocity maxima. Comparisons of the TiO bands to model and template spectra are fully consistent with the temperatures previously derived for this system.  In particular, the reversal of temperatures with mass -- in which the higher-mass primary is cooler than its companion -- is confirmed.  We measure the projected rotation velocities of the compononents; the primary is rotating at least twice as rapidly as the secondary.  At the two radial velocity maxima, H$\\alpha$ emission lines of both components stick out to either sides of the H$\\alpha$ central wavelength, which is dominated by nebula emission. This enables us to model the individual H$\\alpha$ lines of the primary and the secondary. We find that the H$\\alpha$ emission from the primary is at least 7 times stronger than the emission from the secondary. We conclude that the temperature reversal is very likely due to strong magnetic fields inhibiting convection on the primary. ",
        "introduction": "2MASS~J05352184$-$0546085 (hereafter \\twom) is the first known eclipsing binary system comprising two brown dwarfs \\citep{Stassun06}, with dynamically measured masses of $M_1 = 56\\pm 4$ M$_{\\rm Jup}$ and $M_2 = 36\\pm 3$ M$_{\\rm Jup}$. The large component radii of $R_1 = 0.67\\pm 0.03$ R$_\\odot$ and $R_2 = 0.49\\pm 0.02$ R$_\\odot$ suggest that the brown dwarfs are very young, and indeed the system satisfies kinematic and distance criteria for membership in the Orion Nebula Cluster (M42, age $\\sim 1$ Myr). Providing the only direct and precise measurements to date of the fundamental physical properties of young brown dwarfs, \\twom\\ represents an important opportunity to test and calibrate theoretical models of brown dwarf formation and early evolution.  Remarkably, \\twom\\ exhibits a reversal of effective temperatures: The lower-mass secondary is warmer than the higher-mass primary.  The ratio of effective temperatures, $T_2 / T_1$, is $1.064\\pm 0.004$ \\citep[][hereafter SMV07]{Stassun07}. Such a temperature reversal is not predicted by any current models of brown dwarf evolution.  One possible explanation for this temperature reversal considered by SMV07 is that convection in the primary has been somehow suppressed, perhaps through the presence of a strong surface field on the primary brown dwarf. In this Letter we present high-resolution spectra of \\twom\\ which (a) provide independent confirmation of the temperature reversal, and (b) reveal that the primary is indeed much more magnetically active than the secondary. ",
        "conclusions": "We have used high resolution spectra of \\twom\\ taken at radial velocity maxima to investigate the temperature reversal of this enigmatic brown dwarf binary. We used two TiO bands to constrain the absolute temperatures of the two components. In the $\\gamma$-band, we have used PHOENIX model spectra, but it is not clear to what extent these models can accurately reproduce the real TiO spectra. From the model spectra, we find temperatures of 2850\\,K and 3000\\,K ($\\pm 120$\\,K) for the primary and secondary, respectively. In the $\\epsilon$-band, we have used spectra from stellar templates yielding much lower temperatures, but effects of metallicity may introduce differences between template spectra and the spectrum of \\twom. This method yields temperatures between 2600\\,K and 2750\\,K.  Most importantly, the $\\epsilon$-band analysis confirms the finding by SMV07 of a temperature reversal in this system, such that the higher-mass primary is indeed cooler than the lower-mass secondary.  A possible explanation for the temperature reversal may be that large scale magnetic fields are inhibiting convection. In low mass stars, magnetic fields manifest themselves for example as H$\\alpha$ emission, and at least in sun-like stars activity is correlated to the rotation rate. We find that the primary of \\twom\\ shows at least 7 times stronger activity than the secondary, and the former is also rotating at least a factor of 2 more rapidly than the latter.  We can convert the equivalent widths into normalized H$\\alpha$ luminosities \\cite[see][]{Reiners07}. We find $\\log{L_\\textrm{\\small H$\\alpha$}/L_\\textrm{\\small bol}} = -3.47$ and $-4.30$ for the primary and secondary components, respectively.  \\citet{Reiners07} have measured H$\\alpha$ luminosities and magnetic fields in a sample of M dwarfs. They found that for a given spectral type H$\\alpha$ is closely correlated with magnetic field. From the normalized H$\\alpha$ luminosities of \\twom\\ we can estimate the mean magnetic flux on its two componenents. The primary's H$\\alpha$ luminosity of $\\log{L_\\textrm{\\small H$\\alpha$}/L_\\textrm{\\small bol}} = -3.47$ is on the top end of H$\\alpha$ luminosities found among M6--M7 dwarfs and indicate a very high magnetic flux level of $Bf \\approx 4$\\,kG over the entire object \\citep[cp. left panel of Fig.\\,12 in][]{Reiners07}. The secondary's H$\\alpha$ luminosity is only around average for that spectral range \\citep[see also][]{West04} and indicates significantly weaker magnetic flux of $Bf \\approx 2$\\,kG.  Recently, \\citet{Chabrier07} have shown that the effect of inhibited convection could indeed explain the temperature reversal in \\twom. They suggest that magnetic fields with surface values of a few kG might be sufficient for severely inhibiting convection in parts of low mass star interiors. The observations presented here support this scenario."
    },
    "0711/0711.0470_arXiv.txt": {
        "abstract": "Free streaming in a \\emph{mixture} of collisionless non-relativistic dark matter (DM) particles is studied by solving the linearized Vlasov equation implementing methods from the theory of multicomponent plasmas. The mixture includes Fermionic, condensed and non-condensed Bosonic particles decoupling in equilibrium while relativistic, heavy thermal relics that decoupled  when non-relativistic (WIMPs), and sterile neutrinos that decouple \\emph{out of equilibrium} when they are relativistic. The different components interact via the self-consistent gravitational potential that they source.  The free-streaming length $\\lambda_{fs}$ is obtained from the marginal zero  of the gravitational polarization function,  which separates short wavelength Landau-damped from long wavelength Jeans-unstable \\emph{collective} modes. At redshift $z$ we find $ \\frac{1}{\\lambda^2_{fs}(z)} = \\frac{1}{(1+z)}\\, \\big[ \\frac{0.071}{\\textrm{kpc}}\\big]^2  \\sum_{a}     \\nu_a \\, g^\\frac{2}{3}_{d,a}\\,\\left( {m_a}/{\\mathrm{keV}}\\right)^2 I_a  $, where $0\\leq \\nu_a \\leq 1$ are the \\emph{fractions} of the respective DM components   of mass $m_a$  that decouple when the effective number of ultrarelativistic degrees of freedom is $g_{d,a}$, and $I_a$ are dimensionless ratios of integrals of the distribution functions which only depend on the microphysics at decoupling and are obtained explicitly in all the cases considered. If sterile neutrinos produced either resonantly or non-resonantly that decouple near the QCD scale are the \\emph{only} DM component, we find $\\lambda_{fs}(0) \\simeq 7\\,\\mathrm{kpc} \\,(\\mathrm{keV}/m)$ (non-resonant), $\\lambda_{fs}(0) \\simeq 1.73\\,\\mathrm{kpc}\\, (\\mathrm{keV}/m)$ (resonant).   If WIMPs with $m_{wimp} \\gtrsim 100\\,\\mathrm{GeV}$ decoupling at $T_d  \\gtrsim 10 \\,\\mathrm{MeV}$ are present in the mixture with $\\nu_{wimp} \\gg 10^{-12}$, $\\lambda_{fs}(0) \\lesssim 6.5 \\times 10^{-3}\\,\\mathrm{pc}$ is \\emph{dominated} by CDM. If a Bose Einstein condensate is a DM component its free streaming length is consistent with CDM because of the infrared enhancement of the distribution function. ",
        "introduction": "Candidate dark matter (DM) particles are broadly characterized as cold, hot or warm depending on their velocity dispersions. The \\emph{concordance} $\\Lambda\\mathrm{CDM}$ standard cosmological model emerging from CMB, large scale structure observations and simulations favors the hypothesis that DM is composed of primordial particles which are cold and collisionless\\cite{primack}. Although this model is very successful in describing the large scale distribution of galaxies,  recent observations hint at possible discrepancies summarized as the ``satellite'' and ``cuspy halo'' problems. In the $\\Lambda\\mathrm{CDM}$ model the CDM power spectrum favors  small scales, which become non-linear first and collapse in a hierarchical ``bottom-up'' manner and dense clumps survive the mergers in the form of ``satellite'' galaxies. Large-scale simulations within the $\\Lambda\\mathrm{CDM}$ paradigm  lead  to an overprediction of  ``satellite'' galaxies \\cite{moore2}, which is almost an order of magnitude larger than the number of satellites that have been observed in Milky-Way sized galaxies\\cite{kauff,moore,moore2,klyp,will}.  Furthermore, large-scale  N-body simulations of CDM clustering     predict a   density profile monotonically increasing towards the center of the halos\\cite{dubi,frenk,moore2,bullock,cusps}, with asymptotic behavior $\\rho(r)\\sim r^{-\\gamma}$ where $1\\leq \\gamma \\lesssim 1.5$\\cite{moore,frenk,cusps} which describes accurately clusters of galaxies, but indicates a divergent cusp at the center of the halo. In contrast with these results, recent observations seem to indicate central cores in dwarf galaxies\\cite{dalcanton1,van,swat,gilmore,salucci}, sparking the ``cusps vs cores'' controversy.  Warm dark matter candidates (WDM) have been invoked as possible solutions to these potential discrepancies\\cite{mooreWDM,turokWDM}, these particles feature velocity dispersions intermediate between CDM and HDM and can relieve the satellite and cuspy halo problems. The clustering properties of collisionless DM in the linear regime depend on a fundamental length scale: the free-streaming length, that determines a cutoff in the power spectrum. Length scales larger than the free-streaming scale undergo gravitational instability, and shorter scales are damped.  A \\emph{simple estimate} of the free-streaming length $\\lambda_{fs}$ is obtained from the familiar Jeans' length by replacing the speed of sound by the   velocity dispersion of the particle. An equivalent estimate is obtained by computing the distance that the particle travels within a dynamical (Hubble) time\\cite{kt} $\\lambda_{fs}\\sim \\langle \\vec{V}^2\\rangle^\\frac{1}{2}/ H_0$.  However, a   thorough assessment of DM particles and structure formation requires a more detailed and reliable determination of the free-streaming length. The necessity for this  has been recently highlighted by the recent results of ref.\\cite{gao} that suggest that the first stars form in filaments of the order of the free streaming scale.  Perturbations in a collisionless system of particles with gravitational interactions is fundamentally different from fluid perturbations in the presence of gravity. The (perfect) fluid equations correspond to the limit of vanishing mean free path. In a gravitating fluid pressure gradients tend to restore hydrostatic equilibrium with the speed of sound in the medium and short wavelength fluctuations are simple acoustic waves. For large wavelengths the propagation of pressure waves cannot halt gravitational collapse on a dynamical time scale. The dividing line is the Jeans length: perturbations with  wavelengths shorter than this oscillate as sound waves, while perturbations with longer wavelength   undergo gravitational collapse.  In a gas of collisionless particles with gravitational interaction the situation is different since the mean free path is much larger than the size of the system (Hubble radius) and the fluid description is not valid. Instead the Boltzmann-Vlasov  equation for the distribution function must be solved to extract the dynamics  of  perturbations\\cite{peebles,bert}. Just as in the case of plasma physics, the linearized Boltzmann-Vlasov equation describes \\emph{collective} excitations\\cite{plasmas}. In the case of a collisionless gas with gravitational interactions these collective excitations describe particles free-streaming in and out of the gravitational potential wells of which they are the source. The damping of short wavelength \\emph{collective} excitations is akin to \\emph{Landau damping in plasmas}\\cite{plasmas} and is a result of the phenomenon of \\emph{dephasing} via phase mixing in which the particles are out of phase with the potential wells that they produce\\cite{bt}. This situation is similar to Landau damping in plasmas where dephasing between the charged particles and the self-consistent electric field that they produce lead to the collisionless damping of the collective modes\\cite{plasmas}. For a thorough review of collective excitations  and their Landau damping in gravitational systems see ref.\\cite{trema1}.  Gilbert\\cite{gilbert} studied the linearized Boltzmann-Vlasov equation in a matter dominated cosmology for non-relativistic particles described by an (unperturbed) Maxwell-Boltzmann distribution function. In this reference the linearized this equation was cast as a Volterra integral equation of the second kind which was solved numerically. The result of the integration reveals a limiting value of the wavevector below which perturbations are Landau  damped and above which perturbations grow via gravitational instability. Eventually the redshift in the expanding cosmology makes wavevectors that are initially damped  to enter the band of unstable modes and grow\\cite{gilbert}. However, the dividing wavevector  between damped and growing modes emerges at very early times during which the expansion can be neglected\\cite{gilbert}. The results of the numerical study were consistent with replacing the speed of sound by the Maxwellian velocity dispersion in the Jeans length (up to a normalization factor of $\\mathcal{O}(1)$). Gilbert's equations were used also by Bond and Szalay\\cite{bondszalay} in their pioneering study of collisionless damping of density fluctuations in an expanding cosmology. These authors focused primarily on massive neutrinos and solved numerically Gilbert's integral equation but approximated the Fermi-Dirac distribution function and provided a fitting function for the numerical results.  Gilbert's equations were also solved numerically to study dissipationless clustering of neutrinos in ref.\\cite{bran} but with a truncation of their Fermi-Dirac distribution function   and an  analytic fit to the numerical solution of the integral equation was provided. Bertschinger and Watts\\cite{bertwatts,bert} also studied numerically  Gilbert's equation within the context of cosmological perturbations from cosmic strings and massive light neutrinos, and more recently similar integral  equations were solved  approximately for thermal neutrino relics in ref.\\cite{ringwald}.  Most of the studies of the solutions of the linearized Boltzmann-Vlasov equation for collisionless particles addressed one single species of DM candidates\\footnote{The exception being ref.\\cite{bertwatts} wherein massive neutrinos decoupled with a Fermi-Dirac distribution  and a cosmic string source $\\propto \\delta^3(\\vec{r})$ were studied, neglecting any radiation component.} and generally in terms equilibrium distribution functions or approximations thereof, for example Maxwell-Boltzmann for relics that decouple non-relativistically or Fermi-Dirac (without chemical potential) (or truncations of this distribution) for neutrinos.  However, it is   likely that DM may be composed of \\emph{several} species, this possibility is suggested in ref.\\cite{palazzo} and most extensions of the standard model generally allow several possible candidates, from massive weakly interacting   particles (WIMPs) to ``sterile'' neutrinos  ( $SU(2)$ singlets)\\cite{kusenko}.  Furthermore, several possible WDM candidates may decouple \\emph{out of local thermodynamic equilibrium} (LTE) with distribution functions which may be very different from the usual ones in LTE. This is the case for sterile neutrinos produced non-resonantly via the Dodelson-Widrow (DW) mechanism\\cite{dw} or through a lepton-driven MSW (Mikheyev-Smirnov-Wolfenstein) resonance\\cite{este}.  \\medskip  {\\bf{A. Motivation and goals:}} In this article we study free streaming of decoupled collisionless non-relativistic (DM) candidates focusing on \\emph{two} aspects: \\begin{itemize} \\item{ A \\emph{mixture} of (DM) particles including CDM and WDM candidates: typical studies of structure formation invoke either CDM \\emph{or} WDM, but it is likely that \\emph{both} candidates are present with different fractions $\\nu$ of the total DM component. In fact most particle physics extensions beyond the standard model have plenty of room for a variety of CDM, WDM or HDM candidates. Thus we allow all of these components, each one contributing an arbitrary fraction $\\nu$ to the total (DM) content of the Universe. Although the DM candidates are collisionless and do not interact directly with each other, they interact \\emph{indirectly} via the gravitational potential that they source. As a result the free-streaming length of the mixture is a non-trivial function of the individual free-streaming lengths.   }  \\item{Free streaming has mostly been studied in the above references with particles that decoupled either when ultrarelativistic (as is the case for neutrinos) or non-relativic as in the case of weakly interacting massive particles (WIMPs) but generally in local thermodynamic equilibrium (LTE), namely with Fermi-Dirac, Bose-Einstein or Maxwell-Boltzmann distributions respectively. We seek to obtain the corresponding free streaming lengths for particles that decoupled \\emph{in or out} of LTE with \\emph{arbitrary} isotropic distribution functions, without any truncation. This aspect is important for sterile neutrinos either produced non-resonantly\\cite{dw} or resonantly\\cite{este,kusenko} because these particles    decoupled while relativistic but  \\emph{out} of LTE. Therefore, we consider the most general \\emph{mixture} of Fermionic and Bosonic thermal relics that decouple   when relativistic, including the possibility of a Bose-Einstein condensate (BEC)\\cite{coldmatter},   heavy non-relativistic thermal relics, as the case of WIMPs, and sterile neutrinos that decoupled \\emph{out of LTE} when ultrarelativistic. } \\end{itemize}  In order to carry out this program analytically we first neglect the cosmological expansion and solve the linearized Boltzmann-Vlasov equation in the non-expanding case \\emph{exactly} by implementing methods from the theory of multicomponent plasmas\\cite{plasmas}. The neglect of the cosmological expansion is warranted by the detailed numerical study in refs.\\cite{gilbert,bondszalay} wherein it was found that the dividing wavevector between the Landau damped modes and the modes that grow under gravitational instability is insensitive to the expansion, just as in the case of the Jeans instability where the Jeans wavevector can be extracted in the non-expanding case include the redshift dependence of the density, speed of sound and scale factors \\emph{a posteriori}\\cite{kt}. The analytic \\emph{exact} solution for the free-streaming wave-vector \\emph{today} $k_{fs}( 0)=2\\pi/\\lambda_{fs}( 0)$ in terms of the full distribution functions without truncation, \\emph{in or out of LTE} is a main result of this program, one which yields a reliable determination of free-streaming lengths for mixtures of (DM) components that decoupled with arbitrary distribution functions.      \\medskip  This program is carried out by implementing methods from the theory of multicomponent plasmas\\cite{plasmas}, in particular we obtain the ``gravitational polarization'' function\\cite{trema1,bt} for a \\emph{mixture} of (DM) components akin to the dielectric response function of multicomponent plasmas\\cite{plasmas}. The collective excitations are described by the zeroes of this function in the complex frequency plane and the free-streaming wave-vector $k_{fs}$ is identified as that wavevector that separates between the Landau-damped short wavelength modes and the gravitationally Jeans- unstable long-wavelength modes.   Based on the Liouville evolution of the decoupled distribution functions and assuming that the expansion of the universe is slow enough so that it can be treated adiabatically, we provide a scaling argument   that determines  the following dependence of the free streaming length on the redshift,  \\be \\lambda_{fs}(z) = {\\lambda_{fs}(0)}{\\sqrt{1+z}}\\,.\\ee  \\medskip  {\\bf{B. Summary of  Results:} } Our main result for the comoving free streaming length  $\\lambda_{fs}(z)$  at redshift $z$  of mixed (DM) is   \\bea \\frac{1}{\\lambda^2_{fs}(z)} = \\frac{1}{(1+z)}\\, \\Big[\\frac{0.071}{\\textrm{kpc}}\\Big]^2 \\sum_{\\textrm{ species}}&&\\Bigg\\{ \\nu_F \\, g^\\frac{2}{3}_{d,F}\\,\\left(\\frac{m_F}{\\mathrm{keV}}\\right)^2 I_F[u] + \\nu_s \\, g^\\frac{2}{3}_{d,s}\\,\\left(\\frac{m_s}{\\mathrm{keV}}\\right)^2\\, 6.814 +\\nu_B \\, g^\\frac{2}{3}_{d,B}\\,\\left(\\frac{m_B}{\\mathrm{keV}}\\right)^2 I_B[x_d,u_d]+ \\nonumber \\\\&& 10^{12}\\, \\nu_{wimp}\\, \\, g^\\frac{2}{3}_{d,wimp}\\, \\Big(\\frac{m_{wimp}}{100\\,\\textrm{GeV}}\\Big)\\,\\Big(\\frac{T_{d}}{10 \\,\\textrm{MeV}}\\Big) \\Bigg\\}\\,,\\label{lfsintro}\\eea  where $\\nu_a$ is the \\emph{fraction} of (DM) of each component with $\\sum_a \\nu_a=1$, $g_{d,a}$ is the effective number of ultrarelativistic degrees of freedom at decoupling for each species $(a)$ of mass $m_a$,  and the functions $I_F,I_B$ are dimensionless ratios of integrals of the distribution functions of the decoupled particles which are determined by the microphysics at decoupling. Their explicit expressions in the cases considered are given in section (\\ref{sec:deco}). The label $F$ refer to \\emph{all} possible Fermions with chemical potential $\\mu$ decoupled in LTE  at a temperature $T_d$ while ultrarelativistic, and  sterile neutrinos produced non-resonantly via the (DW) mechanism\\cite{dw} for which the chemical potential vanishes and $I_F[0]=2\\ln(2)/3\\zeta(3)=0.3844$, and the label $s$ refers solely to sterile neutrinos produced via a lepton-driven (MSW) resonance via the mechanism described in ref.\\cite{este}. The label $B$ corresponds to condensed or non-condensed Bosons of mass $m$ and chemical potential $\\mu$ that decoupled at temperature $T_d$ while ultrarelativistic. The function $I_B$ features an infrared divergence in the limits $\\mu/T_d ; m/T_d \\rightarrow 0$ or $\\mu = m$ for any value of the mass. This latter case corresponds to the case of a Bose-Einstein Condensate\\cite{coldmatter}. Thus Bosonic particles that decoupled while ultrarelativistic with $\\mu/T_d \\ll 1$ \\emph{or} that formed a BEC lead to small free-streaming lengths and behave as CDM. Finally, (WIMPs) are considered to be decoupled while non-relativistic with a Maxwell-Boltzmann distribution function.  Eq.(\\ref{lfsintro}) clearly shows that (WIMPs) with $m\\sim 100\\,\\mathrm{GeV}$ that decoupled kinetically in LTE  at $T\\sim 10\\,\\mathrm{MeV}$\\cite{dominik} dominate all other contributions to $k_{fs}$ resulting in an extremely small free streaming length $\\lambda_{fs} \\sim 6.5 \\times 10^{-3}\\,\\mathrm{pc}$ \\emph{unless} their fractional abundance $\\nu_{wimp} \\lesssim 10^{-12}$.  If (DM) is dominated by sterile neutrinos (produced either resonantly or non-resonantly) that decoupled near the QCD scale\\cite{dw,este} with $m \\sim \\mathrm{keV}$, we find that the typical free-streaming lengths \\emph{today} are $\\lambda_{fs} \\sim 2-7 \\, \\mathrm{kpc}$, where the larger value corresponds to the non-resonant and the lower value to the resonant production mechanisms respectively. The lower values are  consistent with the recent numerical study of the formation of the first stars in filamentary structures\\cite{gao} and the values for the ``cores'' extracted from the data in ref.\\cite{gilmore} for dwarf spheroidal galaxies (dSphs) ($r_c \\sim 0.5 \\,\\mathrm{kpc}$). The upper values agree with the cores extracted from the data in ref.\\cite{salucci} for spiral galaxies ($r_c \\sim 10 \\,\\textrm{kpc}$). ",
        "conclusions": "In this article we have implemented methods from the theory of multicomponent plasmas to study free streaming of a \\emph{mixture} of non-relativistic DM candidates that include  Fermionic and Bosonic particles that decouple in LTE while relativistic, including the possibility of a Bose-Einstein condensate, heavy thermal relics that decoupled in when non-relativistic (WIMPs), and sterile neutrinos that decouple \\emph{out of LTE} when they are relativistic. We solve \\emph{exactly} the Boltzmann-Vlasov equation for the gravitational perturbations in a non-expanding cosmology and obtain the ``gravitational polarization function'' whose zeroes determine the dispersion relations of the collective excitations of the self-gravitating collisionless gas of particles. The free-streaming wave vector is obtained from the marginal solution that separates Landau damped short wavelength perturbations from unstable collective modes. We obtain the free-streaming length for \\emph{arbitrary} (but isotropic) distributions of the  particles that decoupled \\emph{in or out of LTE} solely in terms of the fractional abundance of the different species and integrals of their distribution functions which depend on the microphysics at decoupling.  Because all of the components are non-relativistic and assuming  that the expansion is slow we provide an adiabaticity argument  that allows us to implement a simple rescaling of densities, velocities dispersion and wavelengths to extract the redshift dependence of the free-streaming length. The validity of this adiabatic approximation is confirmed by the similarity of the result to the Jeans' length for Newtonian perturbations in an expanding cosmology and by the explicit result for the free-streaming wavevector for thermal neutrinos obtained in ref.\\cite{ringwald}.  The main result for the free streaming length as a function of redshift for an arbitrary mixture of DM components is \\be \\frac{1}{\\lambda^2_{fs}(z)} = \\frac{1}{(1+z)}\\, \\Big[\\frac{0.071}{\\textrm{kpc}}\\Big]^2 \\,  \\sum_a \\nu_a\\,g^\\frac{2}{3}_{d,a} \\left(\\frac{m_a}{\\mathrm{keV}}\\right)^2 \\,I_a \\,,\\label{lamfsz} \\ee where \\be I_a = \\Bigg[ \\frac{\\int_0^\\infty f^{(0)}_{d,a}(y)dy}{\\int_0^\\infty y^2\\,f^{(0)}_{d,a}(y)dy}\\Bigg] \\label{Ial}\\ee   is a ratio of integrals of the distribution functions that only depends on the $m_a/T_d,\\mu_a/T_d$ and dimensionless couplings and is completely determined by the  microphysics at decoupling. Evaluating these integrals for thermal Fermionic and Bosonic relics (with or without condensation), WIMPs and sterile neutrinos decoupled out of LTE either resonantly or non-resonantly with the distribution functions obtained in refs.\\cite{este,dw} respectively, we find the general result     \\bea \\frac{1}{\\lambda^2_{fs}(z)} = \\frac{1}{(1+z)}\\, \\Big[\\frac{0.071}{\\textrm{kpc}}\\Big]^2 \\sum_{\\textrm{ species}}&&\\Bigg\\{ \\nu_F \\, g^\\frac{2}{3}_{d,F}\\,\\left(\\frac{m_F}{\\mathrm{keV}}\\right)^2 I_F[u] + \\nu_s \\, g^\\frac{2}{3}_{d,s}\\,\\left(\\frac{m_s}{\\mathrm{keV}}\\right)^2\\, 6.814 +\\nu_B \\, g^\\frac{2}{3}_{d,B}\\,\\left(\\frac{m_B}{\\mathrm{keV}}\\right)^2 I_B[x_d,u_d]+ \\nonumber \\\\&& 10^{12}\\, \\nu_{wimp}\\, \\, g^\\frac{2}{3}_{d,wimp}\\, \\Big(\\frac{m_{wimp}}{100\\,\\textrm{GeV}}\\Big)\\,\\Big(\\frac{T_{d}}{10 \\,\\textrm{MeV}}\\Big) \\Bigg\\}\\,,\\label{lfsconclu}\\eea where $\\nu_a$ is the partial fraction of each component with $\\sum_a \\nu_a =1$, the sum over Fermionic species $(F)$ includes thermal relics \\emph{and} sterile neutrinos produced non-resonantly via the (DW) mechanism\\cite{dw}, for which the chemical potential vanishes and $I_F[0]=0.3844$, the label $s$ is for sterile neutrinos produced via a lepton-driven (MSW) resonance as described in ref.\\cite{este}, and the label $B$ stands for condensed or non-condensed Bosonic thermal relics.  This expression features  several important consequences relevant for large scale structure formation: \\begin{itemize} \\item{A non-negligible $\\nu_{wimp}\\gg 10^{-12}$, for a CDM candidate with $m \\sim 100\\,\\textrm{GeV}$ which decoupled at $T_d \\sim 10 \\, \\mathrm{MeV}$\\cite{dominik} overwhelms all other components (but for a possible BEC) and leads to   small free streaming lengths consistent with CDM  regardless of the presence of WDM or HDM components. For $\\nu_{wimp} \\gg 10^{-12}$ the free-streaming length of \\emph{mixed} DM is completely dominated by WIMPs and is given today by \\be \\lambda_{fs}(0) \\sim \\frac{0.014\\,\\mathrm{pc}}{g^\\frac{1}{3}_d}\\,\\Big( \\frac{100\\,\\mathrm{GeV}}{m}\\Big)^\\frac{1}{2}\\, \\Big( \\frac{10\\,\\mathrm{MeV}}{ T_d}\\Big)^\\frac{1}{2}\\,, \\ee from which it follows that for WIMPs with $m \\sim 100\\,\\mathrm{GeV}$ that decouple kinetically at $T_d \\sim 10 \\,\\mathrm{MeV}$\\cite{dominik} when $g_d \\sim 10$\\cite{kt}  \\be \\lambda_{fs}(0) \\sim   6.5\\times 10^{-3}\\,\\mathrm{pc}\\,.\\label{wimplum}\\ee This cut-off scale might well be related to the smallest non-linear structures found in\\cite{cusps}   unless there is some substantial violent relaxation and  merging. }  \\item{For vanishing chemical potential $u_d =0$, non-BEC Bosonic ultrarelativistic relics feature an infrared enhancement in  $I_B[x_d,0]$ which must be regulated by assuming that the integrals of the distribution function are cutoff of order of the Hubble scale at decoupling $H_d \\sim \\sqrt{g_d}\\, T^2_d/M_{Pl}$. If this is the \\emph{only} DM component, the free-streaming length  today is \\be \\lambda_{fs}(0) \\sim \\frac{14}{g^\\frac{1}{3}_d\\,\\sqrt{I_B[x_d,0]}}\\,\\Big( \\frac{\\mathrm{keV}}{m}\\Big) \\mathrm{kpc}\\,~~;~~I_B[x_d,0]  \\sim \\frac{1}{2\\zeta(3)}\\, \\ln\\Bigg[\\frac{2\\,T_d}{H_d+\\sqrt{H^2_d+m^2}}\\Bigg] \\,   .\\ee  For BEC Bosons $x_d=u_d$ the function $I_B[x_d,x_d]$ is divergent as a consequence of the infrared divergence in the numerator of $I_a$ (eq. \\ref{Ia}). However,  we highlighted that the physics of BEC formation in an expanding cosmology must be assessed in greater detail in order to understand the behavior of superhorizon modes. This observation also holds for the non-condensed Bose gas decoupling when relativistic with $\\mu/T_d \\ll1$ because the distribution function in this case is also infrared sensitive. In both these cases an estimate of the free streaming length may be obtained by introducing an infrared cutoff in the integral in the numerator of $I_a$ of the order of the Hubble scale, since superhorizon modes cannot establish thermal equilibrium via causal processes as discussed above. If a BEC is the \\emph{only} DM component, its free-streaming length today is approximately given by \\be \\lambda_{fs}(0) \\sim \\frac{0.014\\,\\mathrm{pc}}{g^\\frac{1}{12}_d}\\,\\Big(\\frac{\\mathrm{keV}}{m} \\Big)^\\frac{3}{2}\\, \\Big(\\frac{T_d}{\\mathrm{GeV}} \\Big)\\,.  \\ee Therefore even when these relics decoupled when they were ultrarelativistic, they could effectively act as CDM components. This is a consequence of the fact that the distribution functions favor small momenta. }  \\item{If sterile neutrinos that decouple \\emph{out of LTE} near the QCD scale produced either non-resonantly via the (DW)\\cite{dw} mechanism or via a lepton-driven MSW resonance\\cite{este} near the QCD scale are the \\emph{only} DM components we find the following free-streaming lengths today \\bea   \\lambda_{fs}(0)   & \\simeq &  7  \\, \\mathrm{kpc}\\,  \\,\\Big( \\frac{\\mathrm{keV}}{m}\\Big)~~\\mathrm{non-resonant}    \\\\\\lambda_{fs}(0)  & \\simeq &  1.73  \\, \\mathrm{kpc}\\,  \\,\\Big( \\frac{\\mathrm{keV}}{m}\\Big)~~\\mathrm{resonant}\\,.   \\eea  The smaller values of the free streaming length are compatible with those in ref.\\cite{gao} where it is found that first stars form  in filamentary structures with length scales of the order of the free streaming scale and within a factor $\\sim 3-4$ seem to be also consistent with the ``cores'' resulting from the fit of the density profile for dwarf spheroidal galaxies in ref.\\cite{gilmore} $\\sim 0.5\\,\\mathrm{kpc}$. The larger values are consistent with those found in ref.\\cite{salucci} for the cored profiles of spiral galaxies $\\sim 5-10 \\, \\mathrm{kpc}$.    }  \\end{itemize}   We believe that these results  lead to a significant advance in the understanding of collisionless (DM) because trying to obtain a reliable estimate of the free-streaming lengths   via the numerical integration of Gilbert's equations for a combination of \\emph{arbitrary} distribution functions corresponding to particles that decoupled in or out of LTE is undoubtedly a daunting task.  The sensitivity of the free-streaming scale to the details of the distribution function at low momentum and the importance of a reliable determination of the free-streaming length as a measure of the cutoff of the power spectrum of linearized cosmological perturbations require a fundamentally sound understanding of the microphysics of production and decoupling of sterile neutrinos, a program currently underway\\cite{boyho}."
    },
    "0711/0711.2369_arXiv.txt": {
        "abstract": "Observations of the Sagittarius dwarf spheroidal (Sgr dSph) galaxy were carried out with the H.E.S.S. array of four imaging air Cherenkov telescopes in June 2006. A total of 11 hours of high quality data are available after data selection. There is no evidence for a very high energy $\\gamma$-ray signal above the energy threshold at the target position. A 95\\% C.L. flux limit of $\\rm 3.6 \\times 10^{-12} cm^{-2}s^{-1}$ above 250 GeV has been derived. Constraints on the velocity-weighted cross section $\\rm \\langle \\sigma v \\rangle$ are calculated in the framework of Dark Matter particle annihilation using realistic models for the Dark Matter halo profile of Sagittarius dwarf galaxy. Two different models have been investigated encompassing a large class of halo types. A 95\\% C.L. exclusion limit on $\\rm \\langle \\sigma v \\rangle$ of the order of $\\rm 2 \\times 10^{-25} cm^3s^{-1}$  is obtained for a core profile in the 100 GeV - 1 TeV neutralino mass range. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.4779_arXiv.txt": {
        "abstract": "Experiments with thin ZnS(Ag) scintillators provide evidence with C.L. $>$ 99.99{\\%} for the existence of DArk Electric Matter Objects - daemons (presumably negatively charged Planckian particles with $M \\sim $ 10$^{-5}$ g) captured from the Galactic disk into near-Earth, almost circular heliocentric orbits (NEACHOs). Their flux at $V \\approx $ 10-15 km/s was found to be as high as $f_{\\oplus } >$ 10$^{-7}$ cm$^{-2}$s$^{-1}$ and vary with $P$ = 0.5 y, with maxima in March and September. A daemon flux $f_{\\oplus } \\sim $ 10$^{-7}$-10$^{-6}$ cm$^{-2}$s$^{-1}$ is capable of accounting for the Troitsk anomaly in the $^{3}$T $\\beta $-spectrum and suggests its more pronounced manifestation in future KATRIN experiment. In view of the channeling effect on iodine recoil nuclei in the NaI(Tl) crystal, the DAMA/NaI experiment is also apparently detecting a flux of daemons, $f_{\\oplus } \\sim $ 6$\\times $10$^{-7}$ cm$^{-2}$s$^{-1}$, but in this case of those falling with $V$ = 30-50 km/s from strongly elongated, Earth-crossing heliocentric orbits (SEECHOs) oriented in the antapex direction, as a result of which the number of events detected in the 2-6-keV interval varies with $P$ = 1 y. ",
        "introduction": "Our Universe started from Planckian scales, and it appears only reasonable to assume that the larger part of its mass, Dark Matter (DM), remains confined in Planckian objects, elementary black holes whose gravitational radius cannot be smaller than their Compton wavelength. Such objects with $M \\sim $ 10$^{-5}$ g and $r_{g} \\sim $ 10$^{-33}$ cm can carry an electric charge of up to $Z$e $\\approx  G^{1/2}M \\approx $ 10e. While the properties of similar DArk Electric Matter Objects, \\textit{daemons}, were considered by a number of authors [1-8], but the possibility of their detection was perceived skeptically (e.g., [1]).  We started from a notion that negative daemons are nuclear-active particles. Their capture of atomic nuclei should give rise to a release of energy $W \\approx $ 1.8\\textit{ZZ}$_{n}A^{-1/3}$ MeV ($\\sim $10$^{2}$ MeV), with the attendant emission of scintillation-active electrons and nucleons [9]. Becoming confined in the remainder of the nucleus, the daemon should bring about successive decay of nucleons [10], thus lowering the charge of the remainder to $\\vert Z_{n}\\vert -\\vert Z\\vert  <$ 0, which makes it possible for the daemon to capture another nucleus with emission of new particles, etc. We noted also that a negative daemon is capable of catalyzing fusion of light nuclei, up to $Z_{n} \\sim $ 6-9 (an analog of the muon catalysis of deuterons [11]). It should be pointed out that our understanding of the properties and behavior of daemons developed parallel to and in intimate connection with the experiments we were conducting.  Adopting $\\sim $0.3 GeV/cm$^{3}$ for the DM density in the galactic halo, we arrive at $\\sim $5$\\times $10$^{-12}$ cm$^{-2}$s$^{-1}$ for the halo daemon flux intercepted by the Earth. The same applies to the galactic disk DM population with velocity dispersion of 4-30 km/s. The Sun moves relative to the disk objects with $V \\approx $ 20 km/s (in the direction of the apex, which does not lie in the plane of ecliptic [12]). For such a velocity, the effective cross section of the Sun, amplified by gravitational focusing, exceeds its geometric cross section $\\sim $10$^{3}$ times.  In crossing the Sun, daemons suffer slowing down, so that a sizable part of them falls back to move along strongly elongated and rapidly shrinking rosette-shaped trajectories, whose perihelia lie inside the Sun. These daemons settle eventually toward its center to form there a compact daemon kernel, which is capable of accounting for many specific features in solar physics, including generation of non-electron neutrinos [13]. If a daemon propagating along a rosette trajectory crosses the Earth's gravitational sphere of action, the perihelion of this trajectory has a high probability to leave the body of the Sun. This is how daemons enter and build up in stable SEECHOs (until their next encounter with the Earth). Straightforward estimates made in 1996 [14] in the gas-kinetic approximation of the mean free path suggest that the SEECHO daemon flux may reach as high as $f_{\\oplus } \\sim $ 3$\\times $10$^{-7}$ cm$^{-2}$s$^{-1}$ for a velocity of 30-50 km/s.  Recent celestial mechanics calculations [15] showed that SEECHOs crowd primarily in the ``shadow'' zone, i.e., the zone on the antapex side relative the Sun, because the petals of the rosette trajectories emerging into the apex hemisphere can no longer reach the Earth's orbit as a result of the slowing down the daemons undergo after several crossings of the Sun. Whence it follows that the flux of SEECHO daemons should exhibit a distinct one-year modulation with a maximum some time in June (it is appropriate to point out here that the direction to the apex is slightly different for different galactic disk populations, namely, stars of various classes, interstellar gas etc.; therefore, pinpointing this direction relative to the disk daemons is a task for the future).  Subsequent crossings by SEECHO objects of the Earth's gravitational sphere result in their gradual transfer to NEACHOs and accumulation there. It is from NEACHOs that daemons fall on the Earth with $V \\approx $ 10(11.2)-15 km/s. ",
        "conclusions": ""
    },
    "0711/0711.4297.txt": {
        "abstract": "A full description of the 5.5-yr low excitation events in $\\eta$ Carinae is presented. We show that they are not as simple and brief as previously thought, but a combination of two components. The first, the {\\it slow variation} component, is revealed by slow changes in the ionization level of circumstellar matter across the whole cycle and is caused by gradual changes in the wind-wind collision shock-cone orientation, angular opening and gaseous content. The second, the {\\it collapse} component, is restricted to around the minimum, and is due to a temporary global collapse of the wind-wind collision shock. High energy photons (E $>$ 16 eV) from the companion star are strongly shielded, leaving the\tWeigelt objects at low ionization state for $>$6 months. High energy phenomena are sensitive only to the {\\it collapse}, low energy only to the {\\it slow variation} and intermediate energies to both components. Simple eclipses and  mechanisms effective only near periastron (e.g., shell ejection or accretion onto the secondary star) cannot account for the whole 5.5-yr cycle.  We find anti-correlated changes in the intensity and the radial velocity of P Cygni absorption profiles in Fe\\,{\\sc ii}~$\\lambda6455$  and He\\,{\\sc i}~$\\lambda7065$  lines, indicating  that the former is associated to the primary and the latter to the secondary star. We present a set of light curves representative of the whole spectrum, useful for monitoring the next event (2009 January 11). ",
        "introduction": "$\\eta$ Carinae is one of the most luminous and massive stars in the Milky Way. It underwent episodes of large mass ejections in recent centuries, one of them creating the Homunculus bipolar flow with  $\\sim$12~M$_{\\odot}$ \\citep{SGH03_IR}. It continues to lose mass at a rate of $\\sim$10$^{-3}$ $\\dot{\\rm{M}}$ yr$^{-1}$ through a stellar wind \\citep{hillieretal01}, while intervening gas and dust precludes a clear view of the central source by ground-based observations.  Ground-based spectra show a mix of narrow, broad, permitted and forbidden emission lines \\citep{HA92_eta,b29}, some of them displaying P Cygni absorption profiles.. A comprehensive description of the spectrum is found in  \\citet{daminelietal07}, hereafter Paper I, and in the references  cited in that paper. The paradoxical aspect of the spectrum is the presence of lines from high and low energy states. On the basis of  truly periodic variations of the lines, previous work was able to show  that the central source is a binary star.  The primary star is colder and more luminous and the invisible companion is hotter and fainter \\citep{b5}. A tremendous wind-wind collision (WWC) was revealed from the X-ray light curve and spectra \\citep{b2, b70, henley08}.  High resolution images from the ground and from space showed the existence of condensations, named Weigelt objects \\citep{weigelt86}, that are the main source of  narrow lines and of an extended stellar wind, from where broad lines are formed \\citep{b6}.  Although we have a basic picture of the system, many details are not yet understood. This is because the spectrum is incredibly complex and variable. Moreover, the observational properties of variations in lines and continua  have not been presented previously. Some key features, like the permitted broad emission line He\\,{\\sc ii}~$\\lambda$4686 \\citep{b20} have just been discovered  in this star, even if it is one of the most frequently observed in the entire sky. This line eluded detection for  more than 50 years, due to its faintness (EW$<2$~\\AA) and transient  appearance just before minimum. In spite of being faint, it reveals a huge reservoir of high energy photons. This line  originates close to the central source, but the precise location and mechanism are unknown \\citep{b15,SokerBehar06}. Temporal variations of the lines are an important tool to understand the system, but these have not been fully explored.  %to the full extent up to now.  Extensive studies have also been made of the variability of H$\\alpha$. \\citet{kd05} examined the variability of HST STIS observations of H$\\alpha$ and showed that the profile did not repeat at the same phase. In particular, a flat topped H$\\alpha$ profile was observed during the 2003 event, but this was not seen during the 1998/9 event. Extensive UVES observations of the SE lobe were also undertaken for the 2003 event \\citep{b19,Weis05}. As the light is reflected by the Homunculus, this allows the events to be observed from different directions. Reflected variability was seen, although the line profiles are different from what we see  in the direction of the central star. This confirms that the event is not spherically symmetric, a property that could be interpreted as either a latitude dependent shell event, or a consequence of the different viewing angles relative to the binary orbital plane.  A consistent quantitative interpretation of these data set remains to be done.  In Paper I %\\citep{daminelietal07} we presented  the fundamental parameters of the 5.5-yr cycle. They are the period length P~=~2022.7$\\pm1.3$ d, the phase 0 on  $T_0=$JD 2,452,819.8 ( defined by the disappearance of He\\,{\\sc i}~$\\lambda6678$ narrow component), and the high stability of the period during the last 60 years, only understandable in a binary scenario. However, the events are not simple eclipses, and to decipher their nature we need to examine the temporal behavior of the permitted and forbidden lines, and their associated features (broad, narrow and P Cygni components).  The paper is organized as follows. In section \\S \\ref{data} we present the source and quality of data; \\S \\ref{specinterp} we present a general view where individual emission lines originate; in \\S \\ref{binary} we present the general framework of the binary model to interpret the data; in \\S \\ref{min} we discuss the bimodal nature of the event; in \\S \\ref{hei10830} we discuss the peculiarities of the He\\,{\\sc i}~$\\lambda10830$ line; in \\S \\ref{multispectral} we present a multispectral view of the  {\\it collapse} component; and finally, in \\S \\ref{discus}, we set forth a general discussion and our conclusions. ",
        "conclusions": "\\label{discus}  From an analysis of the narrow components of forbidden  emission lines we infer that the event is composed of two parts: the {\\it slow variation}, and the {\\it collapse} components.  The former comprises almost the entire 5.5-yr cycle, is centered at a later time than the {\\it collapse} component and the fading/recovering branches have a smaller degree of asymmetry. These lines are formed in the Weigelt objects and represent the way they `see'  the ionizing source, which is the radiation field of the secondary star that escapes through the shock-cone opening. The gentle fading and recovery indicates that the cone is not completely transparent. The obscuration of the ionizing source, as seen from the Weigelt objects, increases and decreases as the secondary star travels around the orbit and deep into and out of the primary's wind. The {\\it slow variation} component indicates that the shape and gas density inside the shock-cone change along the cycle. Models invoking a shell ejection, accretion onto the secondary  companion, or a simple eclipse by the primary, are unable to explain the {\\it slow variation} component.  The {\\it collapse} component is seen in many spectral features --- X-rays, high excitation Weigelt lines, He{\\,\\sc i} P~Cgyni absorption, and He{\\, \\sc I} emission. It is likely that multiple processes are affecting the collapse phase, with its duration ultimately set by the amount of time the companion spends near periastron in its highly eccentric orbit. Eclipse-like effects (`cone eclipse') could produce a decrease in lines formed in the Weigelt objects, in the inner system and reflected in the dusty Homunculus and features. However, not all directions can be affected at the same time.  Due to uncertainties in the orbital and stellar parameters it is unclear how close the primary and secondary stars approach at periastron. However, because of the much higher densities near periastron, and the rapidly changing orbit, conditions in the shock-cone will change, and it is possible that instabilities will occur. The observed behavior of many spectral features might be explained by a global disruption of the wind-wind collision shock. With this model it is also easy to explain the observed asymmetry seen in the {\\it collapse} light curves around minimum. The bow shock is not symmetrically orientated around periastron and its structure is not identical either side of periastron. In addition to the Coriolis effect, the companion stars are approaching each other before periastron and receding after, changing the relative speeds at which the winds collide.  Although we could not identify unambiguously the  signature of the WWC restoration, it must happen at most 3 months later than $\\phi=0$. At ~3 months, their is the sharp recovery in the intensity of X-rays and the He\\,{\\sc i} broad emission lines. This is either due to the restoration of the WWC, or due to the end of the `cone eclipse' (in which case the restoration of the shock must have occurred earlier).  The doubly ionized forbidden lines show a remarkable behaviour in the {\\it collapse} component. There is a trend of the time delay of line fading with the  ionization potential, indicating that the Weigelt objects see an ionizing spectrum that becomes progressively softer as the minimum is approached, and progressively harder in the recovery phase. This must be caused by optical depth effects of intervening gas in between the ionizing source and the blobs. The slow re-ionization of the blobs ($>$6 months after the minimum for [Ne\\,{\\sc iii}]) is most likely due to obscuration effects. Since the line intensity curves of the high energy forbidden  lines do not flatten out in the high excitation state, it is probable that only a fraction of the gas in the Weigelt objects is ionized to Ne$^{++}$,  S$^{++}$ and Ar$^{++}$.  Very high energy features,such as X-rays and  He\\,{\\sc ii}~$\\lambda4686$, display almost only the {\\it collapse} component. Conversely, at radio cm wavelengths the light curve is dominated by the {\\it slow variation} component. At mm wavelengths the {\\it collapse} component is present at a modest level.  The minimum of the radio 7-mm light curve is centered at $\\sim$42 days after $\\phi=0$. Although the WWC shock may be responsible for a fraction of the radio emission at these energies, most of the emission comes from ionization in the gas outside the binary system. Some effects, such as the light travel time from the source and recombination time of the gas,  can represent a delay to the center of the minimum in order that this epoch is compatible with the center of minimum in X-rays.  He\\,{\\sc i}  P Cygni absorption (and broad emission) lines cannot be formed in the wind of the primary star, since its speed reaches values larger than those of  P Cygni absorption in Fe\\,{\\sc ii}~$\\lambda6455$, which is formed in the external layers of the primary's wind.  As the P Cygni absorption component of Fe\\,{\\sc ii}~$\\lambda6455$ reaches the maximum intensity and maximum negative velocity $\\sim$3 weeks after $\\phi=0$. the wind of the primary star is recombined to a maximum extent at that time. It must occur when the (ionizing) secondary star is at opposition. This interpretation is corroborated by the fact that around $\\phi=0$ the velocity of the P Cygni absorption component of Fe\\,{\\sc ii}~$\\lambda6455$  is shifting to the negative side, and that of He\\,{\\sc i}~$\\lambda7065$ P Cygni absorption to the opposite direction. Other features associated with the shock-cone of the secondary star support this scenario, such as the He\\,{\\sc i}~$\\lambda6678$  broad emission component that also is shifting to the positive side of radial velocities and reaches the minimum at $\\phi$ = 0.01. The X-ray flux reaches a minimum and the column density ($N_{\\rm H}$) reaches a peak 30 days after $\\phi=0$. Moreover, the fluorescent (narrow) line Fe\\,{\\sc ii}~$\\lambda8490$  also reaches a minimum around  $\\phi$ = 0.015.  These data, taken together, indicate that the opposition of  the secondary star occurs 20-30 days after $\\phi=0$. We do not claim this is an eclipse in a classical sense, neither from the primary stellar disk nor from its wind.   The He\\,{\\sc i}~$\\lambda$10830 line indicates that material along our line-of-sight latitude reaches much higher negative velocities ($-$1400~km~s$^{-1}$ ) than the typical velocities we see in the high excitation phases ($-$550~km~s$^{-1}$). Most likely the P Cygni absorption is due to secondary gas that flows through the walls of the shock-cone and is left behind as the secondary star turns around the center of mass at periastron. Variations in the V and R line peaks indicate that there are other emitting regions for this spectral line, in addition to the Weigelt objects and the winds of the companion stars.  Although the observations reported here are compatible with the generic binary scenario presented in section \\ref{binary}, none of the existing models describes the whole data set. Numerical 3-D simulation of two-colliding winds for parameters appropriate to the $\\eta$ Carinae system will be needed to assist in understanding the complex variability that is observed. Work in progress \\citep{Okasaki} is very promising, having been able to reproduce the near-periastron collapse, compatible with data we presented here.  If we are to understand $\\eta$ Carinae and its evolutionary state, the identification of genuine Keplerian velocity-shifts, and the determination of the masses and spectral types of both stars, is a crucial goal. The intensive efforts applied to understand $\\eta$~Carinae, which is a difficult system because of extreme conditions (masses, mass-loss, orbital eccentricity), brings benefits to the field of colliding wind binaries. Binarity in eta Car may not have been responsible for the 1843 giant eruption, and so does not explain the other examples of extreme LBVs, like P Cygni, V1 in N2363, etc. However, the secondary star in $\\eta$~Carinae probes different  regions of the primary's wind and, hopefully, will enable the determination of its mass.  This is helpful to understanding the parameters of evolved massive stars, which are difficult to diagnose when in isolation.  The next minimum is predicted to start on 2009 January 11 ($\\pm$2 d). Fig. \\ref{2009} is useful to plan monitoring campaigns through the {\\it collapse} component. The time frame starts in early 2008 November and extends through 2009 July. This is the best event since 1948 for ground-based observations, since its central core fits entirely in the good observing season. The next favorable event will not occur before 2020. The critical time for narrow components of the high excitation lines and He\\,{\\sc i} encompasses the time interval December 15 to January 20. The P Cygni absorption and the $N_{\\rm H}$ to the X-ray source will reach the maximum around February 10. In order to improve significantly the observations made in the 2003.49 event it is necessary to monitor this critical time interval more frequently than one observation per day.  Since we now have a better picture of the events, we can plan for the next critical observations focused on a few representative lines. Major gains, relative to the present knowledge, will come from higher spatial resolution and signal-to-noise observations. Signatures from high energy phenomena, like X-rays, He\\,{\\sc ii}~$\\lambda$4686 and the Balmer jump  are especially important."
    },
    "0711/0711.1195_arXiv.txt": {
        "abstract": "The X-ray burster \\src\\  shows extremely regular Type I X-ray bursts whose energetics and recurrence times agree well with thermonuclear ignition models. We present calculations of sequences of burst lightcurves using  multizone models which follow the nucleosynthesis ($\\alpha$p and rp-processes) with an extensive nuclear reaction network. The theoretical and observed burst lightcurves show remarkable agreement. The models naturally explain the slow rise (duration $\\approx 5\\ {\\rm s}$) and long tails ($\\approx 100\\ {\\rm s}$) of these bursts, as well as their dependence on mass accretion rate. This comparison provides further evidence for solar metallicity in the accreted material in this source, and gives a distance to the source of $(6.07\\pm 0.18)\\ {\\rm kpc}\\ \\xi_b^{-1/2}$, where $\\xi_b$ is the burst emission anisotropy factor. The main difference is that the observed lightcurves do not show the distinct two-stage rise of the models. This may reflect the time for burning to spread over the stellar surface, or may indicate that our treatment of heat transport or nuclear physics needs to be revised. The trends in burst properties with accretion rate are well-reproduced by our spherically symmetric models which include chemical and thermal inertia from the ashes of previous bursts. Changes in the covering fraction of the accreted fuel are not required. ",
        "introduction": "The basic physics of Type I X-ray bursts as thin shell flashes on the surfaces of accreting neutron stars was understood many years ago (e.g.~\\citealt{fuj81,LE95}). Nonetheless, detailed comparisons of observations and theory are often less than successful \\citep{fuj87,vplewin}. The burster \\src\\ (aka Ginga 1826-24) is an important exception. \\cite{Ubertini99} dubbed it the ``clocked burster'' because of its extremely regular bursting behavior. They found a burst recurrence time close to 6 hours with a dispersion of approximately 6 minutes. \\cite{Bildsten00} dubbed it the ``textbook burster'' because of the good agreement with theory. He noted that the burst energetics and recurrence times were as expected for the inferred accretion rate $\\dot M\\approx 10^{-9}\\ M_\\odot\\ {\\rm yr^{-1}}$, and proposed that the long burst tails, lasting $\\approx 100\\ {\\rm s}$, were powered by rp-process hydrogen burning \\citep{WW81,HF84}. The rp-process involves a series of proton captures and beta-decays on heavy nuclei close to the proton drip line, for which there are considerable uncertainties in beta decay rates, reaction rates, and nuclear masses. Type I X-ray bursts offer an important test of our understanding of this process.  \\cite{gal04} (hereafter G04) studied 24 bursts observed from \\src\\ by the {\\em Rossi X-ray Timing Explorer (RXTE)} between 1997 November and 2002 July, and carried out a detailed comparison of the observed recurrence times and energetics with theory. During this period, the accretion rate (assumed to be proportional to the persistent X-ray luminosity) increased by $\\approx 50$\\%, while the burst recurrence time decreased from $\\Delta t\\approx 6$ hours to $\\approx 4$ hours. At the observed accretion rate $\\dot M\\approx 10^{-9}\\ M_\\odot\\ {\\rm yr^{-1}}$, hydrogen burns stably between bursts by the beta-limited hot CNO cycle \\citep{HF65}, heating the accumulating layer of hydrogen and helium at a rate that depends only on the mass fraction of CNO elements. The resulting  ignition column depth of $\\approx 2\\times 10^8\\ {\\rm g\\ cm^{-2}}$ is nearly independent of $\\dot M$ (e.g.~\\citealt{CB00}), in agreement with the observed scaling which is close to $\\Delta t\\propto 1/\\dot M$.  G04 used the variation of burst energetics with accretion rate to constrain the composition of the accreted material. The energetics are measured by the ratio $\\alpha$ of integrated persistent flux between bursts to the burst fluence, equivalent to the ratio of the gravitational energy release from accretion to the nuclear energy release in bursts. The observed value $\\alpha\\approx 40$ is close to the value expected for the $4$--$5\\ {\\rm MeV}$ per nucleon energy release in hydrogen burning \\citep{Bildsten00}. G04 found that $\\alpha$ decreased from $\\approx 44$ in 1997 to $\\approx 40$ in 2002. This  is consistent with the change in composition of the fuel layer expected from hot CNO burning between bursts, if the accreted material has solar metallicity. For a shorter recurrence time, less hydrogen burns between bursts, resulting in a larger burst energy and smaller $\\alpha$. In addition, G04 found that whereas the ignition models predicted that the ignition column depth should increase slightly with $\\dot M$, the observed trend was the opposite.  The study of G04 did not include models of the burst lightcurves. In this paper, we compare observations of \\src\\ with multizone models \\citep{woo04}, which include a large nuclear reaction network to follow the rp-process in detail. We show that these models naturally explain the slow rise (duration $\\approx 5\\ {\\rm s}$) and long tails ($\\approx 100\\ {\\rm s}$) of these bursts. We show that our time-dependent models reproduce the scalings of the recurrence time with accretion rate for a metallicity near solar.      \\begin{figure} \\epsscale{1.6} \\plotone{f1.eps} \\caption{Comparison of observed and calculated lightcurves. The histogram shows the average lightcurve from the bursts observed during the year 2000 when the recurrence time was $\\approx 4$ hours (G04, Figure 2). The error bars are the 1$\\sigma$ variations from burst to burst. The solid and dashed curves are the average burst profiles from models A3 ($Z=0.02$) and B3 ($Z=0.001$), which have $\\Delta t=3.9$ and $4.0$ hours respectively. The inset magnifies the rise and the initial part of the decay. The grey bands indicate the 1$\\sigma$ variation of the burst profiles about the average.} \\label{fig:lc} \\end{figure} ",
        "conclusions": "\\tablenotemark{a}} \\tablewidth{0pt} \\tablehead{ \\colhead{Model} & \\colhead{Number} & \\colhead{Z} & \\colhead{$\\dot M$} & \\colhead{$\\Delta t$} & \\colhead{$E_{\\rm burst}$} & \\colhead{$\\alpha$} & \\colhead{$\\Delta M$} \\nl \\colhead{} & \\colhead{of bursts} & \\colhead{} & \\colhead{($10^{-9}\\ M_\\odot\\ {\\rm yr^{-1}}$)} & \\colhead{(h)} & \\colhead{($10^{39}\\ {\\rm ergs}$)} & \\colhead{} & \\colhead{($10^{21}\\ {\\rm g}$)}\\nl } \\startdata A1 & 19 & 0.02 & 1.17 & 5.4 (0.1) & 4.67 (0.20) & 57.4 (2.8) & 1.14 (0.03)\\nl A2 & 18 & 0.02 & 1.43 & 4.3 (0.1) & 4.67 (0.11) & 55.6 (1.2) & 1.11 (0.03)\\nl A3 & 30 & 0.02 & 1.58 & 3.85 (0.06) & 4.73 (0.07) & 55.0 (0.9) & 1.10 (0.02)\\nl A4 & 13 & 0.02 & 1.75 & 3.48 (0.06) & 4.84 (0.06) & 53.6 (0.8) & 1.11 (0.02)\\nl B1 & 12 & 0.001 & 1.17 & 12.8 (0.6) & 13.3 (0.7) & 47.8 (0.4) & 2.71 (0.14)\\nl B2  & 17 & 0.001 & 1.43 & 6.04 (0.41) & 7.74 (0.49) & 47.4 (1.0) & 1.57 (0.11)\\nl B3  & 15 & 0.001 & 1.75 & 3.98 (0.28) & 6.26 (0.32) & 47.3 (2.4) & 1.27 (0.09)\\nl \\enddata \\tablenotetext{a}{See text for definitions of quantities.} \\end{deluxetable*}     \\begin{figure} \\epsscale{1.2} \\plotone{f2.eps} \\caption{Burst parameters for the $Z=0.02$ models A1--A4 compared with individual observed bursts. The properties of the observed bursts are plotted as diamonds. The properties of individual model bursts are plotted as crosses, with average values for each model plotted as squares joined by solid lines. } \\label{fig:alpha} \\end{figure}    Once the bursts have been normalized in this way, the agreement between the observed and theoretical lightcurves for $Z=0.02$ (model A3) is remarkable. Model A3 fits the observed decay exceedingly well out to about 40s, falling a little below the observed flux between 40 and 120s. The burst to burst variations in the models are also of comparable magnitude to the burst to burst variations in the data. The most significant difference is that the theoretical model shows a distinct two-stage rise which is not apparent in the observed lightcurve (see the lower panel of Fig.~1).  Model B3, which has a low metallicity, does not reproduce the observed lightcurve. Given the uncertainty in the distance to the source, the normalization may be adjusted to bring the observed and predicted peak luminosities into agreement, but the shape of the decay provides an additional constraint. In model B3, the lower metallicity leads to very little hydrogen burning between bursts, giving less helium and more hydrogen at ignition than in the solar metallicity model. The result is a  burst with a much longer tail, inconsistent with the observed profile for any distance.     \\begin{figure} \\epsscale{1.6} \\plotone{f3.eps} \\caption{Change in the lightcurve as a function of accretion rate. Averaged lightcurves from models A1 (dotted curve), A2 (solid curve), and A3 (dashed curve). The grey shaded areas indicate the standard deviation of the burst to burst variation  around the average lightcurve. As accretion rate increases, the burst becomes less luminous, with a longer rise time and slower decay.} \\label{fig:lc3} \\end{figure}    \\begin{figure} \\epsscale{1.6} \\plotone{f4.eps} \\caption{Fitted exponential decay timescales for observed bursts (open symbols) and theoretical lightcurves from models A1, A2, A3 and A4 (solid symbols). For the models, we choose a representative sample of 5 bursts (burst numbers 7 to 11 from each run), and displace them slightly on the horizontal axis for clarity. Note that some of the theoretical lightcurves show a three stage decay. The uncertainties in the fitted decay times are comparable to the scatter in the points. We adopt the same relation between $F_X$ and $\\dot M$ as in Figure \\ref{fig:alpha}.} \\label{fig:times} \\end{figure}  In Figure \\ref{fig:alpha}, we compare the burst recurrence times, energies, $\\alpha$ values, and ignition masses as a function of accretion rate for the observed bursts and the solar metallicity models A1--A4. The observed burst fluence $E_b$ is related to the burst energy $E_{\\rm burst}$ by $E_b=4\\pi d^2\\xi_bE_{\\rm burst}$, and we adopt the same value of $\\xi_bd^2$ that we determined earlier by comparing the burst peak luminosity with the observed peak flux. We convert the bolometric flux $F_X$ to accretion rate $\\dot M$ using the relation $4\\pi d^2\\xi_pF_X=\\dot Mc^2 z/(1+z)$, where $\\xi_p$ is a factor that accounts for possible anisotropy in the persistent emission (e.g.~Fujimoto 1988). For the same value of distance, we find that $\\xi_p/\\xi_b=1.55$ is required for the recurrence times to match in Figure \\ref{fig:alpha}. The theoretical alpha values have been divided by this same factor to give the correct prediction for alpha including anisotropic emission.  The agreement between the models and observations in Figure \\ref{fig:alpha} is excellent. The main difference is that the burst energy is overpredicted by approximately 5\\%, and the $\\alpha$ value underpredicted by the same factor. The low metallicity models B1--B3 are not consistent with the observed trends. For example, the burst energies show a factor of 2 variation in these models over the observed accretion rate range (Table 1), much larger than the $\\approx 10$\\% variations seen in the data.  Figure \\ref{fig:lc3} shows how the burst lightcurves change with accretion rate (models A1 to A3). G04 pointed out that the burst lightcurves changed with recurrence time. They fitted the decay of the bursts to an exponential profile with a break in the exponential decay time. The theoretical curves show a similar two timescale decay (this is particularly apparent in the lower panel of Fig.~1). G04 noted a significant increase in the first exponential decay timescale as $\\Delta t$ decreased, from $14.7\\pm 0.7$s when $\\Delta t\\approx 6$ hours to $17.5\\pm 1.1$s when $\\Delta t\\approx 4$ hours. The second exponential decay timescale $43\\pm 1$s showed no significant change. G04 speculated that this change was a result of the composition of the fuel, since more hydrogen is depleted between bursts for larger $\\Delta t$.  Figure \\ref{fig:lc3} shows that the model lightcurves show a similar change with accretion rate to the observed lightcurves, presumably relating to the changing hydrogen/helium ratio at ignition. This is quantified in Figure \\ref{fig:times}, which shows fitted exponential decay times for the observed and model bursts.  Overall, the agreement is good, although the decay times of the model bursts are generally slightly shorter than observed (consistent with the difference in the burst tails in Figure 1), and they show much larger burst to burst variations than observed. Also, whereas the observed bursts are well fit by a two-timescale decay, some of the model bursts show a more complex decay profile, which we have fitted with three timescales."
    },
    "0711/0711.4854_arXiv.txt": {
        "abstract": "In this paper, we apply the statefinder diagnostic to the cosmology with the Abnormally Weighting Energy hypothesis (AWE cosmology), in which dark energy in the observational (ordinary matter) frame results from the violation of weak equivalence principle (WEP) by pressureless matter.  It is  found that there exist closed loops in the statefinder plane, which is an interesting characteristic of the evolution trajectories of statefinder parameters  and can be used to  distinguish AWE cosmology from the other cosmological models. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.4912_arXiv.txt": {
        "abstract": " ",
        "introduction": "The formation and evolution of massive ($>10$M\\sun) stars plays a key role in the final fate of their parental molecular clouds, and in the appearance and evolution of their host galaxies. They inject large amounts of mechanical and radiative energy creating, either by a single star or a stellar cluster, the most spectacular gaseous nebulae in the Cosmos. Also, they generate fast shocks which heat up the surrounding plasma to temperatures above 10$^6$ K, and carve large interstellar \"holes\" that continuously stir the general interstellar medium.  From the moment they are born their powerful outflows begin to plough into the surrounding molecular material. Similarly, the output of Lyman continuum radiation, that sets them apart from lower mass stars, ionize the cloud and form a dense and hot H~II region. The expansion of this photoionized nebula drives a shock wave that both compress and may trigger further star formation on the one hand, and disperse a large portion of the molecular cloud on the other. As the new OB star enters the field population its winds and ultra-violet radiation continue to influence the general interstellar medium. These effects strengthen after their short main sequence lifetime, through the supergiant and Wolf-Rayet phases, and culminates in the most energetic of stellar events, the supernova explosion. Thus, the energy input from massive stars, via the combined effects of expanding H II regions and supernova remnants, can shape the interstellar medium of gaseous galaxies, creating large, expanding structures that may even vent mass and energy into the halo. It is this litany of energetic phenomena that gives massive stars such a pivotal role in astrophysics. Hence, it is not surprising to find that the dynamics of these phenomena have figured strongly in John's work over the years. ",
        "conclusions": ""
    },
    "0711/0711.1530_arXiv.txt": {
        "abstract": "The evolution over time of the magnetic activity and the resulting X-ray and UV coronal and chromospheric emissions of main-sequence dG, dK, and dM stars with widely different ages are discussed. Young cool stars spin rapidly and have correspondingly very robust magnetic dynamos and strong coronal and chromospheric X-ray -- UV (XUV) emissions. However, these stars spin-down with time as they lose angular momentum via magnetized winds and their magnetic generated activity and emissions significantly decrease. Studies of dK--dM stars over a wide range of ages and rotations show similar (but not identical) behavior. Particular emphasis is given to discussing the effects that XUV emissions have on the atmospheres and evolution of solar system planets as well as the increasing number of extrasolar planets found hosted by dG--dM stars. The results from modeling the early atmospheres of Venus, Earth and Mars using recently determined XUV irradiances and inferred winds of the young Sun are also briefly discussed. For example, the loss of water from juvenile Venus and Mars can be explained by action of the strong XUV emissions and robust winds of the young Sun. We also examine the effects of strong X-ray and UV coronal and chromospheric emissions (and frequent flares) that dM stars have on possible planets orbiting within their shrunken habitable zones (HZs) -- located close to the low luminosity host stars (HZ $\\leq$ 0.4 AU). Dwarf M stars make interesting targets for further study because of their deep outer convective zones (CZs), efficient dynamos, frequent flares and strong XUV emissions. Furthermore, a large fraction of dM stars are very old ($>$5 Gyr), which present intriguing possibilities for the development of highly advanced modes of intelligent life on planets that may orbit them. ",
        "introduction": "\\vspace{-2mm}  Over the last 15 years, as part of the ``Sun in Time'' program, we have been carrying out an in-depth study of solar-type stars with different ages, rotation rates and correspondingly vastly different levels of magnetic activity. The main goals of this program are to investigate the solar dynamo, and to determine the dependence of coronal X-ray/EUV emissions and chromospheric/transition region (TR) FUV/UV emissions on stellar age and rotation. The latter goal was expanded to include the contruction of XUV spectral irradiance tables for the Sun over its main sequence lifetime (e.g. Guinan et al. 2003; Ribas et al. 2005). These spectral XUV irradiances are very important in studies of the effects of the young Sun's increased radiation on paleoplanetary atmospheres and environments. We use a homogeneous sample of $\\sim$15 single, nearby dG0--5 stars as proxies for the Sun (and solar-type stars), covering ages from $\\sim$100 Myr to 8.5 Gyr. The ages of the younger program stars are inferred from memberships in various clusters / moving groups, while the ages for the older stars are secured from isochronal fits. The rotation periods are obtained chiefly from ground-based photometry (star spot modulation) or from chromospheric Ca {\\sc ii} rotationally induced variability. Excellent relationships are found among rotation, age and activity (see Guinan et al. 2003). The program stars have been observed from X-ray--UV using {\\em ROSAT}, {\\em ASCA}, {\\em Chandra}, {\\em XMM}, {\\em EUVE}, {\\em FUSE}, {\\em IUE} and {\\em HST}.  The strength of our sample is its homogeneity. In essence, we use a $\\sim$1 M$_\\odot$ star as a laboratory to study the effects of the stellar dynamo (with similar convective zone (CZ) depths) by varying the only free parameter, P$_{\\rm rot}$. Our studies have shown that the Sun's magnetic field has steadily declined as its rotation slowed due to magnetic breaking. Also, the study of the young Sun's (as well as other dG stars) XUV fluxes using solar proxies reveals unprecedented diagnostics for the state of the younger solar system and the physics of the much more active early Sun (Ayres et al. 1996; Guinan et al. 2003; Ribas et al. 2005). Fig. 1 shows a plot of rotation period vs. stellar age for most of our ``Sun in Time'' sample, with a power law fit, showing the spin-down of the Sun and other solar-type stars with age. Fig. 2 shows the related decrease in X-ray luminosity (L$_{\\rm x}$) with age for the ``Sun in Time'' sample. As shown in these figures, the young Sun was rotating over 10x faster than the present Sun and had correspondingly high (5--1000x) levels of XUV coronal and chromospheric emissions. Most of the observed ranges in L$_{\\rm x}$ for the older, less active stars shown in Fig. 2 arise from solar-like activity cycles. Some results of the ``Sun in Time'' program have been summarized by G\\\"{u}del et al. (1997), Guinan et al. (2003), Guinan \\& Ribas (2004) and Ribas et al. (2005).  \\begin{figure} \\includegraphics[width=4in]{GStars.eps} \\vspace{-2mm} \\caption{Plot showing the increase in P$_{\\rm rot}$ for dG0--5 stars with increasing age. This spin-down with age arises from magnetic breaking from angular momentum loss via magnetized winds.}\\label{fig1} \\vspace{-4mm} \\end{figure}  \\begin{figure} \\includegraphics[width=4in]{IAUS_GX.eps} \\vspace{-2mm} \\caption{The variation of X-ray luminosity (L$_{\\rm x}$) for solar-type dG0--5 stars are shown and plotted against age. The ranges in L$_{\\rm x}$ for the Sun and other older solar-type stars arise mainly from activity cycles.}\\label{fig2} \\vspace{-8mm} \\end{figure}  The ``Sun in Time'' program also bears on the crucial question of the influence of the young Sun's strong XUV emissions on the developing planetary system -- in particular on the photochemical and photoionization evolution of early planetary atmospheres. The constructed spectral irradiance tables are of interest to researchers in paleoplanetary atmospheres and for studies on the atmospheric evolution of the large number of extrasolar planets found orbiting other G-type stars. The irradiance data have been published by Ribas et al. (2005) and examples in the FUV/UV region are given in Fig. 3. As shown, there are dramatic decreases in the FUV ({\\em FUSE}) and FUV/UV ({\\em IUE}) emission line strengths (and irradiances) as the star loses angular momentum and spins down with age. Fig. 4 shows the averaged XUV irradiance changes over time for the Sun and other solar-type (dG0--5) stars from solar proxies of different ages. The plot illustrates how emissions associated with hotter plasmas diminish more rapidly as the stars rotate slower with age. The coronal X-ray/EUV emissions of the young main sequence Sun were approximately 100--1000x stronger than the present Sun. Similarly, the TR and chromospheric FUV \\& UV emissions of the young Sun are 10--100x and 5--10x stronger, respectively, than at present. Also shown in Fig. 4 is the slow increase ($\\sim$8\\%/Gyr) in the bolometric luminosity of the Sun over the past 4.6 Gyr. Over this time, the Sun's luminosity increased from $\\sim$0.7 L$_\\odot$ to 1.0 L$_\\odot$. This change in solar luminosity arises from the acceleration of nuclear fusion in its core that causes its radius and surface temperature to slowly increase with time.  \\begin{figure} \\includegraphics[width=5in]{IAUS_G2.eps} \\vspace{-3mm} \\caption{\\small Plots of flux density at 1 AU for the {\\em FUSE} O {\\sc vi} (1030--1040\\AA) region (left) and the FUV--UV 1200-1700\\AA~region from {\\em IUE} low dispersion {\\em SWP} spectrophotometry (right). As shown in the figures the flux densities decrease rapidly with stellar age. The stars are plotted from top to bottom in order of increasing age and decreasing activity.}\\label{fig3} \\vspace{-5mm} \\end{figure}  \\begin{figure} \\includegraphics[width=4.5in]{IAUS_G3.eps} \\vspace{-3mm} \\caption{\\small The smoothed XUV spectral irradiances of solar-type stars for three spectral intervals plotted against age. The flux scales were normalized to unity at the Sun's age. Also shown is the variation of bolometric solar luminosity (L/L$_\\odot$) over time, where L$_\\odot$ = 1 is the luminosity of the present 4.6 Gyr Sun (L$_\\odot$).}\\label{fig4} \\vspace{-5mm} \\end{figure}  We have been collaborating with planetary scientists and astrobiology groups to study the effects of the young Sun's  (and other dG0--5 stars) strong XUV irradiance on the loss of water from Mars and its implications for the oxidation of the Martian soil (Lammer et al. 2003). It has been assumed, from topographic and geological studies, that Mars was originally warmer and much wetter than at present, and likely possessed a $\\sim$1 bar atmosphere. Lammer et al. considered ion pick-up sputtering, as well as dissociative recombination processes. The loss of H$_2$O from Mars over the last 3.5 Gyr was estimated to be equivalent to a global martian H$_2$O ocean with a depth of about $\\geq$12 m. If ion momentum transport, a process to be studied in detail by Mars Express, is significant on Mars, the water loss may be enhanced by a factor of about 2. For this study it has been assumed that, for the first billion years, Mars had a hot liquid iron core and, through rotation, possessed a significant magnetic field and resulting magnetosphere. This magnetosphere essentially protected the early Martian environment from the combination of high levels of XUV radiation and strong winds of the younger Sun that would have otherwise removed its atmosphere. However, Mars is a smaller, less massive planet than the Earth, and lost heat at a faster rate. Thus its iron core solidified $\\sim$1 billion years after the planet's formation. Without the magnetic field, the outer Martian atmosphere was exposed to the ionizing effects and strong winds of the early Sun, and thus partially eroded. Photolysis of water ensued, with a preferential loss of the lighter Hydrogen over the heavier Oxygen. The loss of water and water vapor from the atmosphere resulted in a greatly diminished Greenhouse Effect and a rapid cooling of the lower Martian atmosphere. This rapid cooling (dropping below the freezing point of water) permitted some water to remain behind, possibly as permafrost trapped below the Martian soil. This scenario is illustrated in Fig. 5.  A recent study (Kulikov et al. 2006) has been published using solar-proxy data (from Ribas et al. 2005 and Wood et al. 2002; 2005) to study the atmosphere and water loss from early Venus. This study shows that water on early Venus (only 0.71 AU from the Sun) was essentially all lost during the first $\\sim$0.5 Gyr after its formation from the vigorous action of strong (massive) winds and high XUV fluxes from the young Sun. In another study, by Grie$\\ss$meier et al. (2004), ``Sun in Time'' irradiance date were used to investigate the atmospheric loss of extrasolar planets resulting from XUV heating, which can lead to the evaporation of ``Hot Jupiters.''  \\begin{figure} \\includegraphics[width=4.5in]{IAUS_G5.eps} \\vspace{-3mm} \\caption{Cartoon depicting the early evolution and ensuing erosion/loss of the Martian atmosphere. (Based on Lammer et al. 2003)}\\label{fig5} \\vspace{-3mm} \\end{figure}  \\vspace{-2mm} ",
        "conclusions": ""
    },
    "0711/0711.4123_arXiv.txt": {
        "abstract": "We measure the distribution of velocities for prograde and retrograde satellite galaxies using a combination of published data and new observations for 78 satellites of 63 extremely isolated disc galaxies (169 satellites total). We find that the velocity distribution is non-Gaussian ($>$ 99.9\\% confidence), but that it can be described as the sum of two Gaussians, one of which is broad ($\\sigma = 176 \\pm 15$ km s$^{-1}$), has a mean prograde velocity of  $86\\pm30$ km s$^{-1}$, and contains $\\sim$ 55\\% of the satellites, while the other is slightly retrograde with a mean velocity of $-21 \\pm 22$ km s$^{-1}$ and $\\sigma = 74 \\pm 18$ km s$^{-1}$ and contains $\\sim$ 45\\% of the satellites. Both of these components are present over all projected radii and found in the sample regardless of cuts on primary inclination or satellite disc angle.  The double-Gaussian shape, however, becomes more pronounced among satellites of more luminous primaries.  We remove the potential dependence of satellite velocity on primary luminosity using the Tully-Fisher relation and still find the velocity distribution to be asymmetric and even more significantly non-Gaussian. The asymmetric velocity distribution demonstrates a connection between the inner, visible disc galaxy and the kinematics of the outer, dark halo. The reach of this connection, extending even beyond the virial radii, suggests that it is imprinted by the satellite infall pattern and large-scale effects, rather than by higher-level dynamical processes in the formation of the central galaxy or late-term evolution of the satellites. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.2639_arXiv.txt": {
        "abstract": "In these lecture notes, the role of strangeness in relativistic astrophysics of compact stars is addressed. The appearance of strange particles, as hyperons, kaons, and strange quarks, in the core of compact stars is examined and common features as well as differences are presented. Impacts on the global properties of compact stars and signals of the presence of exotic matter are outlined for the various strange phases which can appear in the interior at high densities. ",
        "introduction": "Neutron stars are the final endpoint in the evolution of stars more massive than eight solar masses. They are born in a spectacular explosive event, a core-collapse supernova, which can outshine an entire galaxy in its brightness. A surprisingly good estimate about the characteristic masses and radii can be made by adopting the argument of Landau \\cite{Landau32}. The delicate balance between gravitational energy and Fermi energy can only be achieved up to a certain critical number of nucleons (fermions) which is $N\\approx (M_P/m_N)^3 \\approx 10^{57}$, where $M_P$ stands for the Planck mass and $m_N$ for the nucleon mass. The corresponding maximum mass for neutron stars amounts to $M_{\\rm max} \\approx M_p^3/m_N^2 \\approx 1.8 M_\\odot$. This mass estimate is exactly in the region of maximum masses presently discussed in the literature with much more refined theoretical models of neutron star matter. There are excellent modern textbooks on compact stars available, which discuss the physics of neutron stars and its various phases in great detail. We refer the interested reader to the books by Norman Glendenning \\cite{Glen_book} and by Fridolin Weber \\cite{Weber_book} for a thorough treatment of this exciting field in the interplay of physics and astrophysics. Most recently, a new textbook on compact stars appeared which updates and complements the existing ones \\cite{Haensel_book}. A recent review article on the relation between the nuclear equation of state and neutron stars can be found in \\cite{Lattimer:2006xb}, a nice popular article on neutron stars in \\cite{Lattimer:2004pg}.  In these lecture notes, we discuss the global properties of neutron stars, masses and radii, and how those change for different compositions in their interior. In particular, we address the appearance of new and exotic phases in dense neutron star matter and how they can modify the mass-radius diagram for compact stars. In section~\\ref{sec:hyperons}, the baryons of the SU(3) octet are treated together in a chiral SU(3)$\\times$SU(3) effective Lagrangian which is motivated from the approximate symmetries of the underlying theory of strong interactions, quantum chromodynamics (QCD). When the properties of hyperons are tied to the known experimental data on hypernuclei, hyperons are present in neutron star matter above twice normal nuclear matter density affecting the properties of massive neutron star configurations. In section~\\ref{sec:kaons}, strange mesons, the $K^-$, and its presence in neutron star matter by forming a Bose-Einstein condensate are addressed. A relativistic field-theoretic approach is utilised to outline the basic features of kaon condensation for compact stars. Ultimately, hadrons have to be described in terms of quark degrees of freedom at high densities. The phase transition from the chirally broken hadronic phase to the approximately chirally restored quark matter phase is examined within perturbative QCD and the MIT bag model to illustrate the main features of a strong first order phase transition in compact star matter in section~\\ref{sec:quarks}. Characteristics of the mass-radius relation are highlighted which signal the presence of exotic matter in the core of compact stars. ",
        "conclusions": "Compact stars can have a rich structure with new and exotic phases being present in their cores. Hyperons, heavy baryons with strangeness, can appear in neutron star matter. As a new fermionic degree of freedom, hyperons lower the overall pressure so that the maximum mass of a compact star with hyperons is substantially smaller than the one of a neutron star consisting of nucleons and leptons only. Also, negatively charged Bose-Einstein condensate could be formed via kaon condensation at high densities. As the Bose-Einstein condensate does not give a direct contribution to the pressure of the system, the stable sequence of compact stars in the mass-radius diagram simply stops as soon as the kaon condensed phase appears or extremely dense compact stars are generated with unusually small radii. Finally, at high densities quark matter can be present. Depending on the onset of the phase transition and the strength of the first-order phase transition from the hadronic to the quark matter phase, a third family of compact stars is found in the mass-radius diagram. The compact stars of this third family have a core of pure quark matter and smaller radii than ordinary hybrid stars. The physics of dense quark matter in compact star opens exciting new perspectives in exploring the phase diagram of QCD at high densities. The new research facility FAIR at GSI Darmstadt will explore this fascinating regime of strong interaction physics in the near future with relativistic heavy-ion collisions with bombarding energies tuned so as to create compressed baryonic matter at highest baryon densities."
    },
    "0711/0711.2852_arXiv.txt": {
        "abstract": "We investigate the stability of black hole accretion disks in a primordial environment (POP III disks for short), by solving the vertical structure of optically thick disks, including convective energy transport, and by employing a one-zone model for optically thin isothermal disks. Because of the absence of metals in POP III disks, we find significant differences in stability associated with ionization between POP III disks and the disks of solar metallicity. An unstable branch in S-shaped equilibrium curves on the $\\dot{M}-\\Sigma$ (mass accretion rate - surface density) plane extends to a larger surface density compared with the case of disks of solar metallicity. The resulting equilibrium loci indicate that quasi-periodic oscillations in luminosity can also be driven in POP III disks, and their maximal luminosity is typically by an order of magnitude larger than that of the disks of solar metallicity. Such a strong outburst of POP III disks can be observed by future huge telescopes, in case that the mass is supplied onto the disks at the Bondi accretion rates in typical virialized small dark halos. ",
        "introduction": "According to recent advances in the formation theory of population III (POP III) stars, it has been understood that they could be very massive $(10^2-10^3M_{\\odot })$ \\citep{ON98,Abel00,BCL02}.  Those stars end up as black holes as massive as $10-10^3 M_\\odot$ (here after POP III BHs), in the case that the mass of the progenitor stars is within the range $40M_\\odot \\lsim M_* \\lsim 140M_\\odot$ or $M_*\\gsim 260M_\\odot$ \\citep{HW02}. The accretion disks formed around such black holes can emit a significant amount of soft X-rays, which can contribute to reionization of the universe \\citep{Ricotti04}. Moreover, such black holes could be candidates for the progenitors of ultra-luminous X-ray sources \\citep{Mii05} or the building blocks of super-massive black holes \\citep{Rees84}.  The accretion disks surrounding POP III BHs are formed from the accreting gas around the progenitor stars.  In the case that the mass of a POP III star is within the range $40M_\\odot \\lsim M_* \\lsim 140M_\\odot$ or $ M_* \\gsim 260M_\\odot$, the star forms a black hole directly without a SN \\citep{HW02}. In this case, the accreting gas is expected to contain a small amount of heavy elements. In addition, recent numerical simulations on low-mass galaxy formation, taking into account the effects of SNe feedback, suggest that the metallicity of the gas remain in such galaxies during their early star-burst phase is as low as $\\sim 10^{-4} Z_\\odot$ \\citep{wada03}. Therefore, the metallicity of the accretion disks surrounding POP III BHs is expected to be very low ($Z \\lsim 10^{-4}Z_\\odot$). We hereafter call these accretion disks as POP III disks.  The thermal instability of the standard accretion disk has been investigated in detail (e.g, \\authorcite{SH75} \\yearcite{SH75}; \\authorcite{SS76} \\yearcite{SS76}; \\authorcite{Pringle76} \\yearcite{Pringle76}). By assuming solar metallicity (hereafter POP I disks), Mineshige and Osaki (\\yearcite{MO83}) obtained the S-shaped equilibrium loci on the $\\dot{M}-\\Sigma$ (mass accretion rate v.s. surface density) plane, implying that the unstable disk undergoes limit-cycle oscillation between the upper, hot, ionised branch and the lower, cool neutral branch (see also \\authorcite{Hoshi79} \\yearcite{Hoshi79}; \\authorcite{Meyer81} \\yearcite{Meyer81}; \\authorcite{Smak82} \\yearcite{Smak82}; \\authorcite{aobon} \\yearcite{aobon} for a review). It is so-called dwarf-nova type instability, which successfully explains the observed properties of dwarf nova \\citep{Meyer84}. An X-ray nova is also thought to be caused by such a disk instability, although its central object is expected to be a black hole \\citep{HW89,MW89}. The limit-cycle behavior for active galactic nuclei has been investigated by \\citet{LS86} (see also \\authorcite{MS90} \\yearcite{MS90}).   Mayer and Duschl (\\yearcite{MD05a}) investigated this instability on the POP III disks in which they also obtained the equilibrium loci on the $\\dot{M}-\\Sigma$ plane. However, their results based upon a one-zone approximation, even if the disk is optically thick, in which the energy transfer due to convection along the direction perpendicular to the equatorial plane of the disk is not included. As a result, the equilibrium loci are not S-shaped, which does not directly lead to the limit-cycle oscillations driven by thermal and secular instabilities. Thus, we need a more sophisticated approach in order to understand the nature of the stability of POP III disks. In this paper, we re-examine stability of POP III disks, by solving the vertical structure of optically thick disks, including the effect of convection. For optically thin disks, we adopt a one-zone approximation. In addition, we evaluate the amplification of the luminosity of accretion disks surrounding POP III BHs by the limit-cycle oscillation. We also discuss the possibility to detect the POP III disks by calculating their emergent spectra in the burst phase.  This paper is organised as follows: we describe basic equations in section 2. We compare the stability of POP III disks with POP I disks in section 3. In section 4, we discuss the observational possibility of POP III black hole accretion disks, and summarise in the final section. ",
        "conclusions": "By solving the vertical structure of the optically-thick accretion disks, including the convective energy transport, and by employing the one-zone model for optically thin isothermal disks, we investigate the stability of the disks in primordial environments.  POP III disks ($Z\\lsim 10^{-4}Z_\\odot$) are thermally and secularly unstable, and exhibit limit-cycle oscillations, like the POP I disks ($Z\\sim Z_\\odot$). The outbursts of POP III disks are stronger than those of POP I disks. The maximal luminosity in the burst state is expected to be an order of magnitude larger in POP III disks than in POP I disks.   The expected flux densities of POP III disks surrounding black holes of $M=10^3M_\\odot$ are $\\sim 10^{-34} {\\rm erg~s^{-1}~cm^{-2}~Hz^{-1}}$ around $1\\rm{\\mu m}$ in the case of $z\\sim 10$, whereas the expected number of sources is $O(10^3)$ per $10'\\times 10'$ field of view. Those disks are within reach of future facilities, such as OWL.   \\bigskip \\bigskip  Numerical calculations were carried out at Rikkyo University. This work was supported in part by a special postdoctoral researchers program in RIKEN (KO), and by Ministry of Education, Culture, Sports, Science, and Technology (MEXT) Young Scientists (B) 17740111(KO) and 17740110 (HS)."
    },
    "0711/0711.2580_arXiv.txt": {
        "abstract": "Highly relativistic jets are a key element of current gamma-ray burst models, where the jet kinetic energy is converted to radiation energy at optically thin shocks. High-energy neutrinos are also expected, from interactions of protons accelerated in the same shocks. Here we revisit the early evolution of a relativistic jet, while the jet is still inside the star, and investigate its neutrino emission. In particular we study propagation of mildly relativistic and ultrarelativistic jets through a type Ib progenitor, and follow reverse shocks as the jets cross the star. We show that protons can be accelerated to $10^4$--$10^5$ GeV at reverse shocks, and efficiently produce mesons. The mesons experience significant cooling, suppressing subsequent neutrino emission. We show, however, that the neutrino yield from the reverse shock is still reasonably large, especially for low-luminosity and long-duration jets, where meson cooling is less severe. We discuss implications of our results in the context of neutrinos from choked jets, which are completely shock heated and do not break out of the star. From a choked jet with isotropic equivalent energy of $10^{53}\\,\\mathrm{erg}$ at 10 Mpc, we expect $\\sim$20 neutrino events at IceCube. ",
        "introduction": "The leading model of gamma-ray bursts (GRBs) involves a relativistic ($\\Gamma_j\\gtrsim 100$, where $\\Gamma_j$ is the jet Lorentz factor) fireball jet, where the observed gamma rays are produced by radiation from Fermi-accelerated electrons in optically thin shocks (for reviews see e.g.~\\cite{Zhang:2003uk,Piran:2004ba,Meszaros:2006rc}). In the most popular model for the more common long-duration GRBs, the so-called collapsar model, the core of a massive star collapses to a black hole or neutron star, driving a highly relativistic jet which breaks out of the star~\\cite{Woosley:1993wj,MacFadyen:1998vz}. Within this scenario, high-energy neutrinos from photomeson interactions of accelerated protons have been studied~\\cite{Waxman:1997ti,Rachen:1998fd,Waxman:1998yy,Waxman:1999ai,Bahcall:1999yr,AlvarezMuniz:2000st,Meszaros:2001ms,Guetta:2002du,Dermer:2003zv,Li:2002dw,Murase:2006mm}. With $1\\,\\mathrm{km^3}$ size neutrino detectors such as IceCube~\\cite{Ahrens:2003ix} and KM3Net~\\cite{Katz:2006wv} being constructed, we are entering an age when we can test various predictions for high-energy neutrinos. Waxman and Bahcall predicted $\\epsilon_\\nu > 10^5\\,\\mathrm{TeV}$ neutrinos from external reverse shocks~\\cite{Waxman:1999ai}, and similar predictions have been made in Refs.~\\cite{Dermer:2003zv,Li:2002dw} for the forward shock. In addition, internal shocks can occur while the relativistic jet is still in the star. M\\'esz\\'aros and Waxman predicted a $\\epsilon_\\nu \\gtrsim 5 \\,\\mathrm{TeV}$ neutrino precursor burst from such a set up~\\cite{Meszaros:2001ms}. It is expected that a precursor neutrino burst is present even if the jet does not successfully break out of the star, producing a choked jet that is dark in gamma rays and bright in neutrinos.  Inferring from the observed rate of GRBs, highly relativistic ($\\Gamma_j \\gtrsim 100$) jets occur in perhaps $\\lesssim 10^{-3}$ of core-collapse supernovae~\\cite{Berger:2003kg}. However, the jet signature may be more common, and a significantly higher portion of supernovae may be endowed with slower, mildly relativistic jets; according to late-time radio observations, perhaps as high as a few $\\%$~\\cite{Totani:2003rk,Berger:2003kg,Soderberg:2004ma,Granot:2004rh,van Putten:2004dh}. The recent detection of low-luminosity GRB 060218 also suggests from detection rates that low-luminosity GRBs, with mildly relativistic ejecta, are more common than conventional high-luminosity GRBs by factors as large as $10^2$~\\cite{Campana:2006qe,Cobb:2006cu,Pian:2006pr,Soderberg:2006vh,Liang:2006ab,Toma:2006iu,Waxman:2007rr}. The occurrence rate is high enough that despite their lower luminosity, their contribution to the neutrino background is comparable to or larger than conventional GRBs~\\cite{Murase:2006mm,Gupta:2006jm}.  In addition to being more frequent, mildly relativistic jets are more baryon rich, which work positively for neutrino emission. Indeed, the detectability of neutrinos from internal shocks of mildly relativistic jets has been studied by Razzaque, M\\'esz\\'aros, and Waxman~\\cite{Razzaque:2004yv} and extended by Ando and Beacom~\\cite{Ando:2005xi} (see also~\\cite{Razzaque:2005bh,Koers:2007je}). In these scenarios, internal shocks occur at $r_s\\approx 10^{10}$--$10^{11}\\,\\mathrm{cm}$, smaller than the typical radius of type Ib progenitors. Accelerated protons interact via proton-photon ($p\\gamma$) and proton-proton ($pp$) interactions to produce mesons, which consequently produce neutrinos on decay.  In this paper we investigate neutrinos arising from protons that are accelerated in the reverse shock, while the relativistic jet is still propagating inside the star. We adopt a type Ib progenitor as our model star, but we discuss how our results apply to type Ic and II progenitors also. The reverse shock is strongly motivated because it is a generic prediction of the initial stages of any jet moving through the star, regardless of its Lorentz factor, and whether it breaks out of the star or not. Therefore, we consider two types of supernovae: one is the perhaps more numerous supernova containing a mildly relativistic ($\\Gamma_j \\sim 10$) jet, and the other is a supernova containing an ultrarelativistic ($\\Gamma_j \\sim 100$) jet that can ultimately cause the observed GRBs. In both cases we consider choked as well as successful jets, depending on its duration.  We show that protons can be accelerated to $10^4$--$10^5$ GeV in the reverse shock. The low and high-energy protons produce mesons efficiently by $pp$ and $p\\gamma$ interactions respectively. For energies in between, cooling mechanisms are faster, and result in a characteristic suppression feature in the meson spectrum. Mesons that are produced experience significant cooling due to the high proton and photon densities in the shocked jet. However, we still expect a reasonable number of neutrino events. For instance, for a choked mildly relativistic jet with isotropic energy $10^{54}$ erg at 10 Mpc, we expect $\\sim2$ neutrino events at IceCube from the reverse shock. Interestingly, assuming the same jet Lorentz factor, we find that a lower isotropic energy of $10^{53}$ erg produces more neutrino events ($\\sim20$ events). This is due to meson cooling. Specifically, we find that neutrinos are favored by high-Lorentz factor, low-luminosity, and long-duration jets, as meson cooling is less significant for such jets. The rate of core-collapse supernova within 10 Mpc is expected to be $\\sim$1--3 yr$^{-1}$~\\cite{Ando:2005ka}, implying particle acceleration in reverse shocks can potentially be tested in the next decades.  Neutrino emission from the reverse shock is potentially important for short-duration jets\\footnote{We note that the word ``short-duration'' here does not mean another class of GRBs, which is often referred to as short-duration and hard-spectrum GRB, and is believed to be associated with compact-star mergers~\\cite{Ando:2004pc,Nakar:2007yr}.} that are completely shocked and do not break out of the star, i.e., choked jets. We argue that while it is possible that internal shocks occur prior to jet shocking (in particular for low $\\Gamma_j$ jets),  the allowed range of radii for internal shocks to occur is constrained from above by the reverse shock, and below by the fact that the stellar material is opaque to high-energy neutrinos. This suggests that neutrino emission from choked jets are of similar magnitude to our estimates.  This paper is organized as follows. In Sec.~II we discuss the evolution of the jet as it propagates through the star. In Sec.~III we treat proton acceleration, and suppression of neutrino emission due to proton and meson cooling. In Sec.~IV we utilize the results of the previous sections to investigate the detectability of high-energy neutrinos from reverse shocks. Finally, we finish with discussions in Sec.~V and conclusions in Sec.~VI. ",
        "conclusions": "We showed that neutrino emission from reverse shock accelerated protons are strongly suppressed due to meson cooling in the hot jet and hot stellar material surrounding the reverse shock. Despite this suppression, we still expect a fairly promising number of events at IceCube, ranging from $2$ to $20$ for a close supernova of $10\\,\\mathrm{Mpc}$ hosting a choked mildly relativistic jet of total isotropic energy $10^{53}$--$10^{54}\\,\\mathrm{erg}$. We find that for a supernova of given energy $E_{\\rm{iso}}$, neutrino emission is favored by high-Lorentz factor, low-luminosity, and long-duration jets, for which meson cooling is less severe.  Our result sets the scale for neutrino emission from choked jets. For choked jets to provide more neutrinos, internal shocks are required to occur in the unshocked region of the jet. However, this is geometrically constrained from above by the reverse shock, and from below by the fact that the stellar matter is opaque to high-energy neutrinos for small radii.  Regarding neutrino precursor signals which precede GRB photons, we conclude that the reverse shock neutrinos likely plays a small role. However, they will precede those neutrinos from internal shocks, allowing the temporal profile of precursor neutrinos from a close supernova to provide an identification of proton acceleration at the reverse shock."
    },
    "0711/0711.3854_arXiv.txt": {
        "abstract": "The $^6$Li abundance observed in metal poor halo stars exhibits a plateau similar to that for $^7$Li suggesting a primordial origin.  However, the observed abundance of $^6$Li is a factor of $10^3$ larger and that of $^7$Li is a factor of 3 lower than the abundances predicted in the standard big bang when the baryon-to-photon ratio is fixed by WMAP.  Here we show that both of these abundance anomalies can be explained by the existence of a long-lived massive, negatively-charged leptonic particle during nucleosynthesis.  Such particles would capture onto the synthesized nuclei thereby reducing the reaction Coulomb barriers and opening new transfer reaction possibilities, and catalyzing a second round of big bang nucleosynthesis. This novel solution to both of the Li problems can be achieved with or without the additional effects of stellar destruction. ",
        "introduction": "It has recently been pointed out (e.g.,~\\cite{Asplund:2005yt})  that the abundances of both $^6$Li and $^7$Li observed in metal poor halo stars (MPHSs) are not in agreement with those predicted from standard big-bang nucleosynthesis (BBN). Specifically, the $^6$Li abundance as a function of metallicity exhibits a plateau similar to that for $^7$Li in very metal-poor stars, suggesting a primordial origin for both isotopes. This $^6$Li abundance plateau, however, is a factor of $\\sim 10^3$ larger than that predicted by BBN.  A less severe problem exists for $^7$Li; the BBN value based upon the baryon-to-photon ratio fixed by the WMAP analysis~\\cite{Spergel:2006hy} of the cosmic microwave background is roughly a factor of three higher than is observed.  Moreover, a longstanding effort in cosmology has been the search for evidence of unstable particles that might have existed in the early universe.  It is thus natural to ask whether the two lithium abundance anomalies might be a manifestation of the existence of such a particle.  In this context a number of possible solutions to the $^6$Li problem have been proposed which relate the Li anomalies to the possible existence of unstable particles in the early universe~\\cite{decaying}.  This has been extended in several recent studies~\\cite{Pospelov:2006sc,Cyburt:2006uv,Kaplinghat:2006qr,Kohri:2006cn,Bird:2007ge,Hamaguchi:2007mp} to consider heavy negatively charged decaying particles that modify BBN, but in rather different ways.  In these latter studies, the heavy particles, here denoted as $X^-$, bind to the nuclei produced in BBN to form $X$-nuclei.  The massive $X^-$ particles would be bound in orbits with radii comparable to those of normal nuclei.  Hence, they would reduce the reaction Coulomb barriers thereby enhancing the thermonuclear reaction rates and extending the duration of BBN to lower temperatures.  Pospelov~\\cite{Pospelov:2006sc} suggested that a large enhancement of the $^6$Li abundance could result from a transfer reaction involving an $X^-$ bound to $^4$He (denoted $^4$He$_X$), i.e. $^4$He$_X$($d$,$X^-$)$^6$Li.  Although this was an intriguing idea, Hamaguchi et al.~\\cite{Hamaguchi:2007mp} have recently pointed out via a more complete quantum mechanical calculation that Pospelov's estimate for the $^4$He$_X$($d$,$X^-$)$^6$Li cross section was too large; this leads to too high a $^6$Li abundance.  Cyburt et al.~\\cite{Cyburt:2006uv} further motivated this hypothesis by identifying the $X^-$ as the supersymmetric partner of the tau lepton, i.e., a stau, and considered $X^-$ transfer reactions for $^7$Li and $^7$Be production, too.  Although their calculation is based on a fully dynamical treatment of the recombination of nuclei with $X^-$ particles and also of BBN, they used cross sections for all $X^-$ transfer reactions involving $^6$Li, $^7$Li, and $^7$Be that were too large, as we discuss below.  Therefore, the calculated abundances are to be viewed with caution.  Kaplinghat and Rajaraman~\\cite{Kaplinghat:2006qr} observed that the decay of an $X^-$ when bound to $^4$He would occasionally knock out a proton or neutron to produce $^3$He or $^3$H, thereby enhancing their abundances and the abundance of $^6$Li through reactions with other primordial $^4$He nuclei at higher energies.  Kohri and Takayama~\\cite{Kohri:2006cn} studied the recombination of nuclei with $X^-$ particles, and suggested the possibility of solving the $^7$Li problem.  However, they did not carry out dynamical calculations involving recombination processes and BBN simultaneously.  This forced them to introduce artificial parameters for the fractions of the captured nuclei, which turn out to be different from the fractions obtained by solving the recombination plus BBN fully dynamically.  A new resonant reaction $^7$Be$_X$($p$,$\\gamma$)$^8$B$_X$ has recently been proposed by Bird et al.~\\cite{Bird:2007ge} that destroys $^7$Be$_X$ through an atomic excited state of $^8$B$_X$, and the present study identifies another effect in this reaction that might also destroy $^7$Be. Thus there has been a great deal of recent progress in $X^-$ catalyzed BBN in three important aspects: the simultaneous description of recombination and ionization processes of $X^-$ particles with nuclei in the description of BBN, use of updated reaction rates involving the $X$- nuclei, and inclusion of new resonant processes by which $^7$Be is destroyed. No previous calculation has involved all of these effects in a single dynamical calculation of BBN in order to study their effects on the $^6$Li and $^7$Li problems.  In this Article, we present the results of a thorough dynamic analysis of the effects of $X^-$ particles on BBN.  The important difference from previous works is, firstly, that we carried out a fully dynamical BBN calculation by taking account of the recombination and ionization processes of $X^-$ particles by normal and $X$-nuclei as well as the many possible nuclear reactions among them.  Secondly, the reaction rates on normal and $X$-nuclei used in the present study are based on quantum mechanical calculations of the cross section like those of Hamaguchi et al.~\\cite{Hamaguchi:2007mp}, which we believe to be correct.  Thirdly, we have not only included the important $^7$Be destruction mechanism identified by Bird et al.~\\cite{Bird:2007ge}, but have identified another potentially important destruction mechanism involving the reaction channel $^7$Be$_X$+$p$ $\\rightarrow ^8$B$^*$($1^+$, 0.770 MeV)$_X$ $\\rightarrow ^8$B$_X$+$\\gamma$ which has the potential to destroy $^7$Be$_X$ via capture through the $1^+$ nuclear excited state of $^8$B.  We show that when all these effects are included, the single hypothesis of the existence of the $X^-$ particle provides a remarkable solution to both the $^6$Li and $^7$Li abundance anomalies.  We can then use this constraint to place interesting limits on the $X^-$ relic abundance and its decay lifetime and mass. ",
        "conclusions": "\\label{sec4} In summary, we have investigated light-element nucleosynthesis during BBN taking into account the possibility of massive, negatively-charged $X^-$ particles which would bind to the light-nuclei.  When the chemical and kinetic processes associated with such particles are included in a BBN code in a fully dynamical manner, along with the reactions enabled by the $X^-$ particles, the $X^-$ particles are found to enhance the reaction rates in BBN, both by reducing the charge of the resulting $X$-nuclei, and by enabling transfer reactions of the $X^-$ particles.  $X^-$ particles greatly enhance the production of $^6$Li, primarily from the $X^-$ transfer reaction $^4$He$_X$($d$,$X^-$)$^6$Li.  The $^7$Li abundance, however, decreases when the $X^-$ particle abundance is larger than~0.1 times the total baryon abundance.  In this case, the $^7$Li abundance decreases with the $X^-$ particle abundance due to the inclusion of two resonance channels for  $^7$Be$_X$($p$,$\\gamma$)$^8$B$_X$ through the nuclear and atomic excited states of $^8$B$_X$.  It was found to be important to predict precisely the binding energies and excited states of exotic $X$-nuclei in realistic quantum mechanical calculations.  Both abundance ratios of $^6$Li/H and $^7$Li/H observed in MPHSs are obtained with an appropriate choice for the lifetime and abundance of the  $X^-$ particle.  These observational constraints imply a lifetime and abundance roughly in the range of $\\tau_X \\sim 2 \\times 10^3$~s and $Y_X \\sim 0.1$. We deduce that this $Y_X$ value requires that $m_X \\sim 50$~GeV in order to guarantee that this abundance of $X^-$ particles survives to the epoch of nucleosynthesis."
    },
    "0711/0711.4968_arXiv.txt": {
        "abstract": "{We present meridional flow measurements of the Sun using a novel helioseismic approach for analyzing SOHO/MDI data in order to push the current limits in radial depth. Analyzing three consecutive months of data during solar minimum, we find that the meridional flow is as expected poleward in the upper convection zone, turns equatorward at a depth of around 40\\,Mm ($\\sim0.95\\,R_\\odot$), and possibly changes direction again in the lower convection zone. This may indicate two meridional circulation cells in each hemisphere, one beneath the other. } ",
        "introduction": "The first observations of the Sun's meridional flow were made independently by \\citet{Duvall1979}, \\citet{Howard1979}, and \\citet{Beckers1979}, who found a poleward drift of $\\sim$20\\,m/s at the solar surface. By mass conservation, this meridional flow has to change direction somewhere in the solar interior, which may happen in the lower regions of the convection zone.  The meridional circulation plays an important role in many solar models. In mean-field models of angular momentum transport, the meridional circulation is responsible for how differential rotation is established and maintained \\citep{Ruediger1989,Thompson_etal2003,Miesch2005}. Anisotropic Reynolds stresses, arising from random turbulent fluctuations of the plasma motions, can induce a meridional circulation, which in turn drives the differential rotation. In numerical simulations, \\citet{BrunToomre2002} find that the differential rotation profile is established by turbulent convection, and that the Reynolds stresses play a crucial role in transporting meridional momentum toward the equator. The meridional circulation is also an important ingredient in some dynamo models \\citep{Charbonneau2005}. In flux-transport models of the solar dynamo (Babcock-Leighton models) the meridional flow is needed to transport the surface poloidal magnetic field to the bottom of the convection zone, where it is then converted into a toroidal field by rotational shear \\citep{DikpatiCharbonneau1999}.  Using helioseismic techniques, the meridional flow is usually observed in the upper layers of the solar convection zone. Quite good spatial resolution is achieved in this regime, in particular with ring-diagram analysis \\citep[e.g.,][]{GH_etal1999}. Deeper down in the convection zone the horizontal resolution is poorer and to probe there the viewing angle needs to be large. No return flow has conclusively been observed to date. Using time-distance helioseismology, \\citet{Giles1999} measured the radial profile of the meridional circulation down to the bottom of the convection zone and, with the added con\\-straint that the meridional circulation has to be contained within the convection zone, inferred a return flow at around 0.8 solar radii. Applying Hankel analysis, \\citet{BraunFan1998} saw a tentative return flow at around 40 Mm depth (but also consistent with zero within the error margin), whereas \\citet{DuvallKos2001} observed no return flow. ",
        "conclusions": "We analyzed high-resolution full-disk velocity data from SOHO/MDI along the center meridian covering three Carrington rotations from April to June 1997. Using a Fourier transform and measuring the frequency shift of the p-mode ridges in the power spectrum, we are able to obtain a meridional flow profile as a function of radius/depth. Figure~\\ref{fig3} shows that in the Northern as well as the Southern hemisphere the flow is poleward close to the solar surface and turns equatorward at around 40\\,Mm. Such a return flow was previously tentatively observed by \\citet{BraunFan1998}, analyzing data from the same period (but only for 8 days instead of 3 months). The flow may again turn poleward in deeper regions indicating a counter cell, although the data is also consistent with zero flow at lower regions. Such counter cells are consistent with recent numerical simulations \\citep{Miesch2007}. The near-surface flow being somewhat larger in the Southern than in the Northern hemisphere may be because of the inclination of the solar rotation axis \\citep{BeckGiles2005}."
    },
    "0711/0711.3409_arXiv.txt": {
        "abstract": "We present here the VO access to the results of an analysis of the spectra of Sloan Digital Sky Survey (SDSS) galaxies performed with the STARLIGHT code by \\cite{Cidetal05}. The results include for each galaxy the original SDSS spectrum, the best-fit synthetic spectrum, the star formation history, the pure emission line spectrum corrected for the underlying stellar population (in SDSS emission line galaxies) and the intensity of several emission/absorption lines. The database will be accessible from the PGos3 server and it will be released at the end of summer 2007. ",
        "introduction": "The interoperability concept promoted by the Virtual Observatory (VO) is not only limited to data sharing, but it also includes enhancing the possibility to analyze these data (usually by means of comparison with theoretical models) and to infer relevant physical quantities. This task can only be performed if front-end VO interfaces (mostly, applications) are able to provide access to theoretical data (see Rodrigo's et al. contribution in these proceedings) and if analysis tools are implemented in a VO compliant fashion. Theoretical and observational data must not be hardwired in the VO-analysis tools: rather, they should be obtained from VO services. This is the only way to minimize the model-dependent bias in the analyzed data.  The use of VO services to input theoretical and observational data into VO analysis tools creates diversity and improves productivity: since the VO provides a homogeneous structure for data sharing (VOTables), it allows to design automatic tools and workflows that can use different sets of observational and theoretical data without making program drivers for each specific data set. Once data providers distribute their results in VOTables, they can use these analysis tools for further research.  Although this is in principle the most desirable scenario, there are several difficulties that make this kind of workflow difficult to establish, mainly: (i) There is no user-friendly way to recover the credits of the data used in current VO applications; an issue mentioned several times during this meeting and especially relevant for theoretical studies. (ii) The access and exploration of VO-enabled services with theoretical datasets\\footnote{Examples are the PEGASE-HR (see Ph. Prugniel's contribution in these proceedings) and PGos3 databases (see Cervi\\~no's et al. contribution in these proceedings) for evolutionary synthesis models at {\\tt{http://vo.obspm.fr/cgi-bin/siap/pegasehr.pl}} and {\\tt{http://ov.inaoep.mx/pgos3}} respectively.} are not implemented in most VO-applications, with VOSpec and VOSED as the only exceptions.  In this contribution we describe the first stages of the implementation of the analysis tool STARLIGHT \\citep{Cidetal05} in the VO framework. ",
        "conclusions": ""
    },
    "0711/0711.1418_arXiv.txt": {
        "abstract": "We report in this paper phase-resolved spectropolarimetric observations of the rapidly-rotating fully-convective M4 dwarf V374~Peg, on which a strong, mainly axisymmetric, large-scale poloidal magnetic field was recently detected.  In addition to the original data set secured in 2005 August, we present here new data collected in 2005 September and 2006 August.  From the rotational modulation of unpolarised line profiles, we conclude that starspots are present at the surface of the star, but their contrast and fractional coverage are much lower than those of non-fully convective active stars with similar rotation rate. Applying tomographic imaging on each set of circularly polarised profiles separately, we find that the large-scale magnetic topology is remarkably stable on a timescale of 1~yr;  repeating the analysis on the complete data set suggests that the magnetic configuration is sheared by very weak differential rotation (about 1/10th of the solar surface shear) and only slightly distorted by intrinsic variability.  This result is at odds with various theoretical predictions, suggesting that dynamo fields of fully-convective stars should be mostly non-axisymmetric unless they succeed at triggering significant differential rotation. ",
        "introduction": "Most cool stars exhibit demonstrations of activity such as cool surface spots, coronal activity or flares, usually attributed to magnetic fields.  The current understanding is that these magnetic fields are generated at the surfaces and in the convective envelopes of cool stars through dynamo processes, involving cyclonic motions of plasma and rotational shearing of internal layers. In partly-convective Sun-like stars, dynamo processes are thought to concentrate where rotation gradients are supposedly largest, i.e., at the interface layer between their radiative cores and convective envelopes.  Cool stars with masses lower than 0.35~\\msun\\ being fully convective \\citep{Chabrier97}, they lack the thin interface layer presumably hosting solar-type dynamo processes and are thought to rotate mainly as rigid bodies \\citep{Barnes05, Kuker05};  yet they are both very active \\citep[e.g.,][]{Delfosse98} and strongly magnetic \\citep{Saar85, Johns96, Reiners07}.  This led theoreticians to propose that the intense activity and strong magnetism of fully convective dwarfs may be due to non-solar dynamo processes, in which cyclonic convection and turbulence play the main roles while differential rotation contributes very little \\citep[e.g.,][]{Durney93};  the exact mechanism capable of producing such fields is however still a debated issue. The most recent numerical dynamo simulations suggest that these stars are actually capable of producing large-scale magnetic fields, although they disagree on the actual properties of the large scale field.  While some find that fully-convective stars rotate as solid bodies and host purely non-axisymmetric large-scale fields \\citep{Kuker05, Chabrier06}, others diagnose that these stars should succeed at triggerring differential rotation and thus produce a net axisymmetric poloidal field \\citep[e.g.,][]{Dobler06}.  Existing observational data on fully convective dwarfs do not completely agree with any of these models.  Observations of low-mass stars indicate that surface differential rotation vanishes with increasing convective depth, so that fully-convective stars rotate mostly as solid bodies \\citep{Barnes05}.  At first sight, this result confirms nicely the predictions of \\citet{Kuker05}, leading one to expect fully-convective stars to host purely non-axisymmetric large-scale fields \\citep{Chabrier06}. However, spectropolarimetric data of the rapidly rotating M4 dwarf V374~Peg (G~188-38, GJ~4247, HIP~108706) collected with ESPaDOnS for almost three complete rotation periods over nine rotation cycles \\citep[hereafter D06]{Donati06} demonstrated that this fully-convective star hosts a strong mostly-axisymmetric poloidal field despite rotating almost rigidly, in contradiction with theoretical expectations.  In this new paper, we present and analyse spectropolarimetric data of V374~Peg collected with ESPaDOnS at three different epochs, 2005 August (i.e., those used in D06), 2005 September (providing only partial coverage of the rotation cycle) and 2006 August (providing full, though not very dense, coverage of the rotation cycle). First, we briefly describe the new observations and the Zeeman detections we obtained.  We then apply tomographic imaging to our data sets and characterise the spot distributions and magnetic topologies at the surface of V374~Peg in both 2005 and 2006;  from the observed temporal evolution of the magnetic topology, we also estimate both the differential rotation and the long-term magnetic stability of V374~Peg. Finally, we discuss the implications of our results for our understanding of the dynamo processes operating in fully-convective stars.  \\begin{table*} \\begin{center} \\caption[]{Journal of observations.  Subsets of observations collected sequentially on the same night (with no time gaps) are listed on a single line.  Columns 1--7 list the UT date, the number of observations in each subset, the heliocentric Julian date, the UT time, the exposure time, the peak signal to noise ratio (per 2.6~\\kms\\ velocity bin) and the rms noise level (relative to the unpolarised continuum level and per 1.8~\\kms\\ velocity bin) in the average circular polarization profile produced by Least-Squares Deconvolution (see text). The rotational cycle $E$ from the ephemeris of equation (\\ref{eq:eph}) is given in column 8. Whenever several observations are present in a subset, cols.~3, 4 and 8 list the values corresponding to the first and last subset observations, whereas cols.~6 and 7 mention the minimum and maximum values reached within the subset.  } \\begin{tabular}{@{}cccccccc} \\hline Date & $n_{\\rm obs}$ & HJD          & UT      & $t_{\\rm exp}$ & \\sn\\  & $\\sigma_{\\rm LSD}$ & Cycle          \\\\ & & (2,453,000+) & (h:m:s) &   (s)         &       &   (\\ptt) \\\\ \\hline 2005   & & & & & & & \\\\ \\hline Aug 19 & 6  & 601.78613--601.87607 & 06:43:29--08:52:59 & 4 $\\times$ 300.0 & 145--170 & 6.0--7.7 & 0.00--0.20 \\\\ Aug 19 & 12 & 601.94413--602.13790 & 10:30:60--15:10:01 & 4 $\\times$ 300.0 & 99--172 & 5.6--11.3 & 0.35--0.79 \\\\ Aug 21 & 23 & 603.75042--604.19199 & 05:52:00--15:15:52 & 4 $\\times$ 300.0 & 137--208 & 5.7--7.8 & 4.41--5.40 \\\\ Aug 23 & 23 & 605.73892--606.13069 & 05:35:25--14:59:33 & 4 $\\times$ 300.0 & 125--176 & 6.0--8.8 & 8.87--9.75 \\\\ \\hline Sep 08 & 6  & 621.94402--622.02801 & 10:30:49--12:31:46 & 4 $\\times$ 300.0 & 135--161 & 7.1--9.2 & 45.24--45.43 \\\\ \\hline 2006   & & & & & & \\\\ \\hline Aug 05 & 2 & 952.92084--952.93771 & 09:58:04--10:22:22 & 4 $\\times$ 275.0 & 172--181 & 5.7--5.8 & 788.00--788.04 \\\\ Aug 05 & 1 & 953.05250 & 13:07:39 & 4 $\\times$ 275.0 & 180 & 5.6 & 788.30 \\\\ Aug 07 & 2 & 955.02780--955.04651 & 12:31:58--12:58:55 & 4 $\\times$ 300.0 & 192--192 & 5.3--5.3 & 792.73--792.78 \\\\ Aug 08 & 2 & 955.86045--955.87780 & 08:30:57--08:55:55 & 4 $\\times$ 300.0 & 163--182 & 5.5--6.4 & 794.60--794.64 \\\\ Aug 08 & 2 & 956.02162--956.03903 & 12:23:01--12:48:05 & 4 $\\times$ 300.0 & 194--198 & 5.2--5.4 & 794.96--795.00 \\\\ Aug 09 & 2 & 956.82353--956.84159 & 07:37:44--08:03:44 & 4 $\\times$ 300.0 & 185--191 & 5.2--5.5 & 796.76--796.80 \\\\ Aug 09 & 2 & 957.02738--957.04305 & 12:31:15--12:53:50 & 4 $\\times$ 260.0 & 182--187 & 5.3--5.4 & 797.22--797.26 \\\\ Aug 10 & 2 & 957.84510--957.86535 & 08:08:44--08:37:54 & 4 $\\times$ 360.0 & 212--222 & 4.6--4.6 & 799.06--799.10 \\\\ Aug 10 & 2 & 958.03059--958.05071 & 12:35:50--13:04:49 & 4 $\\times$ 360.0 & 191--208 & 4.7--5.1 & 799.47--799.52 \\\\ Aug 11 & 2 & 959.03187--959.05226 & 12:37:38--13:06:59 & 4 $\\times$ 360.0 & 215--220 & 4.7--4.7 & 801.72--801.76 \\\\ Aug 12 & 2 & 960.03589--960.05609 & 12:43:23--13:12:27 & 4 $\\times$ 360.0 & 205--205 & 4.8--4.9 & 803.97--804.02 \\\\ \\hline \\end{tabular} \\end{center} \\label{tab:log} \\end{table*} ",
        "conclusions": "\\label{sec:disc}  Spectropolarimetric observations of V374~Peg were carried out with ESPaDOnS and CFHT at 3 epochs over a full time span of about 1~yr.  Clear Zeeman signatures are detected in most Stokes $V$ spectra;  weak distortions are also observed in the unpolarised profile of photospheric lines.  Using DI, we find that only a few low-contrast spots are present at the surface of V374~Peg, covering altogether 2\\% of the overall surface.  The spot distribution is apparently variable on a timescale of 1~yr.  Even though we may be missing small spots evenly peppering the photosphere \\citep[e.g.,][]{Jeffers06}, V374~Peg nevertheless appears as fairly different from warmer, partly-convective, active stars, whose surfaces generally host large, high-contrast cool spots often covering most of the polar regions at rotation rates as high as that of V374~Peg \\citep[e.g.,][]{Donati00}. \\cite{Barnes04} also report the lack of a cool polar structure for the M1.5 -- largely but not fully convective -- fast rotator HK~Aqr ($\\Prot = 0.431$~d).  A spot occupancy of about 2\\% has been observed for very slow rotators as well \\citep{Bonfils07, Demory07}, suggesting that spot coverage may not correlate with rotation rate in M dwarfs, as oppposed to what is observed in hotter stars \\citep[e.~g.,][]{Hall91}.  Using ZDI, we find that the surface magnetic topology of V374~Peg is mostly poloidal and fairly simple, with most of the energy concentrating within the lowest-order axisymmetric modes.  This confirms and amplifies the previous results of D06.  Again, this is at odds with surface magnetic topologies of warmer active stars, usually showing strong and often dominant toroidal field components in the form of azimuthal field rings more or less encircling the rotation axes \\citep[e.g.,][]{Donati03a}.  As for surface spots, ZDI is mostly insensitive to small-scale magnetic regions (and in particular small-scale bipoles) that can also be present at the surface of V374~Peg;  we can however claim that very little magnetic energy concentrates in spherical-harmonic terms with $\\ell$ ranging from 5 to 25.  We find that the surface magnetic topologies of V374~Peg at epoch Aug05 and Aug06 are very similar apart from an overall shift of about 0.1 rotation cycle.  This is confirmed by our success at fitting the complete Stokes $V$ data set down to almost noise level with a unique magnetic field topology twisted by a very small amount of solar-like differential rotation.  It suggests in particular that magnetic topologies of fully-convective stars are globally stable on timescales of at least 1~yr.  Again, this is  highly unusual compared to the case of warmer active stars, whose detailed magnetic topologies evolve beyond recognition in timescales of only a few weeks \\citep[e.g.,][]{Donati03a}.  The updated rotation period and photospheric shear we find for V374~Peg are respectively equal to $0.445654\\pm0.000002$~d and $0.0063\\pm0.0004$~\\rpd respectively;  the corresponding time for the equator to lap the pole by one complete rotation cycle is equal to 2.7~yr, about 10 times longer than solar.  We speculate that this unusually small differential rotation and long lap time are in direct relationship with the long lifetime of the magnetic topology, suggesting that differential rotation is what mostly controls the lifetime of magnetic topologies.  New data collected on a smaller time span (a few months) and featuring in particular no intrinsic variability (other than that generated by differential rotation) are needed to provide an independent and definite confirmation of this result.  Although the spot distributions at the surface of V374~Peg at epoch Aug05 and Aug06 apparently differ from more than a simple phase shift of 0.1 cycle, both data sets are compatible with a unique brightness map undergoing moderate intrinsic variability (assuming differential rotation is similar to that derived from Stokes $V$ data).  The low spottedness level and weak spot contrast implies that densely sampled data sets at 2 nearby epochs are needed to confirm this result, derive an independent determination of differential rotation parameters from unpolarised spectra and check that spots and magnetic fields trace the same atmospheric layers.  The strong mostly-axisymmetric poloidal field of V374~Peg can hardly be reconciled with the very weak level of surface differential rotation in the context of the most recent dynamo models of fully-convective stars. While \\citet{Dobler06} predicts that rapidly-rotating fully-convective stars should be able to sustain strong poloidal axisymmetric fields with the help of significant (though antisolar) differential rotation, \\citet{Chabrier06} find that these stars should only be able to trigger purely non-axisymmetric magnetic topologies if differential rotation is lacking (or remains below a threshold of about 1/3 of the solar value, Chabrier, private communication).  Both findings are at odds with our observational results, suggesting that our theoretical understanding of dynamo action in fully-convective stars is still incomplete.  We report the detection of strong chromospheric activity in V374~Peg, including quiescent emission in Balmer lines along with sporadic flaring events in H$\\alpha$, H$\\beta$ and the \\caii\\ IRT.  This is noticeably different from what was recently reported by \\citet{Berger07} for the M8.5 brown dwarf TVLM~513-46546, whose rotationally modulated H$\\alpha$ emission identically repeats over several successive rotation cycles. From the width and variability of emission profiles, we conclude that activity occurs mostly close to the stellar surface on V374~Peg, with flares apparently concentrating at specific longitudes. However, there is no clear spatial correlation between the photospheric magnetic field, the starspots and these chromospheric regions, e.g., strongest magnetic fields being observed when chromospheric activity is low.  Simultaneous multiwavelength observations of V374~Peg, including in particular optical spectropolarimetry, X-ray, and radio monitoring, could allow us to investigate the 3D magnetosphere of fully convective stars (e.g., by extrapolating the photospheric field topology up to the corona), and to unravel how photospheric activity relates to chromospheric and coronal activity.  Finally, we note that our \\vsini\\ and rotation period estimates imply that $\\rstar\\sin i=0.32\\pm0.01$~\\rsun, and thus that $\\rstar\\simeq0.34$~\\rsun\\ (assuming $i\\simeq70\\degr$). Given the $M_{\\rm K}$ magnitude of V374~Peg (equal to $7.0\\pm0.3$), we can safely derive that V374~Peg has a mass of $0.28\\pm0.05$~\\msun\\ \\citep{Delfosse00}.  It implies that the radius of V374~Peg is slightly larger than what theoretical models predict for low-mass stars of similar masses.  \\cite{Ribas06} also reports stellar radii larger than the theoretical predictions for very-low mass eclipsing binaries (whose rotation period is very short), whereas the radii of slowly rotating M dwarfs match the predicted values \\citep[e.~g.,][]{Segransan03}.  According to the phenomenological model of \\cite{Chabrier07} the presence of a strong magnetic field may inhibit stellar convection and therefore lead to larger radii with respect to an inactive star of equal mass.   Future observations are thus needed to explore in more details the magnetic topologies, brightness distributions and differential rotation of V374~Peg and other fully-convective stars.  Finding out how these quantities vary with mass and rotation rate could in particular show us the way to more successfull simulations of dynamo processes in low-mass fully-convective stars."
    },
    "0711/0711.4644.txt": {
        "abstract": "Gravitational waves from test masses bound to geodesic orbits of rotating black holes are simulated, using Teukolsky's black hole perturbation formalism, for about ten thousand generic orbital configurations. Each binary radiates power exclusively in modes with frequencies that are integer-linear-combinations of the orbit's three fundamental frequencies. General spectral properties are found with a survey of orbits about a black hole taken to be rotating at 80\\% of the maximal spin.  The orbital eccentricity is varied from 0.1 to 0.9.  Inclination ranges from $20^\\circ$ to $160^\\circ$, and comes to within $20^\\circ$ of polar.  Semilatus rectum is varied from 1.2 to 3 times the value at the innermost stable circular orbits.  The following general spectral properties are found: (i) 99\\% of the radiated power is typically carried by a few hundred modes, and at most by about a thousand modes, (ii) the dominant frequencies can be grouped into a small number of families defined by fixing two of the three integer frequency multipliers, and (iii) the specifics of these trends can be qualitatively inferred from the geometry of the orbit under consideration. Detections using triperiodic analytic templates modeled on these general properties would constitute a verification of radiation from an adiabatic sequence of black hole orbits and would recover the evolution of the fundamental orbital frequencies. In an analogy with ordinary spectroscopy, this would compare to observing the Bohr model's atomic hydrogen spectrum without being able to rule out alternative atomic theories or nuclei. The suitability of such a detection technique is demonstrated using snapshots computed at 12-hour intervals throughout the last three years before merger of a kludged inspiral. The system chosen is typical of those thought to occur in galactic nuclei, and to be observable with space-based gravitational wave detectors like LISA. Because of circularization, the number of excited modes decreases as the binary evolves. A hypothetical detection algorithm that tracks mode families dominating the first 12 hours of the inspiral would capture 98\\% of the total power over the remaining three years. ",
        "introduction": "\\label{Introduction}   The birth of physics in its modern form can arguably be placed at the first successful efforts to monitor the motion of the planets as they orbit the sun, and to model their planar elliptic orbits. Since then, gravitation has transitioned from the best to the least tested among the fundamental laws of physics.  Gravitational spectroscopy, monitoring power spectra of the outputs from gravitational wave detectors, can verify the existence of two-body systems with a dramatically different character. Athough characterized by mass ratios similar to those for planetary systems, these systems are so distorted by the strong gravity near black hole event horizons that their orbits look more like balls of twine than planar ellipses.  These orbits differ significantly from those with relativistic precessions which have already been observed and which may soon be measured in our galactic center \\cite{will 2007}.  Given the relative paucity of gravitational experiments, and the enigmatic status of gravitation compared to the other fundamental interactions, opportunities to observe these systems are highly valuable.  Though unlikely candidates for Brahe-Kepler-like catalysts of another revolution in physics, such observations would help to elevate tests of gravity to a level that better compares with those for the other fundamental physical interactions.  General relativity predicts that black holes will radiate gravitational waves characterized by discrete frequency spectra in at least two broad classes of realistically observable scenarios: black hole ringdowns, and during extreme mass ratio inspirals (EMRIs).  In both, the radiation is a consequence of a single black hole having been slightly perturbed. Observations of this radiation are therefore ideally suited for studying phenomena governed completely by the physics of isolated black holes.  In the case of black hole ringdown, the source of the black hole's perturbation is an arbitrary transient event, and the radiation is emitted as the hole settles back into a quiescent state.   Ringdown radiation consists of a superposition of exponentially damped monochromatic waves, with frequencies and damping timescales that are determined by a pair of continuous parameters, the black hole's mass and spin.  Potentially observable ringdown events include binary systems that merge to form single black holes, and supernovae that result in black hole formation.  Ringdown phenomena have been studied in detail elsewhere \\cite{flanagan hughes 1998, dreyer et al, berti et al 2006, berti et al 2007 ringdown}, and are not the focus of this work.  This work instead focuses on EMRIs---black hole binary systems in which the mass of the black hole is much greater than the mass of the companion. The planned space-based interferometer LISA is anticipated to observe anywhere from tens to thousands of EMRIs produced by the capture of compact stellar mass black holes (and the occasional white dwarf) by the megamassive black holes found in galactic nuclei (masses ranging from $10^5~M_\\odot$ to $10^7~M_\\odot$) with sufficient sensitivity to determine the mass, spin, and quadrupole moment of the larger black hole to within a fraction of a percent \\cite{gair et al 2004, hopman alexander 2006, miller et al 2005, sigurdsson 2003, barack cutler 2004, barack cutler 2007}. In the recent mock LISA data challenge, a proof of principal for a variety of detection algorithms using simulated signals and noise, three groups using independent detection algorithms recovered EMRI masses and spins to within a few to a few tenths of a percent \\cite{gair mandel wen, gair babak porter barack, mldc}.  Advanced ground-based detectors (with a lower frequency cutoff of about 10 Hz \\cite{advanced ligo}) might also observe similar capture events by intermediate black holes in globular clusters \\cite{brown et al 2006}. Observation of these intermediate mass ratio inspirals (IMRIs) are more speculative, might occur with an estimated upper limit of a few to a few tens per year \\cite{mandel et al 2007}, and due to their less extreme mass ratios, may require theoretical waveform models that are more sophisticated than those based on black hole perturbation theory alone. See Ref.~\\cite{amaro-seoane et al} for a review of EMRI science with both ground-based and space-based detectors.  Unlike radiation from black-hole ringdown, for EMRIs the source of the black hole perturbation is persistent. On timescales that are short compared to the observable lifetime of the binary, the companion behaves as a test mass orbiting on a geodesic of the background spacetime.  Such orbits are in general characterized by three fundamental frequencies \\cite{schmidt 2002, mino 2003}, and the corresponding radiation is tri periodic, or rather, a superposition of modes that oscillate exclusively at frequencies that are integer-linear combinations of the orbit's three fundamental frequencies \\cite{drasco hughes 2004} \\begin{subequations}\\label{waveform model} \\begin{align} h_{+} - i h_{\\times} &= \\sum_{m kn } h_{mkn} e^{-i \\omega_{mkn}t}~, & \\\\ \\omega_{mkn} &= m \\omega_\\phi + k \\omega_\\theta + n \\omega_r ~, \\label{frequencies}& \\end{align} \\end{subequations} where I am using the convention that all unconstrained summation indices range from $-\\infty$ to $\\infty$. Here $\\omega_\\phi$, $\\omega_\\theta$, and $\\omega_r$, are the fundamental orbital frequencies associated with motion in the (Boyer-Lindquist) coordinate shown as a subscript, $h_+$ and $h_\\times$ are the two independent components of the metric perturbation measured by a distant observer, and $h_{mkn}$ are complex amplitudes that depend on the observer's position, and on the parameters of the binary (masses, spin, and orbit geometry).  The radiation described by Eq.~(\\ref{waveform model}) is only a snapshot of a complete EMRI waveform.  While work toward precision waveforms for non-test-mass motion throughout the lifetime of EMRIs remains an active field (see the discussion of Capra waveforms in Ref.~\\cite{drasco 2006}), for short enough times these waveform snapshots are as exact as any envisioned. For example, for the case of a mass $\\mu$ on a circular orbit of initial radius $r$ about a nonrotating hole with mass $M \\gg \\mu$, the normalized overlap of the radiation from geodesic motion alone and that from the true inspiraling motion will be greater than 95\\% for times shorter than \\cite{drasco 2006} \\begin{equation} \\label{dephasing time} t_\\text{dephase} \\approx (1 \\text{ day}) \\left(\\frac{M}{10^6 M_\\odot}\\right) \\left(\\frac{r}{6M}\\right)^{11/4} \\left(\\frac{M/\\mu}{10^5}\\right)^{1/2} ~, %t =  M\\left[ \\sqrt{\\frac{5 \\pi}{108}} \\left( \\frac{r}{M} \\right)^{11/4} \\right] \\sqrt{\\frac{M}{\\mu}} ~, \\end{equation} whereas the orbital period is \\begin{equation} t_\\text{orbit} = (8 \\text{ minutes}) \\left(\\frac{M}{10^6 M_\\odot}\\right) \\left(\\frac{r}{6M}\\right)^{3/2}~. %t = 2\\pi M \\left( \\frac{r}{M} \\right)^{3/2} \\end{equation} Understanding the radiation from simple test mass motion may prove a sufficient basis for useful observations, however crude.  This is the motivating principal behind the work described here.  This paper has two main results.  The first is an observation of trends in a survey of numerically simulated EMRI-snapshots, for thousands of different generic orbital configurations, using the code described in Refs.~\\cite{hughes et al 2005} and \\cite{drasco hughes 2006}.  The observed trends are summarized as follows: (i) When the modes are sorted in order of decreasing power, their power decreases fast enough to be primarily confined to a relatively small number of modes compared to the number of modes computed when using the algorithm introduced in Refs.~\\cite{hughes et al 2005} and \\cite{drasco hughes 2006}. (ii) The dominant modes can be grouped into a small number of families defined by fixing two of the integer frequency multipliers.  (iii) The specifics of these trends (for example, the precise falloff of power as a function of mode index, or the identity of the dominant mode families) can be qualitatively inferred from the geometry of the orbit under consideration, as illustrated by the following two examples.  Simple orbits that are very nearly circular are expected to be typical of IMRIs that could be found with ground-based detectors.  Mandel \\emph{et al}.~estimate that the most likely IMRI formation mechanism should result in an eccentricity $e < 10^{-4}$ by the time the orbital frequency is brought into the observable band above 10 Hz for advanced LIGO \\cite{advanced ligo}, and that even for the less likely formations mechanisms that can yield higher eccentricities, 90\\% of the systems should have $e < 0.1$ at 10 Hz \\cite{mandel et al 2007}. For these orbits power falls off as a power law with mode index, and 99\\% of the radiated power is confined to a few to $\\sim 10$ modes with frequencies (\\ref{frequencies}) in the following families: \\begin{equation} \\label{easy modes} \\omega = \\left\\{ \\begin{array}{ll} m\\omega_\\phi & \\text{ dominant power at } m=2 \\\\ \\pm m\\omega_\\phi + \\omega_r & \\text{ dominant power at } m=2 \\\\ \\pm m\\omega_\\phi + \\omega_\\theta & \\text{ dominant power at } m=1 \\end{array} \\right. ~. \\end{equation} where the upper (lower) sign of $\\pm$ refers to prograde (retrograde) orbits, and mode-power decreases exponentially as $m$ is varied away from its value for the dominant mode.  Note that for power spectra, there is no observational consequence for a change in a frequency's overall sign.  For any given spectrum, I always define mode indices ($m,k,n$) in such a way that the frequency is positive: $\\omega_{mkn} > 0$.  More eccentric orbits are expected to be typical of the EMRIs observed by space-based detectors.  EMRIs that fall into LISA's frequency band are thought to be born with such large initial eccentricities ($10^{-6} \\lesssim 1 - e \\lesssim 10^{-3}$) that, although radiation circularizes them, even at the time of merger as many as half of the systems should have residual eccentricities of about $e \\gtrsim 0.2$ \\cite{barack cutler 2004}. Increasing orbital eccentricity dramatically slows the rate of power falloff and ultimately results in a spectrum dominated, albeit much less so, by the following mode families: \\begin{equation} \\label{hard modes} \\omega = \\left\\{ \\begin{array}{l} \\pm 2\\omega_\\phi + n\\omega_r \\\\ \\pm \\omega_\\phi + \\omega_\\theta + n\\omega_r  \\\\ \\pm 2\\omega_\\theta + n\\omega_r \\end{array} \\right. ~. \\end{equation} where again $\\pm$ is $+$ for prograde orbits and $-$ for retrograde orbits, and mode power decreases as $n$ is varied away from its value $n_{\\max}$ at which the spectra are peaked.  The value of $n_{\\max}$ can be very crudely predicted, with an accuracy on the order of 10\\%, from eccentricity $e$ according to the formula \\begin{equation} \\label{PM peaks} n_{\\max} \\approx \\exp(1/2) \\left( 1-e\\right)^{-3/2}~, \\end{equation} that was fitted \\cite{drasco hughes 2006} to the peaks in the spectra derived by Peters and Mathews,  who in 1963 used an analytical treatment of Newtonian orbits with the quadrupole formula for gravitational radiation \\cite{peters mathews}. For these eccentric orbits, the rate of power falloff as a function of mode index is such that as many as $\\sim 10^2$ to $10^3$ modes are needed to capture 99\\% of the power radiated.  The second main result of the paper is a proposal for combining the analytical waveform model (\\ref{waveform model}) with the observed spectral trends so as to create a means for verifying fundamental aspects of black hole physics. Precision measurements that will test the fundamental physics of black holes have long been among the primary motivations for gravitational wave observatories \\cite{abramovici et al, lisa science}. However, whether or not there are any feasible means for implementing one of the more grand measurements, measuring the multipolar decomposition of spacetime near black hole candidates in such a way as to identify general relativity as the only valid theory, has been a subject of debate \\cite{hughes 2006, psaltis et al 2007}. In its most pure form, that measurement will require more than the current understanding of what exactly is predicted by general relativity and its alternatives.  I propose a related, but less ambitious measurement that will constitute a minimal verification of one of general relativity's fundamental predictions for black hole perturbation:  that black holes perturbed by bound test masses radiate according to the above waveform model (\\ref{waveform model}) for times that are long compared to the spectrum's fundamental periods $2\\pi / \\omega_{\\phi, \\theta, r}$ but short compared to the inspiral time.  The measurement would also recover the evolutionary sequence of those three fundamental frequencies. This is a minimal verification in that it only provides a means for confirming general relativity's prediction under the assumption that the perturbed spacetime geometry is that of a rotating black hole. Observations that provided such a verification may ultimately be subjected to the more grand tests based on generalizations \\cite{li lovelace 2007} of Ryan's theorem \\cite{ryan}. In the absence of a greater theoretical understanding of alternatives to general relativity and black holes ,however, the verification alone would not disqualify alternatives for either the background spacetime or the theory of gravity.  The remaining sections of this paper are outlined as follows.  In Sec.~\\ref{EMRI-snapshot spectra}, I review the relevant equations that define generic black hole orbits and that describe the radiation they produce.  In Sec.~\\ref{Sample spectra}, the radiative spectra for three sample configurations are studied and general trends are discussed. In Sec.~\\ref{Survey of many spectra} I simulate spectra from a grid of orbital configurations and discuss how the general trends from the previous section vary across the grid. Section \\ref{Spectra from a kludged inspiral} describes how the spectral trends evolve throughout the lifetime of an approximated inspiral. In Sec.~\\ref{Verification of black hole orbits}, I outline a practical means for extracting signals characterized by these general spectral trends from data collected by gravitational wave observatories. Section \\ref{Conclusion} summarizes the paper's main points. ",
        "conclusions": "The number of significant modes in generic EMRI-snapshot spectra has been shown generally to be much more manageable than one might have guessed from earlier truncation algorithms \\cite{drasco hughes 2006}. This should lead to improved truncation algorithms, which will reduce the cost of future data analysis efforts. Such improvements should exploit the trends observed here in the relationship between orbit geometry and spectral signature.  The ability to predict the multiplier $n \\approx n_{\\max}$ of the radial frequency for the dominant modes using a formula based on such simple approximations \\cite{peters mathews} is encouraging.  It suggests that many of the trends in these spectra might be understood analytically using more recent tools \\cite{moreno-garrido mediavilla buitrago, poisson, ganz et al}.  The detection algorithms that are suggested here for verifying minimal aspects of relativity and black hole physics may ultimately be used in future gravitational wave detections.  However, more work is needed to determine whether or not they are cost-efficient and science-efficient alternatives to traditional search techniques. Another possibly interesting area for future work is to explore the dependence of the snapshot spectra on the spin of the larger black hole.  It has been implicitly assumed that the large values of spin considered are somehow representative of observable EMRIs.  This is reasonable since the few existing measurements due to modeling x-ray spectra from galactic nuclei suggest near-maximal spins \\cite{miller, reis et al}. Still, future work that tests the generality of these parameter values by simulating other EMRIs would be worthwhile."
    },
    "0711/0711.4521_arXiv.txt": {
        "abstract": "Conservation of \\textit{canonical} angular momentum shows that charged particles are typically constrained to stay within a poloidal Larmor radius of a poloidal magnetic flux surface. However, more detailed consideration shows that particles with a critical charge to mass ratio can have zero canonical angular momentum and so be both immune from centrifugal force and not constrained to stay in the vicinity of a specific flux surface. Suitably charged dust grains can have zero canonical angular momentum and in the presence of a gravitational field will spiral inwards \\textit{across} poloidal magnetic surfaces toward the central object and accumulate. This accumulation results in a gravitationally-driven dynamo, i.e., a mechanism for converting gravitational potential energy into a battery-like electric power source. ",
        "introduction": "To the best of the author's knowledge, no plasma has ever been observed to make a Kepler orbit around the Sun, a planet, or a moon. For example, the Earth's magnetotail does not make Kepler orbits around the Earth nor does the solar wind make a Kepler orbit around the Sun. This is puzzling because in the astrophysical literature plasmas are routinely presumed to make near-Kepler orbits in the presence of the gravitational field of a central object, e.g., accretion disks are modeled as MHD\\ plasmas in a near-Kepler orbit around a central object. The qualifier `near' is used because it is conventionally presumed for gases and plasmas that a radial gradient \\ of an isotropic pressure exists and that this pressure gradient provides a modest outward force which slightly reduces the amount of centrifugal force required to balance the inward force of gravity and so achieve a stable circular orbit (e.g., see Ref. \\cite{Weidenschilling:2000} for a discussion of the implications of this effect in the context of a system consisting of gas and solid particles).  If Kepler orbiting plasmas are so ubiquitous in astrophysics, then why is there not even a single example in the great variety of plasma/gravitational situations in our own solar system? A Kepler orbit is essentially a property of a single point particle -- it is not a property of a collection of independent point particles. For example, Earth and Mars are individually in Kepler orbits around the Sun, but the center of mass of Earth and Mars is not in a Kepler orbit around the Sun. If one replaced Earth and Mars by some statistically large number of point particles, then each could be in its own Kepler orbit around the Sun, but the center of mass of this configuration would not be in a Kepler orbit.  A collisional gas in a container could be in a Kepler orbit around the Sun because the walls of the container \\textit{bind} the particles to stay within a fixed distance of the center of mass so that the entire system can be considered as a point particle located at the center of mass. The transition from single point particle behavior to the behavior of a group of particles can be seen by considering a moon in close orbit around a massive planet. The tidal forces resulting from the gradient of the planet's gravitational force can be so strong as to overcome the binding forces and fracture the moon. The fragments would then follow distinct individual Kepler orbits and the center of mass of these orbits would not follow a Kepler orbit. Similarly, in order for a gas (or plasma) to have its center of mass follow a Kepler orbit there would have to be some binding force, such as would be provided by container walls, that would prevent the particle constituents of the gas (or plasma) from separating onto distinct Kepler orbits. Collisions might provide binding at the interior of a gas cloud, but not at the periphery because particles at the periphery moving away from the center would not encounter other particles with which they could collide. The periphery would simply expand into vacuum if there is no wall to prevent this expansion and the particles constituting the periphery would make Kepler orbits substantially different from the Kepler orbit calculated for the center of mass.  \\bigskip  Accretion disks are composed of dust and gas and the dust to gas mass ratio is estimated \\cite{Greaves:2000,Augereau:1999} to range from $\\sim10^{-2}$ to $\\sim1.$ UV radiation photo-ionizes \\cite{Wolf:2003,Lee:1996,Sickafoose:2000} the dust and gas so the accretion disk can be considered as a dusty plasma \\cite{Verheest:2000} consisting of charged dust grains, electrons, and ions. The magnetorotational instability (MRI) \\cite{BalbusHawley:1991} and the unipolar induction dynamo (UID) \\cite{Lovelace:1976} assume accretion disks are axisymmetric ideal MHD\\ plasmas and neglect dusty plasma physics effects. The MRI\\ and UID\\ additionally assume that accretion disks obey both Kepler dynamics and ideal MHD. Thus accretion disks are considered to be ideal MHD\\ plasmas in a circular Kepler orbit about a central object; they are supposed to conserve angular momentum and have frozen-in magnetic flux. One can then ask what does the trajectory of an individual particle in the accretion disk look like. Since MHD is assumed and MHD\\ is based on the assumption that all particles make cyclotron orbits, it seems that the individual particle in question should be making a cyclotron orbit. On the other hand since the whole plasma is supposed to be making a circular Kepler orbit around the central object, then presumably this cyclotron-orbiting particle must also be making a Kepler orbit around the central object. A Kepler-orbiting cyclotron orbit is not an obvious concept to visualize, at least to this author. This conceptual difficulty suggests that instead of assuming that a particle is simultaneously Kepler-orbiting and cyclotron orbiting, one should go back to first principles to investigate how Kepler and cyclotron orbits relate to each other. Perhaps it will then become obvious how to visualize a Kepler-cyclotron orbiting particle, or perhaps not. ",
        "conclusions": ""
    },
    "0711/0711.3071_arXiv.txt": {
        "abstract": "We report on the discovery of a new quantitative relationship between X-ray gas morphology and radio and X-ray AGN luminosities in normal elliptical galaxies. This is the second paper in a series using data on 54 objects from the {\\it Chandra} public archive and builds on the findings of Paper I, which demonstrated that hydrostatic equilibrium in elliptical galaxies holds, at best, only approximately and that the shape of the X-ray isophotes is unrelated to the shape of the gravitational potential. Instead, the gas is almost always asymmetrically disturbed. In this paper, we quantify the amount of asymmetry and study its correlation with other galaxy properties. We also determine radio powers and derive X-ray AGN luminosities for our galaxy sample. We find that the amount of asymmetry in the gas is correlated with both measures of AGN activity, in the sense that the hot gas is more disturbed in galaxies with higher radio and X-ray AGN luminosities. We find no evidence that galaxy density has significant effects on gas morphology. We do however find evidence for a correlation between gas asymmetry and the presence of hot ambient gas, which we interpret as a signature of hydrodynamic interactions with an external ambient medium. Surprisingly, the AGN--morphology connection persists all the way down to the weakest AGN luminosities in rather X-ray faint galaxies. This is strong morphological evidence that supports the general importance of AGN feedback, even in normal elliptical galaxies. ",
        "introduction": "In the first paper of this series \\citep[][hereafter Paper I]{DiehlGallery} we conducted a morphological analysis of the hot interstellar medium (ISM) of normal elliptical galaxies, derived from archived {\\it Chandra} observations of 54 objects. This paper introduced a new technique to isolate the diffuse hot gas emission from the contaminating effects of unresolved point sources, and presented a gallery of adaptively binned gas-only images. By fitting elliptical isophotes to these images, we demonstrated that the apparent flattening of the X-ray gas is completely uncorrelated with, and often significantly larger than, that of the starlight. The absence of any correlation, at radii inside the optical effective radius, implies that the hot gas cannot be in perfect hydrostatic equilibrium, unless the potential is dominated by a dark matter component completely unrelated to the stellar mass distribution. This is unlikely, since within the effective radius gravitational potentials are generally expected to be stellar-mass dominated \\citep[e.g.][]{MamonDarkmatterI,HumphreyDarkmatter}. Consequently, efforts to measure the shapes of dark matter halos using X-ray gas isophotes may be fruitless in most cases \\citep[e.g.][]{BuoteGeomtest,BuoteNGC720}, and radial mass profiles derived from X-ray data \\citep{HumphreyDarkmatter,FukazawaMassprofiles} could be in error by factors of order unity.  This paper takes up the question left unaddressed by Paper I: what {\\it is} the origin of the gas flattening, if not gravitational? In Paper I, we already excluded rotational flattening as the dominant cause. A qualitative analysis of gas morphologies showed that they are often asymmetric and disturbed. In this paper we will quantify this asymmetry, argue that it and the isophotal ellipticity reflect the same underlying disturbance, and show strong evidence that the root causes are the central active nucleus---even in those systems with very weak active galactic nuclei (AGN)---and interactions with ambient external gas.  Asymmetries are well known in X-ray studies of galaxies, groups, and clusters. {\\it Chandra} and {\\it XMM-Newton} observations of clusters of galaxies reveal strong deviations from smooth surface brightness distributions. Many of these features are believed to be due to cold fronts, caused by infalling substructure into the cluster center \\citep[e.g.][]{SarazinColdfront}, or gas sloshing in the gravitational potential \\citep[e.g.][]{MarkevitchSloshing}. In other cases, depressions in the surface brightness distribution are found to be coincident with extended radio emission \\citep[e.g.][]{McnamaraHydraA2}. These depressions, or ``bubbles'', are generally believed to be pockets of low-density extremely hot plasma, inflated by jets powered by the central AGN \\citep[see][for a recent review]{McNamaraAGNReview}. However, there are many cases where the depressions have no obvious radio counterparts, which has brought them the nickname ``ghost cavities'' \\citep{McnamaraGhostCavities}. These cavities are now interpreted to be relic bubbles that have detached from the radio source and are buoyantly rising radially outward in the cluster gas.  Unexpectedly, {\\it Chandra} and {\\it XMM} observations of normal elliptical galaxies have also revealed a wealth of highly disturbed gas morphologies with qualitatively similar properties to clusters. Many individual observations have been analyzed in detail and a variety of different explanations have been proposed. Sharp, one-sided drops in surface brightness are usually interpreted as signs of ram-pressure, distorting the gas in elliptical galaxies as they move through the ambient intracluster or intragroup medium, as in NGC~4472 \\citep{IrwinNGC4472,BillerNGC4472} and NGC~1404 \\citep{Machacek}. To explain the very asymmetric emission in NGC~7618, \\citet{KraftNGC7618} even argue for a major group-group merger. Other observations are not interpreted as results of environmental effects, but rather as signatures of the central AGN. For a few cases, such as NGC~4374 \\citep{Fin01} or NGC~4472 \\citep{BillerNGC4472}, the association with the AGN is indicated clearly by the correspondence between the radio source and depressions in the X-ray gas distribution. While these cavities are generally surrounded by cool rims \\citep[][and references therein]{McNamaraAGNReview}, two low-power radio galaxies (Centaurus~A and NGC~3801) have recently been identified with shocks surrounding the radio lobes \\citep{KraftCentaurusA,KraftCentaurusA2,CrostonShocks}. In other objects, one can directly detect the X-ray counterpart of the radio jet \\citep[e.g.][]{SambrunaXrayjets, HarrisXrayJet, HardcastleXrayJet}. However, most cases are less clear-cut. For example, the origin of the shock-like features seen in NGC~1553 \\citep{BlantonNGC1553} or NGC~4636 \\citep{JonesNGC4636} is still a mystery. The morphologies in other galaxies qualitatively resemble ghost cavities very close to the core \\citep[e.g. NGC 5044,][]{BuoteNGC5044XMM}, without evidence for an active radio source.  These individual cases each come with their own analysis techniques and interpretations, making general statements about elliptical galaxies problematic. Detailed analyses are generally limited to the X-ray brightest galaxies with clearly identifiable features, which are often dismissed as special cases. However, our morphological survey of normal elliptical galaxies in Paper I shows that disturbances and asymmetries in the gas are actually the norm rather than the exception, even for relatively X-ray faint galaxies. Most galaxies in the {\\it Chandra\\/} archive lack the signal, and therefore the contrast, to reliably identify cavity or shock features individually. In this paper, we introduce a statistical measure of asymmetry (the asymmetry index, $\\eta$) that allows us to treat all objects on an equal footing. We will show that the asymmetry index is strongly correlated with two different measures of nuclear activity, and with the presence of ambient hot gas.  The analysis in this paper begins with the elliptical isophotal fits of Paper I, and is organized as follows. In \\S\\ref{s4.analysis} we review the techniques of Paper I and describe the sample. We then calculate smooth functional fits to the isophotal profiles, use these fits to generate symmetric surface brightness models, and use the residuals from these models to define the asymmetry index $\\eta$ for each galaxy. We show by simulations that $\\eta$ is a monotonic measure of the sort of asymmetry expected from multiple cavities, and that it is independent of non-morphological parameters. We describe our techniques to extract the central AGN X-ray luminosity and to determine radio power and environment. In \\S\\ref{s4.results} we show that the asymmetry and isophotal ellipticity are correlated, and thus probably measure the same underlying disturbance. We show that neither of these quantities depends on neighbor galaxy density, but that they are correlated with measures of AGN power and the presence of ambient gas. The consequences for AGN feedback in galaxies and interactions with the ambient medium are discussed in \\S\\ref{s4.discussion}, before we conclude with a reiteration of our major results in \\S\\ref{s4.conclusions}. ",
        "conclusions": "\\label{s4.conclusions}  We have extended our earlier work from Paper I on normal elliptical galaxies, which showed that optical and X-ray gas ellipticities are uncorrelated in the regions where stellar mass dominates, contrary to what is expected from gas in hydrostatic equilibrium. Instead, the hot gas is generally disturbed. To elucidate the nature of these asymmetries,  we have introduced the asymmetry index $\\eta$, which measures the statistical deviation from a symmetric surface brightness model. This allows a systematic study of asymmetry even in relatively low-luminosity objects, in which asymmetric features are barely resolved.  We find a strong correlation between the asymmetry index and two independent measures of AGN activity: the radio continuum power at $20\\cm$ from NVSS \\citep{NVSS} and the X-ray AGN luminosity extracted from {\\it Chandra} data. The observed AGN--asymmetry correlation persists all the way down to the weakest AGN, where the NVSS survey reaches its detection limit. This is quite surprising, since these objects generally lack extended jet signatures in their radio images and are mostly detected as weak central point sources, if at all. We also find the asymmetry index to be correlated with the outer temperature gradient $\\alpha_{24}$, a proxy for the presence of a hot intragroup or intracluster medium. The strength of this correlation is comparable to that of the AGN--asymmetry correlation, indicating that hydrodynamic interaction with the ambient medium is comparably important to AGN in determining the X-ray morphology of normal ellipticals. However, we find no such correlation with the density of neighboring galaxies, and no evidence that galaxy-galaxy interactions play a significant role in shaping the hot ISM in these systems. Alternatively, the presence of an outer temperature gradient may also indicate the availability of a fuel reservoir for the central AGN. In this case, the increased amount of asymmetries in the X-ray gas could indicate a greater duty cycle for AGN outbursts in these objects.  The emerging picture is consistent with the AGN persistently stirring up the interstellar medium through intermittent outbursts, and strengthens the case for the AGN to be at least partly responsible for offsetting cooling in normal elliptical galaxies. We will address the impact of the central AGN on temperature and entropy profiles in detail in Paper III of this series \\citep{DiehlEntropy}."
    },
    "0711/0711.2288_arXiv.txt": {
        "abstract": "We show that candidate contact binary asteroids can be efficiently identified from sparsely sampled photometry taken at phase angles $\\alpha>60$\\degr. At high phase angle, close/contact binary systems produce distinctive lightcurves that spend most of the time at maximum or minimum (typically $>$$1\\,$mag apart) brightness with relatively fast transitions between the two.  This means that a few ($\\sim$5) sparse observations will suffice to measure the large range of variation and identify candidate contact binary systems.  This finding can be used in the context of all-sky surveys to constrain the fraction of contact binary near-Earth objects. High phase angle lightcurve data can also reveal the absolute sense of the spin. ",
        "introduction": "For most purposes, small solar system bodies are best observed at solar opposition. Measurements taken at low phase angle (the Sun-target-observer angle, denoted $\\alpha$; see Fig.~\\ref{Fig.Geometry}) benefit from strong backscattering, thus rendering the target brighter. In astrometric or surface-averaged measurements, the enhanced signal-to-noise due to the low phase angle geometry is desirable.  However, if the goal is to derive information on the shape of the target, then high-$\\alpha$ data become useful. This has been noted in the past, in the context of lightcurve inversion problems \\citep{2003A&A...404..709D}. As a body rotates, the amount of sunlight it reflects to the observer depends on the time varying projected cross-section, which in turn depends on its shape.  If well-sampled lightcurve information is obtained at multiple observing geometries, then the shape of the object can be derived. At low phase angles the non-convex features of an object are not apparent from its lightcurve.  This means that in the absence of high-$\\alpha$ data the non-convex figure of a contact binary can be wrongly interpreted as being composed of a single convex component.  One known effect of high-$\\alpha$ observations on lightcurves is the increase in the peak-to-peak range of variation \\citep{1956ApJ...123..331G}.  For single ellipsoidal bodies, and for phase angles $\\alpha < 40\\degr$, the lightcurve range $\\Delta m$ has been shown to vary roughly linearly with $\\alpha$ \\citep{1990A&A...231..548Z}. This is generally known as the amplitude-phase relationship (APR). When observing from Earth, the maximum phase angle at which a small solar system body can be imaged depends essentially on its heliocentric distance.  The distant Kuiper belt objects can be observed up to $\\alpha\\sim2\\degr$, Jovian Trojan asteroids at $\\alpha\\lesssim10\\degr$, and main-belt asteroids at $\\alpha\\lesssim30\\degr$.  Only near-Earth objects (NEOs) can be observed at phase angles $\\alpha>30\\degr$ which makes them the best candidates for high-$\\alpha$ studies.  On the order of twenty binary NEOs have been discovered, both photometrically \\citep[e.g.,][]{1997Icar..127..431P,1998Icar..133...79P} and from radar observations \\citep[e.g.,][]{2000IAUC.7518....2N,2002aste.conf..151O}. From the known sample, it is possible to note a few characteristic features \\citep{2006Icar..181...63P}.  The primaries have mean diameters of $\\sim 1\\,$km (with a spread of one order of magnitude), usually two to four times those of their secondary companions.  One case, 69230 Hermes, has a size ratio close to unity \\citep{2003IAUC.8227....2M}.  The separation distances are small, usually a few primary radii, and the mutual orbits are nearly circular.  Most measured primaries spin (around the shortest physical axis) close to the critical break-up rate for strengthless bodies with mean density $\\sim 2\\,$g$\\,$cm$^{-3}$. The secondaries have synchronized spin and orbital periods. In the few ($\\sim$10) cases where shapes can be inferred, the primaries are well described by oblate spheroids, while the secondaries have elongated shapes along the line of centers. One NEO, 4769 Castalia, is suspected from radar data to be a contact binary \\citep{1990Sci...248.1523O}.  In recent years, explanations to how binary NEOs may form have mostly converged on the idea of the splitting of a parent body.  \\citet{1980Icar...44..807W} proposed that binaries may form by rotational fission of an object effected by an off-center impact.  \\citet{2002Sci...296.1445M} argued that spin-up due to tidal interactions during close encounters with planets was the most likely formation process for NEO binaries, and \\citet{2006AREPS..34..157B} were the first to suggest that radiaton forces (YORP effect, Yarkovsky-O'Keefe-Radzievskii-Paddack) provide the torque needed for rotational fission.  The main features of the NEO binary sample described above have led \\citet{2007ApJ...659L..57C} to propose the following formation mechanism: due to YORP, $\\sim 1\\,$km objects are spun up close to break-up rates, which leads to the accumulation of regolith around the equator. As YORP continues to spin up the body, the surface material is eventually stored in orbit around the (now) primary, where it coagulates into a secondary; the latter grows into an elongated, nearly Roche equilibrium shape balancing gravitational and inertial accelerations. The subsequent evolution of the mutual orbit is driven by BYORP, the binary version of YORP, which may lead the components away from each other but may also bring them together to mutual contact \\citep{2005Icar..176..418C,2007Icar..189..370S}.  To obtain a well-sampled lightcurve of an NEO, capable of unequivocally establishing its nature, would require extensive observations at multiple geometries. For a sparse survey such as will be undertaken by Pan-STARRS (typically revisiting each target $\\sim4$ times per lunation), it may take years before the shape of the object can be determined. The method presented here can single out potential contact binaries after $\\sim5$ measurements at high-$\\alpha$.  In this letter, we present the first lightcurve simulations of contact binaries at high phase angle, and use them to address the detectability of such systems.  \\FigGeometry  \\begin{deluxetable}{ccccccc} \\tablecaption{Four close/contact binary systems considered.} \\label{Table.Binaries} \\tablewidth{0pt} \\tablehead{ \\colhead{\\#\\tna} & \\colhead{$q$\\tnb} & \\colhead{($B/A, C/A$)\\tnc} & \\colhead{($b/a, c/a$)\\tnc}  & \\colhead{$l$\\tnd} & \\colhead{H.E.\\tne} } \\startdata 1 & 1.00 & (0.66, 0.60) & (0.66, 0.60) & 1.00 & yes \\\\ 2 & 0.67 & (0.77, 0.69) & (0.53, 0.49) & 1.00 & yes \\\\ 3 & 0.25 & (0.92, 0.83) & (0.51, 0.48) & 1.19 & yes \\\\ 4 & 0.13 & (1.00, 0.90) & (0.51, 0.48) & 1.00 & no \\\\ \\enddata \\tablenotetext{a}{System number} \\tablenotetext{b}{Mass ratio} \\tablenotetext{c}{Primary and secondary axis ratios} \\tablenotetext{d}{Orbital distance in units of $A+a$} \\tablenotetext{e}{Components in hydrostatic equilibrium} \\end{deluxetable} ",
        "conclusions": "We have presented the case for using high phase angle observations to find extreme lightcurves produced by contact binaries in the solar system. Sparse observations of contact binaries at phase angles $\\alpha>60$\\degr\\ are extremely efficient at detecting the large range of brightness variation characteristic of these objects' lightcurves. From Earth, this method is most relevant to observations of near-Earth objects as they attain the largest phase angles.  At this point, a survey of high-$\\alpha$ measurements alone should probably not be used to estimate the fraction of contact binaries in the NEO population.  Deviations from hydrostatic equilibrium shape and the precise observing geometry ($\\theta$ and $\\phi$) affect the exact detection probability, thereby introducing uncertainty in the derived abundance. The technique is extremely efficient at detecting equator-on, relatively smooth contact binaries with symmetric components, but confirmation of the exact configuration of a system that shows large $\\Delta m$ in just a few high-$\\alpha$ measurements should always be sought using follow-up observations.  The sparse quality of planned all-sky surveys such as Pan-STARRS implies that a full shape solution can be obtained only after several years of data have been collected. It is therefore extremely useful to develop fast techniques that allow the identification of potentially interesting targets from a Pan-STARRS-type survey (or even from dedicated surveys using sub-1m telescopes), which can be followed up on other telescopes."
    },
    "0711/0711.0843_arXiv.txt": {
        "abstract": "We investigate the clustering of galaxies selected in the 3.6~\\micron\\ band of the {\\it Spitzer\\/} Wide-area Infrared Extragalactic (SWIRE) legacy survey.  The angular two-point correlation function is calculated for eleven samples with flux limits of $S_{3.6}\\ge 4$--400~\\mujy, over an 8~square degree field.  The angular clustering strength is measured at $>5$-$\\sigma$ significance at all flux limits, with amplitudes of $A=(0.49$--$29)\\times10^{-3}$ at one degree, for a power-law model, $A\\theta^{-0.8}$.  We estimate the redshift distributions of the samples using phenomological models, simulations and photometric redshifts, and so derive the spatial correlation lengths.  We compare our results with the GalICS (Galaxies In Cosmological Simulations) models of galaxy evolution and with parameterized models of clustering evolution.  The GalICS simulations are consistent with our angular correlation functions, but fail to match the spatial clustering inferred from the phenomological models or the photometric redshifts.  We find that the uncertainties in the redshift distributions of our samples dominate the statistical errors in our estimates of the spatial clustering.  At low redshifts (median $z\\le0.5$) the comoving correlation length is approximately constant, $r_0=6.1\\pm0.5h^{-1}$~Mpc, and then decreases with increasing redshift to a value of $2.9\\pm0.3h^{-1}$~Mpc for the faintest sample, for which the median redshift is $z\\sim1$.  We suggest that this trend can be attributed to a decrease in the average galaxy and halo mass in the fainter flux-limited samples, corresponding to changes in the relative numbers of early- and late-type galaxies.  However, we cannot rule out strong evolution of the correlation length over $0.5<z<1$. ",
        "introduction": "Galaxies are not distributed randomly across the sky.  At least at low redshifts, they appear to trace distinct patterns: galaxy clusters are connected to each other by long, filamentary structures of galaxies, interspersed with large voids in which few or no galaxies are seen. A plausible theoretical motivation has arisen for the formation of such large-scale structures (LSS) of galaxies, namely that the galaxies are tracing an underlying distribution of dark matter. In their most modern form, models for the formation of these large scale structures postulate that the evolution of the dark matter density field is inextricably linked to the formation and evolution of the galaxies themselves \\citep[e.g.,][]{Cole00,Granato00,Hatton03}. These models generally invoke some variation of the biased hierarchical paradigm, in which overdensities, or `halos', in the dark matter distribution undergo successive mergers over time to build halos of increasing mass, with galaxies forming from the baryonic matter in these halos.  From these models, and recent observations, it is clear that the relationship between the properties of galaxies, and the properties of the dark matter halos in which they reside, is subtle, and is an area in which observational constraints are particularly valuable in constraining models. We would like to know, observationally, what sort of galaxy occupies what sort of halo as a function of redshift, and how the properties of galaxies change with both redshift and the masses of their parent halos.  One method that has proven especially useful in providing such observational constraints is measuring clustering amplitudes. Fundamentally, the biased hierarchical paradigm requires that overdensities of dark matter should themselves cluster together on the sky, with the strength of clustering depending on their mass \\citep{Kaiser84,Bardeen86}. In principle, we can make great strides in understanding the relationship between galaxies and the underlying dark matter distribution by measuring the clustering strength of galaxies selected in a particular way, and relating this to theoretical predictions for halo clustering \\citep{Benson01}.  Such observations are, however, not straightforward to perform. Reliable clustering measurements require large, homogeneous samples of sources selected over large enough areas of sky to sample a range of dark matter density regimes. Ideally, we would like such observations to be performed in the near- and mid-infrared: the (restframe) near-infrard is most sensitive to evolved stars and so can pick up large samples of passively evolving systems, whereas the mid-infrared is sensitive to the dusty, active sources in which the stars and central black holes in (at least some) passively evolving systems are thought to form. Infrared observatories available up to now, however, have not been capable of mapping large enough areas to the required depths to find sufficient numbers of sources, or in enough bands to even crudely discriminate between different populations.  The launch of the {\\it Spitzer\\/} Space Telescope \\citep{Werner04} offers the potential to overcome these problems, due to its ability to map large areas of sky in the infrared to greater depths than any previous observatory, and in multiple bands so that dusty, active systems can be differentiated from passively evolving systems. The {\\it Spitzer\\/} Wide-area Infrared Extragalactic (SWIRE) survey \\citep{Lonsdale03,Lonsdale04} is the largest of the {\\it Spitzer\\/} Space Telescope's six Cycle~1 legacy programmes.  The survey covers a total area of 49~square degrees, split between six fields, in all seven of {\\it Spitzer's\\/} imaging bands (3.6--160~\\micron). The area and depth of SWIRE combine to produce a survey of significant comoving volume, 0.2$h^{-3}$~Gpc$^3$ over $0<z<2$, and spatial scales of $\\sim100h^{-1}$~Mpc at $z\\ge1$.  A principal goal of SWIRE is to study the clustering behaviour of a variety of extragalactic populations.  In \\citet{Oliver04} we presented the first detection of galaxy clustering in the survey, measuring a two-point angular correlation function at 3.6~\\micron\\ from our validation data, and in \\citet{Farrah06} we presented results on the clustering of Ultraluminous Infrared Galaxies (ULIRGs) at $z>1$.  \\citet{Fang04} presented angular correlation functions at 3.6--8.0~\\micron\\ from the 4-square degree {\\it Spitzer\\/} First Look Survey.  In this paper we extend the analysis of the 3.6-\\micron\\ clustering to larger scales and fainter flux limits (higher redshifts).  We begin with a summary of definitions and formalisms in section~\\ref{sec:defs}, then in section~\\ref{sec:sample} we discuss the sample selection, including star/galaxy separation and the angular selection function.  In section~\\ref{sec:corrfunc} we present our measurements of the two-point angular correlation function.  We compare our results with previous measurements in the K-band (section~\\ref{sec:kband}) and with the GalICS semi-analytical simulations (section~\\ref{sec:galics}).  The angular clustering amplitudes are used to estimate the spatial correlation lengths, which are then compared with simple parameterized models of clustering evolution, the GalICS simulations and results from the literature (section~\\ref{sec:r0}).  Section~\\ref{sec:conclude} draws together some conclusions from our analysis.  We use $H_0=100h^{-1}$~km~s$^{-1}$~Mpc$^{-1}$ with $\\Omega_M=0.3$ and $\\Omega_\\Lambda=0.7$.  Magnitudes are in the AB system unless otherwise noted. ",
        "conclusions": "\\label{sec:conclude}  The SWIRE survey has allowed us to measure the angular clustering of sources selected at 3.6~\\micron\\ over larger scales and with greater significance than previously has been possible.  The two-point angular correlation function has been measured in eleven flux-limited samples down to $S_{36}=4$~\\mujy, corresponding to median redshifts $z\\le1$. These angular results are in good agreement with the GalICS simulations and with $K$-band clustering measurements.  We have used Limber's equation to infer the spatial correlation lengths from the angular measurements and four estimates of the redshift distributions of each sample.  We find that the systematic uncertainty in the $dN/dz$ distribution dominates the statistical errors, and adopt the photometric redshift distribution from {\\sc ImpZ} as the best estimate of $dN/dz$, extrapolating to higher redshifts using the GalICS results.  The comoving correlation length varies slowly around $r_0=6.1\\pm0.5h^{-1}$~Mpc out to a median redshift of $z=0.5$ and then decreases with increasing redshift.  If this change at higher redshifts is due to evolution in the clustering, then the required evolution is very strong (with $r_0=27\\pm6h^{-1}$~Mpc and $\\epsilon=4.8\\pm0.8$) and these SWIRE galaxies would be in environments that will evolve into massive clusters by the present day.  There is some indication from the \\citet{Franceschini06} and \\citet{Xu03} phenomological models that the change in correlation length is due to a decrease in the average 3.6-\\micron\\ luminosity (or equivalently, stellar mass) of the fainter samples rather than evolution in the clustering.  This decrease in average luminosity is the result of an increased fraction of late-type galaxies in the faint samples.  Comparing these SWIRE data with the halo models of \\citet{Matarrese97} suggests that the fainter samples are selecting sources in lower mass halos.  These latter two results provide a consistent picture where lower-mass late-type galaxies are preferentially found in lower mass halos and thus have weaker clustering.  The comparison of our data with the GalICS simulations has shown that these models are a good match to the angular clustering, but their redshift distributions differ markedly from the phenomological models and the photometric redshift distributions.  This highlights the need for a better understanding of the spatial clustering and its evolution.  Higher-resolution numerical models and improved semi-analytic models will allow us to directly explore the relation between the clustering of the galaxies and the dark matter, i.e.\\ the bias.  For example, the dark matter Millenium Simulation \\citep{Springel05} has a lower mass resolution limit than GalICS, and combined with semi-analytic models of the galaxies it is a better match to the area and depth of SWIRE than the current models.  Future work will explore the comparison between these models and the SWIRE data presented here."
    },
    "0711/0711.0378_arXiv.txt": {
        "abstract": "We report spectroscopic observations of the nova M31N-2007-06b, which was found to be spatially coincident with the M31 globular cluster Bol 111. This nova is the first out of more than 700 discovered in M31 over the past century to be associated with one of the galaxy's globular clusters. A total of three spectra of the nova were obtained 3, 6, and 36 days post discovery. The data reveal broad (FWHM$\\sim3000$~km~s$^{-1}$) Balmer, NII, and NIII emission lines, and show that the nova belongs to the He/N spectroscopic class. The He/N class of novae are relatively rare, making up roughly 15\\% of the novae with measured spectra in M31, and roughly 20-25\\% of the Galactic novae for which spectroscopic data are available. The implications of a nova, particularly an He/N nova, occurring in a globular cluster are discussed. ",
        "introduction": "Classical Novae are semi-detached binary stars where the Roche-lobe-filling component (typically a cool, near-main-sequence star) transfers mass to a white dwarf companion (Warner 1995). If the rate of accretion is low enough to allow the accreted material to become sufficiently degenerate, a thermonuclear runaway (TNR) eventually ensues, driving substantial mass loss from the system, and leading to a nova eruption. Luminosities as high as M$_V \\simeq -9$ are often observed at the peak of eruption, making novae second only to gamma-ray bursts and supernovae in the energetics of their outbursts. Their high luminosities ($-6 \\lessim$ M$_V \\lessim -9$) and frequencies of occurrence ($\\sim50$~yr$^{-1}$ in a galaxy like M31), make them potentially powerful probes of the evolution of binary systems in different (extragalactic) stellar populations.  The standard picture for the formation of a nova system is through the common envelope evolution of an initially detached progenitor binary leading to the formation of a short-period, semi-detached system with the red dwarf transferring mass to its white dwarf companion (Paczy\\'nski 1976; Meyer \\& Meyer-Hofmeister 1979). Population synthesis studies based on this formation mechanism have predicted that the proportion of fast and bright novae -- like Type Ia SNe, which are thought to have similar progenitors -- is expected to be higher in younger stellar populations that contain on average more massive white dwarfs (Yungelson et al. 1997). Since novae with massive white dwarfs are expected to have shorter recurrence times between eruptions, the luminosity-specific nova rate (LSNR) should also be higher in younger populations. Despite these predictions, the weight of observational evidence suggests the opposite; that the LSNR is at least as high (and perhaps higher) in ellipticals and in spiral bulges, compared with early-type galaxies (e.g., Williams \\& Shafter 2004).  The most throughly studied extragalactic system is M31, where more than 700 novae have been discovered since the pioneering work of Hubble early in the 20th century (e.g. see Darnley et al. 2006; Pietsch et al. 2007; Shafter 2007a, and references therein). A principal conclusion reached by these studies is that, contrary to the predictions of population synthesis models, M31 novae appear to belong primarily to the galaxy's bulge population. This surprising result led Ciardullo et al. (1987) to speculate that the bulge nova rate may be enhanced by nova binaries that were spawned in M31's globular cluster system, and subsequently injected into the bulge through 3-body encounters in clusters (McMillan 1986), or by tidal disruption of entire clusters, or both.  To test the possibility that white dwarf accretors are also enhanced in the clusters themselves, searches for novae in M31's globular cluster system were undertaken by Ciardullo et al. (1990a) and by Tomaney \\& Shafter (1992). In the former study, Ciardullo et al. examined 54 globular clusters that fell within their prior M31 bulge nova surveys (Ciardullo et al. 1987; Ciardullo et al. 1990b). They found no nova eruptions, and were able to show that novae were several times less likely to be found in globular clusters than were high-luminosity X-ray sources, where roughly $\\sim$20\\% are associated with globular clusters (Crampton et al. 1984). In the latter study, Tomaney \\& Shafter used a multi-fiber spectrograph to monitor more than half ($\\grtsim$200) of M31 globular clusters over an effective survey time of almost a year. Once again, no novae were found suggesting that any enhancement of novae in M31's globular clusters was less than that of the high-luminosity X-ray sources.  The situation changed recently when as part of the ROTSE-IIIb program, Quimby et al. (2007) discovered a nova, M31N-2007-06b, that was spatially coincident with the M31 globular cluster system Bol 111. In this Letter we report spectroscopic observations of M31N-2007-06b, which clearly establish that the nova belongs to the He/N spectroscopic class proposed by Williams (1992). After describing the spectroscopic observations in the next section, we conclude by discussing the implications of an He/N nova arising in a globular cluster. ",
        "conclusions": "Our principal conclusions can be summarized as follows:  1) The nova M31N-2007-06b is spatially coincident to within observational errors ($\\pm 0.4''$ in each coordinate) with one of the galaxy's globular clusters (Bol 111). This discovery marks the first time a nova has been associated with a globular cluster in M31, and only the second extragalactic nova to arise in a globular cluster (the other being the nova found by Shara et al. in one of M87's globular clusters.)  2) Our spectroscopic observations clearly establish that M31N-2007-06b is a member of the He/N type in William's (1992) classification system. Although such novae are expected to be more common in younger stellar populations like that associated with M31's disk, the shorter recurrence times and the higher luminosities expected for these novae strongly enhance their discovery probability in any stellar population.  3) Taken at face value, the association of a single nova with a globular cluster out of the more than 700 novae discovered in M31 over the past century is consistent with the hypothesis that novae in globular clusters are not enhanced relative to the field. However, it is almost certain that a significant number of novae in M31's globular cluster system have been missed over the past century given the difficulty in detecting globular cluster novae, making an enhancement over the field seem likely. However, a more thorough and systematic monitoring campaign of M31's globular cluster system over many years will be required before a definitive conclusion can be reached."
    },
    "0711/0711.0652_arXiv.txt": {
        "abstract": "{% Azimuthal magnetorotational instability is a mechanism that generates nonaxisymmetric field pattern. Nonlinear simulations in an infinite Taylor-Couette system with current-free external field show, that not only the linearly unstable mode $m=1$ appears, but also an inverse cascade transporting energy into the axisymmetric field is possible. By varying the Reynolds number of the flow and the Hartmann number for the magnetic field, we find that the ratio between axisymmetric ($m=0$) and dominating nonaxisymmetric mode ($m=1$) can be nearly free chosen. On the surface of the outer cylinder this mode distribution appears similarly, but with weaker axisymmetric fields. We do not find significant differences in the case that a constant current within the flow is added. } ",
        "introduction": "MHD-instabilities in radiation zones of differentially rotating stars influence the rotation profile of the star due to angular momentum transport (Denissenkov \\& Pinsonneault 2007; Yoon, Langer \\& Norman 2006). Tayler instability (Tayler 1973) or helical magnetorotational instability (Hollerbach \\& R\\\"udiger 2005) are examples for such instabilities. According to Spruit (1999) the Tayler instability could possibly support a dynamo. In their model the instability should create a significant poloidal field that due to differential rotation is wound up into a toroidal one and thus close the loop for a self-sustained dynamo (Braithwaite \\& Spruit 2004; Braithwaite 2006). While this idea sounds promising, there are doubts about its general functionality. Simulations in very similar geometry with $z$-dependent differential rotation did not show a dynamo (Gellert, R\\\"udiger \\& Elstner 2007). Zahn, Brun \\& Mathis (2007) claim that such a process is not possible in this simple way. The authors argue that the nonaxisymmetric modes wound up by differential rotation can not contribute to the axisymmetric field, because the azimuthal wave\\-number is always $m\\neq0$. As we show in the following, azimuthal magnetorotational instability (AMRI) creates also an axisymmetric component in the nonlinear regime. For a scenario based on AMRI as instability mechanism that might overcome the arguments of the incompatible wavenumber and would make field regeneration at least possible. It does not necessarily mean, however, that it supports a dynamo.  This work consists of two aspects. First we describe the development of nonaxisymmetric modes in the nonlinear regime of AMRI-unstable Taylor-Cou\\-ette flow. Especially the observed inverse cascade transporting energy into the axisymmetric field and its dependence on the shear is analyzed. As second aspect we point out that weak currents within the flow do not basically influence the instability. Field structure and energy distribution remain almost unchanged. ",
        "conclusions": "In an infinite Taylor-Couette system we have shown, that nonlinear simulations of AMRI lead to an unstable nonaxisymmetric mode $m=1$ of the magnetic field if a certain strength of the external magnetic field is reached. Furthermore the instability is suppressed if the external field becomes too strong (for given Reynolds number). Equivalent is the point of view of an instability of the magnetic field due to shear. Here also a certain strength of the shear or value of the Reynolds number has to be reached to trigger the instability of the $m=1$ mode. And if the shear exceeds a certain value, the instability is not observable anymore.  The simulations reveal in the saturated state a distribution of magnetic energy over several modes, not only the $m=1$ is excited. However, the energy transport works in both directions, not only in the small-scale direction to larger $m$, but also to lower $m$. A significant amount of energy is shifted inversely into the axisymmetric mode ($m=0$). The shear strength is responsible for the ratio between magnetic energy in axisymmetric and nonaxisymmetric mode $m=1$. Near the stability boundaries, the nonaxisymmetric mode carries nearly all energy. In the central region of the unstable parameter space the energy ratio between modes $m=0$ and $m=1$ is roughly one.  Adding a small positive or negative current within the flow does not affect the instability, the field pattern and final energy distribution do not change compared to the current-free case. A small deviation from the profile $B^0 \\propto R^{-1}$ does not visibly influence the occurrence of the instability."
    },
    "0711/0711.0742_arXiv.txt": {
        "abstract": "{We present $B$ and $V$ light curves of a large stellar flare obtained with the Wide Field Camera at the Isaac Newton 2.5-m telescope (La Palma). The source object is a faint ($m_V=21.38$) foreground star in the field of the Andromeda galaxy, with its most probable spectral type being dM4. We provide an estimate of the total flare energy in the optical range and find it to be of the order of $10^{35}$ erg. The cooling phase of the large flare shows three additional weak flare-like events, which we interpret as results of a triggering mechanism also observed on the Sun during large coronal mass ejections.} ",
        "introduction": "Flares are known as sudden and violent events releasing magnetic energy and hot plasma from the stellar atmospheres. We observe them on magnetically active stars and, much more closely, on the Sun. Electromagnetic radiation is emitted across the entire spectrum, from radio waves through the optical range to X-rays and $\\gamma$ rays. The total energy released during a typical solar flare is of the order of $10^{30}$~erg, while the largest solar two-ribbon flares can emit up to $10^{32}$~erg.  According to the accepted model of solar flares, in the upper atmosphere, between oppositely oriented magnetic field lines, a current sheet forms and magnetic reconnection takes place, which results in acceleration of particles, and thus produces electromagnetic radiation and plasma heating. On the basis of the solar paradigm one could simply expect that modelling stellar flares is just a question of scaling (e.g., as a function of the released energy, size, duration, etc.). However, besides general similarities, stellar flare observations also unveil problems that cannot be explained with the extended canonical solar flare model, arising from the different spectral distribution of the emitted energy, the role of age and spectral type of the host star, the multiform magnetic field topologies on stars, the tidal forces in active binaries, etc. The only way we can come closer to drawing up similarities and differences between solar and stellar flares is to continuously broaden the ensemble of analyzed flares (solar and stellar as well), thus making an observational base of different types of flaring activities available for theoretical purposes. This is how solar flare observations can help in the correct interpretation of stellar flares, and reversely.  Flare activity in cool stars is a very common phenomenon. Flares in stars with spectral types earlier than M are observed mainly in UV and X-rays and optical flares (like the most energetic white light flares in the Sun; Hudson, Wolfson, \\& Metcalf \\cite{huetal06}) are rare. However, in the less luminous low-mass dM stars optical flares often occur. It is also known that red dwarfs from $\\approx0.3 M_\\odot$ down to the hydrogen burning limit can release a considerable amount of energy via flare eruptions, even when, in most cases, these objects show very low magnetic activity in their quiescent state. In this paper we present photometric observations of a large optical flare event observed on 2000 September 25 that is associated with one such low-mass dM stars. In Sect.~\\ref{data} a short description of the data is given, and in Sect.~\\ref{star} we summarize the stellar properties of the host star. In Sect.~\\ref{flare} we give an estimate of the energy released by the large flare and, finally, in Sect~\\ref{postflare} an additional two weaker short-term flares are identified and a plausible scenario is set up for the three weak post-flare events that followed the large flare.  \\begin{figure} \\includegraphics[width=0.48\\textwidth]{finding.eps} \\caption{Finding chart from the M31 field with our target (INT WFC image).} \\label{finding} \\end{figure}  \\begin{table*} \\caption{Catalog data on the observed target} \\label{T1} \\begin{tabular}{lccc}\\hline & INT WFC & DIRECT &\t2MASS \\\\ \\hline object ID \t& M31\\_J00453912+4140395 & D31J004539.1+414039.5 & J00453911+4140396 \\\\ $\\alpha$ (J2000)\t\t& 00\\,45\\,39.12 &  00\\,45\\,39.11\t& 00\\,45\\,39.12 \\\\ $\\delta$ (J2000)\t\t& +41\\,40\\,39.5 & +41\\,40\\,39.46\t& +41\\,40\\,39.7 \\\\ $B$\t\t\t& 23.18 & \t&\t\\\\ $\\sigma_B$\t\t& 0.11\t&\t&\t\\\\ $V$\t\t\t& 21.38\t& 21.62\t&\t\\\\ $\\sigma_V$\t\t& 0.03\t& 0.07\t&\t\\\\ $I$\t\t\t& \t& 18.18\t&\t\\\\ $\\sigma_I$\t\t& \t& 0.07\t&\t\\\\ $J$\t\t\t& \t&\t& 16.35\t\\\\ $\\sigma_J$\t\t&\t&\t& 0.11\t\\\\ $H$\t\t\t&\t&\t& 15.98\t\\\\ $\\sigma_H$\t\t&\t&\t& 0.16\t\\\\ $K$\t\t\t&\t&\t& 15.63\t\\\\ $\\sigma_K$\t\t&\t&\t& 0.18\t\\\\ \\hline \\end{tabular} \\end{table*} ",
        "conclusions": ""
    },
    "0711/0711.4482_arXiv.txt": {
        "abstract": "We readdress the problem of finding a simultaneous description of the pion form factor data in $e^+e^-$ annihilations and in $\\tau$ decays. For this purpose, we work in the framework of the Hidden Local Symmetry (HLS) Lagrangian and modify the vector meson mass term by including the pion and kaon loop contributions. This leads us to define the physical $\\rho$, $\\omg$ and $\\phi$  fields as linear combinations of their ideal partners, with coefficients being meromorphic functions of $s$, the square of the 4--momentum flowing into the vector meson lines. This allows us to define a dynamical, {\\it i.e.} $s$-dependent, vector meson mixing scheme. The model is overconstrained by extending the framework in order to include the description of all meson radiative ($VP\\gamma$  and $P\\gamma \\gamma$ couplings) and leptonic ($Ve^+e^-$ couplings) decays and also the isospin breaking ($\\omg/\\phi \\ra \\pi^+ \\pi^-)$ decay modes. The model provides a simultaneous, consistent and good description of the $e^+e^-$ and $\\tau$ dipion spectra. The expression for pion form factor in the latter case is derived from those in the former case by switching off the isospin breaking effects specific to $e^+e^-$ and switching on those for  $\\tau$ decays. Besides, the model also provides a good account of all decay modes of the form $VP\\gamma$, $P\\gamma \\gamma$  as well as the isospin breaking decay modes. It leads us to propose new reference values for the $\\rho^0 \\ra e^+ e^-$ and $\\omg \\ra \\pi^+ \\pi^-$ partial widths which are part of our description of the pion form factor. Other topics ($\\phi \\ra K \\overline{K}$, the $\\rho$ meson mass and width parameters) are briefly discussed. As the $e^+e^-$ data are found perfectly consistent with $\\tau$ data up to identified isospin breaking effects, one finds no reason to cast any doubt on them and, therefore, on the theoretical estimate of the muon anomalous moment $a_\\mu$ derived from them. Therefore, our work turns out to confirm the relevance of the reported  3.3 $\\sigma$ discrepancy between this theoretical estimate of $a_\\mu$ and its  direct BNL measurement. ",
        "introduction": "\\label{introduction} \\indent \\indent In order to study the phenomenology of light flavor mesons below 1 GeV, like partial decay widths or meson form factors, one needs a framework which includes in a well defined manner the lowest mass nonets of pseudocalar (P) and vector (V) mesons. Such a framework is well represented by the Hidden Local symmetry (HLS) model \\cite{HLSOrigin,HLSRef}. In this approach, vector mesons are gauge bosons of a spontaneously  broken  hidden local symmetry which generates their (Higgs--Kibble, HK) masses. Besides the non--anomalous sector, this model has an anomalous sector, hereafter called FKTUY Lagrangian \\cite{FKTUY}, which aims at describing couplings of the form $VVP$, $VP \\gamma$ $P \\gamma \\gamma $, $VPPP$ for light flavor mesons.  In its original form, the full (non--anomalous and anomalous) bare HLS Lagrangian fulfills the $U(3)$ symmetry, as it possesses both Nonet Symmetry and $SU(3)$ flavor symmetry. As such, the HLS model covers a limited phenomenological scope. In order to broaden this scope, especially in order to account for $VP \\gamma$ and $P\\gamma \\gamma$ couplings as derived from measured partial widths, the original HLS model should be supplemented with symmetry breaking mechanisms.  Several  $SU(3)$ symmetry breaking schemes have been proposed \\cite{BGP,BGPbis,Heath1} following the original idea of Bando, Kugo and Yamawaki (BKY) \\cite{BKY}. It has been shown \\cite{rad,mixing}  that the most successful  variant is the so--called ``new scheme\" of Ref. \\cite{Heath1} briefly recalled in Appendix \\ref{CC}. However,  breaking only the $SU(3)$ symmetry is insufficient \\cite{rad,mixing} in order to reach a satisfactory description of the data on  $VP \\gamma$ and $P\\gamma \\gamma$ couplings which requires one to also break  Nonet Symmetry. This was first performed in an {\\it ad hoc} manner in \\cite{rad} with the aim of recovering the radiative decay couplings of  O'Donnell \\cite{ODonnell} which is the most general set  fulfilling only $SU(3)$ symmetry~; therefore, the model developed in \\cite{rad} was indeed in agreement with general group theoretical considerations with additionally a mechanism smoothly breaking $SU(3)$ flavor symmetry.  Slightly later, it was shown \\cite{mixing,WZWChPT} that an appropriate mechanism for Nonet Symmetry breaking can be produced by adding determinant terms \\cite{tHooft} to the HLS Lagrangian~; the result was shown to meet all properties of Extended Chiral  Perturbation Theory (EChPT) \\cite{leutw,leutwb,feldmann} at leading order in the breaking parameters. Additionally, it was also proved \\cite{WZWChPT} that the Nonet Symmetry breaking mechanism proposed in \\cite{rad} was indeed an appropriate approximation of the (rigorous) mechanism derived from adding the determinant terms to the bare HLS Lagrangian.  However, breaking  $SU(3)$ and Nonet symmetries is still not enough to be in position of describing fully the whole set of radiative decays  $VP \\gamma$~; indeed, a process like $\\phi \\ra \\pi^0 \\gamma$ requires including a $\\omg-\\phi$ mixing scheme~; additionally, any global fit of all available $VP \\gamma$ transitions cannot be successful without introducing such a mixing. Traditionally, the   $\\omg-\\phi$ mixing is described \\cite{rad,mixing} by rotating the fields $\\omg_I$ and $\\phi_I$ which are the entries of the bare vector field matrix, generally called  ideal fields. An angle involving the mixing of the $\\eta$ and $\\eta^\\prime$  mesons is also required which has been shown \\cite{WZWChPT} to vanish in the limit of exact $SU(3)$ symmetry of the HLS Lagrangian.  A brief account of the HLS model in its anomalous (FKTUY) and non--anomalous sectors is given in Appendices \\ref{AA} and \\ref{BB}, mostly focused on the subject of this work. Appendix \\ref{CC} describes shortly but fully the various symmetry breaking procedures and the field renormalization scheme, except for the $SU(2)$ breaking mechanism which is the main subject of this paper. Appendix \\ref{YM} deals with some specific features of the Yang--Mills piece of the HLS Lagrangian which provides the kinetic energy part of the vector meson fields and should also undergo SU(3) symmetry breaking.  These Appendices illustrate  that the model depends on a very few number of free parameters as clear from Appendix \\ref{DD} where they all appear. Among these free parameters, some are specific to our HLS based model, but some others are constrained. For instance, one of the two SU(3) breaking parameters is the ratio of decay constants $f_K/f_\\pi$ and can be constrained by the corresponding measured value.   On another hand, decay processes like $\\omg/\\phi \\ra \\pi^+ \\pi^-$ can hardly be understood without some scheme for $\\rho^0-\\omg$ and  $\\rho^0-\\phi$ mixings. Moreover, the decay behaviour of $\\rho^0$ and $\\rho^\\pm$ cannot differ at the coupling constant level. It happens that the HLS model at one loop order provides a mechanism which allows us to perform a full $\\rho^0-\\omg-\\phi$ mixing starting from the corresponding ideal (bare) fields. Within the non--anomalous HLS Lagrangian, the precise mechanism, kaon loop contributions, has been already described in \\cite{su2}. However, we shall see that the (FKTUY) anomalous sector provides $K^* K$ loops as additional mechanism and that another one ($K^*\\overline{K}^*$ loops) is provided by the Yang--Mills piece.   Within the HLS model, as recalled below, $ V_i \\leftrightarrow V_j$ transitions among the ideal $\\rho^0_I$, $\\omg_I$ and $\\phi_I$  fields are generated by loop effects~; these transitions generally have not a constant amplitude but rather depend on the squared 4--momentum flowing through the $V_i$ (and $V_j$) line(s). We show in Section \\ref{OneLoop} that charged and neutral kaon, $K^* K$ and $K^*\\overline{K}^*$ loops come through their sum in  the $ \\omg_I\\leftrightarrow \\phi_I$ amplitude, while they come through their difference in $\\rho^0_I \\leftrightarrow \\phi_I$ and $\\rho^0_I \\leftrightarrow \\omg_I$ amplitudes. This also means that the $\\omg - \\phi$ mixing proceeds from quantum effects rather than symmetry breaking effects, in contrast with the $\\rho^0-\\omg$ or $\\rho^0-\\phi$ mixings. These transition amplitudes are given by Dispersion Relations which should be subtracted  in order to make the integral convergent (see, for instance, Appendix A in \\cite{mixing}). This gives rise to polynomials with real coefficients to be fixed using external renormalization conditions.  In the exact SU(3) symmetry limit, charged and neutral kaons carry the same mass and one can expect (or require) the polynomials associated with the charged and neutral kaon loops to coincide. Likewise, the charged and neutral $K^* K$ loop functions can be made equal. In this case, the $ \\omg_I\\leftrightarrow \\phi_I$ amplitude survives with its renormalization polynomial, while the $\\rho^0_I \\leftrightarrow \\phi_I$ and $\\rho^0_I \\leftrightarrow \\omg_I$ transition amplitudes exactly vanish. If one breaks the SU(3) flavor symmetry leaving conserved the $(u,d)$ quark sector, the same conclusion holds. However, if one introduces a breaking of $SU(2)$ flavor symmetry, the mass difference beween $u$ and $d$ quarks generates a  mass difference between the charged and neutral kaons (and $K^*$'s). Then, the three possible transition amplitudes do not identically vanish any longer and, moreover, they depend on the invariant mass  associated by the 4-momentum flowing through the vector  meson lines. Stated otherwise~: the $\\rho^0_I$, $\\omg_I$ and $\\phi_I$ mixing into the physical $\\rho^0$, $\\omg$ and $\\phi$ fields should also be invariant mass dependent. As already noted in \\cite{mixing,su2} this implies that the vector squared mass matrix, which has to be diagonalized in order to define the physical fields, is also invariant mass dependent and that the notion of mass for the physical vector fields becomes unclear as soon as one goes beyond tree level.  The property that $\\rho^0-\\omg$ mixing should be invariant mass dependent has been the subject of several studies  in the framework of general local effective field theories \\cite{RhoOmg0}, then in Vector Dominance Models \\cite{RhoOmg1,RhoOmg2} where it was pointed out that the mixing amplitude should vanish  at $s=0$, as the $\\rho^0$ and $\\omg$ self--masses, in order to preserve gauge invariance. One may also quote other studies going in the same direction \\cite{RhoOmg3,RhoOmg4} attributing the mixing to  finite width effects, or quark loops (and pion loops). Using only the pion form factor data available at that time, Ref. \\cite{ RhoOmg5} derived an approximate expression for the $\\rho^0-\\omg$ mixing amplitude~; however, limiting that much the  kind  of data used, one cannot really observe a clear mass dependence effect. On the other hand, one  should also note that Ref. \\cite{RhoOmg6} proved  that isospin violation effects describing the $\\rho^0-\\omg$ mixing are not accounted for by the low energy constants (LEC) of Chiral Perturbation Theory (ChPT) but are generated near threshold by the difference between charged and neutral kaon loops~; it was also shown that these effects are tiny at the two--pion threshold ($10^{-4}$) while they are known to be at the percent level in resonance peak region. This illustrates the invariant mass dependence of the $\\rho^0-\\omg$ mixing. However, one should note that the absence of the LECs in this calculation suggests a one-loop result may be unreliable.  Therefore, the question we address is to consider the effects of loops on $\\rho^0,~\\omg,~\\phi$ mixing. However, in order to have some chance to single out the mass dependent behaviour, one clearly has to treat simultaneously the pion form factor data in the spacelike and timelike regions together with the largest possible set of light meson decays (radiative, leptonic and  isospin violating strong decays) within a single framework. As argued above, the HLS model, equiped with suitable symmetry breaking mechanisms, seems able to provide such a framework. We shall not try to include Isospin Symmetry breaking effects into the coupling constants  which would help by providing more parameter freedom in the fit procedures. If really needed, it can certainly be done along the lines of the BKY breaking scheme \\cite{BKY,Heath1} as illustrated by \\cite{Hashimoto}.    \\vspace{0.5cm} \\indent \\indent The question of having a unified description of the largest possible set of low energy data is by itself interesting. However, this also addresses the puzzling and long--standing problem of the difference between the (estimated) isospin 1 part of the pion form factor in $e^+e^-$ annihilations and in $\\tau$ decays which are related through the Conserved Vector Current assuption (CVC). The importance of the problem is enhanced by its implication for the predicted value  of the muon anomalous magnetic moment $a_\\mu$ to be compared with the direct BNL  measurement \\cite{BNL}~; referring to the latest account by M. Davier \\cite{Davier2007} the estimate of the hadronic vacuum polarization (which enters the theoretical estimate) derived from $e^+e^-$ data provides a 3.3 $\\sigma$ disagreement between the theoretical estimation of $a_\\mu$ and the BNL direct measurement \\cite{BNL}~;  moreover, the $\\tau$ data estimate of the hadronic vacuum polarization provides a value of $a_\\mu$ very close to its direct measurement \\cite{Davier2007}.  Except for an experimental problem with $e^+e^-$ annihilation data (which seems by now unlikely)  in the  data recently collected at Novosibirsk \\cite{CMD2-1995corr,CMD2-1998-1,CMD2-1998-2,SND-1998}, or some new (or unidentified) physics effect, the disagreement between $e^+e^-$ and $\\tau$ data \\cite{Aleph,Opal,Cleo} is hard to explain. Indeed, {\\it a priori} the single difference between these two channels, should be due to Isospin Symmetry breaking (IB) effects. However, the comparison has been already performed with  IB effects accounted for in both $e^+e^-$ and $\\tau$ data. This includes \\cite{DavierPrevious1,DavierPrevious2,DavierPrevious3} pion mass values in kinematical factors, (a parametrization of the) $\\rho-\\omg$ mixing, charged and neutral $\\rho$  mass and width differences, short range \\cite{Marciano} and long range \\cite{Cirigliano1,Cirigliano2,Cirigliano3,Mexico1,Mexico2,Mexico3} IB effects in the $\\tau$ partial decay width to two pions.  This persistent disagreement may point towards new physics effects \\cite{Fred1}~; however, one should also note that the way some IB effects are accounted for has been questioned several times. For instance, effects due to the charged and neutral $\\rho$ pole positions \\cite{Fred2,Fred3} were considered, but have not been found sufficient in order to solve the observed discrepancy \\cite{DavierPrevious2,DavierPrevious3}~;  $\\rho-\\omega$ mixing effects may also have been poorly estimated \\cite{Maltman1}. However, based on sum rules derived  using an OPE input, K. Maltman \\cite{Maltman2,Maltman3} concluded there is inconsistency of the (presently) estimated isospin 1 part of $e^+e^-$ data with expectation, while $\\tau$ data provide a nice agreement.   We plan to address this question once again by building an effective model relying on the (symmetry broken) HLS model. In this approach, we plan to have a framework giving simultaneously an account of the partial decay widths of light mesons decays (radiative, leptonic, isospin breaking decay modes), of the pion form factor in $e^+e^-$ data (both timelike and close spacelike regions) and in $\\tau$ decay. By construction, the corresponding expressions of the pion form factor will be such that they will solely differ from each other by Isospin Symmetry breaking effects, mostly located in the $(\\rho,~\\omega,~\\phi)$ mixing scheme which, of course, has no counterpart in  $\\tau$ decay. Stated otherwise, our model is built in such a way that going from the pion form factor expression in $e^+e^-$ annihilations to the expression valid for $\\tau$ decays is performed by switching off IB effects specific to $e^+e^-$ and switching on those specific to $\\tau$ decays.  At start, the model we built is rendered complicated by the large number of possible loops involved. Fortunately, it can be somewhat simplified without loosing too much physics insight. Of course, this model depends on some parameters to be fixed in a fitting procedure~; we should define a fitting procedure flexible enough that $\\tau$ decay data can be removed or kept. Stated otherwise, the light mesons decays and the pion form factor in $e^+e^-$ data are expected to fix practically the U(3)/SU(3)/SU(2) breaking model. In this approach, we may get a prediction of the $\\tau$ decay 2--pion spectrum which can be compared  with the existing measurements. Including $\\tau$ decay data should only refine the values of the fitted parameters.  \\vspace{0.5cm} \\indent \\indent The paper is organized as follows~: In Section \\ref{HLS1}, we derive the Lagrangian pieces of relevance in order to deal with the pion form factor in $e^+e^-$ annihilations and $\\tau$ decays, while in Section \\ref{FFunbrk} we give the pion form factor expressions without loop corrections and symmetry breaking effects, mostly for illustration.  In Section \\ref{OneLoop} we discuss the loop corrections which modify the vector meson mass matrix and perform already some simplifications. The modified vector meson squared mass matrix we propose is given in Section \\ref{diagonal} with the diagonalization procedure and the relation between physical and ideal vector meson fields. The method used in order to renormalize the loop functions defining the self--energies and the transition amplitudes is sketched in Section \\ref{Renormalization}. The form factor functions used for $\\tau$ decays and $e^+e^-$ annihilations are given in resp. Sections \\ref{ModelFFtau} and \\ref{ModelFFee}.  A necessary ingredient affecting the  pion form factor in $e^+e^-$ annihilations is the photon vacuum polarization (VP), which is discussed in Section \\ref{ModelFFgVP}. Fitting with the partial width expressions is briefly discussed in Section \\ref{DecayWidths}~; more details can be found in \\cite{su2}, where a practically identical method is used with constant mixing functions, however. The way to deal with the various kinds of data used in the fit procedures is described in Section \\ref{FullDataSet}, with a special emphasis on our dealing with some correlation phenomena present in the existing data. In Section \\ref{DataFitting}, we fully describe the fit procedures we worked out under various conditions and the results and conclusions we reach~; we also comment on the numerical and physical properties of our model. Finally, Section \\ref{summary} is devoted to a summary of our conclusions.  As already commented upon in course of the Introduction, several Appendices gather results of ours or others already published. They are given in an attempt to be as self contained as possible~; they are placed outside the main text in the interest of clarity and ease of reading. However, Appendix \\ref{YM} contains new results concerning the Yang--Mills piece of the HLS Lagrangian. ",
        "conclusions": "\\label{summary} \\indent \\indent Within the context of the HLS model, as for other models, at tree level the so--called ideal fields ($\\rho_I$, $\\omg_I$, $\\phi_I$) are mass eigenstates. This simple picture vanishes at one--loop order. In this case,  kaon loops generate non--zero amplitudes allowing $\\omg_I \\longleftrightarrow \\phi_I$ transitions. Breaking of Isospin Symmetry in the pseudoscalar sector  generates a mass difference between kaons and, besides,  the transition amplitudes $\\omg_I \\longleftrightarrow \\rho_I$  and $\\phi_I \\longleftrightarrow \\rho_I$, even if of small magnitude, are no longer vanishing. Additionally, these amplitudes exhibit a dependence on $s$, the square of the momentum flowing through the vector meson lines. As the physical fields $\\rho$, $\\omg$, $\\phi$ are expected to be eigenstates of the squared mass matrix, this unavoidably leads to define them as linear combinations of their ideal partners. However, as the transition amplitudes are $s$--dependent, it is clear that these combinations should also be $s$--dependent.  We substantiate these considerings starting from the  HLS Lagrangian, modified by including in the squared mass matrix of the neutral vector mesons all  self--energies and transition amplitudes. Making the assumption that the physical neutral vector fields should be eigenstates of the loop modified squared mass matrix of the (ideal) neutral vector meson, we solve the eigenvalue problem perturbatively. This leads to physical vector meson fields expressed as linear combinations of their ideal partners with definite $s$--dependent coefficients, which are actually analytic -- or, rather, meromorphic -- functions of $s$. Of course, this algebra leaves unchanged the charged and/or open strangeness sector, as the starting fields acquire a running mass but no transition from one to another meson field.  The main mechanism producing the vector meson field mixing is the occurence of neutral and charged kaon loops in transitions between the ideal neutral vector meson fields. We have also shown that the anomalous HLS--FKTUY sector and the Yang--Mills lagrangian piece provide as supplementing mechanisms the  $K^*K$ and $K^*\\overline{K}^*$ loops, occuring in transitions in nice correspondence with kaon loops. Within this framework, the $\\omg \\phi$ mixing has been shown to proceed from quantum (loop) effects, while the $\\rho^0\\omg $ and $ \\rho^0\\phi$ mixings follow from Isospin Symmetry breaking effects and vanish when this symmetry is restored. From a numerical point of view, the bulk of the effect is carried more by the subtraction polynomials than the loop expression themselves, as these are logarithmic functions of small amplitude in the physics region of interest.  The vector meson mixing allows to dynamically generate isospin violating couplings $\\phi \\pi^+ \\pi^-$ and $\\omg \\pi^+ \\pi^-$ at the (modified HLS) Lagrangian level. With this at hand, we have been able to construct the pion form factor expression at one--loop order modified in order to account for Isospin Symmetry breaking through only a dynamically generated $\\rho_I-\\omg_I-\\phi_I$ mixing scheme. {\\it A priori}, this fully affects the $e^+e^-$ physics but, by no way, the $\\tau$ physics~; in this sector, as a second order refinement, we have been led to accept shifting with respect to each other the Higgs--Kibble masses of the $\\rho^0$ and $\\rho^\\pm$ mesons. This provides a tiny effect, nevertheless clearly visible in ALEPH $\\tau$ data, but not  obvious in CLEO data. New data expected from B-factories may confirm the need for this mass shift. As a final result, we end up with structureless residuals in the fitted regions, which confirm that our dynamical mixing scheme is appropriate.  The mixing properties have been introduced in the anomalous decay amplitudes $V \\ra P \\gamma$ and $P \\ra V \\gamma$. These processes actually represent our main lever arm  while defining numerically our isospin breaking scheme.  Beside a good description of the $V \\ra P \\gamma$ and $P \\ra V \\gamma$ decay data, this allows a  very good simultaneous description of all pion form factor data from the close spacelike region to the $\\phi$ mass, in $e^+e^-$ annihilations as well as in $\\tau$  decays. The physical  ground of this result can be traced back to the fact that the $\\rho^0$ meson is a mixture of isospin 0 and 1 states (as the $\\phi$ and $\\omg$ meson), in contrast to the $\\rho^\\pm$ meson which is purely isospin 1. Actually, extracting the isospin 1 part of the pion form factor in $e^+e^-$ annihilations requires to split up the $\\rho^0$, $\\phi$ and $\\omg$ contributions. This is done automatically by our model, and one can claim that such a splitting cannot be done without some model.  The net result of this model is to prove that the lineshapes for the pion form factor in $e^+e^-$ annihilations  and in $\\tau$  decays are perfectly consistent with each other, without any further breaking of CVC than a possible tiny $\\rho^0-\\rho^\\pm$ mass difference. In a further step, one may include the data from KLOE as well as data expected to come from BaBar and Belle concerning the $\\tau$ spectrum on the one hand and the $e^+ e^-$ Initial State Radiation samples on the other hand. However, our comparison of $e^+e^-$ and $\\tau$ data does not seem to leave room for any kind of new physics.  \\vspace{0.5cm} \\indent \\indent On the other hand, we have shown that this model allows a good account of all decays of the form  $V \\ra P \\gamma$ and $P \\ra V \\gamma$. The case of the $\\omg \\ra \\eta \\gamma$ and $\\omg \\ra \\pi^0 \\gamma$ partial widths, where some disagreement is observed with the so--called ``fit\" values proposed by the PDG, has been discussed and we have argued that the real situation is somewhat unclear.  Our dealing with the pion form factor data has led us to propose improved values for data sharply related with the $e^+e^- \\ra \\pi^+ \\pi^-$ annihilation process, namely the $\\rho^0 \\ra e^+e^-$ and  $\\omg \\ra \\pi^+ \\pi^-$ partial width decays. In both cases, we find that the reference values should be significantly modified and we propose for these new reference data.  Finally, we have briefly commented upon the mass and width of the $\\rho^0$ and $\\rho^\\pm$ mesons and argued that the best motivated definition should rely on the pole position in the complex $s$--plane and related definitions.  \\vspace{0.5cm} \\indent \\indent Having shown that  $e^+e^-$ data and $\\tau$ data are perfectly consistent with each other provided one uses  an appropriate model of Isospin Symmetry breaking, we can conclude that there is no reason to question the $e^+e^-$ data. This result is important as these data serve to estimate numerically the hadronic photon vacuum polarization used in order to  predict  the value of the muon anomalous magnetic moment $g-2$ . Therefore our model indirectly confirms the 3.3 $\\sigma$ discrepancy between the BNL direct measurement of the muon anomalous magnetic moment and its theoretical estimate."
    },
    "0711/0711.1572_arXiv.txt": {
        "abstract": "{Our position inside the Galaxy requires 3D-modelling to obtain the distribution of the Galactic magnetic field, cosmic-ray (CR) electrons and thermal electrons. } {Our intention is to find a Galactic 3D-model which agrees best with available radio observations.} {We constrain simulated all-sky maps in total intensity, linear polarization, and rotation measure (RM) by observations. For the simulated maps as a function of frequency we integrate in $15\\arcmin$ wide cones the emission along the line of sight calculated from Galactic 3D-models. We test a number of large-scale magnetic field configurations and take the properties of the warm interstellar medium into account.} {From a comparison of simulated and observed maps we are able to constrain the regular large-scale Galactic magnetic field in the disk and the halo of the Galaxy. The local regular field is 2~$\\mu$G and the average random field is about 3~$\\mu$G. The known local excess of synchrotron emission originating either from enhanced CR electrons or random magnetic fields is able to explain the observed high-latitude synchrotron emission. The thermal electron model (NE2001) in conjunction with a proper filling factor accounts for the observed optically thin thermal emission and low frequency absorption by optically thick emission.  A coupling factor between thermal electrons and the random magnetic field component is proposed, which in addition to the small filling factor of thermal electrons increases small-scale RM fluctuations and thus accounts for the observed depolarization at 1.4~GHz.} {We conclude that an axisymmetric magnetic disk field configuration with reversals inside the solar circle fits available observations best. Out of the plane a strong toroidal magnetic field with different signs above and below the plane is needed to account for the observed high-latitude RMs. The large field strength is a consequence of the small thermal electron scale height of 1~kpc, which also limits the CR electron extent up to a height of 1~kpc not to contradict with the observed synchrotron emission out of the plane. Our preferred 3D-model fits the observed Galactic total intensity and polarized emission better than other models over a wide frequency range and also agrees with the observed RM from extragalactic sources.}  \\keywords {polarization -- Radiation mechanisms: non-thermal -- Radiation mechanisms: thermal -- ISM: magnetic field -- ISM: structure -- Galaxy: structure}  \\titlerunning{Galactic 3D-models} \\authorrunning{X.H.~Sun et al.} ",
        "introduction": "Synchrotron radiation is the major component of the Galactic radio continuum emission from centimetre to longer wavelengths. The diffuse Galactic synchrotron emission originates from relativistic electrons gyrating in the magnetic field. An additional smaller contribution is observed from supernova remnants (SNRs) and pulsar wind nebulae located in the Galactic plane, where cosmic ray (CR) electrons are accelerated to high energies. Synchrotron emission is intrinsically highly polarized, but Faraday rotation effects in the magnetized interstellar medium along the line of sight and also beam averaging effects reduce the observed percentage polarization to much smaller values. It is crucial to understand the regularity, strength and spatial structure of the magnetic field in our Galaxy, which are still under discussion although the first detection of Galactic polarization was several decades ago \\citep{wsp62,wsb+62}. Our position inside the Galactic disk at about 8.5~kpc distance from the Galactic centre requires modelling of the magnetic field and CR electrons to interpret observed synchrotron maps. Polarization maps are influenced by Faraday rotation effects along the line of sight and thus require to take the distribution of thermal electrons in the Galaxy into account as well. Important information on the Galactic magnetic field strength and direction comes from new rotation measure (RM) data. All these information must be commonly considered to constrain Galactic 3D-models in order to find the most realistic one. This paper presents a general attempt in that direction.  Previous 3D-models for the magnetic field, CR electrons and thermal electrons were derived individually from selected data sets. Their general relevance can be established by combining these models to explain other, in particular new observational data, for instance from WMAP.  Here we consider RMs from extragalactic sources (EGSs), thermal and non-thermal intensities and linearly polarized emission together, which proves to give stringent constraints to combined 3D-models. We note that simulated maps are also required as a robust reference for the technical design of future telescope projects like the SKA or as inputs for data reduction pipelines like that for the PLANCK mission.  The paper is organized as follows: we review the status of Galactic models in Sect. 2, describe the method of our simulations in Sect. 3, summarize available observations used to constrain the 3D-models in Sect. 4 and briefly describe numerical details of the simulations in Sect. 5. The 3D-model results are presented in Sect. 6. Some general problems relevant to the modelling are discussed in Sect. 7. Conclusions are presented in Sect. 8. ",
        "conclusions": "Galactic 3D-models of the distribution of thermal electrons, cosmic-ray electrons and magnetic fields were used to calculate via the HAMMURABI code all-sky total intensity, polarization and RM maps. A comparison with observed maps over a wide frequency range is made to adapt the 3D-model parameter in the best possible, although still qualitative way (Table~1).  The main conclusions are : \\begin{itemize} \\item The NE2001 thermal electron model combined with a filling factor $f_e$ as derived by \\citet{bmm06} reproduces the optical thin 22.8~GHz free-free emission template and also the observed spectral flattening in the Galactic plane between 408~MHz and 22~MHz caused by absorption of optically thick thermal gas. We find indications that the filling factor might be larger towards high $z$ than that found by \\citet{bmm06}.  \\item An ASS field configuration plus a reversal inside the solar circle either in a ring or between the inner edges of the Sagittarius-Carina arm and the Scutum-Crux arm is favoured by RMs of EGSs observed along the Galactic plane. A BSS field configuration also fits these RMs, however, it then fails to fit the observed RM gradient in latitude direction. The strength of the regular magnetic field component at the solar radius is 2~$\\mu$G. The field in the Galactic centre region is constant at 2~$\\mu$G in a direction opposite to that at the solar radius.  \\item The strength of the Galactic random field component is 3~$\\mu$G.  \\item A regular halo magnetic field with a maximum of 10~$\\mu$G at a height of 1.5~kpc is required to account for the observed high-latitude RMs. This toroidal field is asymmetric to the Galactic plane and decreases sharply towards the plane. The variation of the field towards high $z$ is not a sensitive parameter in our simulation.  \\item The CR electron distribution has an exponential scale height of 1~kpc and a radial scale length of 8~kpc. The disk is truncated at a height of 1~kpc. The local CR flux density at 10~GeV is taken twice the directly measured value at earth. The energy power law spectral index $p$ is taken to be constant at $-$3 for frequencies above 408~MHz and $-$2 below.  \\item The observed local excess of synchrotron emission is taken into account. The excess is either caused by an enhanced local CR electron density or by an increase of the local turbulent magnetic field. Our modelling can not distinguish between both possibilities.  \\item In addition to the filling factor we introduce a coupling factor between the thermal electron density and the random magnetic field component to reproduce the observed large RM fluctuations and the observed depolarization at 1.4~GHz. \\end{itemize}  However, a number of caveats remain to the above conclusions.  In our simulation we use the NE2001 model as input for the thermal electron distribution and density. This requires a large halo field strength and hence a truncation of the CR electrons at $z$~=~1~kpc. This might be not very realistic. Indeed, there are indications as discussed in Sect. 7.3 that the scale height of thermal electrons might be larger than 1~kpc. When the thermal electron scale height is increased by a factor of 2 we get a regular halo magnetic field of 2~$\\mu$G and the truncation of the CR electron density at 1~kpc is avoided. We conclude that the high-latitude thermal electron density needs more investigations and thus the halo field parameters remain open.  For large areas along the Galactic plane RM data are still missing or very rare. A final answer on the magnetic field configuration depends on a much more complete RM data set. Pulsar RM data are not yet included in our model, although they are potentially better tracers of the regular Galactic magnetic field along the line of sight than those from EGSs as they show no intrinsic RM. Pulsar RMs are clearly important to detail the disk magnetic field throughout the Galaxy on smaller scales than in our model and might also reveal deviations from a general regular pattern being assumed here.  Despite of significant observational progress over the last years \\citep[e.g. ][]{hml+06} more RM data are clearly needed to be conclusive and pulsar distances need to be better constrained in addition.  We note that our 3D-model differs from previous models, which were derived from selected observational data.  These models then fail to properly reproduce other observations, which, however, trace relevant information of the interstellar medium, which can not be ignored on the way towards a physical relevant model."
    },
    "0711/0711.1434_arXiv.txt": {
        "abstract": "{Structural parameters and the total \\mv\\ magnitude are important properties to be used in the characterization of individual globular clusters. With the increasing statistics, especially of the intrinsically faint objects, collective properties of the Galactic globular cluster system will be better defined, with reflexes on our understanding of the formation history of the Galaxy.} {The determination of structural parameters of 11 faint Galactic globular clusters that, in most cases, had not been previously studied in this context. The clusters are IC\\,1257, Lyng\\aa\\,7, Terzan\\,4, Terzan\\,10, BH\\,176, ESO\\,452-SC11, ESO\\,280-SC08, 2MASS-GC01, 2MASS-GC02, GLIMPSE-C01 and AL\\,3, which are projected not far from the central region of the Galaxy. Field-star contamination is significant in the colour-magnitude diagrams. Half of the sample have an absorption $\\aV\\ga7$, reaching in some cases $\\aV\\ga15$.} {Stellar radial number-density and surface-brightness profiles are built with 2MASS photometry that, for the present clusters, corresponds basically to giant-branch stars. Field-star decontamination is essential for clusters in dense fields, so an algorithm that we previously developed for open clusters in rich fields is employed to better define cluster sequences. With decontaminated photometry we also compute the total \\mv\\ of four such globular clusters, using M\\,4 as a template. King-like functions are fitted to the radial profiles, from which the core, half-light, half-star count and tidal radii are derived, together with the concentration parameter. Parameters derived here are compared to the equivalent ones of other Galactic globular clusters available in the literature.} {Compared to massive globular clusters, those in the present sample have smaller tidal and larger core radii for a given Galactocentric distance, which is consistent with rather loose structures. Globular cluster radii derived from number-density and surface-brightness profiles have similar values. The present magnitude estimates are $\\mv\\approx-4.9$ (ESO\\,280-SC08), $\\mv\\approx-5.8$ (2MASS-GC01) and $\\mv\\approx-5.6$ (2MASS-GC02). We also estimate $\\mv\\approx-6.1$ for GLIMPSE-C01, which results somewhat less luminous than previously given. The density profiles of Tz\\,4 and 2MASS-GC01 present evidence of post-core collapse clusters.} {Structural parameters and luminosity of most of the faint globular clusters dealt with in this paper are consistent with those of Palomar-like (low-mass and loose structure) globular clusters. This work helps to improve coverage of the globular cluster parameter space.} ",
        "introduction": "\\label{intro}  Astrophysical parameters of globular clusters (GCs) are of great interest, first because they can be used to characterize GCs individually, what eventually will lead to a deeper understanding of their properties as a class in the Galaxy. On a broad perspective, such parameters may provide important clues to the investigation of the formation and evolution processes of the GC system of the Galaxy and elsewhere, e.g. Harris (1996, and the 2003 update\\footnote{\\em http://physun.physics.mcmaster.ca/Globular.html} - hereafter H03), \\citet{MvdB05} and \\citet{GCProp}. On the observational perspective, the depth of this kind of analysis depends directly on the statistical completeness of the sample, especially in the faint-GC tail, and the reliability of the derived parameters.  In general terms, the standard picture of the GC structure assumes a isothermal central region and a tidally truncated outer region, usually described by a roughly spherical symmetry (e.g. \\citealt{Binney1998}). In this context, the structure of most GCs is well approximated by a rather dense core and a sparse halo, but with a broad range in concentration. The most commonly used structural parameters are the core (\\rc), half-light (\\rhL) and tidal radii (\\rt), as well as the concentration parameter $c=\\log(\\rt/\\rc)$, related to the isothermal sphere, single-mass model introduced by \\citet{King1962} to fit the surface brightness profiles (SBPs) of Galactic GCs.  As they age, GCs are continually harassed by external processes such as tidal stress and dynamical friction (from the Galactic bulge and disk, and giant molecular clouds), and internal ones such as mass loss by stellar evolution, mass segregation and low-mass star evaporation (e.g. \\citealt{Khalisi07}; \\citealt{Lamers05}; \\citealt{GnOs97}). Over a Hubble time, these effects tend to decrease cluster mass, which may accelerate the core collapse phase for some clusters (\\citealt{DjMey94}, and references therein). As a dramatic end-result, low-mass clusters may disrupt completely over the course of a few disruption time scales (\\tdis). Observational evidence for the dependence of \\tdis\\ on cluster mass are given by, e.g. \\citet{JA82}, \\citet{Wielen91} and \\citet{Lamers05}. The latter authors found $\\tdis\\sim M^{0.62}$, which agrees with the N-body estimate of \\citet{BM03} for the disruption time scale due to the Galactic tidal field. Theoretically, the dependence of \\tdis\\ on cluster mass was investigated by, e.g. \\citet{Spitzer58} and \\citet{GnOs97}. Most of the above processes are related to the large-scale mass distribution of the Galaxy. Thus, their strength should depend on Galactocentric distance, and they are expected to be more effective on the less-massive clusters (e.g. \\citealt{DjMey94}; \\citealt{vdBergh91}; \\citealt{CheDj89}). Consequently, the distribution of structural parameters among the Galactic GC population, as well as their dependence on Galactocentric distance, may also be related to the physical conditions prevailing at the early formation phase (e.g. \\citealt{vdBergh91}).  All these arguments considered, it is natural to assume that the long-term (more than 10\\,Gyr, in most cases) dynamical evolution of the collective GC population (as well as the individual GCs) may be reflected on the statistical properties of their structural parameters. It is in this context that intrinsically faint GCs play an important r\\^ole, since they are more sensitive to the processes discussed above, and have shorter dynamical-evolution time scales than the massive ones.  In most cases, SBPs of faint GCs may contain significant noise owing to the intrinsically-low surface brightness outside the central region, and  because of the random presence of relatively bright stars, either from the field or cluster members. The same applies to the outer parts, where background starlight is usually a major contributor even in luminous GCs.  Structural parameters measured in such SBPs would certainly be affected in varying degrees. In such cases, the alternative is to work with stellar radial density profiles (RDPs), in which only the projected number-density of stars is taken into account, irrespective of the individual stellar brightness. RDPs are particularly appropriate for the outer parts, provided that a statistically significant, and reasonably uniform, comparison field is available to tackle the background contamination. Such wide fields are provided by 2MASS (e.g. \\citealt{DetAnalOCs}; \\citealt{BB07a}). However, the 2MASS photometric depth limit implies that basically the giant branch can be accessed in GCs more distant than a few kpc from the Sun (\\citealt{BB07b}). Thus, it is necessary to take into account depth-limited effects when considering such structural parameters in the context of absolute comparisons with parameters obtained with deeper photometry. This issue has been studied in \\citet{BB07b} by means of model GCs built assuming different conditions, which allowed us to investigate relations among radii measured in RDPs and SBPs, together with their dependence on photometric depth.  \\citet{TKD95} presented a catalogue of SBPs for 125 Galactic GCs with structural parameters derived mostly with \\citet{King66} profile. Based mostly on the latter work, H03 provided structural parameters for 141 (of the total catalogued 150 at the time) GCs. Recently, \\citet{Cohen07} used surface brightness profiles in \\jj, \\hh\\ and \\ks\\ with 2MASS\\footnote{{\\em http://www.ipac.caltech.edu/2mass/releases/allsky/}} photometry to derive core radii for 104 GCs, by means of fits of a power-law with a core function. The 9 remaining GCs in H03 are faint and were not included in \\citet{Cohen07}.  The compilation of H03 contains 150 GCs but, as a consequence of recent deeper surveys, the number of known GCs in the Galaxy has been increasing. For instance, \\citet{KMB2005} discovered the far-IR GC GLIMPSE-C01. \\citet{Carraro05} identified Whiting\\,1 as a young halo GC. Two stellar systems detected with the Sloan Digital Sky Survey (SDSS) in the outer halo, Willman\\,1 (SDSS\\,J1049$+$5103) and SDSS\\,J1257$+$3419, might be GCs or dwarf galaxies (\\citealt{Willman05}; \\citealt{Sakamoto06}). \\citet{OBB06} identified AL\\,3 as a new GC in the bulge with a blue horizontal branch. Next, \\citet{FMS07} found evidence that FSR\\,1735 is an inner Galaxy GC, and \\citet{Belokurov07} found the faint halo GC SEGUE\\,1 using SDSS. More recently, \\citet{Koposov07} reported the discovery of two very-low luminosity halo GCs (Koposov\\,1 and 2) detected with SDSS. The latter three GCs are beyond the scope of the present paper because they are too distant and/or exceedingly faint. Another object, FSR\\,0584, was recently identified as a nearby Palomar-like halo GC or an old open cluster (OC) by \\citet{FSR584}. Finally, \\citet{FSR1767} identified FSR\\,1767 as a bulge-projected ($\\ell=352.6^\\circ$, $b=-2.17^\\circ$) Palomar-like GC ($\\mv\\approx-4.7$) that appears to be the nearest one ($\\ds\\approx1.5$\\,kpc) so far discovered. Fundamental parameters and updates to the GCs in H03 were provided by \\citet{GCProp}, also available as a Vizier on-line catalogue\\footnote{\\em http://vizier.cfa.harvard.edu/viz-bin/VizieR?-source=J/A$+$A/450/105}.  The main goal of the present work is to derive structural parameters, using RDPs and SBPs built with 2MASS star-counts and photometry, for the remaining 9 faint GCs in H03, which lack this kind of data. The targets are IC\\,1257, Lyng\\aa\\,7, Terzan\\,4, Terzan\\,10, BH\\,176, ESO\\,452-SC11, ESO\\,280-SC08, 2MASS-GC01, and 2MASS-GC02. We also include in the analyses the more recently discovered GCs GLIMPSE-C01 (\\citealt{KMB2005}) and AL\\,3 (\\citealt{OBB06}).  This paper is structured as follows. In Sect.~\\ref{TargetGCs} we present previous results on the present GCs that are relevant for this work. In Sect.~\\ref{PhotPar} we detail the 2MASS photometry, including extraction and errors, build the near-infrared colour-magnitude diagrams (CMDs) for the present GCs, briefly discuss the tools and algorithms employed to decontaminate the CMDs of field stars, and estimate the absolute $V$ magnitudes of ESO\\,280-SC06, 2MASS-GC01, 2MASS-GC02 and GLIMPSE-C01. In Sect.~\\ref{Struc} we build RDPs and SBPs and derive the core, half-light, half-star count and tidal radii, together with the concentration parameter for the present GC sample. In Sect.~\\ref{Disc} we discuss potential effects of depth-limited photometry in the derivation of structural parameters from radial profiles, and compare such parameters with those of other globular clusters in the Galaxy. Finally, concluding remarks are given in Sect.~\\ref{Conclu}. ",
        "conclusions": "\\label{Conclu}  Compilations of astrophysical parameters of globular (e.g. H03) and open (e.g. WEBDA\\footnote{\\em obswww.unige.ch/webda}) clusters represent a fundamental step in the study of several aspects related to the clusters themselves and the Galactic subsystems they belong to.  The majority of the GCs are rather populous and concentrated, which results in high contrast with respect to the background/foreground stellar field. The bottom line is that such GCs can be rather easily observed in most regions of the Galaxy, from near the center to the remote halo outskirts. Thus, through the statistics of observables such as the age, metallicity, luminosity, dynamics and kinematics, GCs can be taken as excellent tracers of the formation history of the Galaxy. Obviously, the same applies to extra-Galactic GCs and their host galaxies.  Faint globular clusters, in particular, are important because, compared to massive clusters, they are more severely affected by the many Galaxy tidal effects and have shorter dynamical-evolution time scales. Consequently, the distribution of their structural parameters throughout the Galaxy may provide clues to better understand the dynamical evolution of individual GCs and their system.  With the present work we derived structural parameters of 9 faint GCs listed in H03, and 2 others recently identified, by means of stellar radial number-density and surface-brightness profiles. These GCs either lacked such parameters completely or, in some cases, had first-order estimates based on CCD observations covering fields smaller than the angular diameters of clusters. In the present work, the profiles were built with wide-field 2MASS photometry (\\jj\\ band), whose depth-limit together with the distances of the present GCs, result basically in giant-branch stars. In all cases, the radial profiles reached at least to the respective tidal radius. We also estimated total \\mv\\ values, which lacked for ESO\\,280-SC06, 2MASS-GC01 and 2MASS-GC02, and provided an independent \\mv\\ estimate for GLIMPSE-C01 (which resulted $\\approx2$\\,mag fainter than previously given), based on field-star decontaminated photometry of the giant stars, using M\\,4 as a template GC. Our estimates are in the range $-6.1\\la\\mv\\la-4.9$, which confirms them as low-luminosity GCs.  The stellar radial profiles were fitted with King-like functions from which the core, half-star count, half-light and tidal radii, as well as the concentration parameter, were derived. Most of the present objects turned out to be so-far neglected Palomar-like (low-mass, loose structure) GCs. The distributions of core and tidal radii, with respect to Galactocentric distance, are also consistent with the Palomar-like nature.  Besides a better understanding on the individual structure of some faint globular clusters, what also results from this work is an improved coverage of the GC parameter space."
    },
    "0711/0711.4027_arXiv.txt": {
        "abstract": "We report the results of CCD $V$ and $r$ photometry of the globular cluster NGC 5466. The difference image analysis technique adopted in this work has resulted in accurate time series photometry even in crowded regions of the cluster enabling us to discover five probably semi-regular variables. We present new photometry of three previously known eclipsing binaries and six SX Phe stars. The light curves of the RR Lyrae stars have been decomposed in their Fourier harmonics and their fundamental physical parameters have been estimated using semi-empirical calibrations. The zero points of the metallicity, luminosity and temperature scales are discussed and our Fourier results are transformed accordingly. The average iron abundance and distance to the Sun derived from individual RR Lyrae stars, indicate values of [Fe/H]=$-1.91 \\pm 0.19$ and D = $16.0 \\pm 0.6$ kpc, or a true distance modulus of $16.02 \\pm 0.09$ mag, for the parent cluster. These values are respectively in the Zinn \\& West  metallicity scale and in agreement with recent luminosity determinations for the RR Lyrae stars in the LMC.  The $M_V$-[Fe/H] relation has been recalibrated as $M_V=+(0.18\\pm0.03)[Fe/H]+(0.85\\pm0.05)$ using the mean values derived by the Fourier technique on RR Lyrae stars in a family of clusters. This equation predicts $M_V=0.58$ mag for [Fe/H]=$-1.5$, in agreement with the average absolute magnitude of RR Lyrae stars calculated from several independent methods. The $M_V$-[Fe/H] relationship and the value of [Fe/H] have implications on the age of the globular clusters when determined from the magnitude difference between the horizontal branch and the turn off point (HB-TO method). The above results however would not imply a change in the age of NGC~5466, of 12.5$\\pm$0.9 Gyr, estimated from recent isochrone fitting. ",
        "introduction": "Determination of the age, distance and metallicity of globular clusters has enormous relevance in several astrophysical problems such as the construction of a model of the formation of the Galaxy, and the stellar structure and evolution on the Horizontal Branch (HB). The distribution of globular cluster ages as a function of galactocentric distance, the dispersion of cluster ages at a given metallicity and the role of a \"second-parameter\" (other than metallicity) in the structure of the HB, are aspects that may be constrained by our knowledge of these fundamental quantities.  It is known that the HB morphology in globular clusters of similar metallicities varies, and the identification of a second parameter responsible of this has originated much discussion in the astronomical literature. Recent findings that the more massive clusters tend to have HBs more extended to higher temperatures led Recio-Blanco et al. (2006) to suggest the cluster total mass may be  the second parameter. Considering the fact that the HB morphology varies due to stellar evolution, the age is still a strong candidate for the second parameter status (e.g. Stetson et al. 1999; Catelan 2000 and references therein).  Two populations of galactic globular clusters have been identified on the basis of their metallicities, kinematics and horizontal-branch structure, which has implications on the history of the formation of our Galaxy (see e.g. Zinn, 1993; van den Bergh 1993; Zinn 1996). The determination of accurate absolute ages of globular clusters is, however, linked to our knowledge of the cluster distance scale (Gratton et al. 1997; Gratton 1998; Carretta et al. 2000; VandenBerg 2000; Gratton et al. 2003). Nevertheless, calculations of relative ages, i.e. cluster age differences relative to well studied nearby clusters (e.g. M92 and M5), at a fixed metallicity, provide an age scale  with an uncertainty $\\leq$ 1 Gyr, which allows one to conclude that the outer-halo globular clusters, with galactocentric distances $R_G \\geq$ 50 kpc, are 1.5-2 Gyr younger than the inner-halo clusters, $R_G \\leq$ 10 kpc,  of the same metallicity (see Stetson et al. 1999; VandenBerg 2000 and references therein). As for the relative ages and their connection with globular cluster metallicities, it has been found that all clusters with [Fe/H] $\\leq -1.7$ are old and coeval and on average 1.5 Gyr older than intermediate-metallicity clusters ($-1.7 \\leq [Fe/H] \\leq -0.8$). All the clusters with [Fe/H] $\\geq -0.8$ are ~1 Gyr younger than the most metal-poor ones. The age dispersions in each metallicity group are $\\leq$ 1 Gyr (De Angeli et al. 2005). There is no correlation of the cluster age nor of metallicity with the galactocentric distance.  Direct tests to calibrate the cluster chronology should involve however, direct measurements of the metallicities of individual stars in the clusters and direct estimates of the distances and ages of the systems. These efforts have been made recently, based on detailed spectroscopic analysis (James et al. 2004; Gratton et al. 2005; Yong et al. 2005; Cohen \\& Melendez, 2005) and  deep photometric measurements of the main sequence and subgiant branch in clusters of different  metallicities and galactocentric distances (e.g. Stetson et al. 1999).  An alternative and efficient method to estimate these relevant physical parameters takes advantage of the fact that the light curve morphology of RR Lyrae stars is connected with fundamental stellar physical parameters, as first suggested by Walraven (1953). The mass, effective temperature, luminosity, helium fraction and iron content for RR Lyrae stars can be calculated from the parameters of the Fourier decomposition of their light curves using the semiempirical calibrations calculated in recent years (e.g. Simon \\& Clement 1993, Jurcsik \\& Kov\\'acs 1996, Morgan et al. 2007). The present capabilities of CCD photometry and difference image analysis (Alard \\& Lupton 1998, Alard 2000, Bramich et al. 2005, Bramich 2007) enables us to perform accurate photometry of very faint stars in remote systems, even in the crowded central regions of globular clusters. There are now numerous works in which the physical parameters of RR Lyrae stars have been estimated from their Fourier light curve decomposition and consequently the mean distances and metallicities of the parent clusters have been determined (for a summary of these works see L\\'azaro et al. 2006).  In the present work we report standard $V$ and instrumental $r$ CCD photometry of NGC 5466, perform the Fourier light curve decomposition of known RR Lyrae stars and report the discovery of new variable stars in the cluster. The globular cluster NGC 5466 (R.A.(2000)$=14^h05^m27^s$.3, DEC(2000)$=+28^{\\circ} 32'04''$) is located in the intermediate galactic halo (l=42.15, b=+73.59, $Z=$ 15.9 kpc, R$_G$=16.2 kpc). It is a rather diffuse system that, up until 1918, was not considered a globular cluster but more likely an open cluster (Shapley 1918). Soon after, however, it was listed for the first time as \"certainly a globular cluster\" (Shapley 1919) and then described as a loose globular cluster (Shapley 1930). Its concentration class is XII and its spectral type F5 (Buonanno et al. 1984). From the periods of its RR Lyrae stars the cluster is of the Oosterhoff type II (Oosterhoff 1939). Its diffuse nature is most likely due to the effects of Galactic tides, an idea supported by the recent discovery of tidal tails around the cluster (Belokurov et al. 2006; Grillmair \\& Johnson 2006).  In the catalogue of variable stars in globular clusters (Clement et al. 2001; 2002) there are reported 32 variables in NGC 5466; 13 RRab, 6 RRc; 2 possible second overtone pulsators or RR2, 1 possible double mode pulsator or RRd, 1 anomalous cepheid; 3 eclipsing binaries and 6 SX Phe stars.  Photoelectric photometry of the cluster dates back to the work of Cuffey (1961) in the $PV$ system. Buonanno et al. (1984) performed $BV$ photographic photometry aimed at producing a colour magnitude diagram and to measure accurate positions of the stars in the cluster. CCD photometry has been reported by Corwin et al. (1999). These authors determine the ephemerides of 20 RR Lyrae stars. In the present paper, we aim to revisit the light curves of known RR Lyrae stars using the relatively new difference imaging technique and to Fourier decompose them into their harmonics, so as to estimate individual physical parameters of astrophysical importance. We also search for new variables.  In Sect. 2 we describe the observations and data reductions. In Sect. 3 individual objects are discussed and new variables are presented. In Sect. 4 we calculate the physical parameters using the Fourier light curve decomposition method and discuss the transformations to homogeneous scales. In Sect. 5 we discuss the results in the wider context of the RR Lyrae distance scale and the age of the cluster. In Sect. 6 we summarize our conclusions.  \\begin{table*} \\footnotesize{ \\begin{center} \\caption[{\\small Standard $V$ and mean instrumental magnitudes $v$ per observing run for 11 standard stars taken from Buonanno et al. (1984). {\\it n} is the number of images considered in the average.}] {\\small Standard $V$ and mean instrumental magnitudes $v$ per observing run for 11 standard stars taken from Buonanno et al. (1984). {\\it n} is the number of images considered in calculating the mean and standard deviation.} \\label{tab_estandar} \\hspace{0.01cm} \\begin{tabular}{|lc|ccc|ccc|ccc|ccc|} \\hline & & $v$ & & & $v$ & & & $v$ & & &$v$& &  \\\\ Star$^*$ & $V$&H04& $\\sigma_{\\rm H04}$ & n &H05& $\\sigma_{\\rm H05}$ &n&K05& $\\sigma_{\\rm K05}$ & n &SPM05& $\\sigma_{\\rm SPM05}$& n \\\\ \\hline F,21 &   14.11 & 15.012 &  0.003& 24 & 15.104 & 0.007& 36 &        &     &   &16.618 &  0.001&4 \\\\ H,3 &   16.66 & 17.677 &  0.008 &41& 17.731 & 0.016& 50 &        &       & &       & &\\\\ I,7 &   15.78 & 16.806 &  0.007& 42& 16.895 & 0.007& 48 &        &      &  &18.457  & 0.010&5\\\\ J,11 &   16.87 & 17.879 &  0.011 &44& 17.974 & 0.022& 52 &        &       & &19.414  &  0.014&5\\\\ M,28 &   17.64 & 18.641 &  0.038& 41& 18.767 & 0.027& 53 & 19.742 &  0.015&3 & 20.278 & 0.008&5\\\\ O,231 &   14.28 & 15.255 &  0.004 &44& 15.385 & 0.006& 51 & 16.617 &  0.008&2 &16.866  & 0.004&5\\\\ P,259 &   14.13 & 15.039 &  0.005& 43& 15.152 & 0.010& 51 & 16.585 &  0.010&3&16.719 &  0.005&5 \\\\ Q,263 &   16.37 & 17.375 &  0.008 &44& 17.480 & 0.017& 53 & 18.667 &  0.007&3 &19.028 &  0.004 &5\\\\ R,266 &   16.67 & 17.702 &  0.039& 43& 17.836 & 0.019&53 & 18.905 &  0.031&3 &19.271 &  0.007 &5\\\\ S,F285 &   14.88 & 15.904 &  0.013 &43& 15.999 & 0.011&53 &        &       & &       & &\\\\ Z,299 &   17.57 & 18.506 &  0.013& 43& 18.635 & 0.022 &51 &        &       & & 20.128 &  0.010 &5\\\\ \\hline \\end{tabular}  \\end{center} * Star identifications with capital letters are from Cuffey (1961) while numbers are the equivalents in Buonanno et al. (1984). } \\end{table*} ",
        "conclusions": "The technique of difference imaging has proven to be a powerful tool in order to find new variables even in crowded image regions. Here we report the discovery of five red semiregular variables, and the independent detection of three known eclipsing binaries and of faint SX Phe stars.  Physical parameters of astrophysical relevance, $log~(L/L_{\\odot})$, $log~T_{\\rm eff}$, $M_V$ and [Fe/H], have been derived for the RR Lyrae stars in NGC~5466 using the Fourier decomposition of their light curves. Special attention has been paid to the conversion of these parameters into broadly accepted scales. It has been found that $T_{\\rm eff}$ for the RRc stars from the calibration of Simon \\& Clement (1993), and for RRab stars from the calibration of Jurcsik (1998), cannot be reconciled with the theoretical predictions of the instability strip bounds. Also, these temperatures produce unexpected distributions of both RRc and RRab stars on the HB, confirming the conclusion of Cacciari et al. (2005) and Di Fabrizio et al. (2005) that the temperatures derived from the above calibrations of the Fourier parameters are uncertain.  Proper transformations are required to set the mean values of $M_V$ and [Fe/H] of the RR Lyraes in adequate scales. Once this is done, the parameters derived for a family of globular clusters define a  $M_V$-[Fe/H] relationship consistent with the metallicity and distance scales determined most recently from the RR Lyraes in the LMC (Gratton et al. 2004; Clementini et al. 2003). The values for NGC 5466 are [Fe/H]=$-1.91 \\pm 0.18$ and $D=16.0 \\pm 0.6$ kpc, or a true distance modulus of $16.02 \\pm 0.09$. The above results and the isochrones of Salarsis \\& Weiss (2002) imply an age of NGC~5466 slightly younger, by about 0.8 Gyr, than previous estimates.  \\begin{table} \\footnotesize{ \\begin{center} \\caption[Nuevas variables] {\\small Standard $V$ magnitudes of RR Lyrae stars in NGC 5466 (portion).} \\label{tab8} \\hspace{0.01cm} \\begin{tabular}{ccc} \\hline Star & HJD & $V$\\\\  \\hline V2&   2453104.2887&\t17.614\\\\ V2&   2453104.3231&\t17.664\\\\ V2 &  2453104.3295&\t17.655\\\\ V2 &  2453104.3357&\t17.661\\\\ V2 &  2453104.3419&\t17.657\\\\ ...&&\\\\ \\hline \\end{tabular} \\end{center} } \\end{table}  \\begin{table} \\footnotesize{ \\begin{center} \\caption[Nuevas variables] {\\small Instrumental $r$ magnitudes of RR Lyrae stars in NGC 5466 (portion).} \\label{tab9} \\hspace{0.01cm} \\begin{tabular}{ccc} \\hline Star & HJD & $r$\\\\  \\hline V2 &  2453104.3262&16.815\\\\ V2 &  2453104.3326&16.819\\\\ V2  & 2453104.3388&16.829\\\\ V2 &  2453104.3526&\t16.812\\\\ V2 &  2453104.3657&\t16.806\\\\ ...&&\\\\ \\hline \\end{tabular} \\end{center} } \\end{table}  \\begin{table} \\footnotesize{ \\begin{center} \\caption[Nuevas variables] {\\small Standard $V$ magnitudes of new variables in NGC 5466 (portion).} \\label{tab10} \\hspace{0.01cm} \\begin{tabular}{ccc} \\hline Star & HJD & $V$\\\\  \\hline V37 & 2453104.3388&\t15.730\\\\ V37&  2453104.3526&\t15.732\\\\ V37 & 2453104.3657&\t15.730\\\\ V37 & 2453104.3789&\t15.737\\\\ V37 & 2453104.3851&\t15.738\\\\ ...&&\\\\ \\hline \\end{tabular} \\end{center} } \\end{table}  \\begin{table} \\footnotesize{ \\begin{center} \\caption[Nuevas variables] {\\small Standard $V$ magnitudes of three known eclipsing bianries (V28, V29 and V30) and six SX Phe (NH29, NH30, NH35, NH38, NH39, NH49) stars in NGC 5466 (portion).} \\label{tab11} \\hspace{0.01cm} \\begin{tabular}{ccc} \\hline Star & HJD & $V$\\\\  \\hline V28 & 2453104.3975&\t18.554\\\\ V28 & 2453104.4037&\t18.511\\\\ V28 & 2453104.4098&\t18.546\\\\ V28 & 2453104.4160&\t18.540\\\\ V28 & 2453104.4400&\t18.551\\\\ ...&&\\\\ \\hline \\end{tabular} \\end{center} } \\end{table}"
    },
    "0711/0711.3607_arXiv.txt": {
        "abstract": "The stellar population, metallicity distribution and ionized gas in the elliptical galaxies NGC\\,6868 and NGC\\,5903 are investigated in this paper by means of long-slit spectroscopy and stellar population synthesis. Lick indices in both galaxies present a negative gradient indicating an overabundance of Fe, Mg, Na and TiO in the central parts with respect to the external regions. We found that $Mg_2$ correlates both with ${\\rm FeI}_{\\lambda5270}$ and ${\\rm FeI}_{\\lambda5335}$, suggesting that these elements probably underwent the same enrichment process in NGC\\,6868. However, only a marginal correlation of $Mg_2$ and ${\\rm FeI}_{\\lambda5270}$ occurs in NGC\\,5903. The lack of correlation between computed galaxy mass and the $Mg_2$ gradient suggests that these elliptical galaxies were formed by merger events. The stellar population synthesis shows the presence of at least two populations with ages of  13\\,Gyr and 5\\,Gyr old in both galaxies.  We have estimated the metallicity of the galaxies using SSP models. The central region of NGC\\,6868 ($|R|\\la0.5$\\,kpc) presents a deficiency of alpha elements with respect to iron and solar metallicity. The external parts,  present a roughly uniform distribution of $\\rm [\\alpha/Fe]$ ratios and metallicities ranging from $[Z/Z_\\odot] = -0.33$ and solar. A similar conclusion applies to NGC\\,5903.  Concerning the emitting gas conspicuously detected in NGC\\,6868, we test three hypotheses as ionizing source: an H\\,II region, post-AGB stars and an Active Galactic Nucleus (AGN). Diagnostic diagrams involving the ratios $[NII]_{\\lambda6584}/H\\alpha$, $[OI]_{\\lambda6300}/H\\alpha$ and $[SII]_{\\lambda6717,31}/H\\alpha$, indicate that values measured in the central region of NGC\\,6868 are typical of LINERs. Together with the stellar population synthesis, this result suggests that the main source of gas ionization in NGC\\,6868 is non-thermal, produced by a low-luminosity AGN, probably with some contribution of shocks to explain ionization at distances of $\\sim3.5$\\,kpc from the nucleus. ",
        "introduction": "\\label{Intro}  A question still open to debate is how elliptical galaxies are formed. Three scenarios have been proposed to address this issue: {\\em (i)} monolithic collapse of gas clouds linked to high energy dissipation rates \\citep{arim87}; {\\em (ii)} galaxy merger \\citep{burk00}; and {\\em (iii)} hybrid models that involve an early monolithic collapse followed by merger events at later times \\citep{ogan05}. A consequence of the monolithic collapse (associated to a large-scale star formation) is that a conspicuous metallicity gradient is expected to establish along the full range of galactic radius (e.g. \\citealt{car93a}; \\citealt{ChioCar02}). However, if merger events are important, the gradient is expected to disappear. A few ellipticals present strong signs of recent dynamical perturbation \\citep{schw90}, which may provide important clues to their formation processes and evolution.  Important information on the above issues has been derived from metal line-strength indices (e.g. \\citealt{davis87}) and their radial variation within galaxies. $Mg2$ line-strength distribution in early-type galaxies, for example, can vary considerably, ranging from essentially featureless to structured profiles showing, e.g., changes of slope possibly associated with kinematically decoupled cores, or anomalies in the stellar population of some ellipticals \\citep{car93b}.  In addition, the origin of the ionized gas in ellipticals is another controversial issue, because until recently it was believed that the interstellar medium (ISM) in these galaxies was not significant. However,  studies by, e.g. \\citet{phi86} have shown that more than 50\\% of the elliptical and lenticular galaxies contain significant amounts of ionized gas. More recently, \\citet{macc96} have traced ionized gas in a sample of 73 elliptical and SO galaxies by means of narrow-band filters centered on ${\\rm H}\\alpha$ and [NII]. They detected emission gas in about 3/4 of the galaxies, with masses amounting to $10^3$ to $10^5\\,\\ms$.  In the present work we investigate properties of the elliptical galaxies NGC\\,6868 and NGC\\,5903, because previous studies have shown that they contain important interstellar medium (\\citealt{macc96}; \\citealt{ferr99}), have complex stellar and gas kinematics (\\citealt{ca00}). Besides, both ellipticals belong to galaxy groups, which might imply past close encounters (and even mergers) with companions. Potential consequences of such past interactions are the triggering of discrete star formation events with stellar populations differing in metallicity and age (\\citealt{wor92}) and the presence of large-scale dynamical perturbations that may affect galaxy morphology by building subsystems such as a stellar and/or dust disk (\\citealt{ca00}). In this context, NGC\\,6868 and NGC\\,5903 can be taken as interesting candidates where the stellar population and metallicity distribution can be investigated, which are important parameters to understand the formation and evolution of these galaxies.  The main goal of the present paper is to investigate, by means of long-slit spectroscopy, the origin, star-formation history, metallicity and ionized gas distribution in NGC\\,6868 and NGC\\,5903. The analyses will be based mostly on a stellar population synthesis method (based on templates built with star clusters and H\\,II regions), and the radial distribution of absorption and emission lines.  This paper is organized as follows. In Sect.~\\ref{Observations} we describe the observations. Lick indices are measuremed are discussed in Sect.~\\ref{Lick}. The stellar population synthesis and results are described in Sect.~\\ref{StellarPop}. Metallicity and ionized gas are discussed in Sect.~\\ref{MetGas}. Discussions are in Sect.~\\ref{Dicussion}. ",
        "conclusions": "\\label{Dicussion}  The stellar population, metallicity distribution and ionized gas in NGC\\,6868 and NGC\\,5903 have been investigated in this paper by means of long-slit spectroscopy and stellar population synthesis. Lick indices  of both galaxies  consistently present a negative gradient that indicates an overabundance of Fe, Mg, Na and TiO in the central parts with respect to the external regions. We found that $Mg_2$ correlates both with ${\\rm FeI}_{\\lambda5270}$ and ${\\rm FeI}_{\\lambda5335}$, which suggests that these elements probably underwent the same enrichment process in NGC\\,6868. However, in NGC\\,5903 only a marginal correlation of $Mg_2$ and ${\\rm FeI}_{\\lambda5270}$ occurs.  Galaxy mass and the $Mg_2$ gradient computed in the present work are consistent with previous results that suggest that NGC\\,6868 and NGC\\,5903 were formed by merger events. In addition, the stellar population synthesis clearly shows the presence of at least two populations of different ages in both galaxies. Particularly in the central regions of NGC\\,6868 and NGC\\,5903, the 13\\,Gyr population contributes with $70\\pm16$\\% and $56\\pm12$\\%, respectively. The 5\\,Gyr population contributes with $28\\pm15$\\% and $25\\pm12$\\%, respectively in flux fraction at $\\rm5870\\,\\AA$.  The central regions of NGC\\,6868 and NGC\\,5903 present a deficiency of alpha elements with respect to the external parts, which suggests different star-formation histories in both regions. With respect to the metallicity, the central regions of both galaxies have higher $[Z/Z_\\odot]$ values than the external parts. The range in metallicity spanned by the sampled regions of NGC\\,6868 appears to be larger than in NGC\\,5903.  Concerning the ionized gas in NGC\\,6868, the ratios $[NII]_{\\lambda6584}/H\\alpha$, $[OI]_{\\lambda6300}/H\\alpha$ and $[SII]_{\\lambda6717,31}/H\\alpha$ consistently suggest the presence of a LINER at the galaxy center. These results, together with the stellar population synthesis, suggest that the main source of gas ionization in NGC\\,6868 is non-thermal, produced by a low-luminosity active nucleus, probably with some contribution of shocks to explain ionization at distances of $\\sim3.5$\\,kpc from the nucleus."
    },
    "0711/0711.3868.txt": {
        "abstract": "A recently developed spectral-element adaptive refinement incompressible magnetohydrodynamic (MHD) code [Rosenberg, Fournier, Fischer, Pouquet, J. Comp. Phys. 215, 59-80 (2006)] is applied to simulate the problem of MHD island coalescence instability (\\ci) in two dimensions. \\ci\\ is a fundamental MHD process that can produce sharp current layers and subsequent reconnection and heating in a high-Lundquist number plasma such as the solar corona [Ng and Bhattacharjee, Phys. Plasmas, 5, 4028 (1998)]. Due to the formation of thin current layers, it is highly desirable to use adaptively or statically refined grids to resolve them, and to maintain accuracy at the same time. The output of the spectral-element static adaptive refinement simulations are compared with simulations using a finite difference method on the same refinement grids, and both methods are compared to  pseudo-spectral simulations with uniform grids as baselines. It is shown that with the statically refined grids roughly scaling linearly with effective resolution, spectral element runs can maintain accuracy significantly higher than that of the finite difference runs, in some cases achieving close to full spectral accuracy. \\\\ ",
        "introduction": "\\label{Intro}  In many hydrodynamic or magnetohydrodynamic  (MHD) applications in astrophysics or space physics, it is essential that a numerical simulation resolve the development of sharp spatial structures accurately. While pseudo-spectral methods generally can maintain high accuracy, they are mainly applied on more regular geometry and require more uniform grids, which can make it difficult to reach high resolution in order to resolve sharp isolated structures especially in flows dominated by such structures. Static or adaptive mesh refinement (AMR) methods can put more grid points in and around isolated structures in order to resolve them, but may not achieve similar high-accuracy if low order spatial discretizations are used \\cite{rosenberg2007}. They are particularly useful in bounded flows, where pseudo--spectral methods are often not optimized. Therefore, it is of great interest to develop numerical schemes that can combine high accuracy and high spatial resolution. Adaptive spectral element methods have the potential to do just that, providing spectral-like accuracy that can be applied efficiently to resolve isolated structures. In this paper, we concentrate on comparing the accuracy of simulation results from a spectral element based AMR code (SE)\\cite{rosenberg2006, rosenberg2007} and a finite difference based AMR code (FD) \\cite{astronum2006,b2005} on an astrophysical problem that requires high spatial resolution as well as high accuracy, the so--called  MHD island coalescence instability (\\ci) problem \\cite{nb1998}. In order to make meaningful comparisons, we let each code run on essentially the  same non-uniform grid (i.e. with the total degrees of freedom in the problem fixed) that is refined a priori in regions of the grid where the current sheets will form in the \\ci. A separate pseudo-spectral code (PS), running on uniform grids, is also used to provide a baseline for the comparisons.  There is little in the literature about attempts to compare FD and high order methods. Often, it is simply accepted that the fixed lower order FD methods will be less accurate, but that, due to their simplicity and efficiency, {\\it h}--type grid refinement can always be carried out to improve the accuracy for any real problem. We do not attempt in this work to compare performance metrics of the methods, preferring instead to focus on their ability to produce accurate solutions. To an extent, then, it is clear that {\\it h}-refinement will improve solution accuracy for many problems; however, previous work \\cite{rosenberg2007} suggests that local high order may be {\\it required} in certain instances. We consider this issue of {\\it h}-refinement in FD briefly in this work. Previous {\\it comparative} work in one dimension \\cite{basdevant86} demonstrates that spectral methods, including SE, are more accurate than FD methods due primarily to dispersion problems in the latter. Nevertheless, this work concludes that the SE method is not well suited to the calculation of thin internal layers (structures), particularly when their location is unknown. This assertion is made because, while accuracy is found to be good, Basdevant et al. 1986 find that polynomial orders are required to be inordinately high even in one dimension. This conclusion is refuted in later work \\cite{mavriplis94}, which found that the SE method is indeed well suited to this type of problem provided adaptivity is used (see also \\cite{rosenberg2006}). More recent work \\cite{st-cyr2007} also demonstrates that the SE method exhibits smaller errors than a low-order finite volume scheme at comparable resolutions for nearly all of a series of tests of a shallow water model on a cubed--sphere grid that is either adaptively or statically refined. However, to the best of our knowledge, ours is the first work that performs a quantitative comparison of FD with SE methods in divergence--free nonlinear flows in two space dimensions with nonconforming (statically) adaptive grids.  We continue our introductory remarks on the \\ci\\ in \\Sec{sec_island_instability}, where we present details of the problem, and provide a motivation for our work. In the process, we discuss important properties of the \\ci\\ as well as some aspects of MHD flows that will serve our later discussion. We provide details about the simulation set up as well as some diagnostic measures used in the comparisons in \\Sec{sec_props}. In \\Sec{sec_methods} we present the numerical methods used in sufficient detail to elucidate the results. Our results are presented in \\Sec{sec_results}, where we offer several types of accuracy comparisons relevant to MHD flows, and the \\ci, in particular. We conclude in \\Sec{sec_conclusion} with a summary of our findings, and some discussion about the relative advantages of the methods. ",
        "conclusions": "\\label{sec_conclusion}  It is worth noting that at the present stage, the spectral element method described in this paper is more costly in terms of computational time than the FD and PS methods with which we compare our results, the latter being optimal for periodic boundary conditions. We have not tried to compare CPU usage of the different codes due to the fact that the SE and FD codes are running on different computational platforms; furthermore, both codes are still under continuous development, which can change efficiency greatly. Instead, we concentrate on comparing accuracy obtained by the two schemes, mainly to answer the question of whether using SE scheme has any advantage that deserves putting more effort into further development. We believe the results of this paper have shown that indeed the SE method has substantial advantages and needs to be investigated further.  We have shown that using statically refined grids with DOF of roughly linear scaling, SE method can produce results with high accuracy that can sometimes even approach the spectral accuracy of the PS method. This can be potentially very helpful in the investigation of the 3D \\ci\\ problem in the ideal limit, which will require adaptive grids that can resolve distinct features. The higher accuracy of the SE method can potentially provide a reliable definitive answer to the important question of whether a true current sheet forms in the Parker's problem of solar and stellar coronal heating.  Our main conclusion that SE methods can produce simulations with accuracy somewhat in between PS and FD methods is not surprising in itself. However, our results yield important quantitative information in the context of statically refined grids in problems with distinct spatial structures.  The accuracy of FD runs presented here is mainly for comparison purposes, since we have imposed the restriction of using the same static grids with the same DOF. This is by no means a suggestion that FD method cannot be used in the investigation of the \\ci, or similar problems. Indeed we have also studied briefly that accuracy in FD method can also be increased by using more DOF. However, we have not studied the trade off of doing this with respect to the increase of the usage of computational resources. This may be an important topic for future investigation."
    },
    "0711/0711.2398_arXiv.txt": {
        "abstract": "We report on the results of a time-series photometric survey of NGC 2362, carried out using the CTIO 4m Blanco telescope and Mosaic-II detector as part of the Monitor project.  Rotation periods were derived for $271$ candidate cluster members over the mass range $0.1 \\la M/\\msun \\la 1.2$.  The rotation period distributions show a clear mass-dependent morphology, qualitatively similar to that in NGC 2264, as would be expected from the age of this cluster.  Using models of angular momentum evolution, we show that angular momentum losses over the $\\sim 1-5\\ {\\rm Myr}$ age range appear to be needed in order to reproduce the evolution of the slowest rotators in the sample from the ONC to NGC 2362, as found by many previous studies.  By incorporating \\spitzer IRAC mid-IR measurements, we found that $3-4$ objects showing mid-IR excesses indicative of the presence of circumstellar discs were all slow rotators, as would be expected in the disc regulation paradigm for early pre-main sequence angular momentum evolution, but this result is not statistically significant at present, given the extremely limited sample size. ",
        "introduction": "\\label{intro_section}  NGC 2362 is a very young open cluster, at a moderate distance.  We adopt the parameters of \\citet{moi2001} for the remainder of this work: an age of $5^{+1}_{-2}\\ {\\rm Myr}$, distance $1480\\ {\\rm pc}$ and reddening $E(B-V) = 0.10\\ {\\rm mag}$.  Note however that there is some controversy in the literature regarding the age of this cluster, with estimates ranging from $3 - 9\\ {\\rm Myr}$.  \\citet{may07} have examined the relative ages of young open clusters, finding that NGC 2362 does indeed seem to be older than the ONC, but that it may be younger than IC 348, and they suggest an age of $\\sim 3\\ {\\rm Myr}$.  Nevertheless, the extremely young age, and lack of substantial nebulosity or differential reddening in this cluster make it an ideal testing ground for theories of pre-main sequence stellar evolution. \\citet{dahm05} carried out an extensive spectroscopic survey, concentrating on the H$\\alpha$ $6563\\ {\\rm \\AA}$ and Li\\ {\\sc i} $6707.8\\ {\\rm \\AA}$ features, finding that a large fraction ($\\sim 91$ per cent) of the T Tauri star (TTS) population in NGC 2362 is composed of weak-line emitters (WTTS), as expected given the canonical $5\\ {\\rm Myr}$ age for the cluster.  They found that $\\sim 5-9$ per cent of the TTS population are likely to still be undergoing accretion on the basis of the measured H$\\alpha$ equivalent widths, which is comparable to the fraction of objects found to exhibit L-band excess emission indicative of an inner disc by \\citet{hll2001}.  NGC 2362 has also been observed more recently with \\spitzer using the IRAC (Infra-Red Array Camera) instrument in the $3.6$, $4.5$, $5.8$ and $8.0\\ {\\rm \\mu m}$ bands.  \\citet{dh07} used these observations to derive spectral energy distributions for candidate NGC 2362 members, and to characterise disc emission in objects exhibiting infra-red excesses.  They found an upper limit to the fraction of objects exhibiting optically thick discs of $\\sim 7 \\pm 2$ per cent, with an additional $\\sim 12 \\pm 3$ per cent of objects exhibiting weak or optically thin discs.  \\subsection{Evolution of stellar angular momentum} \\label{amevol_section}  A long-standing issue in our understanding of the evolution of stellar rotation rates on the pre-main sequence is that significant loss of angular momentum appears to be required in order to reproduce the observed rotation rates of low-mass stars.  During the first few Myr of their evolution, low-mass pre-main sequence (PMS) stars contract by factors of $\\sim 2-3$.  Hence, if angular momentum is conserved, all sub-solar mass stars, which are thought to have initial rotation periods of $\\sim 10\\ {\\rm days}$, should rotate at periods shorter than $2\\ {\\rm days}$ after only $2-3\\ {\\rm Myr}$.  Observations at these ages show that some stars rotate much slower than this rate, and indeed the slowest rotators do not appear to spin up at all over the first few Myr of their evolution.  The most popular scenarios for solving this problem invoke additional angular momentum losses resulting from the interaction of the star with a circumstellar disc (e.g. \\citealt{k91}).  In the ``disc locking'' scenario, this interaction results in transfer of angular momentum from the star to the disc, which maintains a constant stellar rotation rate until the disc dissipates sufficiently to break the coupling.  Alternatively, the angular momentum may be lost through an enhanced stellar wind, powered by the accretion of material from the disc \\citep{mp05}.  In both cases, the most basic prediction of the models is that there should be a correlation between rotation rate and the presence of discs, in the sense that the slow rotators should still be surrounded by discs, whereas for the rapid rotators, the discs should have dispersed.  Much of the recent literature on rotation has been devoted to the search for a correlation between rotation rate and disc indicators, including diagnostics of accretion (UV excess, spectral line diagnostics such as H$\\alpha$, etc.) and of the thermal emission from the disc itself (infra-red excess).  Many of the early studies were inconclusive, due to a combination of effects including small sample sizes, biases, dependence of the adopted disc indicators on other stellar parameters (predominantly mass), and the use of ambiguous disc indicators.  Two recent studies (\\citealt{rebull06}; \\citealt{cb07}) have claimed statistically significant detections of the expected correlation between slow rotators and circumstellar discs, in the ONC ($1 \\pm 1\\ {\\rm Myr}$; \\citealt{h97}) and NGC 2264 ($2-4\\ {\\rm Myr}$; \\citealt{park2000}), respectively.  \\citet{cb06} have also examined IC 348 ($2-3\\ {\\rm Myr}$; \\citealt{luh2003}), but the results were ambiguous due to small number statistics.  These studies have all used the IRAC instrument on the \\spitzer space telescope, which provides mid-IR photometry in four broad bands: $3.6$, $4.5$, $5.8$ and $8.0\\ \\micron$.  In particular, $8.0\\ \\micron$ infra-red flux excesses provide a much more reliable indicator of the presence of circumstellar discs than shorter-wavelength IR excesses (e.g. in the $K$-band).  There is thus mounting evidence in favour of the disc regulation paradigm, but nevertheless, it is important to verify that this result extends to older open clusters.  \\subsection{The survey}  We have undertaken a photometric survey in NGC 2362 using the CTIO 4m Blanco telescope and Mosaic-II imager.  Our goals are two-fold: first, to study rotation periods for a sample of low-mass members, covering K and M spectral types, down to $\\sim 0.1\\ \\msun$, and second, to look for eclipsing binary systems containing low-mass stars, to obtain dynamical mass measurements in conjunction with radial velocities from follow-up spectroscopy.  Such systems provide the most accurate determinations of fundamental stellar parameters (in particular, masses) for input to models of stellar evolution, which are poorly constrained in this age range.  We defer discussion of our search for occultations to another publication (Miller et al., in preparation).  These observations are part of a larger photometric monitoring survey of young open clusters over a range of ages and metalicities (the Monitor project; \\citealt{hodg06} and \\citealt{a2007}).  The remainder of the paper is structured as follows: the observations and data reduction are described in \\S \\ref{odr_section}, and the colour magnitude diagram (CMD) of the cluster and candidate membership selection are presented in \\S \\ref{memb_section}.  The method we use for obtaining photometric periods is summarised in \\S \\ref{period_section} (see \\citealt{i2006} for a more detailed discussion).  Our results are given in \\S \\ref{results_section}, and \\S \\ref{conclusions_section} summarises our conclusions.  In several of the following sections, we use mass and radius estimates for the cluster members.  These were derived from the $5\\ {\\rm Myr}$ NextGen models, using the $I$-band magnitudes, rather than $V - I$ colour or $V$ magnitude, for reasons discussed in \\S \\ref{cmd_section}. ",
        "conclusions": "\\label{conclusions_section}  We have reported on results of an $i$-band photometric survey of NGC 2362, covering $\\sim 0.4\\ {\\rm sq. deg}$ of the cluster.  Selection of candidate members in a $V$ versus $V-I$ colour-magnitude diagram using an empirical fit to the cluster sequence found $1800$ candidate members, over a $V$ magnitude range of $16 < V < 26$ (covering masses from $1.2\\ \\msun$ down to below the brown dwarf limit).  The likely field contamination level was estimated using a simulated catalogue of field objects from the Besan\\c{c}on Galactic models \\citep{r2003}, finding an overall contamination level of $\\sim 65$ per cent, implying that there are $\\sim 630$ real cluster members over this mass range in our field-of-view.  We derived light curves for $\\sim 85\\,000$ objects in the NGC 2362 field, achieving a precision of $< 1$ per cent per data point over $15 \\la i \\la 19$.  The light curves of our candidate cluster members were searched for periodic variability corresponding to stellar rotation, giving $271$ detections over the mass range $0.1 < M/\\msun < 1.2$.  The rotation period distribution as a function of mass was found to show a clear mass-dependent morphology, as seen in our earlier NGC 2547, NGC 2516 and M34 studies, and qualitatively similar to that in NGC 2264, which is thought to have a similar age.  Using rotational evolution models from our earlier NGC 2516 work \\citep{i2007b}, we were able to show that distribution of rotation rates for the rapid rotators in NGC 2362 could be reproduced from the distributions in the ONC and NGC 2264 by evolving them forward in time to $\\sim 5\\ {\\rm Myr}$.  However, for the slow rotators at masses $\\ga 0.6\\ \\msun$, the predicted rotation rates from the models were too rapid, with the observations indicating that the rotation periods of these stars are in fact approximately constant over the $\\sim 1-5\\ {\\rm Myr}$ age range.  This confirms the need for additional angular momentum losses on this part of the PMS, e.g. as would be expected in the disc regulation paradigm.  Finally, we have used the \\spitzer observations of \\citet{dh07} to test the disc regulation hypothesis, by searching for a correlation between rotation rate and mid-IR excesses indicative of circumstellar discs.  No statistically significant correlation was found with the present small sample sizes, but the $3-4$ objects with significant disc signatures in the sample are indeed all slow rotators."
    },
    "0711/0711.0391.txt": {
        "abstract": "We present results from 1351 high resolution spectra of 1215 stars in the Orion Nebula Cluster (ONC) and the surrounding Orion 1c association, obtained with the Hectochelle multiobject echelle spectrograph on the 6.5m MMT. We confirmed 1111 stars as members, based on their radial velocity and/or \\halp~ emission. The radial velocity distribution of members shows a dispersion of $\\sigma=3.1 $ \\kms. We found a substantial north-south velocity gradient and spatially coherent structure in the radial velocity distribution, similar to that seen in the molecular gas in the region. We also identified several binary and high velocity stars, a region exhibiting signs of triggered star formation, and a possible foreground population of stars somewhat older than the ONC.  Stars without infrared excesses (as detected with the IRAC instrument on the {\\em Spitzer Space Telescope}) exhibit a wider spread in radial velocity than the infrared excess stars; this spread is mostly due to a blue-shifted population of stars that may constitute a foreground population.  We also identify some accreting stars, based on H$\\alpha$, that do not have detectable infrared excesses with IRAC, and thus are potential transitional disk systems (objects with inner disk holes). We propose that the substructure seen both in stellar and gaseous component is the result of non-uniform gravitational collapse to a filamentary distribution of gas. The spatial and kinematic correlation between the stellar and gaseous components suggests the region is very young, probably only $\\sim1$ crossing time old or less to avoid shock dissipation and gravitational interactions which would tend to destroy the correlation between stars and gas. ",
        "introduction": "Because most stars are formed in clusters (e.g., \\citet{lada03}), the processes responsible for cluster formation are important to include in any consideration of the mechanisms of star formation. Observations of very young clusters can provide clues to the initial conditions of cluster formation if the cluster has not dynamically relaxed. The increasing sensitivity of infrared studies with both ground-based instruments and the Spitzer Space Telescope has made it more feasible to search for substructure in embedded populations, as seen for example in the young cluster NGC 2264 \\citep{teix06}.  In addition, gas and dust can be cleared away rapidly by stellar energy input, revealing many young cluster members at optical wavelengths.  In the case of NGC 2264, spatially-coherent kinematic substructure in the stellar population now has been detected using optical spectroscopy by \\citet{fure06}.  The Orion Nebula Cluster (ONC) is a touchstone for studies of cluster formation, as it is the closest, relatively populous ($\\gtrsim 2000$ members) young cluster containing an O star.  While many stars in the ONC are heavily extincted (\\citet{ali95}; \\citet{mcca94}; Carpenter, \\citet{carp01}; \\citet{mcca02}; see also \\citet{odel01} and references therein), evaporation of molecular gas by the O6-7 central star $\\theta^1$~C Ori has revealed many of the members of the ONC at visible wavelengths. Optical studies of the stellar population show that it is quite young, with a median age of roughly 1 Myr or less \\citep{hill97}, suggesting that it is reasonable to search for substructure in the cluster related to the initial conditions of formation.  Hillenbrand \\& Hartmann (1998, HH98 hereafter) conducted a preliminary study of the structure and kinematics of the ONC. HH98 constructed azimuthally-averaged stellar source counts and showed that they could be fit with a simple, spherically-symmetric, single-mass King cluster model of core radius $\\sim 0.2$~pc and a central density of $\\sim 2 \\times 10^4$~stars~pc$^{-3}$ (see also \\citet{mcca02}). However, HH98 pointed out that the cluster is not circular in projection; rather, it is elongated north-south along the same direction as the molecular gas filament, suggesting that the cluster might not be completely relaxed.  This elongation is particularly apparent on scales larger than 0.5 pc from the center. The filamentary structure is well expressed in CO observations of the region \\citep{ball87}, as well as in the distribution of pre-main sequence (PMS) stars, shown by recent infrared (IR) observations (Megeath et al 2007, in preparation). According to the models of \\citet{burk04}, self-gravitationally collapsing gaseous sheets can form such elongated, filament-like clouds. The numerical simulations show that density and velocity dispersion becomes larger at the ends, inducing star formation to start at the tip of the filament, just like the position of ONC within the Orion A cloud. The collapsing sheet model also predicts an overall velocity gradient along the filament, in agreement with CO observations of ONC \\citep{ball87}. For such a young star forming region as ONC, the stars are still very close to their birthplace; therefore, the stellar kinematics may reflect initial conditions rather than being relaxed.  We therefore conducted a radial velocity (RV) survey of more than 1200 stars in the northern region of the Orion A cloud to search for clues to the initial conditions of the ONC's formation.  In this paper we present results based on high-resolution spectroscopic observations in the $\\lambda \\sim 6560$~\\AA~ region, which allowed us to measure RVs and identify T Tauri stars (TTS) based on their \\halp~ emission. We find that the stellar population exhibits a strong spatial and kinematic correlation with the $^{13}$CO gas in the region; the observed substructure conclusively demonstrates that the ONC is not relaxed, but instead reflects the initial conditions involved in its formation.  We also comment on the detailed structure, distributional differences seen among stellar groups distinguished by infrared, \\halp, and spatial and kinematic properties. In a subsequent paper we will incorporate followup spectroscopy we have obtained in the Mg triplet region (centered at 5225 \\AA) to examine rotational velocities and spectral typing, along with observations of the 6708 \\AA~ Li line to use as an age estimator. ",
        "conclusions": "\\subsection{Cluster in Formation}  The significant position-velocity substructure seen in both the stars and the molecular gas clearly demonstrates that the ONC is not dynamically relaxed; qualitatively, it appears to be more consistent with dynamical models of cluster formation such as those of Hartmann \\& Burkert (2007) and Bate, Bonnell, \\& Bromm (2003). The curvature of the $^{13}$CO emission seen in the northern region of the position-velocity plot (Figure 4), with a possible corresponding southern feature, is suggestive of gravitational acceleration of material toward the cluster center, where the gravitational potential well should be deepest. Gravitational acceleration is also consistent with the evidence for the largest velocity distribution in the gas located near the center. Many stars exhibit the same position-velocity distribution as the $^{13}$CO, which suggests that these stars are mostly following the motion of the dense gas within which they formed. If these motions were primarily the result of stellar energy input through winds or photoionization, it is not clear that the gas and many stars should show the same kinematics, as such energy input could easily blow the gas away from the stars - some possible examples of this are discussed in the following section.  The observed correlation of stars and dense gas in space and in velocity indicates that the inner regions cannot have experienced more than about one crossing or collapse time; otherwise the velocity structure would be erased as the infalling gas shocks and dissipates its kinetic energy, with the stars passing through the shocked gas. Overall, the spatial coherence in the $^{13}$CO gas kinematics, the evidence that many younger stars follow this motion, and the strongly filamentary structure of the stellar distribution outside the main core of the ONC suggest that the ONC still exhibits features traceable to its initial conditions of formation.  \\subsection{Subgroups and stellar energy input}  While there is a strong general spatial and kinematic correlation between stars and dense gas, there are exceptions.  One especially clear case can be seen on the bottom left panel of Fig. \\ref{fig:ra_chanel} at $DEC\\simeq-5\\degr27\\arcmin$: a tight, linear string of stars  sticking out of the main cloud as a ``nose'' extending between $RV=21$ and 25 \\kms. The left-hand panel of \\ref{fig:nose_proj} shows a declination-velocity plot zoomed into this region, in an RA range shifted slightly westward of the Trapezium region (vertical lines in the right-hand, RA-DEC panel of Fig. \\ref{fig:nose_proj}).  Clearly there is a small group of stars blueshifted by about 1-2 \\kms~ from the gas in the declination range $-5^{d}~20^{m}~>DEC>~-5^{d}~30^{m}$.  Selecting ``off-cloud'' stars by means of an velocity envelope for the gas (thin curve in left panel) results in demonstrating that these stars are spatially concentrated in a region south-west of the Trapezium which is relatively evacuated, as indicated by the weak $8 \\mu$m emission seen in the IRAC map (right panel, grayscale).  This region is also known for a large number of Herbig-Haro objects shaped like bow-shocks pointing back toward the Trapezium region (see description in \\citet{odel01}), consistent with the idea that outflows from the region of most current star formation is blowing out material.  We therefore suggest that the molecular gas in this region originally extended to more negative velocities, but that outflows have cleared away the gas, making an ``indentation'' in the position-velocity plot at $\\sim -5^{d}~26^{m}$, RV $\\sim 24$ \\kms, and leaving the recently-formed stars behind.  There are other interesting small scale structures, like the tight, dense sub-cluster group seen in the middle of the second panel of Fig.  \\ref{fig:dec_chanel} (upper row, 2nd panel from left to right, also somewhat apparent on Fig. \\ref{fig:ra_chanel} upper row, rightmost panel). These can be remnants of small building blocks, traces of density variations in the primordial cloud.   \\subsection{Other Populations?} \\label{foreground}  Based on Fig. \\ref{fig:irac_ccd}, we divided our sample into \\textit{IR-excess sources}, with $([3.6]-[4.5]) > 0.2$ and $([5.8]-[8.0]) > 0.5$; and \\textit{non-IR excess sources} with $([3.6]-[4.5]) < 0.2$ and $([5.8]-[8.0]) < 0.5$. This division of sources is shown on the velocity-declination map with full projection along RA in Figure \\ref{fig:irac_wtts}: the left panel for excess and the middle one for non-excess stars.  Looking at the left panel in Figure \\ref{fig:irac_wtts}, southward of $-6\\degr$ there are very few CTTS. This abrupt change is an unfortunate selection effect. The targets selected for our second observing run (S fields) had a southern declination limit of $-6\\degr$, exactly where the sudden change of excess/non-excess ratio happens. And since targets for the first observing run (F fields) were selected from the 2MASS color--magnitude diagram, our selection was biased against stars with IR excesses because of the $(H-K)<0.5$ criterion.  While many of the non-excess stars (middle panel) follow the general distribution of gas and excess stars, a modest number clearly exhibit a broader distribution in radial velocity, with more blue- than redshifted stars. This can be seen on the right panel, where we present RV histograms of excess (solid bars) and non-excess (gray-shaded boxes) stars, for three declination ranges (north, with $DEC>-5.2$; mid, with $-5.2>DEC>-6.0$; and south, $DEC<-6.0$ - as represented by the location of the histograms as well). We divided this way because RV bahavior of the gas is very different in these regions: the northern velocity is significantly higher than the mean velocity in the south, and the middle section exhibits a strong RV gradient. Regardless, the velocity distribution of non-excess stars is wider because of a secondary peak at lower velocities.  One possibility is that the lower-velocity stars have been ejected from the main cluster.  Close encounters in multiple star systems might result in ejecting stars while stripping their disks; the red-shifted stars might then eventually plunge into the dense, opaque regions behind, leaving more optically-visible stars with blueshifts.  However, this seems unlikely because it requires a surprisingly large fraction of stars to be ejected, we do not detect a larger density of such objects nearer the center of the ONC, where the high stellar density would be more favorable to ejection. In addition, south of $-6^{\\circ}$ the velocity distribution is clearly bimodal, with a very tight velocity spread for the on-cloud stars; it is not clear why ejection would lead to such distinct structures.  Another possibility is that we are detecting a foreground, older population. Stars in the foreground Orion 1a association (Brown \\etal 1994) may have a heliocentric radial velocity near $\\sim 20$ \\kms, as suggested by measurements of the 25 Ori group \\citet{bric07}.  If this is the case, it would mean that the Orion 1a association is much more spread spatially towards the ONC (in the foreground) than has previously been realized.  A third possibility is that stellar energy input blew out material, forming a small proportion of stars in a bubble wall moving toward us.  To explore this idea further, we plot only the stars with RV values less than 22 \\kms in the RA--DEC plot of Figure \\ref{fig:foreground}. Since there is no gas here, we display the most blueshifted 1 \\kms~ wide channel of the gas at 23 \\kms~ heliocentric velocity. The northern foreground stars seems to be separated from the older southern group, by an almost empty gap running east--west at $-6\\degr00\\arcmin$. However, this might be the result of a combination of our bias against selecting infrared-excess stars mentioned above and accidental fiber coverage. Still, there is some structural similarity between the distribution of these southern foreground stars and the gas ``behind'', redshifted to it. Note how these southern stars are spread on top of two-three denser gas clumps. If this structure is interpreted as a local ``bubble'', this could have been created at the very beginning of star formation by the early emergence of some OB stars in the region. These stars than could have blown away most of the gas from us, some smaller amount towards us, and so other stars could have formed at this early stage in this perturbed southern region.  This picture would require some OB stars in the questioned region, and the list of \\citet{brow94} contains six of them interestingly close to $-6$ degree declination (see Fig. \\ref{fig:foreground}). However, obviously none of these is in close vicinity of the gas now, as there are no apparent HII regions at this part of the molecular cloud. An additional problem is that one would expect such a population to be somewhat older than the main cloud, which is not obvious in the J vs. J-H diagram.  This is not a particular problem for the foreground population hypothesis, as the Orion 1a stars will appear higher in the color-magnitude diagram at a given age than the ONC.    In conclusion, the reasons for the additional velocity spread of the non-excess stars are not clear.  Further monitoring for radial velocity variations as well as more detailed studies of the stellar properties are needed to determine which of the above possibilities are correct.   \\subsection{Comparison with Other Observations and Models}  If the ONC is not dynamically relaxed, why did HH98 find a fairly good fit to a King cluster model?  (Note that \\citet{scal05} point out difficulties with physically interpreting the ONC in this way). The central region of radius $\\sim 0.5$~pc does appear to exhibit a smooth distribution, and thus may be closer to having relaxed, though it is somewhat elongated. On larger scales, the elongation is sufficiently large that azimuthal averaging is not appropriate.  To take an extreme example of what this averaging can do, imagine a cylindrical distribution of material in the plane of the sky with uniform surface density and width $W$. Azimuthal averaging of this distribution results in an average surface density $\\Sigma \\sim$~constant for $r < W$ and $\\Sigma \\propto r^{-1}$ for $r \\gg W$. Interpreting this surface density in terms of a spherically-symmetric density distribution results in $\\rho \\sim$~constant for $r < W$ and $\\rho \\propto r^{-2}$ for $r \\gg W$, which is roughly consistent with a King model. The ONC is not as extreme in structure as this simple example, but it illustrates the potentially misleading nature of azimuthal averaging of an intrinsically highly elongated structure.  Comparing azimuthally averaged stellar densities and global velocity dispersions to N-body calculations, \\citet{krou00} was unable to distinguish whether the ONC is in equilibrium, or is collapsing, or is expanding; a similar result was obtained \\citet{scal05}.  This demonstrates the importance of avoiding azimuthal averaging and obtaining detailed kinematic observations of gas and stars in developing an understanding of the dynamical state of the region.  Tan, Krumholz, \\& McKee (2006; TKM06) have argued for a picture in which rich star clusters take several dynamical times to form, are quasi-equilibrium structures during formation, and thus initial conditions are not very important. In particular, TKM06 discuss the ONC and argue that it is several crossing times old, of order 3 Myr, larger than assumed here.  Here we discuss why we arrive at different conclusions.  Following \\citet{scal02}, TKM06 argue that the smoothness of the spatial distribution of the stars in the ONC argues for long formation times, which allow clumps of stars to disperse.  As pointed out above, the azimuthal averaging done by \\cite{hill98} was somewhat misleading, helping to smooth out filamentary structure.  In addition, the gas kinematics, and to a lesser extent the stellar kinematics, show substructure which is less evident in the spatial distribution alone.  TKM06 also use the analysis of stellar ages in the $0.4 - 6 \\msun$ mass range by \\citet{pall99} to argue for a large age spread in the ONC. Although they note that \\citet{hart03} pointed out problems of contamination by non-members, TKM06 do not take sufficient notice of the mass-dependence of isochrones. As was clear from the original work by \\citet{hill97}, and is evident in Figure 4 of \\citet{pall00}, standard isochrones imply that stars of 1 to 5 $\\msun$ in the ONC have systematically older ages by one to a few Myr than both the lower mass stars and the higher mass stars. This effect is seen in virtually every star-forming region, at least comparing the intermediate mass stars to the lower mass $< 1 \\msun$ stars (see HR diagrams in \\citet{pall00}).  In other words, the ``tail'' of older stars in the distribution of stellar ages found by Palla \\& Stahler (1999, 2000) is strongly populated by a {\\em different} mass range (intermediate mass stars) than the peak of the age distribution (low mass stars).  As pointed out by \\citet{hart03}, if one takes these age determinations at face value, it means that the ONC existed for a few Myr as a cluster forming mostly intermediate mass stars, with very few low mass stars (or high-mass stars). In other words, the ONC (and other regions) begin their first few Myr of existence with extremely non-standard stellar initial mass functions (IMFs).   There is no observational evidence for any young clusters dominated by intermediate-mass stars by number. Therefore, there must be something wrong with the isochrones.  Pre-main sequence ages are contraction ages; one must know the starting radius of the star to determine the age.  TKM06 argue that corrections for starting radii (the ``birthline'' position) are only important for stars with ages $< 1$~Myr. However, \\citet{hart03} showed that this is not necessarily true for intermediate-mass stars; changing the accretion rates at which stars form over plausible ranges can have very big effects on the birthline for $\\gtrsim 1 \\msun$ stars.  Thus, the apparent ages of the intermediate-mass stars can be systematically overestimated due to incorrect birthline assumptions; this allows the ONC to maintain a roughly typical IMF over its formation. With this correction, the fraction of stars older than 2 Myr decreases dramatically."
    },
    "0711/0711.1164_arXiv.txt": {
        "abstract": "The most commonly accepted explanation for the origin of hypervelocity stars in the halo of the Milky Way is that they are the result of tidal disruption of binaries by the massive black hole at the center of the Galaxy. We show that, if this scenario is correct, and if the original binary properties are similar to those in the local stellar neighbourhood, then the hypervelocity stars should rotate with velocities measureably lower than those for field stars of similar spectral type. This may prove to be a more direct test of the model than trying to predict the position and velocity distributions. ",
        "introduction": "The discovery of a population of young, high velocity stars in the halo of the Galaxy (Brown et al. 2005), termed hypervelocity stars, is quite naturally explained in terms of ejection from the region around the massive black hole at the center (e.g. Hills 1988). The original proposal by Hills was that stellar binaries on grazing orbits are tidally disrupted by the black hole, leaving one star bound and one unbound with a high ejection velocity. The connection is further enhanced by the fact that the hypervelocity stars are late-B type stars and thus of similar spectral type to the population of peculiarly young stars present in the central parsec (Ghez et al. 2003, Eisenhauer et al. 2005), although this is partly the result of sample selection (Brown et al. 2006).  Examining the dynamical link between the hypervelocity stars and the so-called Sag~A$^*$ cluster may help to explain the puzzling origins of this latter population. The modern realisation of Hills' scenario invokes the scattering (Perets, Hopman \\& Alexander 2007) of young stellar binaries onto nearly radial orbits by massive perturbers (usually giant molecular clouds) from Galactocentric radii $\\sim $ several parsec (where such stars can form more naturally than in the strong tidal field of the central black hole). The dynamical disruption of these binaries then produces the observed hypervelocity stars and a bound population that results in the Sag~A$^*$ cluster (Gould \\& Quillen 2003, Ginsburg \\& Loeb 2006). However, some issues remain. The apastra of the cluster stars do not extend as far out as the Hills scenario would suggest and the scenario does not naturally explain the second population of young stars in the Galactic center, the so-called He-emission line stars (e.g. Paumard et al. 2006), as these stars are too far from the black hole to be the result of tidal disruption.  At least two alternative scenarios exist for the origin of the hypervelocity stars. They may also be ejected through two-body encounters with massive, dark objects orbiting the central black hole (Yu \\& Tremaine 2003). Indeed, the existence of a second, intermediate mass, black hole ($\\sim 10^3$--$10^4 M_{\\odot}$) in the Galactic center has been suggested (Hansen \\& Milosavljevic 2003) as a mechanism for explaining the presence of the young stars as the tidally stripped remnants of an inspiralling stellar cluster (Gerhard et al. 2001). In this scenario there is no direct one-to-one link between hypervelocity stars and Sag~A$^*$ cluster stars, although both are assumed to be drawn from the same parent population.  In a related fashion, stars may be ejected by two-body encounters with stellar mass black holes $\\sim 10 M_{\\odot}$ (O'Leary \\& Loeb 2006), a population of which is expected to form a cusp around the central black hole as a result of mass segregation in the Galactic potential (Morris 1993; Miralda-Escude' \\& Gould 2003). The efficiency of this process is lower than scattering by an IMBH, and  more subject to stellar destruction by tidal disruption or physical collisions (Ginsburg \\& Loeb 2007). However, it may still be significant and may be applicable if the observed young stars had their origins in a gravitationally unstable accretion disk (Levin \\& Beloborodov 2003) rather than a tidally disrupted cluster.  The original scenario due to Hill requires a binary, whereas the scenarios in which stars scatter off black holes do not. Furthermore, the orbital separations of the binaries that give rise to hypervelocity stars are small enough (Hills 1988; Gualandris, Portegies Zwart \\& Sipior 2005; Bromley et al. 2006) that the stellar rotation of the star is affected by the rotation of hypervelocity stars to test whether they did indeed spend some part of their lifetime in a binary. If the rotation velocities of hypervelocity stars are similar to those of binary stars of the same spectral type, and measureably different to equivalent single stars, then we may interpret this as evidence in favour of the binary disruption origin of hypervelocity stars.  In \\S~\\ref{Bin} we examine the empirical trends of stellar rotation with binary separation, for early-type stars, and how this translates into a relation between stellar rotation and ejection velocity, via the auspices of  Hills' scenario. Observational confirmation would provide a strong argument in favour of this model, as the alternative models do not require binary stars and thus do not imply the same consequence. ",
        "conclusions": "The currently most widely accepted model for the origin of hypervelocity stars is that they correspond to the break-up of binaries that interact with the massive black hole at the Galactic center. We have shown that, based on the empirically determined properties of early-type stars in binaries, and the kinematic requirements of the ejection scenario,  one should expect the hypervelocity stars to rotate slower than the average for single stars of their type. This is the result of the fact that, in order for a star to attain the velocities observed, it must come from  a system whose initial orbital separation is small enough that the rotation is affected by the presence of a binary companion.  In quantitative terms, we have shown that the rotation velocities for the observed late-type B stars are expected to be $\\sim 80$km/s, compared with $\\sim 130$km/s one finds for most single B-type stars. If the observational selection can be pushed into the regime of the early-type A stars, the prediction gets even stronger, with expected rotation velocities $\\sim 40$km/s. Furthermore, it is worth noting that the discovery of hypervelocity Am stars would be a direct signature of initial binarity, since the peculiar spectral signatures of Am stars are the result of radiative levitation in atmospheres that are very quiescent due to their slow rotation. This is why Am stars are almost always found in binaries. (Abt 1965, 2000)  The preceding assertions assume that the properties of the initial binaries are accurately reflected by the studies of those in the solar neighbourhood. However, this is implicit in  Hills' scenario, in that it invokes stars which formed far enough from the black hole that there is no reason to expect the star formation properties to be any different (unlike many scenarios for the origins of the young stars in the central parsec, where the physics is likely quite different). Another underlying assumption that one might question is whether the binaries have had enough time to circularise before being disrupted (i.e. are the solar neighbourhood binaries, in some sense, older)? To answer this, we note Abt \\& Boonyarek's empirical result that the regime of quasi-synchronisation extends to orbital periods well beyond that expected by simple tidal theory. Thus, it is likely the result of processes at work during the formation of the binary, and which then imprint the rotational properties of the stars well before the binary is disrupted. Even in the case where the synchronisation is the result of classical tidal interaction, for early type stars the tidal forces are quite sensitive to the stellar structure, and most of the tidal synchronization takes place in the early stages of the star's evolution (Zahn 1977).  Thus, we consider it a robust prediction that hypervelocity stars resulting from Hills' scenario should be slowly rotating, and measureably different from the rotational properties of field stars of similar spectral type. If hypervelocity stars are instead found to be rotating like normal field B or A stars, then this would favour one of the other ejection scenarios, such as scattering by an intermediate mass black hole (Levin 2006, Baumgardt, Gualandris \\& Portegies Zwart 2006; Sesana, Haardt \\& Madau 2006). Such a measurement should be within reach of current generations of telescopes. Indeed, the spectra of Brown et al. have a velocity resolution of $\\sim 100$ km/s, so that only a modest increase of a factor of 2 in spectral resolution is required to start putting some real constraints on the ejection scenario. In fact, the hypervelocity star HE~0437-5439 (Edelmann et al. 2005), already has a measured rotation velocity of $V \\sin i = 54 \\pm 4$ km/s. This low value is nicely consistent with the expectations of the binary disruption scenario, although its possible origin in the LMC is a complicating factor.  In almost all cases, if one binary member is ejected, the other remains bound to the black hole in a highly eccentric orbit. Thus, the above considerations apply to these stars as well. Indeed, Hills' scenario has been invoked to explain the origin of the so-called Sag~A$*$ cluster. However, the best-studied member of this group, the star SO-2, does not agree well with the predicted properties from the Hills' scenario. The eccentricity is not large enough to be the result of a binary disruption, and it is rotating normally, in contrast to the above prediction. In principle, the presence of other (possibly dark) bodies in the same vicinity may lead to additional perturbations which might explain the present observed parameters. However, the short age of the star renders this explanation problematical. As always, we await further observation."
    },
    "0711/0711.1352_arXiv.txt": {
        "abstract": "It is usually assumed that dark matter is produced during the radiation dominated era. There is, however, no direct evidence for radiation domination prior to big-bang nucleosynthesis. Two non-standard thermal histories are considered. In one, the low-temperature-reheating scenario, radiation domination begins as late as $\\sim 1$ MeV, and is preceded by significant entropy generation. Thermal axion relic abundances are then suppressed, and cosmological limits to axions are loosened. For reheating temperatures $T_{\\rm rh}\\lsim 35~{\\rm MeV}$, the large-scale structure limit to the axion mass is lifted. The remaining constraint from the total density of matter is significantly relaxed. Constraints are also relaxed for higher reheating temperatures. In a kination scenario, a more modest change to  cosmological axion constraints is obtained. Future possible constraints to axions and low-temperature reheating from the helium abundance and next-generation large-scale-structure surveys are discussed. ",
        "introduction": "The Peccei-Quinn (PQ) solution to the strong-CP problem yields the axion, a dark-matter candidate \\cite{crewther,baluni,pq}. If the axion mass $m_{\\rm a}\\gsim 10^{-2}~{\\rm eV}$, axions will be produced thermally, with cosmological abundance \\begin{equation}\\Omega_{{\\rm a}}h^{2}=\\frac{m_{{\\rm a}}}{130~{\\rm eV}}\\left(\\frac{10}{g_{*_{\\rm S},\\rm F}}\\right)\\label{thermalfreeze},\\end{equation} where $g_{*_{\\rm S},{\\rm F}}$ is the effective number of relativistic degrees of freedom when axions freeze out \\cite{changchoi,notinvisible,kt,murayama}. Axions with masses in the $\\sim {\\rm eV}$ range would contribute to the total density in roughly equal proportion to baryons.  Axions in the $\\sim{\\rm eV}$ mass range are relativistic when they decouple at $T_{{\\rm F}}= 30-50~{\\rm MeV}$ \\cite{changchoi}. Free streaming then erases density perturbations, suppressing the matter power spectrum on scales smaller than the axion free-streaming length \\cite{hu1,hu2,kt,haneutrino5}. Light axions would also contribute to the early integrated Sachs-Wolfe (ISW) effect \\cite{isw}. Data from large-scale structure (LSS) surveys and cosmic microwave-background (CMB) observations have been used to impose the constraint $m_{\\rm a}\\lsim1~{\\rm eV}$ to light hadronic axions \\cite{hanraffelt,raffeltwong,slosar}. These arguments apply to any particle relativistic at matter-radiation equality or cosmic microwave background (CMB) decoupling, thus imposing the similar constraint $\\sum_{i} m_{\\nu,i}\\lsim 1~{\\rm eV}$ to the sum of neutrino masses \\cite{crotty,tegmarksdss,barger,seljakneutrino,haneutrino5,fogli2,kristiansen,fukugitawmap3,spergelwmap3,elgaroy,pierpaoli}.  These constraints rely on abundances computed assuming that radiation domination began earlier than the chemical freeze-out of light relics. There is, however, no direct evidence for radiation domination prior to big-bang nucleosynthesis (BBN) \\cite{kamion}. The transition to radiation domination may be more gradual than typically assumed. In such a modified thermal history, two effects may cause relic abundances to change. First, the Hubble expansion rate scales differently with temperature $T$ until radiation domination begins, leading to a different freeze-out temperature. Second, entropy may be generated, suppressing relic abundances.  The universe could have reheated to a temperature as low as $1~{\\rm MeV}$, with standard radiation domination beginning thereafter \\cite{kawasaki,kawasaki2,ichikawa3,giudice2,hannestadreheat,kolbnotario}. This low-temperature reheating (LTR) scenario may be modeled simply through the entropy-generating decay of a massive particle $\\phi$ into radiation, with fixed rate $\\Gamma_{\\phi}$ and initial value $H_{\\rm I}$ of the Hubble parameter. The scalar $\\phi$ may be the inflaton, oscillating as inflation ends and decaying into standard-model particles, or it might be a secondary scalar, produced during preheating \\cite{dolgovlinde,afwise,kofman1,kofman3,shtanovreheat}. This decay softens the scaling of temperature $T$ with cosmological scale factor $a$, increasing the Hubble parameter $H(T)$ and leading to earlier freeze-out for certain relics. Entropy generation then highly suppresses these relic abundances.  Kination models offer another alternative to the standard thermal history, invoking a period of scalar-field kinetic-energy dominance \\cite{kination}, but no entropy production. Without entropy generation, abundances change more modestly.  Past work has shown that cosmological constraints to neutrinos, weakly interacting massive particles, and non-thermally produced axions are relaxed in LTR \\cite{giudice1,giudice2,yaguna}. Non-thermally produced axions ($m_{\\rm a}\\lsim 10^{-2}~{\\rm eV}$) would be produced through coherent oscillations of the PQ pseudoscalar \\cite{notinvisible, kt,torig}. In this paper, we obtain constraints to thermally-produced hadronic axions in the kination and LTR scenarios. While kination modestly loosens limits, LTR dramatically changes the cosmologically allowed range of axion masses.  We begin by reviewing these modified thermal histories and calculating axion relic abundances. We then generalize cosmological constraints to axions, allowing for low-temperature reheating and kination. For reheating temperatures $T_{\\rm rh} \\lsim 35~{\\rm MeV}$, LSS/CMB limits to the axion mass are lifted; constraints are also relaxed for higher $T_{\\rm rh}$. Constraints from the total matter density are also relaxed, but not completely lifted. For $T_{\\rm rh}\\simeq 10~{\\rm MeV}$, the new constraint is $m_{\\rm a}\\lsim 1.4~{\\rm keV}$, while for $T_{\\rm rh}\\simeq 35~{\\rm MeV}$, we find that $m_{\\rm a}\\lsim 43~{\\rm eV}$. If $T_{\\rm rh}\\gsim170~{\\rm MeV}$, standard results are recovered. After estimating the ability of future large-scale-structure surveys to further constrain axion masses for a variety of reheating temperatures, we derive modestly relaxed constraints to axions in the kination scenario. We conclude by considering future possible constraints to the relativistic energy density of axions in a low-temperature reheating model. ",
        "conclusions": "The lack of direct evidence for radiation domination at temperatures hotter than $1~{\\rm MeV}$ has motivated the introduction of kination, low-temperature reheating, and other scenarios for an altered pre-BBN expansion history. In the case of kination, the change in axion abundances and thus cosmological constraints is modest. Low-temperature reheating will suppress the abundance of thermally-produced hadronic axions, once the reheating temperature $T_{\\rm rh}\\sim 50~{\\rm MeV}$. This is rather intuitive once we recall that the axion freeze-out temperature in a radiation dominated cosmology is $\\sim 50~{\\rm MeV}$. If the reheating temperature crosses this threshold, axion densities are severely reduced by dramatic entropy production during reheating.  Total density, large-scale structure, and microwave background constraints to axions are all severely loosened as a result, possibly pushing the the axion mass window to very high values; for $T_{\\rm rh}\\simeq 10~{\\rm MeV}$, the new constraint is $m_{\\rm a}<1.4~{\\rm keV}$. For $T_{\\rm rh}\\gsim 170~{\\rm MeV}$, standard radiation dominated results are recovered. The inclusion of information on smaller scales will probe higher axion masses and lower reheat temperatures. More precise measurements of the matter power spectrum on all scales will probe lower axion masses. Kination also relaxes constraints to axions, though much less markedly. Future probes of primordial helium abundance will either lead to further constraints on axion properties, or, if axions are directly detected, provide a new view into the thermal history of the universe during the epoch $10~{\\rm MeV}\\lsim T\\lsim 170~{\\rm MeV}$."
    },
    "0711/0711.4431_arXiv.txt": {
        "abstract": "We use a large sample of galaxy groups selected from the SDSS data release  4 with an  adaptive halo-based group  finder to probe  how the clustering  strength of  groups  depends  on their  masses  and colors.   In particular, we determine the relative  biases of groups of different masses, as  well as that  of groups  with the  same mass  but with  different colors (either  that  of the  central  galaxy,  or the  total  color  of all  group members).  In  agreement with  previous studies, we  find that  more massive groups are more strongly clustered,  and the inferred mass dependence of the halo  bias  is in  good  agreement  with  predictions for  the  $\\Lambda$CDM concordance cosmology.  Regarding the  color dependence, we find that groups with red centrals  are more strongly clustered than groups  of the same mass but with  blue centrals.  Similar results  are obtained when the  color of a group is  defined to be the total  color of its member  galaxies.  The color dependence   is  more  prominent   in  less   massive  groups   and  becomes insignificant in groups with masses $\\gta 10^{14}\\msunh$.  These results are consistent with those obtained by Yang \\etal from an analysis of the 2dFGRS, but inconsistent with those obtained by Berlind {\\etal}, who also used an SDSS group  catalogue.  We construct  a mock  galaxy redshift  survey from the large Millennium $N$-body simulation that is populated with galaxies according to  the semi-analytical model  of Croton \\etal Applying  our group finder to this mock survey, and  analyzing the mock data in exactly the same way as the  true data, we are able to accurately  recover the intrinsic mass and color  dependencies of the halo  bias in the  model.  Interestingly, the semi-analytical model reveals the same  color dependence of the halo bias as we  find in  our  group  catalogue.  In  halos  with $M\\sim  10^{12}\\msunh$, though, the strength  of the color dependence is much  stronger in the model than  in  the  data. We discuss these results in  light of  the assembly bias  of dark  matter halos and  the star formation histories of galaxies. ",
        "introduction": "\\label{sec:intro}   In  the  standard  cold   dark  matter  (CDM)  paradigm  of  structure formation,  virialized CDM  halos are  considered to  be  the building blocks of  the mass  distribution in the  Universe. The  properties of dark matter halos, as well as their formation histories and clustering properties,  have been studied  in great  detail using  both numerical simulations  and analytical  approaches such  as the  (extended) Press Schechter (1974) formalism. These studies have shown that halo bias is mass dependent, in that more massive halos are more strongly clustered (e.g., Mo  \\& White 1996;  Sheth \\& Tormen  1999; Sheth, Mo  \\& Tormen 2001; Seljak \\& Warren 2004;  Tinker \\etal 2005). This mass dependence of the halo bias has played a crucial role in our understanding of the correlation  function  of  both  dark  matter and  galaxies,  via  the halo model (e.g., Cooray \\& Sheth 2002), the halo occupation model (e.g., Jing, Mo \\&  B\\\"orner 1998; Seljak 2000; Peacock \\& Smith 2000;  Berlind  \\&  Weinberg  2002; Magliocchetti  \\&  Porciani  2003; Berlind \\etal 2003; Zheng  \\etal 2005), and the conditional luminosity function (CLF)  model (e.g., Yang, Mo  \\& van den Bosch  2003; van den Bosch, Yang  \\& Mo 2003; Vale  \\& Ostriker 2006; Cooray  2006; van den Bosch \\etal 2007).  Observationally, the clustering strength of galaxy systems (groups and clusters) is  found to increase  with their mean separation  $d \\equiv n^{-1/3}$, with  $n$ the number  density of objects (e.g.   Bahcall \\& West  1992; Zandivarez  \\etal  2003; Padilla  \\etal  2004; Yang  \\etal 2005b). Assuming that richer  groups are statistically associated with more massive halos, one can  estimate the halo masses of galaxy groups using a theoretical  halo mass function, such as the  one given by the current $\\Lambda$CDM model (see Yang \\etal 2005b; Berlind \\etal 2006a; Yang  \\etal 2007).   Thus  the relation  between  mean separation  and clustering strength can be converted into a relation between halo mass and halo bias.  As shown in Yang \\etal (2005b), the mass dependence of the  halo  bias  obtained this  way  from  galaxy  groups is  in  good agreement with the prediction of the current $\\Lambda$CDM model.  In addition to the mass dependence, recent simulations have shown that halo  bias  also depends  on  the  assembly  history.  Using  $N$-body simulations,  Sheth   \\&  Tormen  (2004)   studied  the  environmental dependence of halo assembly and  found that dark matter halos in dense environments assemble at slightly earlier times than halos of the same mass  in low density  regions.  Gao  \\etal (2005)  used a  very large, high-resolution  numerical  simulation  of  structure formation  in  a $\\Lambda$CDM  cosmology to reexamine  the mass  and age  dependence of halo bias.  They  found that, for halos of $M  \\leq 10^{13} \\msunh$ at redshift $z=0$, the bias factor depends not only on halo mass but also on the halo assembly time. This  result has been confirmed by a number of subsequent investigations (Harker  \\etal 2006; Wechsler \\etal 2006; Wetzel \\etal  2006; Gao \\&  White 2007; Jing,  Suto \\& Mo  2007).  The origin of this assembly bias has recently been investigated in a number of papers (Wang, Mo \\& Jing 2007; Keselman \\& Nusser 2007; Hahn \\etal 2007).  If the  properties of galaxies that reside  in haloes of  a given mass correlate with  the assembly history  of that halo, the  assembly bias will impact the galaxy bias, $b_{\\rm gal}$, to the extent that $b_{\\rm gal}$ not only depends on halo mass, as generally assumed in HOD and CLF modeling, but also on the formation history of that halo.  Using a very crude method  for  assigning galaxy  luminosities to dark  matter haloes,    Reed    \\etal  (2007)    conclude,    however,   that   the luminosity-weighted  galaxy ages  cannot   closely trace the  assembly epoch of  their    dark   matter  hosts,  otherwise     the   observed color-dependence of  the clustering  can  not be reproduced.  Using  a galaxy formation model grafted    on  to the   Millennium  Simulation, Croton, Gao  \\&  White  (2007) compared   the original   simulation to `shuffled' versions  in   which the galaxy  populations   are randomly swapped among  haloes of similar  mass. They  found that  the shuffled versions have   correlation functions  that  are  different  from  the original ones at  the 5 to   10 percent level, indicating  that galaxy properties are (weakly)   correlated with the halo assembly   history. Unfortunately, since their    semi-analytical model is not a   perfect representation of the true  universe  (see e.g., Weinmann \\etal  2006b and Baldry \\etal 2006), it is  unclear whether such a correlation also exists in reality.  To test this, various authors  have examined whether there is any sign of assembly bias  in the large scale structure  of the observed galaxy distribution.   Using a  galaxy group  catalogue constructed  from the 2-degree Field  Galaxy Redshift  Survey (2dFGRS; Colless  \\etal 2001), Yang  \\etal (2006)  examined  how the  clustering  strength of  galaxy groups depends on the star formation activity of the central galaxies. At fixed mass  they found that the bias of  galaxy groups decreases as the SFR of the central galaxy increases. Assuming that galaxies with a more active SFR are bluer, this implies that groups with a red central are more  strongly clustered than groups  of the same mass  but with a blue central.  Surprisingly, a  completely opposite trend was found by Berlind  \\etal   (2006a) using  the  group  catalogue constructed  by Berlind \\etal (2006b) from the Sloan Digital Sky Survey (SDSS; York et al.  2000).   They found that  massive groups with bluer  centrals are {\\it more}  strongly biased  on large scales  than groups of  the same mass  but  with  redder   centrals.   Tinker  \\etal  (2007)  used  the (projected) galaxy correlation function  and galaxy void statistics to test  whether   the  galaxy  content   of  halos  of  fixed   mass  is systematically different in low density environments.  Using data from the SDSS combined  with HOD models, they find  that the luminosity and color of  field galaxies are  determined predominantly by the  mass of the halo in which they reside and have little direct dependence on the environment in which the host halo formed. Since, as shown by Sheth \\& Tormen (2004), the  assembly history of a halo  is correlated with its environment,  this  therefore implies  that  there  is no  significant correlation between the color of  a galaxy and the assembly history of its dark matter halo.   Finally, Blanton \\& Berlind (2006)compared the actual  relationship between  galaxy  colors and  large scale  density field to that predicted by  the null hypothesis that galaxy colors are only dependent on the mass of  the halo in which they reside.  To this extent, they use a similar shuffling technique as Croton \\etal (2007), but using galaxy groups in the SDSS rather than dark matter halos in a semi-analytical model  for galaxy formation.  They  find that shuffled and unshuffled color-density relations agree to better than 5 percent, indicating that the large scale environment of a group has only a very mild impact  on the properties  of its member galaxies.   Clearly, the current status as to whether galaxy properties are correlated with the halo assembly history is still very confusing.  In this paper, we re-examine the  dependence of the halo bias on group mass  and on the  colors of  member galaxies,  using the  galaxy group catalogue recently constructed by Yang  \\etal (2007) from the SDSS DR4 (Adelman-McCarthy  \\etal  2006).   An  important improvement  of  this catalogue over that of Yang \\etal (2006) and Weinmann \\etal (2006a) is that group masses are estimated both from the total luminosity and the total stellar mass of member  galaxies, allowing us to examine how the results may be affected by  the uncertainties in the estimate of group masses.  Furthermore, in order to test the impact of the uncertainties in the  group finder and in the  methods used to assign  masses to our groups,  we  use  a  mock  galaxy  redshift  survey  (hereafter  MGRS) constructed   from  the   Millennium  semi-analytical   galaxy  sample (hereafter SAM) of  Croton \\etal (2006) to test  how accurately we can recover the true trends in the semi-analytical model.  This paper is  organized as follows.  Section~\\ref{sec:data} describes the galaxy  sample and  group catalogue used  for our  analysis, while Section~\\ref{sec:method}  presents  our   method  for  estimating  the group-galaxy  cross-correlation function.   Our  observational results for the dependence of the group-galaxy cross-correlation on group mass and  on the  color  of the  central  galaxy are  presented in  Section \\ref{sec:obs}.  The same results. but for the semi-analytical MGRS are presented   in   Section    \\ref{sec:MGRS}.    Finally,   in   Section \\ref{sec:summary}, we summarize  and discuss our findings.  Throughout this  paper, we  adopt the  $\\Lambda$CDM `concordance'  cosmology with $\\Omega_m    =    0.24$,    $\\Omega_{\\Lambda}=0.76$,   $h=0.73$    and $\\sigma_8=0.77$.  Distances are quoted in units of $\\mpch$.  \\begin{deluxetable*}{lcccc} \\tablecaption{Volume-limited Samples of Galaxies\\label{tab1}} \\tablewidth{0pt} \\tablehead{Sample & z &$\\rmag$ & $N_{galaxy}$ (obs.) & $N_{galaxy}$ (mock) }  \\startdata V1 & (0.064, 0.127) & (-22.0, -20.5] &  54116 & 66484 \\\\ V2 & (0.064, 0.157) & (-22.0, -21.0] &  39101 & 51627 \\enddata  \\tablecomments{Column 1  indicates the volume-limited  galaxy sample ID.  Column 2  gives the redshift range of  each sample.  Column 3 is the absolute-magnitude range.  Columns  4 and 5 list the galaxy number in each sample  for observational and mock SDSS catalogues, respectively.} \\end{deluxetable*}   \\begin{deluxetable*}{lccccccccc} \\tabletypesize{\\scriptsize} \\tablecaption{Galaxy Groups in the Observation \\label{tab2}} \\tablewidth{0pt} \\tablehead{ID & $\\log M_{\\rm h}$ & z & $N_{group}$ &\\multicolumn{3}{c}{C0 ($M_L$/$M_S$)} &&\\multicolumn{2}{c}{C1 ($M_L$/$M_S$)} \\\\ \\cline{5-7}\\cline{9-10}& & & ($M_L$/$M_S$) & Red & Green & Blue && Red & Blue \\\\ (1)&(2)&(3)&(4)&(5)&(6)&(7)&&(8)&(9)}  \\startdata G1& $(12.0,12.5]$&(0.064,0.127)& 43233/41314& 0.99/0.99 &0.87/0.88 &0.60/0.66 &&0.94/0.96 &0.65/0.77\\\\ G2& $(12.5,13.0]$&(0.064,0.127)& 15614/15450& 1.01/1.01 &0.91/0.92 &0.67/0.72 &&0.96/0.99 &0.73/0.84\\\\ G3& $(13.0,13.5]$&(0.064,0.157)& 9909/10104& 1.03/1.03 &0.94/0.95 &0.71/0.75 &&0.99/1.01 &0.80/0.88 \\\\ G4& $(13.5,14.0]$&(0.064,0.157)& 2976/2826& 1.04/1.04 &0.96/0.97 &0.73/0.75 &&0.99/1.01  &0.86/0.89 \\enddata  \\tablecomments{Column  1 indicates  the  group sample  ID. Column  2 indicates the  mass range of each  group sample in  terms of $\\log [M_h/\\msunh]$.  Column 3 gives the  redshift range of each sample. Column 4 lists the group  number in the corresponding redshift and mass  ranges, where $M_L$  and $M_S$  mean group  masses estimated from the  ranking of group luminosity $L_{19.5}$  and stellar mass $M_{\\rm  stellar}$, respectively (Yang  \\etal 2007).   Columns 5-7 list the means  of the central galaxy colors  (C0) for groups with red, green  (galaxies between red  and blue Gaussian  peaks), blue central galaxies, respectively.  Columns 8 and 9 list the means of the total  group colors (C1)  for groups which are  separated into red and blue subsamples in similar numbers.  See the text for more details. } \\end{deluxetable*}   \\begin{deluxetable*}{lccccccccc} \\tabletypesize{\\scriptsize} \\tablecaption{Galaxy Groups in the MGRS \\label{tab3}} \\tablewidth{0pt} \\tablehead{ID & $\\log M_{\\rm h}$ & z & $N_{group}$ &\\multicolumn{3}{c}{C0 ($M_L$/$M_S$)} &&\\multicolumn{2}{c}{C1 ($M_L$/$M_S$)}\\\\ \\cline{5-7}\\cline{9-10}& & & ($M_L$/$M_S$) & Red & Green & Blue && Red & Blue\\\\ (1)&(2)&(3)&(4)&(5)&(6)&(7)&&(8)&(9)} \\startdata S1& $(12.0,12.5]$& (0.064,0.127) & 46274/44761& 0.98/0.99 &0.86/0.86 &0.53/0.58 &&0.89/0.90 &0.54/0.61 \\\\ S2& $(12.5,13.0]$& (0.064,0.127) & 16929/16538& 0.98/0.98 &0.87/0.89 &0.57/0.61 &&0.90/0.93 &0.59/0.68 \\\\ S3& $(13.0,13.5]$& (0.064,0.157) & 11139/11231& 0.97/0.97 &0.89/0.91 &0.55/0.60 &&0.92/0.94 &0.61/0.69 \\\\ S4& $(13.5,14.0]$& (0.064,0.157) & 3311/3453& 0.97/0.97 &0.92/0.93 &0.47/0.63 &&0.93/0.94 &0.53/0.68 \\enddata  \\tablecomments{Similar to  Table~\\ref{tab2}, but for  galaxy groups in the semi-analytical Mock Galaxy Redshift Survey (MGRS).} \\end{deluxetable*}   \\begin{figure*} \\plotone{f1.eps}   \\caption{The  color-magnitude   distributions  of galaxies in  the SDSS (left-hand  panel) and the  MGRS (right-hand panel). The dotted- and dashed-  lines indicate the centers of the two Gaussian used  to describe the color distributions  at a given luminosity bin (see text for details). These are used to split the galaxy populations  in red (above  the red sequence),  blue (below the  blue  sequence)  and  green  (in between  the  red  and  blue sequences). For clarity,  only 10\\% of all galaxies in the SDSS and MGRS are shown. } \\label{col:gau} \\end{figure*} ",
        "conclusions": "\\label{sec:summary}  In this paper we have used a large galaxy group catalogue constructed from the SDSS Data Release 4 (DR4) to study how the clustering of galaxy groups depends on their masses and, at fixed mass, on the color of their member galaxies.  We use  the projected  and redshift-space  two-point cross  correlation functions between  groups and  galaxies to  measure the  relative biases  for  groups of different masses and colors.  Our main results can be summarized as follows: \\begin{enumerate} \\item The  correlation amplitude  of galaxy groups  depends strongly  on their masses, and  is in good  agreement with the  mass dependence of the  bias of dark matter halos in the $\\Lambda$CDM concordance cosmology. \\item For  a given mass, the  correlation amplitude of groups  also depends on the  characteristic color  of  its  member galaxies.   Red  groups are  more strongly  clustered  than  blue  groups.   This  color  dependence  is  more prominent  in low-mass  groups,  and becomes  insignificant  in groups  with masses  $\\gta  10^{14}\\msunh$.   These  results  are  in  good,  qualitative agreement  with those obtained  by Yang  \\etal (2005b)  from an  analysis of galaxy  groups in  the 2dFGRS,  but disagree  with the  results  obtained by Berlind \\etal (2006a). \\item The observed color dependence in the data is qualitatively reproduced by the semi-analytical  model of Croton  \\etal (2006), but the  predicted color dependence  is much  too  strong for  low-mass  halos with  $M \\sim  10^{12} h^{-1}\\Msun$. \\item The systematic errors that can be induced by the group finder or by the method of assigning masses to the groups are discussed. And we conclude that our finding of the color dependence of the relative bias in the SDSS observation is robust. \\end{enumerate} Unlike   the  mass   dependence  of   the  group   clustering  which   can  be straightforwardly  linked  to  the  mass  dependence  of  the  halo  bias,  an interpretation of the color dependence  at fixed mass is less straightforward. It  is tempting  to  link  it to  the  assembly bias  that  has recently  been discovered in  numerical simulations, and which  shows that haloes  of a given mass that assemble earlier are more strongly clustered. The assembly time of a dark matter halo, $t_{\\rm main}$, is typically defined as the lookback time at which its  main progenitor reaches a mass  that is half of  the halo's present day mass.  A straightforward, but naive, interpretation  of our results therefore is that halos  that assemble  earlier contain  redder  galaxies. This  would not  only explain the sign of the color-dependence at fixed halo mass detected here, but since the assembly bias becomes less significant for more massive halos (e.g., Gao  \\etal  2005;  Jing  \\etal  2007),  it also  naturally  explains  why  the color-dependence  is weaker  for more  massive groups.   One way  to  invoke a positive,  causal correlation between  the assembly  histories of  dark matter haloes  and the star  formation histories  of the  associated galaxies,  is to assume that the  epoch of the last major merger,  which is strongly correlated with the assembly  time (Li \\etal 2007b), also signals the  time at which star formation  is terminated.  This  may come  about if  the feedback  from merger induced  AGN activity expels  (or heats)  the gas  associated with  the merger remnant,  thus shutting  off the  star  formation activity  (e.g., Di  Matteo, Springel \\& Hernquist 2005; Hopkins  \\etal 2006, 2007).  Such a merger induced truncation  mechanism would ensure  that halos  that assemble  earlier contain redder galaxies.  Alternatively,  the   observed  color  dependence   may  be  related   to  the `archeological' downsizing.  It is well known that,  in hierarchical structure formation, more  massive halos assemble later  (e.g., Lacey \\&  Cole 1993; van den  Bosch 2002;  Wechsler  \\etal  2002).  Therefore,  if  indeed haloes  that assemble earlier  contain redder  galaxies, it would  imply that  more massive haloes (i.e., clusters) contain bluer  galaxies than low mass haloes.  This is in clear  conflict with observations,  which show that more  massive galaxies, which typically  reside in  more massive halos,  on average have  formed their stars earlier.   An attractive solution to  this problem has  been proposed by Neistein \\etal  (2006), who introduced  an alternative `halo  formation time', the  time at  which  the  sum over  the  masses of  {\\it  all} the  virialized progenitors with masses  above a given minimum mass,  reaches half the present day halo  mass.  As shown by  Neistein \\etal (2006), contrary  to the assembly time,  $t_{\\rm  main}$,  defined  as  the  accretion time  of  the  last  main progenitor, $t_{\\rm all}$  {\\it increases} with increasing halo  mass, so that it provides  a natural explanation  for the observed  archeological downsizing that galaxies  in more massive haloes  have older stellar  populations. If the star formation history of a galaxy  is correlated with this formation time, it may also  explain the color-bias  at fixed halo  mass. The reason is  that, as demonstrated in  Neistein \\etal (2006),  although the formation  time, $t_{\\rm all}$, and  the assembly  time, $t_{\\rm main}$,  of halos  are anti-correlated when considering the entire halo  population, at fixed halo mass $t_{\\rm all}$ and $t_{\\rm main}$ are positively correlated.  Thus, at fixed halo mass, a halo that  forms  earlier will  contain  a redder  galaxy  and  will also  assemble earlier."
    },
    "0711/0711.1813_arXiv.txt": {
        "abstract": "For purposes of designing targeted cataclysmic variable (CV) detection surveys and interpreting results of other projects with many CV detections such as the ChaMPlane Survey, we have created a model of the CV distribution in the Galaxy.  It is modeled as a warped, flared exponential disk with a gaussian vertical distribution.  Extinction is based on a detailed Galactic dust and gas model. A luminosity function for CVs is also incorporated, based on a smoothed version of published data.  We calculate predicted field detection rates as a function of the limiting magnitude expected for the detecting system (i.e. WIYN/Hydra or NOAO 4m/Mosaic). Monte-Carlo techniques are used to assess statistical fluctuations in these rates.  We have created maps of the expected CV distribution for the full non-bulge Galactic plane ($20^\\circ<l<340^\\circ$, $|b|<15^\\circ$) for use in both the ChaMPlane Survey and future CV surveys. Assuming a CV distribution with a scale height of 160 pc, the ChaMPlane observational result of 5 CVs in 13 northern fields is best fit by a CV local space density of $0.9^{+1.5}_{-0.5}\\times10^{-5}\\rm{pc}^{-3}$, with the range representing the 95\\% confidence interval. ",
        "introduction": "The ChaMPlane Survey \\citep{grin05} has as a primary goal the determination of the Galactic cataclysmic variable star (CV) distribution to check if CVs can provide a significant contribution to the diffuse X-ray background of the Galaxy. In the survey design, the local space density of $10^{-5}$ CVs per cubic parsec \\citep{patt98, warn95} was used to predict detection rates.  Now, with followup WIYN\\footnote{The WIYN Observatory is a joint facility of the University of Wisconsin, Indiana University, Yale University, and the National Optical Astronomy Observatories.}/Hydra spectroscopy acquired for 13 fields, 11 of which are in the second and third Galactic quadrants \\citep{roge04}, we may begin to examine the accuracy of the published local CV space density. Also, we seek to make more precise predictions for expected detection rates along various lines of sight by using a more realistic model for the CV distribution within the Galaxy and a detailed model for the gas distribution to compute extinction along the lines of sight. ",
        "conclusions": "Based on the model of the disk CV distribution presented here, it appears that the ChaMPlane Survey CV detections reported in \\citet{roge04} for 13 northern WIYN-surveyed fields are consistent with a CV density of $0.6\\times10^{-5}\\rm{pc}^{-3}$ and the range of scale heights proposed by \\citet{patt84}.  This also supports the suggestion of \\citet{grin05} that earlier \\emph{ROSAT} and \\emph{Einstein} derived CV space densities of $3\\times10^{-5}\\rm{pc}^{-3}$ are too high by a factor of $\\sim3$. Completion of spectroscopic identification of potential targets in the ChaMPlane fields will result in much smaller viable ranges for the local space density and scale heights of CVs, narrowing the allowable parameter space.  The model will also allow future survey projects to selectively target potentially high-output fields to significantly boost the number of new CV detections, thus allowing for further investigations of both the spatial distribution of CVs and also the CV luminosity function."
    },
    "0711/0711.3033_arXiv.txt": {
        "abstract": "We present here the prospects for the GLAST Large Area Telescope (LAT) detection of the signature of the lightest Kaluza-Klein particle (LKP). It decays by direct annihilation into electron-positron pairs that may be detectable in the high energy electron flux. We discuss the LAT capability for detecting the high energy (20 GeV - ~1 TeV) cosmic ray electron flux and we analyze the LAT sensitivity to detect LKP-produced electrons for various particle masses. We include an analysis of the diffusive propagation of the electrons in the galaxy. ",
        "introduction": "The nature of dark matter and dark energy is one of the most exciting and critical problems in modern astrophysics. A number of theoretical models predict the existence of dark matter in different forms; so the experimental detection will be crucial. A large number of experiments are ongoing and planned  to detect dark matter, directly and indirectly, both at accelerators and in  space, where they search for dark matter signatures in cosmic radiation (see \\cite{Baltz} and references therein). In this paper we explore the capability of GLAST Large Area Telescope (LAT), scheduled for launch in the beginning of 2008, to detect the signature of dark matter in the high energy cosmic ray electron flux. We have previously demonstrated that LAT will be a  powerful detector of cosmic ray electrons, and will provide measurement of their flux with high statistical confidence \\cite{moiseevICRC}. We should mention that LAT is not designed to distinguish electrons from positrons, so we refer to their  sum as electrons for simplicity. LAT will detect  $\\approx 10^{7} $ electrons per year above 20 GeV with the energy resolution 5-20\\%. Such good statistics permits detection of spectral features, among which could be ones caused by  exotic sources such as Kaluza-Klein particles which manifest higher spatial dimensions. The possibility of the annihilation of the Lightest Kaluza-Klein Particles (LKP), which can be a stable and viable dark matter candidate, directly into electron-positron pairs, was investigated in, e.g. \\cite{Baltz} and \\cite{Cheng}. They estimated that electron-positron pairs are produced in approximately 20\\% of the annihilation cases, which makes the observations viable within the model assumptions.  There are some indications of spectral features in the electron spectrum observed by ATIC \\cite{ATIC} and PPB-BETS \\cite{BETS} around 300-500 GeV, as well as in the positron spectrum measured by HEAT \\cite{HEAT}, encouraging us for  measurements with the LAT. ",
        "conclusions": "We analyzed the capability of the LAT to detect high energy cosmic ray electrons and applied it to the model of LKP direct annihilation into electron-positron pairs. Using the estimates for this model as given in \\cite{Baltz} as an example to demonstrate the detection feasibility, we estimated that within this model, the LAT will be able to recognize the LKP-caused spectral edge in the electron spectrum up to a LKP mass of 600-700 GeV. The results obtained  can be applied to any dark matter model where  electrons are produced in order to estimate the LAT sensitivity. The important feature is that the dominating background in these measurements will be only the \"conventional\" electron flux.  We want to thank all LAT team members, and especially the members of LAT Dark Matter Science Working Group for their support and valuable suggestions. We are grateful to Robert Hartman and Jan Conrad for their comments on this paper."
    },
    "0711/0711.3343_arXiv.txt": {
        "abstract": "{} {A major goal of our survey is to significantly increase the number of detected extra-solar planets in a magnitude-limited sample to improve our knowledge of their orbital elements distributions and thus obtain better constraints for planet-formation models.} {Radial-velocity data were taken at Haute-Provence Observatory (OHP, France) with the ELODIE echelle spectrograph.}% {We report the presence of a planet orbiting {\\small HD}\\,196885\\,A, with an orbital period of \\perd~days. This star was previously suggested to host a $ 386 $-day planet, but we cannot confirm its existence. We also detect the presence of a stellar companion, {\\small HD}\\,196885\\,B, and give some constraints on its orbit. } {} ",
        "introduction": "The {\\it {\\small ELODIE} Planet Search Survey} was an extensive radial-velocity northern survey of dwarf stars at the Haute-Provence Observatory (OHP, France) using the {\\small ELODIE} high-precision fiber-fed echelle spectrograph \\citep{Baranne_etal_1996} mounted on the Cassegrain focus of the 1.93-m telescope. It started to operate at the end of 1993 and acquired data until the summer of 2006, when it was replaced by the new echelle spectrograph {\\small SOPHIE} \\citep{Bouchy_etal_2006}.   This survey is part of a large effort aiming at an extra-solar planet search through radial-velocity measurements, in order to characterize the zoo of exoplanets and to bring strong constraints on their processes of formation and evolution. {\\small ELODIE} was responsible for the findings of many extra-solar planets, among them the first hot Jupiter, 51\\,Peg\\,b \\citep{Mayor_Queloz_1995}. The original sample consisted of 142 stars, but a new sample of 330 stars was defined in 1997. Details about the program and the surveyed sample can be found in \\citet{Perrier_etal_2003}.  In this paper we present the detection of two bodies around the star {\\small HD}\\,196885\\,A, one in the mass range of planets and the other believed to be a stellar companion. This star was previously described to harbor a planet with a period of about 386~days (although no published reference is known), but we could not confirm its presence in our observations (our analysis shows a planet with an orbital period of about \\perd~days instead). The stellar companion was also recently observed using NACO adaptative optics by \\citet{Chauvin_etal_2006,Chauvin_etal_2007}. The stellar properties of {\\small HD}\\,196885\\,A are briefly recalled in Sect.\\,2 and the radial velocities with the two detected companions are described in Sect.\\,3. In Sect.\\,4 we discuss the possibility of existence of more companions, in particular around $ 368 $~days. Last section is devoted to our conclusions on this system. ",
        "conclusions": "In this paper we report the presence of a planet orbiting the {\\small HD}\\,196885\\,A star, with an orbital period of \\perd~days. We also detect the presence of the stellar companion {\\small HD}\\,196885\\,B, that was recently observed by \\citet{Chauvin_etal_2006,Chauvin_etal_2007}. Its orbit is not completely determined, but we are able to provide some constraints: $ P > 40 $~yr, $ a > 14 $\\,AU and $ M \\sin i > 0.28 \\, M_\\odot$. It is possible to find a large set of orbits with much longer orbital periods, that fit the observational data as good as the minimum mass orbits (Tab.\\ref{T3}). In particular, we find a solution at $ P = 120 $~yr with $ M \\sin i = 0.6\\,M_\\odot $, the mass estimated by \\citet{Chauvin_etal_2007}.  The {\\small HD}\\,196885\\,A system was previously described as a star orbited by a planet at 386~days by the California \\& Carnegie Planet Search Team, detected with data acquired with the Lick survey. Although the system is reported on their webpage and many other places, no publication describing the system is yet known. The data acquired with the {\\small ELODIE} survey also present some signal around $ 368 $~days, but a more detailed analysis shows that this signal corresponds to an aliasing of the observational data, probably due to the similarities with the Earth orbital period. We thus do not confirm the existence of the planet obtained with measurements form the Lick survey.  The {\\small ELODIE} radial velocity residuals after subtracting the signal from the two companions still presents a $O-C$ (Fig.\\ref{F1}) slightly above the precision of the instrument. This suggests that the system may still hide additional planetary companions. However, we were unable to find them, even when using a genetic algorithm. The excess in the residuals can also be a contamination from the  spectrum of the stellar companion {\\small HD}\\,196885\\,B. We have searched for its presence in our {\\small ELODIE} data using multi-order TODCOR, a two-dimensional cross-correlation algorithm \\citep{Zucker_etal_2004}, but did not find anything convincing. We have also searched for correlations between line bisectors and radial velocity or residuals. % No significant correlation could be detected. These negative results do not allow us to formally discard the blend scenario, but they render it an unlikely possibility.  The planet {\\small HD}\\,196885\\,Ab is the fourth to be discovered in a close binary with a separation smaller than 25\\,AU. The other cases already known are {\\small Gliese}\\,86 \\citep{Queloz_etal_2000,Lagrange_etal_2006}, $\\gamma$ Cep \\citep{Hatzes_etal_2003,Neuhaeuser_etal_2007} and {\\small HD}\\,41004\\,A \\citep{Zucker_etal_2004}. These systems are ideal for carrying out combined astrometric and radial-velocity observations to constrain the binary dynamic properties and the possible impact of a close binary companion on planet formation and evolution."
    },
    "0711/0711.1170_arXiv.txt": {
        "abstract": "Dusty hyperluminous galaxies in the early universe provide unique environments for studying the role of massive stars in the formation and destruction of dust. At redshifts above ~ 6, when the universe was less than ~ 1 Gyr old, dust could have only condensed in the explosive ejecta of Type~II supernovae (SNe), since most of the progenitors of the AGB stars, the major alternative source of interstellar dust, did not have time to evolve off the main sequence. We present analytical models for the evolution of the gas, dust, and metals in high redshift galaxies, with a special application to SDSS J1148+5251, a hyperluminous quasar at $z = 6.4$ hereafter referred to as \\jay. We show that an average SN must condense at least 1~\\msun\\ of dust to account for the mass of dust in this object, when grain destruction by supernova remnants (SNRs) is taken into account. This required yield is in excess of $\\sim 0.05$~\\msun, the largest mass of dust inferred from infrared observations of Cas~A. If the yield of Cas~A is typical, then other processes, such as accretion onto preexisting grains in molecular clouds are needed to produce the mass of dust in J1148+5251. For such process to be effective, SNR must significantly increase, presumably by non-evaporative grain-grain collisions during the late stages of their evolution, the number of nucleation centers onto which refractory elements can condense in molecular clouds. ",
        "introduction": "The detection of massive amounts of dust in hyperluminous infrared (IR) galaxies at redshifts $z > 6$ raises challenging questions about the sources capable of producing such large amount of dust during the relatively short lifetime of these galaxies (e.g. \\cite{ref:maiolino06,ref:beelen06}). For example, the galaxy SDSS J1148+5251 (hereafter \\jay) located at $z = 6.4$ was observed at far-IR and submillimeter wavelength \\cite{ref:bertoldi03,ref:robson04,ref:beelen06}. Figure \\ref{dustspec} depicts the observed fluxes and spectral fits of various dust emission models through the data. The average IR luminosity of the source is $L_{IR} \\sim 2\\times 10^{13}$~\\lsun, and the average dust mass is $M_d \\sim 2\\times 10^8$~\\msun. Using the Kennicutt relation \\cite{ref:kennicutt98}, one can derive a star formation rate (SFR) of $\\sim 3000$~\\msun~yr$^{-1}$ from the observed far-IR luminosity. For comparison, the Milky  Way galaxy is about 10~Gyr old, has an average SFR of $\\sim 3$~\\msun~yr$^{-1}$, and contains about $5\\times 10^7$~\\msun\\ of dust, a significant  fraction of which  was produced in AGB stars.  \\begin{figure}[htbp] \\begin{center} \\includegraphics[width=4.0in]{dustspec.eps} \\end{center} \\caption{Spectral fits of several possible dust compositions to the observed far-IR and submillimeter observations of \\jay. Wavelengths are in the observer's frame of reference. Dust masses range from $(1-5)\\times10^8$~\\msun\\ and luminosities range from $(2-3)\\times 10^{13}$~\\lsun, depending on dust composition. References to the observations can be found in \\cite{ref:dwek07}.} \\label{dustspec} \\end{figure}  At $z=6.4$ the universe was only 890~Myr old, using standard $\\Lambda$CDM parameters ($\\Omega_m = 0.27$, $\\Omega_{\\Lambda} = 0.73$, and $H_0 = 70$~km~s$^{-1}$~Mpc$^{-1}$). If \\jay\\ formed at $z = 10$ then the galaxy is only 400~Myr old. If the SFR had occured at  a constant rate over the lifetime of the galaxy, its initial mass should have been about 10$^{12}$~\\msun, which is significantly larger than the dynamical mass $M_{dyn} \\approx 5\\times 10^{10}$~\\msun\\ of the galaxy \\cite{ref:walter04}. The high observed SFR may therefore represent a recent burst of star formation that has lasted for only about 20~Myr. The galaxy \\jay\\ is therefore at most $\\sim 400$~Myr old, and probably significantly younger with an age of only $\\sim 20$~Myr. Adopting a current gas mass of $M_g = 3\\times 10^{10}$~\\msun\\ for this galaxy we get that the gas mass fraction at 400~Myr is about 0.60. The dust-to-gas mass ratio is given by $Z_d \\equiv M_d/M_g = 0.0067$.  A significant  fraction of the dust in the Milky Way was produced in AGB stars. However, these stars are not likely  to  contribute significantly to the formation of dust in very young galaxies, since the low mass stars ($M \\approx 2$~\\msun) that  produce most of the dust did not have time to evolve off the main sequence \\cite{ref:dwek98,ref:morgan03,ref:dwek07}. In contrast, core collapse supernovae (SNe) ($M > 8$~\\msun) inject their nucleosynthetic products back into the ISM shortly ($t < 20$~Myr) after their formation, resulting in the rapid enrichment of the interstellar medium (ISM) with the dust that formed in their explosive ejecta. But can SNe account for the large amount of dust seen in this object? The answer to this question is complicated by the fact that SNe are also the main source of grain destruction during the remnant phase of their evolution \\cite{ref:jones96,ref:jones04}. The problem can therefore only be addressed with a detailed model for the evolution of dust in these systems. ",
        "conclusions": "The early universe is a unique environment for studying the role of massive stars in the formation and destruction of dust. The equations describing their chemical evolution can be greatly simplified by using the instantaneous recycling approximation, and by neglecting the delayed contribution of low mass stars to the metal and dust abundance of the ISM. Neglecting any accretion of metals onto pre-existing dust in the interstellar medium, the evolution of the dust is then primarily determined by the condensation efficiency of refractory elements in the ejecta of Type~II supernovae, and the destruction efficiency of dust by SN blast waves.  We applied our general results to \\jay, a dusty, hyperluminous quasar at redshift $z = 6.4$ and found that the formation of a dust mass fraction of $Z_d = 0.0067$ in a galaxy with an ISM mass of $3\\times 10^{10}$~\\msun, requires an average SN to produce between 0.5 and 1~\\msun\\ of dust if there was no grain destruction (Fig. \\ref{snyield}). Such large amount of dust can be produced if if the condensation efficiency in SNe is about unity. Observationally, the required dust yield is in excess of the largest amount of dust ($\\sim 0.054$~\\msun) observed so far to have formed in a SN. This suggests that accretion in the ISM may play an important role in the growth of dust mass. For this process to be effective, SNR must significantly increase, presumably by non-evaporative grain-grain collisions during the late stages of their evolution, the number of nucleation centers onto which refractory elements can condense in molecular clouds."
    },
    "0711/0711.2804_arXiv.txt": {
        "abstract": "{The discoveries of the GZK cutoff with the HiRes and Auger Observatories and the discovery by Auger of clustering  of $\\gtrsim 60$ EeV ultra-high energy cosmic rays (UHECRs) towards nearby ($\\lesssim 75$ Mpc) AGNs along the supergalactic plane establishes the astrophysical origin of the UHECRs. The likely sources of the UHECRs are gamma-ray bursts and radio-loud AGNs because: (1) they are extragalactic; (2) they are sufficiently powerful; (3) acceleration to ultra-high energies can be achieved in their relativistic ejecta; (4) anomalous X-ray and $\\gamma$-ray features can be explained by nonthermal hadron acceleration in relativistic blast waves; and (5) sources reside within the GZK radius. Two arguments for acceleration to UHE are presented, and limits on UHECR ion acceleration are set. UHECR ions are shown to be able to survive without photodisintegrating while passing through the AGN scattered radiation field, even if launched deep in the broad line region.  UHECR injection throughout cosmic time fits the measured energy spectrum of UHECRs, at least for protons.  Local UHECR proton and ion interaction and energy-loss mean free paths are calculated using an empirical fit to the extragalactic background light (EBL) at IR and optical energies. Minimum intergalactic magnetic (IGM) fields $\\approx 10^{-11}$ G are derived from clustering assuming specific source origins, e.g., Cen A, nearby AGNs, or GRBs for the super-GZK CRs seen with Auger. Besides distinct cosmic-ray induced $\\gamma$-ray signatures that should be observed with the {\\it Gamma ray Large Area Space Telescope (GLAST)}, source and GZK neutrino detections and the arrival distribution of UHECR in direction and time can finally decide the sources of cosmic rays at the highest energies. }   \\email{charles.dermer@nrl.navy.mil}   \\begin{document} ",
        "introduction": "A high-significance steepening in the UHECR spectrum at energy $E\\cong 10^{19.6}$ eV was reported earlier this year by the HiRes collaboration \\cite{hires07}, and here at the 2007 M\\'erida ICRC based on observations taken with the Auger Observatory \\cite{yam07}.  This result confirms the prediction of the Greisen-Zatsepin-Kuzmin (GZK) cutoff \\cite{gre66,zk66} in the UHECR spectrum due to photohadronic interactions of UHECRs with photons of the cosmic microwave background radiation (CMBR), and favors astrophysical bottom-up vs. particle physics top-down scenarios for the UHECRs, provided that sources are found within the GZK radius.  At the same time, Auger data shows \\cite{ung07} evidence for mixed composition with substantial ion content in UHECRs with energies as high as a $few \\times 10^{19}$ eV, based on studies of the depth of shower maxima. With hybrid fluorescence detectors and shower counters, Auger provides the strongest evidence yet for metals in the UHECRs, possibly with mean atomic mass $\\langle A \\rangle \\sim 8$ -- 26, significantly different from pure proton and pure Fe composition. This result depends on the accuracy of the nuclear interaction physics used to model showers, but points to the importance of nuclei in the UHECRs, and the meaning of this for GZK physics \\cite{nw00,ta04,alo07}.   The GZK radius of a proton with energy $E= E_{par} = 10^{20}E_{20}$ eV can be estimated by noting that the product of the cross section and the inelasticity in a pion-producing reaction is $K_{\\phi\\pi}\\sigma_{\\phi\\pi} \\approx 70~ \\mu$b \\cite{ad01}, so that the photopion energy-loss pathlength $r_{\\phi\\pi}(E)$ is given by $n_{ph}(E)(K_{\\phi\\pi}\\sigma_{\\phi\\pi}) r_{\\phi\\pi}(E) = 1$, where $n_{ph}(E)$ is the $E$-dependent number density of photons above the pion-production threshold.  There are $\\approx 400$ CMBR photons per cm$^3$ at the present epoch, and ions with Lorentz factor $\\gamma=E/Am_pc^2$ satisfying $\\gamma(h\\nu/m_ec^2)\\gtrsim 2m_\\pi/m_e \\cong 500$ will interact with most of the photons of the CMBR. The mean dimensionless photon energy in the CMBR is $h\\langle \\nu\\rangle/m_ec^2 \\cong 2.70 k_{\\rm B}T_{CMBR}/m_ec^2\\cong 1.3\\times 10^{-9}$, noting that $T_{CMBR} = 2.72$ K at the present epoch; hence UHECR protons with Lorentz factor $\\gamma\\gtrsim 4\\times 10^{11}$, or $E\\gtrsim 4\\times 10^{20}$ eV will have a photo-pion energy-loss pathlength $r_{\\phi\\pi}(E\\gtrsim 4\\times 10^{20}$ eV) $\\approx 12$ Mpc.  By considering the number of photons above threshold in a blackbody distribution, the energy-loss mean free path of an UHECR proton with energy $E$ is found to be \\begin{equation} r_{\\phi\\pi}(E_{20}) \\cong {13.7 \\exp[4/E_{20}] \\over [1 + 4/E_{20}]}\\;{\\rm Mpc}\\;. \\label{eq1} \\end{equation} Figure 1 shows the photopion energy-loss pathlength from eq.\\ (\\ref{eq1}) for cosmic-ray protons interacting with CMBR photons, in comparison with a more detailed calculation \\cite{sta00}.  Auger results show a mixed UHECR composition extending from $\\approx 4.5\\times 10^{17}$ eV -- $4.5\\times 10^{19}$ eV, with an indication of increased average mass at the highest energy datum \\cite{ung07}. The ionic content of the UHECR provides new information to understand UHECR source properties, especially if the distribution of atomic mass $A$ can be obtained with better statistics.  Figure 1 shows expansion, photohadronic, and total energy-loss mean free paths (MFPs) for p and Fe on the CMBR. Calculations of photopair losses follow \\cite{czs92}. The photopion loss calculation can be extended to ions of atomic mass $A$ and charge $Z$ by assuming that the photopion cross-section times inelasticity is at most $= A^{2/3} K_{\\phi\\pi}\\sigma_{\\phi\\pi}$ above threshold. The photopion energy-loss rate of Fe with this assumption for the product of the inelasticity and cross section is shown in Figure 1. Because of the greater mass of Fe than p, far fewer photons are available for photopion interactions than with protons carrying the same energy, so that the corresponding GZK photopion radius is much larger.  %  Ions are in addition subject to losses due to photonuclear interactions that can break up nuclei (for example, via the giant dipole resonance).  The photodisintegration loss length for Fe in the CMBR at $z \\ll 1$, calculated from model fits to photohadronic cross section data \\cite{psb76,ss99}, is plotted in Figure 1.  UHECR Fe is seen to have a comparable GZK-type cutoff, but here at $\\approx 2\\times 10^{20}$ eV, with photopair losses playing a minor role.  A realistic calculation of UHECR ion spectra in the evolving background radiation fields must follow a reaction pathway using Monte Carlo techniques, with the actual effective energy-loss pathlength taking into account the EBL intensity contributed by stars and black holes, then sometimes reprocessed through dust and gas. The exact form of the EBL intensity between $\\approx 1$ and 100 $\\mu$ is poorly known, but can be constrained by empirical galaxy SEDs and $\\gamma$-ray observations.   \\begin{figure}[t] \\center \\includegraphics[scale=0.35]{icrc1160_fig01.eps} \\caption{ Energy-loss mean-free path for UHECR protons (black) and Fe ions (gray) in the CMBR at $z = 0$, as a function of total particle energy. The short-dashed black curve give the analytical approximation for the UHECR proton photopion energy-loss pathlength, eq.\\ (\\ref{eq1}), in comparison with the numerical results, shown by open circles \\cite{sta00}.  Numerical results for the photopair energy-loss pathlength of UHECR protons and Fe are shown by the dotted curves.  The gray dot-dashed curve gives the photodisintegration energy-loss curve of Fe. Expansion losses are shown by the long-dashed line, with $c/H_0 = 4170$ Mpc. } \\label{fig01} \\end{figure}   From the curves shown in Figure 1, we see that for either UHECR protons or Fe ions, the sources of $E \\gtrsim 10^{20}$ eV cosmic rays must be found within a few hundred Mpc, and the $3\\times 10^{20}$ eV cosmic ray detected by the Fly's Eye \\cite{bir95} must have originated from a source within a few tens of Mpc. Both gamma-ray bursts (GRBs) and radio-loud active galactic nuclei (AGNs\\footnote{Blazars are radio-loud AGNs viewed along the jet axis; here the two terms are used interchangeably.})  satisfy this requirement. Besides the availability of extragalactic sources within tens of Mpc, several other conditions must also be met, including \\begin{enumerate} \\item sufficient power to make the UHECRs at the observed intensity; \\item a plausible mechanism to accelerate particles to $\\gtrsim 10^{20}$ eV energies; and \\item a model that can reproduce the observed UHECR spectrum, consistent with astronomical properties of the sources. \\end{enumerate}  The Auger clustering observations \\cite{aug07} suggest new calculations to determine the intergalactic medium (IGM) magnetic field.  If the 2 (out of 27 total) UHECRs with arrival directions within 3$^\\circ$ of Cen A are hypothesized to originate from Cen A, then lower values of the mean IGM field  can be derived, $\\gtrsim 10^{-10}$ G, between us and Cen A, and $\\gtrsim 10^{-11}$ G for sources at 75 Mpc. The clustering observations and the GZK length are tightly coupled concepts \\cite{sta95}, as is established with the Auger data. Here we give new proton and ion energy loss calculations after revisiting the problem of the EBL.  The Auger discovery \\cite{aug07} that the arrival directions of $>56$ EeV are correlated with nearby, $\\lesssim 75 $ Mpc AGNs, in particular, 2 cosmic rays within $3^\\circ$ of Centaurus A, does not mean that AGNs are the sources of UHECRs, any more than O and B stars correlated with galactic $\\gamma$-ray sources mean that cosmic rays are accelerated by high-mass stars \\cite{tor03}. Long duration GRBs trace sites of active star formation, and may be correlated to AGN activity. AGNs trace the local matter distribution due to structure formation, as do radio galaxies and GRBs, or for that matter clusters of galaxies and radio-quiet Sy AGN.  Here we make the case for GRB and radio-loud AGN/blazars as the sources of the UHECRs. UHECR ion acceleration in clusters of galaxies could also make a contribution at $\\lesssim 10^{19}$ eV, or from dim quasars \\cite{bg99}, but this analysis forces us to reject UHECR origin from radio-quiet AGNs.   Effects of UHECR hadron acceleration can be varied, including anomalous $\\gamma$-ray emissions and characteristic behaviors of GRB X-ray light curves.  Detection of high-energy neutrinos from hadronic interactions at the sources will provide the most definitive evidence for the presence of UHECRs \\cite{hal07}. Detection of cosmogenic GZK neutrinos from the decay of pions formed in GZK interactions with CMBR photons gives a calorimetric measure of UHECR power throughout cosmic time, with a characteristic spectrum imprinted by production and propagation effects. ANITA and successor Askaryan-effect detectors should soon be able to constrain a long-duration GRB model of UHECRs, if they were protons.  Metals in the composition of UHECRs change the spectral predictions, as well as the GZK neutrino predictions \\cite{ave04,anc07a}. The result of UHECR ions vs. protons is mainly to enhance neutron $\\beta$-decay neutrinos but not significantly change the $\\sim$ EeV neutrino flux.  Joint analysis of UHECR spectral, composition, and directional information, now given most accurately with the hybrid Auger, with neutrino and $\\gamma$-ray data from GLAST, AGILE, and the ground-based air and water Cherenkov detectors, should lead to a definitive solution to the problem of UHECR origin. A black hole origin is explored here. Crucial to progress are the advances in detector capabilities and instrumentation, multiwavelength observations and multi-disciplinary analyses. ",
        "conclusions": "Auger has already contributed three major discoveries to cosmic-ray physics: \\begin{enumerate} \\item the GZK cutoff, also found with HiRes; \\item Mixed ionic composition in the UHECRs up to a $few\\times 10^{19}$ eV; and \\item statistical demonstration that the clustering of 27 UHECRs with  energies $\\gtrsim 6\\times 10^{19}$ eV  follow the matter distribution as traced by nearby ($\\lesssim 75$ Mpc) AGNs. \\end{enumerate}   These results establish without doubt that UHECRs originate from astrophysical sources. With this information, and based on past theoretical work, I argue that GRBs and radio-loud AGNs, which are classified as blazars when viewed on-axis, are the most probable sources of UHECRs, as can be demonstrated by hadronic $\\gamma$-ray signatures. The most convincing evidence would be direct detection of PeV neutrinos from UHECR sources with IceCube or KM3NET. Detection of GZK EeV neutrinos from photopion interactions of UHECRs as they propagate through EBL with an Askaryan telescope like ANITA will also importantly test astrophysical models.  In this paper, I used $\\gamma$-ray observations to constrain the spectrum of the EBL between 1 and 100 $\\mu$, where it is poorly known. By adopting as a general rule, consistent with observations, that blazar $\\gtrsim 100$ MeV -- TeV $\\gamma$-ray spectra are steeper than $-2$, a low EBL, shown in Figure 9, was favored. With this EBL, we can deconvolve the intrinsic source spectrum of low-redshift sources, and calculate the photodisintegration rate of different UHECR ions. The GZK curves for different ionic species obtained with the low EBL is shown in Figure 16. These results are in accord with the location of the GZK spectral cutoff measured with Auger and HiRes, and a mixed ionic composition, because protons are more difficult to accelerate, and Fe will be broken up.  GRBs and blazars are both viable sites for UHECR acceleration, consistent with the Auger results. What will be important is to associate UHECR arrival directions with a subset of AGNs or galaxies. Small metal-poor bluish star-forming galaxies are likely hosts of long-duration GRBs \\cite{flo03}. Radio galaxies would host misaligned blazars, and slightly misaligned blazars could be $\\gamma$-ray dim and cosmic-ray bright. GLAST will be able to see $\\gamma$-ray dim AGNs (the radio galaxies M87 \\cite{bei07}, Cen A \\cite{har99}, and NGC 6251 \\cite{muk02} are definite, likely, and probable $\\gamma$-ray sources, respectively). Associating radio galaxies or GLAST $\\gamma$-ray galaxies in the nearby universe with arrival directions of UHECRs could decide the question of blazar origin of the UHECRs. With the $\\gamma$-ray, cosmic ray, and neutrino observations, it is likely that the problem of  UHECR origin will soon be solved\\footnote{The theoretical basis for the results presented here are given in my book with Govind Menon entitled ``High-Energy Radiation from Black Holes: $\\gamma$-rays, cosmic rays, and neutrinos,\" to be published by Princeton University Press in 2008.}.      \\vskip0.2in"
    },
    "0711/0711.0801_arXiv.txt": {
        "abstract": "We present an analysis of the 0.95$-$14.5 $\\mu$m spectral energy distributions of nine field ultracool dwarfs with spectral types ranging from L1 to T4.5.  Effective temperatures, gravities, and condensate cloud sedimentation efficiencies are derived by comparing the data to synthetic spectra computed from atmospheric models that self-consistently include the formation of condensate clouds.  Overall the model spectra fit the data well, although the agreement at some wavelengths remains poor due to remaining inadequacies in the models. Derived effective temperatures decrease steadily through the L1 to T4.5 spectral types and we confirm that the effective temperatures of ultracool dwarfs at the L/T transition are nearly constant, decreasing by only $\\sim$200 K from spectral types L7.5 to T4.5.  The condensate cloud properties vary significantly among the L dwarfs in our sample, ranging from very thick clouds to relatively thin clouds with no particular trend with spectral type.  The two objects in our sample with very red $J-K_s$ colors are, however, best fitted with synthetic spectra that have thick clouds which hints at a possible correlation between the near-infrared colors of L dwarfs and the condensate cloud properties.  The fits to the two T dwarfs in our sample (T2 and T4.5) also suggest that the clouds become thinner in this spectral class, in agreement with previous studies.  Restricting the fits to narrower wavelength ranges (i.e., individual photometric bands) almost always yields excellent agreement between the data and models.  In some cases, the models even reproduce the detailed structure of the weak H$_2$O and CH$_4$ absorption features.  Limitations in our knowledge of the opacities of key absorbers such as FeH, VO, and CH$_4$ at certain wavelengths remain obvious, however.  The effective temperatures obtained by fitting the narrower wavelength ranges can show a large scatter compared to the values derived by fitting the full spectral energy distributions; deviations are typically $\\sim$200 K and in the worst cases, up to 700 K.  In spite of the uncertainties remaining in the models, the approach presented here works rather well and should be extended to a much larger sample. ",
        "introduction": "\\label{sec:Introduction}   The emergent spectra of very low-mass (VLM) stars and brown dwarfs, collectively known as ultracool dwarfs, are shaped in large part by gas and condensation chemistry.  The low temperature (500 K $\\lesssim$ $T$ $\\lesssim$ 3000 K), high pressure (1 bar $\\lesssim$ $P$ $\\lesssim$ 10 bar) conditions typical of ultracool dwarf atmospheres favor the formation of molecules (such as CO, CH$_4$, N$_2$, NH$_3$, and H$_2$O) and, at $T$ $<$ 2400 K condensates from the refractory elements (Ti, V, Ca, Al, Fe, Si, and Mg).  Condensation removes gaseous opacity sources (such as TiO, VO, FeH, CrH) from the atmosphere and alters the atmospheric chemistry as the condensates gravitationally settle and form clouds that deplete the upper layers of the atmosphere of both the condensates and their constituent elements \\citep{1994Icar..110..117F,1996ApJ...472L..37F,1999ApJ...519..793L, 2002ApJ...577..974L,1999ApJ...512..843B}.  Condensates also contribute their own opacity, which can be substantial for the more abundant Fe-, Mg-, and Si-bearing species such as forsterite (Mg$_2$SiO$_4$), enstatite (MgSiO$_3$), and solid or liquid Fe \\citep[e.g.,][]{2002ApJ...568..335M,2002ApJ...575..264T}.   This rich chemistry gives rise to complex emergent spectra that are distinctly non-Planckian.  The molecular line opacities dominate over all of the continuum opacity sources (e.g., H$^-$) that are important in the atmospheres of hotter stars \\citep{1995ApJ...445..433A} and as a result the emergent spectra of ultracool dwarfs exhibit strong molecular absorption bands of TiO, H$_2$O, and CH$_4$.  Indeed, it is the strength of these and other molecular bands that are used to define the L and T spectral classes \\citep{1999ApJ...519..802K,2006ApJ...637.1067B}.  L dwarf (2400 K $\\lesssim$ \\teff\\ $\\lesssim$ 1400 K) spectra exhibit weakening TiO and VO bands at red optical wavelengths (0.6$-$1.0 $\\mu$m) due to the formation of Ti-bearing condensates such as perovskite (CaTiO$_3$).  The loss of the gaseous TiO and VO opacity results in a more transparent atmosphere and leads to the emergence of monatomic alkali lines (Na, K, Rb, Cs).  Most prominent are the resonance \\ion{Na}{1} (5890/5896 \\AA) and \\ion{K}{1} (7665/7699 \\AA) doublets, which become increasingly broad with later spectral type.  At near-infrared wavelengths (1$-$2.5 $\\mu$m) the L dwarfs also become progressively redder with later spectral types due to the additional condensate opacity.  As CH$_4$ becomes the dominant carbon-bearing gas in the upper, coolest layers of the atmosphere (CO/CH$_4 <$ 1 for $T\\lesssim$ 1100 K at $P$=1 bar), CH$_4$ absorption bands emerge at near-infrared wavelengths which signals the transition between the L and T (1400 K $\\lesssim$ \\teff\\ $\\lesssim$ 700 K) spectral classes.  The \\ion{K}{1} and \\ion{Na}{1} resonance lines are present in the red optical spectra of T dwarfs \\citep{2003ApJ...594..510B} because the high temperature condensate anorthite (CaAl$_2$Si$_2$O$_8$), a feldspar, has gravitationally settled out of the cooler atmospheric regions.  In the absence of such settling, Na and K would condense into a feldspar in the form of albite (NaAlSi$_3$O$_8$) and orthoclase (KAlSi$_3$O$_8$) \\citep{1999ApJ...512..843B,1999ApJ...519..793L,2002ApJ...568..335M}.   The observed spectra and photometry of ultracool dwarfs can therefore serve as touchstones against which we can test our understanding of atmospheric physics, chemistry, dynamics, and cloud processes as embodied in our atmospheric models.  Shortly after the discovery of the first bona fide brown dwarf, Gliese 229B \\citep{1995Natur.378..463N}, a number of publications appeared comparing the observed spectrum to models \\citep{1996ApJ...465L.123A,1996Sci...272.1919M,1996A&A...308L..29T}. \\citet{2001RvMP...73..719B} presented a review of the substantial subsequent work devoted to understanding the properties of L and T dwarfs.  Although the comparison of models to the spectra of individual objects such as Gl 229B \\citep{2000ApJ...541..374S} and Gl 570D \\citep{2001ApJ...556..373G,2006ApJ...647..552S} has played an important role in the progression of the field, relatively little work has been done on comparing entire classes of models to uniform datasets.  Such comparisons are important as they can reveal the global properties of a set of objects and highlight the systematic strengths and weaknesses of the models.   In this work, we present such a comparison to one particular set of atmospheric models \\citep[][2007, in preparation]{2002ApJ...568..335M}. We aim to understand what the modeling has accomplished to date and to help direct future models and observations along the pathways that seem most promising.  The paper is organized as follows.  Sections \\S2 and \\S3 present the data and model atmospheres used in the analysis, respectively.  In \\S4, we present the results of the comparisons between the data and models and in \\S\\ref{sec:Systematic Uncertainties}, we discuss two systematic uncertainties that affect our analysis. ",
        "conclusions": "We have performed the first fits of the 0.95$-$14.5$\\,\\mu$m SEDs of nine field ultracool dwarfs spanning spectral types L1 to T4.5 to obtain effective temperatures, gravities and values for the cloud sedimentation parameter.  The wide wavelength coverage of the data allows us to easily discriminate between synthetic spectra with different atmospheric parameters because most physical and chemical processes that occur in ultracool dwarf atmospheres leave a signature in one or more of the wavelength regions covered by the data.   In our study of field ultracool dwarfs we have focused on the most significant physical parameters, ignoring the more subtle effects of metallicity variations from the solar value as well as departures from chemical equilibrium caused by vertical mixing in the atmosphere (e.g. \\citet{2006ApJ...647..552S}).  The spectroscopic data come from a variety of instruments and we have estimated the uncertainties in the derived parameters that arise from the absolute flux calibrations of the various spectra that form the SEDs.  Although some wavelength ranges are poorly fitted for some objects, undoubtedly due to remaining inadequacies in the models and, to a smaller extent, to possible variations in metallicity amongst the dwarfs in our sample, overall the model spectra fit the data well.   Given the natural scatter in the observational properties among the objects in our sample and the uncertainties inherent in our procedure, we can identify only general trends in the fitted parameters as a function of spectral type.  The effective temperatures decrease steadily through the L1 to T4.5 spectral types and agree with the determinations of \\citet{2004AJ....127.3516G}. Our results also confirm that \\teff\\ is nearly constant at the L/T transition, decreasing by only $\\sim 200\\,$K from spectral types L7.5 to T4.5.   Two topics of great interest are the evolution of the cloud properties as a function of spectral type and the distribution of cloud properties within each spectral subclass.  Our work provides a first, although somewhat blurry glimpse of this important aspect of ultracool dwarf atmospheres.  We find that the cloud properties, as encapsulated in the most important free parameter of the \\citet{2001ApJ...556..872A} condensation cloud model, \\fsed, vary significantly among L dwarfs, ranging from very thick clouds ($\\fsed=1$) to relatively thin clouds ($\\fsed=3$) with no particular trend with spectral type.  The two L dwarfs in our sample with very red near-infrared colors have best fitting model spectra with thick clouds.  Finally the fits to the two T dwarfs in our sample (T2 and T4.5) suggest that the clouds becomes thinner in this spectral class, in agreement with previous studies.   Fitting individual spectral bands, we almost always obtain excellent fits to the data, in some cases down to the fine structure of the H$_2$O and CH$_4$ bands.  The \\teff\\ obtained in this fashion (Fig. \\ref{fig:TeffSummary}) can show quite a bit of scatter compared to the values derived by fitting the full SED, typically by $\\sim 200\\,$K and in the worst cases, up to 700$\\,$K.  Nevertheless, the trend of decreasing \\teff\\ with spectral type is preserved in all bands, with few exceptions.  However, this exercise shows that in general, \\teff\\ determinations based on the analysis of narrow spectral ranges are not reliable for ultracool dwarfs.   The data point to several directions for improving the models.  First, improved input molecular opacities, particularly for CrH, VO, FeH, and CH$_4$ and the resonance doublet of \\ion{K}{1} (7665, 7699 \\AA) \\citep{2007A&A...474L..21A} are required to improve fits to the early L and all of the T dwarfs, respectively.  The second important area ripe for improvement remains the cloud model, in particular the particles' size distribution, the cloud's vertical profile, and the optical properties of the condensates.   The accurate determination of the physical parameters of isolated field ultracool dwarfs therefore remains an elusive goal.  Our method of SED fitting should be applied to those ultracool dwarfs with well constrained \\teff\\ ($\\pm$20 K) and \\logg\\ ($\\pm$0.05 dex) values (derived from other methods) in order to better assess its internal uncertainties.  Currently, such a sample is limited to a few late-type T dwarfs \\citep{2006ApJ...647..552S,2007ApJ...656.1136S,2007ApJ...660.1507L} that are companions to main sequence stars with known distances, metallicities and well-constrained ages.  Such studies should be extended to include earlier spectral types in order to firmly associate the observable properties of the L and T spectral sequence to fundamental astrophysical parameters."
    },
    "0711/0711.3613_arXiv.txt": {
        "abstract": "We calculate the neutron matter equation of state at finite temperature based on low-momentum two- and three-nucleon interactions. The free energy is obtained from a loop expansion around the Hartree-Fock energy, including contributions from normal and anomalous diagrams. We focus on densities below saturation density with temperatures $T \\leqslant 10 \\mev$ and compare our results to the model-independent virial equation of state and to variational calculations. Good agreement with the virial equation of state is found at low density. We provide simple estimates for the theoretical error, important for extrapolations to astrophysical conditions. ",
        "introduction": "The nuclear equation of state plays a central role in astrophysics, for problems ranging from the structure of neutron stars~\\cite{LP}, neutron star mergers~\\cite{Oechslin} to core-collapse supernovae~\\cite{Mezza,Janka}. Astrophysical applications probe the equation of state at the extremes of isospin and temperature: The mass of a neutron star depends mainly on the equation of state of neutron matter up to densities $\\rho \\sim 4 \\rho_0$~\\cite{KB}, where $\\rho_0 = 0.16 \\fmiq$ is the saturation density of symmetric nuclear matter, while supernova explosions are most sensitive to the properties of nucleonic matter at subnuclear densities and MeV temperatures~\\cite{Mezza}. For many regimes of interest, the equation of state has to be extrapolated from the conditions reached with existing and upcoming experimental facilities. Therefore, reliable theoretical input is needed. In this paper, we present a study of neutron matter at finite temperature, as part of a program to improve the nuclear equation of state input for astrophysics.  Conventional nucleon-nucleon (NN) interactions are nonperturbative as a result of several sources. First, there is a strong short-range repulsion, which leads to bound states of the ``flipped'' potential $\\lambda V_{\\rm NN}$ for small, negative $\\lambda$. This implies that $\\lambda=1$ is far outside the radius of convergence. Consequently, at least the summation of particle-particle ladder diagrams is required~\\cite{BETHE71}. Second, the tensor force, which is singular at short distances, requires iteration in the triplet channels~\\cite{tensor,Fleming}. Finally, there are physical bound and nearly-bound states in the S-waves, which render the perturbative Born series divergent. Recently, it was shown that the first two sources of nonperturbative behavior depend on the choice of NN interaction, and can be removed by evolving nuclear forces to low-momentum interactions $\\vlowk$~\\cite{VlowkReport,VlowkPLB,VlowkRG,Vlowksmooth} with cutoffs around $2 \\fmi$~\\cite{Vlowknucmatt,VlowkWeinberg}. An important additional advantage is that the corresponding leading-order three-nucleon (3N) interactions from chiral effective field theory (EFT) become perturbative in light nuclei for cutoffs $\\lm \\lesssim 2 \\fmi$~\\cite{Vlowk3N}.  At sufficient density ($\\rho \\gtrsim 0.01 \\rho_0$~\\cite{FNN} in nuclear matter), Pauli blocking eliminates the shallow bound states, and thus the particle-particle channel becomes perturbative~\\cite{Vlowknucmatt}. Consequently, the Hartree-Fock (HF) approximation is a good starting point for low-momentum NN and 3N interactions, and perturbation theory (in the sense of a loop expansion) around the HF energy becomes tractable. The perturbative character is due to a combination of Pauli blocking and an appreciable effective range (see also Ref.~\\cite{dEFT}). The 3N interaction is essential for nuclear matter saturation~\\cite{Vlowknucmatt}, while the contributions to the potential energy remain compatible with EFT power-counting estimates. Furthermore, the equation of state becomes significantly less cutoff dependent with the inclusion of the dominant second-order contributions. In this paper, we extend the investigation of Ref.~\\cite{Vlowknucmatt} to neutron matter at subsaturation densities, $\\rho < \\rho_0$, and generalize the perturbative approach to finite temperature.  Based on the work of Kohn, Luttinger and Ward~\\cite{KL,LW}, at finite temperature the loop expansion around the HF free energy can be realized by the perturbative expansion of the free energy. In this paper, we include the first-order NN and 3N contributions, as well as anomalous and normal second-order diagrams with NN interactions. We defer 3N contributions beyond the HF level and higher-order corrections to future work. The pressure, entropy and energy are calculated using standard thermodynamic relations. Since low-momentum interactions are energy independent, the Matsubara sums can be carried out analytically.  Low-momentum interactions $\\vlowk$ and the corresponding 3N forces are defined by sharp or smooth regulators with a variable momentum cutoff $\\lm$. Varying the cutoff is a powerful tool to estimate the theoretical errors due to neglected higher-order many-body interactions and to assess the completeness of the calculations. We use the cutoff dependence to provide simple error estimates, and find that the cutoff dependence is reduced significantly, when second-order contributions are included. The possibility of estimating theoretical errors is an important step towards reliable extrapolations to astrophysical conditions. Finally, we compare our results to the virial equation of state~\\cite{vEOSnuc,vEOSneut} and to variational calculations~\\cite{FP}. The low-density behavior is in good agreement with the virial equation of state. Our results for the energy per particle (see Fig.~\\ref{energy}) highlight the importance of a correct finite-temperature treatment of second and higher-order correlations.  This paper is organized as follows. In Sect.~\\ref{formalism}, we discuss the perturbative expansion at finite temperature and give the expressions for the evaluated diagrams. Our results for the free energy, pressure, entropy and energy are presented in Sect.~\\ref{results}. We conclude and give an outlook in Sect.~\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  This work is part of a program to improve the nuclear equation of state for astrophysics. One of the central objectives is to quantify the theoretical uncertainties in the microscopic nuclear physics input, and to explore the impact on supernovae and neutron stars, for example, through predictions of neutron star masses and radii.  In this first study of neutron matter, we have computed the equation of state at subsaturation densities and temperatures $T \\leqslant 10 \\mev$ based on low-momentum NN and 3N interactions. We have generalized the perturbative approach~\\cite{Vlowknucmatt} to finite temperature, where the free energy is obtained from a loop expansion around the HF energy and the momentum dependence of the self-energy is treated perturbatively. Our results include first-order NN and 3N contributions, as well as anomalous and normal second-order diagrams with NN interactions. The pressure, entropy and energy were then calculated using standard thermodynamic relations. While the HF energy is sizable (and non-perturbative for finite nuclei), the finite-temperature loop expansion around the HF energy seems to be tractable. This is due to a combination of Pauli blocking~\\cite{Vlowknucmatt} and an appreciable effective range~\\cite{dEFT}.  The virial expansion provides a model-independent equation of state for nuclear matter at low density and high temperature~\\cite{vEOSnuc,vEOSneut}, and our perturbative results meet this benchmark. This is very promising, since it will enable us to match the virial equation of state to microscopic calculations based on NN and 3N interactions at higher densities. The comparison of our results to the virial energy per particle highlights the importance of a correct finite-temperature treatment of second and higher-order correlations, which are included only in an average sense in the variational calculations of Ref.~\\cite{FP}. The correct treatment of thermally-excited low-lying states leads to an enhancement of the effective mass at the Fermi surface and consequently to an increase in the entropy, as shown in Fig.~\\ref{entropy}.  We have provided simple estimates for the theoretical error by varying the cutoff in low-momentum interactions. This is a powerful tool to assess theoretical errors due to neglected higher-order many-body forces and due to an approximate many-body treatment. We found that the equation of state becomes significantly less cutoff dependent with the inclusion of second-order contributions, and that the cutoff dependence is small for $\\rho \\lesssim 0.1 \\fmiq$. We note that the errors of the free energy are correlated between different temperatures and grow with increasing density. The first observation implies a relatively small error in the entropy (obtained by a temperature derivative), and consequently similar errors for the energy per particle. The second observation explains the relatively large error band for the pressure (obtained from a density derivative). Finally, we have shown that the uncertainties due to the long-range parts of 3N interactions, the $c_i$ constants, are substantial and overwhelm the error bands from the cutoff variation at this level. We conclude that understanding 3N forces is a frontier in microscopic calculations of the nuclear equation of state, and furthermore that the possibility of estimating theoretical uncertainties is an important step towards reliable extrapolations to astrophysical conditions.  Future work will include systematic studies of the range of validity, quantifying an expansion parameter for the loop expansion, calculations for asymmetric matter, improving the uncertainties at higher densities, how they propagate to astrophysical observables, and understanding the transition to the virial expansion.  \\begin{ack} We thank Scott Bogner, Dick Furnstahl and Chuck Horowitz for useful discussions. AS thanks the GSI Theory Group for the warm hospitality. This work was supported in part by the Virtual Institute VH-VI-041 of the Helmholtz Association, by the BMBF projects ANBest-P and BNBest-BMBF 98/NKBF98, and by the Natural Sciences and Engineering Research Council of Canada (NSERC). TRIUMF receives federal funding via a contribution agreement through the National Research Council of Canada. \\end{ack}"
    },
    "0711/0711.2753_arXiv.txt": {
        "abstract": "The detection by the HESS atmospheric Cerenkov telescope of fourteen new sources from the Galactic plane makes it possible to estimate the contribution of unresolved sources like those detected by HESS to the diffuse Galactic emission measured by the Milagro Collaboration. The number-intensity relation and the luminosity function for the HESS source population are investigated. By evaluating the contribution of such a source population to the diffuse emission we conclude that a significant fraction of the TeV energy emission measured by the Milagro experiment could be due to unresolved sources like HESS sources. Predictions concerning the number of sources which Veritas, Milagro, and HAWC should detect are also given. ",
        "introduction": "The Galactic diffuse $\\gamma$-ray emission is believed to be mostly produced in interactions of cosmic rays with the matter and the radiation fields in the Galaxy, the main production mechanisms being electron non-thermal Bremsstrahlung, Inverse Compton scatterings off the radiation fields and pion decay processes in inelastic collisions of nuclei and matter. Although the standard production mechanisms of $\\gamma$-rays \\citep{Bertsch:1993} explain generally well the spatial and energy distribution of the emission below 1 GeV, the model does not match EGRET observations of the $\\gamma$-ray sky above 1 GeV \\citep{Hunter:1997we}. Many possible explanations have been proposed to account for the GeV excess \\citep{Aharonian:2000iz, Strong:1999hr,Strong:2004,Kamae:2006, deBoer:2005tm, Bergstrom:2006tk,Strong:2006,Stecker}. In any case a major difficulty when studying the diffuse emission is to disentangle the truly diffuse emission from that produced by unresolved sources.  The TeV diffuse $\\gamma$-ray emission from the Galactic Plane \\citep{Atkins:2005} and from the Cygnus region \\citep{Abdo:2006} has been measured by Milagro, a water Cerenkov telescope surveying the northern sky at TeV energies \\citep{Atkins:2000,Atkins:2003}. The diffuse emission measured by Milagro is $(7.3  \\pm 1.5  \\pm 2.3)  \\times {10}^{-11} \\, photons \\, {\\mathrm{{cm}^{-2}}} \\,{\\mathrm{{sr}^{-1}}} \\, {\\mathrm {s^{-1}}} $ for $E>3.5$ TeV in the region ${40}^{o}<l<{100}^{o}$ \\citep{Atkins:2005}. A measurement of the TeV $\\gamma$-ray diffuse flux is important, as it allows us to calculate what fraction of the emission is produced through inverse Compton and what fraction is produced through pion-decay mechanism. Milagro result, which is consistent with EGRET data above 1 GeV, seems to exclude an additional hard spectrum component continuing to above 10 TeV to explain the 60 per cent excess of EGRET flux compared to $\\pi_{0}$-decay production mechanism due to the local cosmic ray flux \\citep{Hunter:1997we}. Recently \\citet{Prodanovic} have argued that no strong signal of pion decay is seen in the $\\gamma$-ray spectrum. The pion decay mechanism would not be able to explain the excess above 1 GeV and the Milagro measurements of the diffuse flux at higher energies reveals a new excess at TeV energies. \\citet{Prodanovic} investigated several possibilities for what might cause the TeV excess, among them possible dark matter decay, contribution from unresolved EGRET sources or sources that are only bright in the TeV range. An extrapolation of the EGRET sources within the range in Galactic longitude of the Milagro observation overestimates the diffuse TeV flux measured. However, there could be a population of sources undetectable by EGRET, but contributing to both the GeV and TeV excess diffuse emission. Recent observations by the TeV observatory, HESS, in fact point to a new class of hard spectrum TeV sources located close to the Galactic Plane, within $\\pm 1$ degrees of latitude, whose average slope is about $E^{-2.3}$ \\citep{Aharonian:2005jn,Aharonian:2005kn}. The HESS detection of high energy $\\gamma$ rays from fourteen new sources has improved significantly the knowledge of both the spatial distribution and the spectra and fluxes of VHE $\\gamma$-ray galactic sources. These HESS results make it possible to estimate with unprecedented precision the contribution of unresolved sources to the Galactic diffuse emission recently extended to TeV energies by Milagro. In fact, the Milagro sensitivity to point sources is about ${10}^{-11}  \\, photons \\, {\\mathrm{{cm}^{-2}}}  \\, {\\mathrm {s^{-1}}} \\, {\\mathrm for} \\quad E > 1 TeV$  and it is of the order of the brightest integral source flux detected by HESS. Sources like the ones detected by HESS are then unresolved by Milagro and contribute to the diffuse emission Milagro measures.  Here, based on HESS results, knowing the sensitivity and the field of view of an experiment, we estimate the number of expected sources and their expected VHE $\\gamma$-ray flux. We then evaluate the contribution of unresolved sources to the diffuse $\\gamma$-ray emission for Milagro. The estimate of the diffuse emission arising from unresolved sources will be a lower limit. In fact, HESS has mostly detected extended sources. For sources larger than the angular resolution of the experiment the source size determines the sensitivity of the instrument. The sensitivity decreases if the extension of the source increases. In Fig.~\\ref{fig1} the HESS source fluxes and the sensitivity as a function of the source extensions are plotted. A clustering of sources about HESS survey sensitivity of 3 percent of the Crab flux reported in \\citep{Aharonian:2005jn} is evident, meaning that there are probably many more to be detected with better sensitivity. \\begin{figure}[ht] \\begin{center} \\includegraphics[angle=0,width=8cm]{f1.eps} \\caption{The HESS source fluxes and the HESS survey sensitivity versus source extensions are plotted \\cite{Aharonian:2005kn}. A clustering of the sources about the HESS sensitivity limit is evident.}\\label{fig1} \\end{center} \\end{figure} ",
        "conclusions": "The number-intensity relation and the luminosity function for the HESS source population were investigated using the assumption that HESS sources are distributed as PSRs and SNRs detected at radio wavelengths. In order for the chosen sample of sources to be complete only the HESS sources with fluxes above 6 percent of the Crab flux were taken into account to derive the number-intensity relation. The contribution of unresolved HESS-like sources to the diffuse emission measured by Milagro was also estimated. Using the logN-logS relation for the HESS sample of Galactic $\\gamma$-ray emitters at least 10 per cent of the diffuse emission at TeV energies is estimated to be due to the contribution of unresolved HESS-like sources. This result is a lower limit for such a contribution because we have taken into account only sources detected above 6 per cent of the Crab flux and because HESS sensitivity gets worse for extended sources, meaning that some extended sources might have been missed by HESS. Using the logN-logS relation we have also predicted the number of HESS-like sources which VERITAS, HESS and HAWC should detect during their survey of the sky. An alternative procedure to evaluate the contribution of unresolved HESS-like sources to Milagro diffuse emission gives the diffuse flux due to unresolved sources comparable to the diffuse emission itself. We finally constrained the slope of the luminosity function. The main uncertainty of this calculation consists in assuming that the distribution of $\\gamma$-ray sources follows the distribution of either pulsars or SNRs observed in the radio. In particular, in order to predict how many PSRs observed in the radio have a PWN and are possible gamma ray emitters we used the result that the spin-down energy loss $dE/dt > dE/dt_c = 4 \\times {10}^{36} erg/s$ for a young energetic pulsar to form a PWN. In this respect we have ignored the existence of pulsars, such as Geminga, which are $\\gamma$-ray loud, yet not observed in the radio.  New observational results support the hypothesis that a population of unresolved sources contribute significantly to the emission at very high energy. Milagro has recently reported the discovery of TeV gamma ray emission from the Cygnus Region of the Galaxy, which exceeds the predictions of conventional models of gamma -ray production \\citep{Abdo:2006} from the same region in the Galaxy where the Tibet Array has detected an excess of cosmic rays \\citep{Amenomori}. Thanks to its improved sensitivity Milagro has also better imaged the whole Northern sky and discovered four sources and four source candidates \\citep{Abdo:2007}. HESS has seen very high energy emission spatially correlated with giant molecular clouds located in the Galactic Center \\citep{Aharonian:nature}. The energy spectrum measured by HESS close to the Galactic Center is $E^{-2.3}$, significantly harder than the $E^{-2.7}$ spectrum of the diffuse emission and equal to the average spectrum of the HESS source population. The emission from the Galactic Center might possibly unveil a cosmic ray accelerator.  To draw more definitive conclusions about the very high energy $\\gamma$-ray sky, new observations are of fundamental importance. New hints will be provide by both MAGIC and VERITAS, which already survey the Cygnus Region. Finally GLAST will investigate the window of energy between 10 MeV to 300 GeV, covering the energy gap left between EGRET and the ground-based low threshold gamma-ray observatories."
    },
    "0711/0711.2109_arXiv.txt": {
        "abstract": "{The \\hhhop\\ molecule probes the oxygen chemistry and the ionization rate of dense circumnuclear gas in galaxies.} {Recent \\hhhop\\ observations show variations in the cosmic-ray ionization rate by factors of $>$10 within our Galaxy.} {Using the JCMT, we have observed the 364 GHz line of p-\\hhhop\\ in the centers of M82 and Arp 220.} {In Arp 220, the line profile suggests that the emission originates in the Western nucleus. In M~82, both the eastern molecular peak and the circumnuclear region contribute to the emission. The derived column densities, abundances, and \\hhhop\\ / \\hho\\ ratios indicate ionization rates similar to or even exceeding that in the Galactic Center.} { Model calculations of the chemistry of irradiated molecular gas indicate a likely origin of this high ionization rate in the extended, evolved starburst of M82. In contrast, irradiation by X-rays from the AGN disk is the most likely model for Arp~220.} ",
        "introduction": "\\label{s:intro}  Starbursts and active galaxies often host large masses of molecular gas and dust in their inner kpc. Gravitational interaction or merging cause the funneling (via bars) of the gas towards the center. The average gas surface- and number densities are significantly higher (by several orders of magnitude) than in the more extended disks. The nuclear molecular gas is often found in tori, feeding the activity and thus controlling its evolution. The triggering and turn-off mechanisms of the activity are dependent on the properties of the gas, which are still poorly known, in particular the fraction of the gas at high temperature and density.  Observations of the molecular gas properties are starting to show great promise not only in allowing us to model the evolution of the starburst, but also in helping to identify and analyze the type of activity lurking in very deeply obscured galactic nuclei. Here, dust extinction is high enough to quench emission even down to the mid-infrared. For reviews of this subject, see \\citet{sanders:ulirgs} and \\citet{gao:ulirgs}.  The star formation rates in the nuclei of nearby galaxies are measured best at \\smm\\ wavelengths, where extinction is negligible and where a rich spectrum of dust continuum and molecular lines allows us to measure fundamental parameters of star-forming matter such as mass, temperature, and density (e.g., \\citealt{aalto:ulirgs}). However, existing observations are mostly at millimeter wavelengths and thus limited to relatively low temperatures and densities. The excellent transmission of sites like Mauna Kea and Chajnantor enables observation of the \\hhhop\\ molecule, which has a twofold importance.  First, the \\hho\\ molecule acts as a natural filter for selecting hot ($\\ga$100\\,K) gas, and therefore traces different regions than CO and other molecules commonly studied from low-altitude sites. The chemistry of this filtering is the evaporation of icy grain mantles at $T\\approx 100$\\,K \\citep{vdtak:h2o}. Other molecules evaporating from grain mantles are not as abundant as \\hho. The \\hhhop\\ molecule is an excellent proxy for \\hho, which cannot be observed in thermal lines from the ground at redshifts below 0.01. In contrast to HDO whose ratio to \\hho\\ strongly depends on ambient conditions, the chemical relation between \\hhhop\\ and \\hho\\ is well-understood and the relevant reactions have been measured in the laboratory \\citep{phillips:h3o+}. Thus, observations of \\hhhop\\ are useful probes of the gas distribution and kinematics in the warm and dense nuclei of active galaxies.  Second, the high proton affinity of \\hho\\ makes \\hhhop\\ a key ion in the oxygen chemistry of dense molecular gas. Combined with information on the \\hho\\ abundance, observations of \\hhhop\\ may be used to trace the ionization rate of dense circumnuclear clouds by cosmic rays (produced in supernovae, i.e., starbursts) and X-rays (from an AGN). The ionization rate of this gas is a fundamental parameter for its dynamics, because it regulates the efficiency of magnetic support against gravitational collapse. Recent APEX observations of \\hhhop\\ show that the cosmic-ray ionization rate in the Sgr~B2 cloud near the Galactic center is $\\sim$10$\\times$ higher than in star-forming regions at 1--4\\,kpc from the Sun \\citep{vdtak:sgrb2}.  Motivated by the results of \\hhhop\\ observations in our Galaxy, we searched for \\hhhop\\ emission in two prototypical active galaxies: M82 and Arp~220. At a distance of 4\\,Mpc, M82 is the prototype evolved starburst, with a large \\hh\\ mass, where many molecules were seen for the first time outside our Galaxy, including \\cop, a tracer of strong irradiation of molecular gas \\citep{fuente:co+}. The molecular emission can be roughly divided into two prominent lobes of molecular emission, separated by 30$''$, which is interpreted as a 400 pc rotating molecular ring. There is also molecular gas in the nucleus, which is best seen in high-$J$ CO lines or in high-density tracer emission, such as HCO (e.g., \\citealt{garcia-burillo:m82}).  The prototypical ultraluminous galaxy Arp~220 ($d$ = 72\\,Mpc), where water lines have been detected \\citep{gonzalez:arp220,cernicharo:arp220}, is a merger of two gas-rich disk galaxies. Significant ($\\gtsim$10$^9$\\,\\msol) molecular gas has gathered in two nuclei (east and west) that have not yet merged into one. The two nuclei are separated by 1$''$ and are surrounded by a 5$''$ extended disk of molecular gas. The bulk of the molecular emission from Arp~220 is emerging from this region (e.g., \\citealt{sakamoto:arp220}), and the JCMT beam covers all of it. Arp 220 has a rich chemistry with surprisingly bright HNC 3--2 line emission, outshining HCN 3--2 by a factor of 2, which indicates unusual chemical or excitational circumstances. High-resolution observations reveal that the HNC emission is emerging from warm gas ($T_{\\rm k} >$ 40\\,K) near the western nucleus \\citep{aalto:madrid}. ",
        "conclusions": "\\label{s:disc}  The column densities of \\hhhop, \\hho, and \\hh\\ in Table~\\ref{t:coldens} and their ratios are compared to chemical models of clouds irradiated by either far-UV or X-ray photons. Far-UV dominated clouds are called photon-dominated regions (PDRs) and the other type X-ray dominated regions (XDRs). The thermal and chemical structures of PDRs and XDRs are quite different, because the absorption cross section for X-rays is much lower than for far-UV. The ionization fraction in XDRs can be as high as $x_e = 10^{-2} - 10^{-1}$, while in PDRs, it is 10$^{-4}$ at most. Furthermore, in XDRs, larger parts of the cloud can be kept at a high temperature. For an elaborate discussion of the differences between PDRs and XDRs we refer to \\citet{meijerink:program} where the codes are described. The grid of models by \\citep{meijerink:grid} is used to look for the best match to our observations. For the PDR models, we also consider an elevated cosmic-ray ionization rate of $\\zeta = 5.0 \\times 10^{-15}$\\,\\rs. As reference, we take the radiation field found by \\citet{spaans:M82} for M82 to model the CO$^+$ abundance found by \\citet{fuente:co+}. Their best-fitting model was an XDR with a density of $n=10^5$\\,\\ccm\\ and an incident X-ray flux of $F_X=5.1$\\,erg\\,\\rs\\,\\scm\\ ($G_0$=10$^{3.5}$). Using this density and radiation field, the model predicts the column densities listed in Table~\\ref{t:model}. The chemical structure of the model is shown in Fig.~\\ref{f:model}. Models for volume densities far below $n=10^5$\\,\\ccm\\ fail to produce appreciable \\hhhop\\ abundances, which makes it unlikely that the origin of the observed emission is in interclump gas, seen e.g.\\ as the $N_H$=10$^{20}$\\,\\scm\\ CO component in M82 \\citep{mao:m82}.  Table~\\ref{t:model} shows that the \\hhhop\\ column density observed in M82 may be produced in either XDR or in the high-$\\zeta$ PDR. The \\hhhop/\\hh\\ ratio in M82 is overproduced in the XDR models and underproduced in the standard PDR, but agrees very well with the high-$\\zeta$ model. None of the models reproduce the observed \\hhhop/\\hho\\ ratio in M82, probably because the observations refer to very different volumes. We conclude that the M82 data are matched best by the PDR model with enhanced cosmic-ray flux. In contrast, the X-ray models do best in reproducing the observations of Arp~220. The \\hhhop\\ column density of 10$^{16}$\\,\\scm\\ is matched by an XDR model with $N_H$=10$^{23}$\\,\\scm, scaled up by a factor 10, for either value of $F_X$. The \\hhhop\\ abundance of \\pow{1}{-8} is matched by the high-$F_X$ model for low $N_H$ or the low--$F_X$ model for high $N_H$. The \\hhhop/\\hho\\ ratio is matched by an XDR of either $F_X$, where low values of $N_H$ are preferred. Nevertheless, a contribution by the eastern nucleus is implausible because of its low mass of dense gas and its low infrared brightness.  In the context of galactic nuclei, a PDR model represents a starburst, where the far-UV is produced by young massive stars and the cosmic rays in supernova remnants. The pure PDR thus stands for a young starburst and the high-$\\zeta$ model for an evolved one. The XDR models represent the much harder radiation environment of an AGN, such as the accretion disk of a supermassive black hole. We find that the observations of M82 are matched by a high-$\\zeta$ PDR, i.e., an evolved starburst. This conclusion is consistent with earlier models of this source \\citep{schreiber:m82}. Our \\hhhop\\ observations do not require X-ray irradiation as in the case of CO$^+$ \\citep{spaans:M82}.  The fact that the XDR models give conflicting requirements on the $N_H$ value toward Arp~220 may well be due to the unusual geometry of this object. The emerging picture of the western nucleus of Arp~220 is that there is a 35 pc optically thick dust disk surrounding what could potentially be an active nucleus. The dust disk is warm ($T_{\\rm D} \\approx$ 170 K) and has a peak H$_2$ column density of $\\approx$10$^{25}$\\,cm$^{-2}$. Hence X-rays, optical, and even infrared emission from the nucleus itself will be absorbed \\citep{downes:arp220}. Surrounding this warm dust disk is a cooler ($T_{\\rm k} \\approx $ 50 K) molecular disk or torus causing CO 2--1 absorption features towards the hotter nuclear dust disk in the line of sight. Most of the observed starburst activity appears to be associated with the rotating disk/torus surrounding the nuclear dust disk. Future modeling efforts of \\hhhop\\ in Arp~220 should take this special geometry into account. However, high-resolution observations are needed to constrain these models. It is possible for example that the disk causes both \\hhhop\\ emission and absorption along the line of sight.  \\begin{table}[tb] \\caption{Results of PDR / XDR models. } \\label{t:model} \\begin{tabular}{rrrrrr} \\hline \\hline \\noalign{\\smallskip} $N_H$ &  $N$(\\hh)  &   $N$(\\hhhop) &  $N$(\\hho) &  \\hhhop/\\hh &  \\hhhop/\\hho \\\\ \\scm  &  \\scm      &   \\scm        &  \\scm      &             &  \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\multicolumn{6}{c}{XDR $F_X$ = 5.1 erg\\,\\rs\\,\\scm} \\\\ \\noalign{\\smallskip} \\pow{1}{22} & \\pow{1.4}{21} & \\pow{6.1}{13} & \\pow{5.6}{15} & \\pow{4.3}{-8} & \\pow{1.1}{-2} \\\\ \\pow{3}{22} & \\pow{7.9}{21} & \\pow{1.9}{14} & \\pow{2.2}{16} & \\pow{2.5}{-8} & \\pow{8.7}{-3} \\\\ \\pow{1}{23} & \\pow{3.9}{22} & \\pow{3.6}{14} & \\pow{5.8}{16} & \\pow{9.2}{-9} & \\pow{6.1}{-3} \\\\ \\noalign{\\smallskip} \\multicolumn{6}{c}{XDR $F_X$ = 16 erg\\,\\rs\\,\\scm} \\\\ \\noalign{\\smallskip} \\pow{1}{22} & \\pow{1.7}{20} & \\pow{3.1}{12} & \\pow{1.7}{14} & \\pow{1.8}{-8} & \\pow{1.8}{-2} \\\\ \\pow{3}{22} & \\pow{2.3}{21} & \\pow{1.5}{14} & \\pow{1.2}{16} & \\pow{6.6}{-8} & \\pow{1.2}{-2} \\\\ \\pow{1}{23} & \\pow{2.5}{22} & \\pow{8.4}{14} & \\pow{8.9}{16} & \\pow{3.4}{-8} & \\pow{9.4}{-3} \\\\ \\noalign{\\smallskip} \\multicolumn{6}{c}{PDR $G_0$=10$^{3.5}$ $\\zeta=5.0\\times10^{-17}$\\,\\rs} \\\\ \\noalign{\\smallskip} \\pow{1}{22} & \\pow{4.6}{21} & \\pow{2.3}{11} & \\pow{1.6}{14} & \\pow{5.0}{-11} & \\pow{1.4}{-3} \\\\ \\pow{3}{22} & \\pow{1.5}{22} & \\pow{1.7}{12} & \\pow{1.6}{15} & \\pow{1.2}{-10} & \\pow{1.1}{-3} \\\\ \\pow{1}{23} & \\pow{5.0}{22} & \\pow{6.8}{12} & \\pow{6.6}{15} & \\pow{1.4}{-10} & \\pow{1.0}{-3} \\\\ \\noalign{\\smallskip} \\multicolumn{6}{c}{PDR $G_0$=10$^{3.5}$ $\\zeta=5.0\\times10^{-15}$\\,\\rs} \\\\ \\noalign{\\smallskip} \\pow{1}{22} & \\pow{4.6}{21} & \\pow{9.2}{12} & \\pow{2.7}{15} & \\pow{2.0}{-9} & \\pow{3.4}{-3} \\\\ \\pow{3}{22} & \\pow{1.4}{22} & \\pow{4.3}{13} & \\pow{1.3}{16} & \\pow{2.9}{-9} & \\pow{3.2}{-3} \\\\ \\pow{1}{23} & \\pow{4.9}{22} & \\pow{1.6}{14} & \\pow{5.1}{16} & \\pow{3.2}{-9} & \\pow{3.1}{-3} \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{table}  \\begin{figure*}[tb] \\centering \\includegraphics[width=5cm,angle=0]{8824-f2a.ps} \\bigskip \\includegraphics[width=5cm,angle=0]{8824-f2b.ps} \\caption{Calculated chemical structure of XDR (top) and PDR (bottom) models}. \\label{f:model} \\end{figure*}  In the near future, the HIFI heterodyne spectrometer onboard ESA's \\textit{Herschel} space observatory will be able to measure \\hho\\ and \\hhhop\\ lines around 1~THz and make accurate estimates of the \\hho/\\hhhop\\ ratios in many sources. These data will allow a much better understanding of the ionization rates of molecular regions inside and outside our Galaxy. Future additions to the model include the effects of freeze-out onto and desorption from dust grains, the effect of metallicity changes, and (especially for the Arp 220 merger) the time dependence introduced by shocks."
    },
    "0711/0711.0019.txt": {
        "abstract": "The SWIFT gamma ray observatory's Burst Alert Telescope (BAT) has detected a sample of active galactic nuclei (AGN) based solely on their hard X-ray flux (14-195\\,keV). In this paper, we present for the first time {\\it XMM-Newton} X-ray spectra for 22 BAT AGNs with no previously analyzed X-ray spectra. If our sources are a representative sample of the BAT AGN, as we claim, our results present for the first time global X-ray properties of an unbiased towards absorption (n$_H < 3 \\times 10^{25}$\\,cm$^{-2}$), local ($<z> = 0.03$), AGN sample. We find 9/22 low absorption (n$_H < 10^{23}$\\,cm$^{-2}$), simple power law model sources, where 4 of these sources have a statistically significant soft component.  Among these sources, we find the presence of a warm absorber statistically significant for only one Seyfert 1 source, contrasting with the ASCA results of \\citet{rey97} and \\citet{geo98}, who find signatures of warm absorption in half or more of their Seyfert 1 samples at similar redshifts. Additionally, the remaining sources (14/22) have more complex spectra, well-fit by an absorbed power law at $E > 2.0$\\,keV.  Five of the complex sources are classified as Compton-thick candidates.  Further, we find four more sources  with properties consistent with the hidden/buried AGN reported by Ueda {\\it et al.} (2007).  Finally, we include a comparison of the {\\it XMM-Newton} EPIC spectra with available SWIFT X-ray Telescope (XRT) observations.  From these comparisons, we find 6/16 sources with varying column densities, 6/16 sources with varying power law indices, and 13/16 sources with varying fluxes, over periods of hours to months.  Flux and power law index are correlated for objects where both parameters vary. ",
        "introduction": "Active galactic nuclei (AGN) surveys are typically dominated by two selection effects: (1) dilution by starlight from the host galaxy and (2) obscuration by dust and gas in the host galaxy and/or the AGN itself (see Hewett \\& Foltz 1994 and Mushotzky 2004).  These factors previously kept an unbiased AGN sample from reach.  However, with the capabilities of SWIFT's Burst Alert Telescope (BAT), this has changed. The BAT surveys the sky in the hard X-ray range of 14-195\\,keV.  As of early 2006, BAT detected $\\sim$150 AGN with a median redshift of 0.03 \\citep{mar05, tue07}.  These AGN were selected purely by their hard X-ray flux, and thus, all but the most heavily absorbed sources (n$_H > 3 \\times 10^{25}$\\,cm$^{-2}$) are not affected by obscuration from gas and dust, which prevents them from being easily detected in UV, optical, or soft band X-ray surveys.  Thus, the BAT survey will provide the information to derive the true distribution of AGN characteristics across the entire electromagnetic spectrum.  The BAT detection limit is a few $\\times 10^{-11}$\\,erg\\,s$^{-1}$\\,cm$^{-2}$.  Thus, the BAT AGNs are powerful sources, with X-ray luminosities over the full range of AGN luminosities. For the closest Seyfert galaxies (z$\\approx0.01$), the BAT flux limit corresponds to a 14 -- 195\\,keV luminosity of roughly $3\\times10^{42}$\\,erg\\,s$^{-1}$. Most of the sources ($\\approx 97$ \\%) have been optically detected as relatively bright nearby objects with $<z> \\approx 0.03$, while most of the sources without bright optical counterparts are blazars, of which there are 15/153 \\citep{tue07}.  Since the BAT sources are mostly bright and nearby, they are good sources with which to study multi-wavelength properties.  Despite the X-ray brightness of the BAT AGN, a number of these sources had yet to be observed spectroscopically in the X-rays.  We present an analysis of {\\it XMM-Newton} EPIC spectra of 23 sources, 2 of which correspond to a pair of interacting galaxies where the BAT source is confused, i.e. likely a combination of the AGN, for which spectra were obtained through discretionary project scientist time.  These sources had no previous X-ray spectrum, were clearly detected by the BAT, and had clear optical counterparts from Digital Sky Survey (DSS) images. Many of these sources now also have SWIFT X-ray Telescope (XRT) spectra available.  Thus, we are able to compare the XRT spectra with the EPIC spectra.  In Section 2, we detail the observations and data analysis for the XMM-Newton and SWIFT observations.  In Section 3, we describe spectral fits to the X-ray spectra, including a search for variability between the XRT and {\\it XMM-Newton} observations.  We summarize our findings in Section 4.  %% How are the BAT AGN selected -- refer to paper ",
        "conclusions": "In this study, we examined the X-ray properties of a sub-sample of BAT detected AGN from the 9-month BAT catalog.  These sources, selected based on their 14-195\\,keV flux, probable optical identifications with the Digital Sky Survey and 2MASS, and their lack of an archival X-ray spectrum, are probably representative of the whole 9-month sample.  Having a bright DSS or 2MASS counterpart does not impart a significant selection effect, considering that all but 20 of the BAT AGN 9-month sample fit this criteria.  The 20 sources without a bright optical/IR counterpart are mostly blazars, of which there are only 15 in the 9-month catalog.  Since blazars are less than 10\\% of the sample, we decided to focus on the majority, non-blazar sources.  For the BAT AGN sample, \\citet{tue07} list 2MASS K$_s$ magnitudes for the sources with available values from NED.  The average value for the BAT sample is 11.47 with a variance of 4.04.  Our sub-sample is slightly dimmer, with an average magnitude of 12.10 and a variance of 1.80, but well within the distribution of BAT AGN magnitudes.  The column density distribution of our sub-sample, shown in Figure~\\ref{fig-bat}, is representative of the larger sample.  In the plot of hydrogen column density versus 14-195\\,keV flux (values from \\citet{tue07}), our 22 {\\it XMM-Newton} follow-up sources span the range of hydrogen column densities. Roughly half of the sources have low column densities (n$_H < 10^{23}$\\,cm$^{-2}$), while half are more heavily absorbed.  We find the same ratio of absorbed to non-absorbed sources in our \\xmm follow-up sample.  While the low absorption sources in our sample span the lower range of hard X-ray fluxes (from BAT), this is expected, since the sources were not previously studied in the X-ray regime.  For the absorbed sources, however, we find that our sample spans the full range of BAT X-ray fluxes.   Given that our sample of sources is a representative sample of the 9-month BAT catalog, it is worthwhile to discuss the general properties of our sources.  To begin, the optical host galaxy classifications of our sources are listed in Table~\\ref{tbl-1}.  Examining the host galaxy classifications, 17/22 of the hosts are classified as spirals or peculiar spirals. This result is interesting, considering that \\citet{gro05} found the hosts of X-ray selected, z $\\approx 0.4-1.3$, Chandra Deep Field sources to be dominated by ellipticals.  If the BAT AGN hosts are predominately spirals, as our sample suggests, this could imply an evolutionary effect in AGN host galaxies between the z$ \\approx 0.03$ and z $\\approx 0.4-1.3$ universe.  From our detailed X-ray spectral fits, we found that 9 of the 22 sources had column densities below $10^{23}$\\,cm$^{-2}$ and spectra well-fit by simple power law or power law with a soft excess models.  Nearly half of these sources showed evidence of having a soft excess.  Optically, all of these sources except for SWIFT J0216.3+5128 and WKK 1263 are Seyfert 1 -- 1.2 sources. As discussed in the detailed spectral fitting section, SWIFT J0216.3+5128 is most likely a blazar.  Though, WKK 1263 is identified as a Seyfert 2 in NED, no optical spectrum is available in the literature to confirm this.  It is possible that this source was simply misclassified.  The remaining sources in our study had more complex X-ray spectra.  For most of these sources, half of the sample, an absorbed power law component model fit is unacceptable.  This was particularly true for ESO 362-G018 and NGC 6860.  The column densities and power law indices computed from the simple model, for these sources, are drastically different than values from more complex models.  Both of these sources, optically Sy 1.5, had optical images indicating dust clouds interspersed with star formation.  Likely, the complex environment contributed to the complexity seen in the X-ray observation.  However, without higher signal-to-noise observations we were unable to resolve the complex spectral components.  These results illustrate the danger of using low quality data/simple models to determine the properties of complex sources.  The remaining 11 sources, half of the sample, had column densities clearly above $10^{23}$\\,cm$^{-2}$.  We classified the observed spectra of these sources as having a ``double power law'' shape, similar to \\citet{tur97}. Optically, these sources are Seyfert 2s.  The exceptions are NGC 612, a weak-lined, giant radio galaxy, and NGC 4992, which have ``galaxy'' spectra or optical spectra showing no emission lines indicative of AGN emission.  NGC 612 was specifically classified as a ``non-LINER'', e.g. having non-AGN line ratios, by \\citet{lew03}, while NGC 4992 is an INTEGRAL source whose optical spectrum led \\citet{mas06} to classify the source as an X-ray bright, optically normal galaxy. Considering the X-ray column densities for these sources, the optical AGN emission could be hidden or obscured by the high column of gas, n$_H > 5 \\times 10^{23}$\\,cm$^{-2}$, in the line of sight.  For all of the complex X-ray spectra sources, we fit the spectra with (1) a partial covering absorption model and (2) a double power law model, where each power law component had a separate absorption model. Most of the sources showed no significant difference in $\\chi^2$ between these two models. This makes it impossible to determine whether the soft flux is the result of scattering of the AGN light, partial covering of the AGN light, or other X-ray sources (such as X-ray binaries or diffuse galaxy emission) contaminating the AGN spectrum.  Particularly for sources, such as NGC 612, with low 0.5 -- 2\\,keV fluxes, e.g., F$_{0.5 - 2 keV} = 2 \\times 10^{-14}$\\,erg\\,s$^{-1}$\\,cm$^{-2}$ corresponding to a luminosity, L$_{0.5 - 2 keV} \\approx 4 \\times 10^{36}$\\,erg\\,s$^{-1}$, within the observed range of Galactic X-ray binaries, since the X-ray luminosity function extends to $\\approx 3 \\times 10^{38}$\\,erg\\,s$^{-1}$ \\citep{gri02}.  For sources with flat spectra (low photon index, $\\Gamma$),  we fit the spectra with a Compton thick model (a reflection dominated model, {\\tt pexrav}).  We added the BAT spectra to the \\xmm data in order to extend the energy range to 200\\,keV.  In addition to providing an adequate fit to the data, replacing the heavily absorbed power law component with a reflection model resulted in higher spectral slopes more consistent with average AGN photon indices (where the BAT spectra were not curved, see Table~\\ref{tbl-10}).  Based on the detailed model fits, NGC 612, ESO 362-G018, MRK 417, ESO 506-G027, and NGC 6860, are classified as Compton-thick candidates.  Two of these sources, ESO 362-G018 and NGC 6860, have measured column densities n$_H << 10^{23}$\\,cm$^{-2}$.  While a column density this low is not expected from reflection in a Compton-thick region, an alternate model where the reflection component is scattered and then absorbed outside of the Compton-thick region, such as employed for MRK 231 \\citep{mah00}, could explain the spectra.  In addition to the Compton-thick candidates, four of the complex spectra sources had high partial covering fractions ($> 0.99$) with the partial covering model and very low ratios of the unabsorbed power law to the absorbed power law component (N$_{\\Gamma_1}$/N$_{\\Gamma_2} < 0.02$) with the double power law model.  Thus, emission from the sources SWIFT J0641.3+3257, SWIFT J0911.2+4533, SWIFT J1200.8+0650, and NGC 4992, was extremely low in the soft band (0.5 -- 2\\,keV) compared to the hard band (2 -- 10\\,keV).  These sources are consistent with the new class proposed by \\citet{ued07} of hidden or buried AGN.  \\citet{ued07} predict that these sources should have lower [\\ion{O}{3}] luminosities than typical Seyfert 2 sources.  Archival optical spectra of NGC 4992 show very weak [\\ion{O}{3}], in fact, so much so that, as discussed above, the spectrum of this source appears as a typical galaxy.  In our optical study (in preparation), we will explore this issue further.  For now, it is important to note that more of these hidden sources exist.  If our sample, showing 4/22 hidden AGN, is representative of the larger BAT sample, we expect that about 1/5th of local AGN have these same properties, making them nearly undetectable in optical samples.  Having classified the sources into categories, we now describe the general properties of our sample as a whole.  To begin, in Figure~\\ref{fig-fluxratio} we plotted the column densities versus two different flux ratios (similar to a diagnostic in \\citet{mal07}).  The column densities we used are listed in Tables~\\ref{tbl-5},~\\ref{tbl-6}, and~\\ref{tbl-8}, where we used the column density of the more heavily absorbed power law component for the complex spectra.  The flux ratios plotted are the ratios of F$_{2 - 10\\,keV}$/F$_{14 - 195\\,keV}$ (medium/hard) and F$_{0.5 - 2\\,keV}$/F$_{2 - 10\\,keV}$ (soft/medium).  In the plots, we represent the three classes of objects (simple power law or power law and blackbody fit sources with low columns (blue), complex heavily absorbed spectra (black), and complex spectra that require more complicated models (red)).  We find that the low absorption sources have average values of medium/hard and soft/medium flux of 0.38 and 0.48, respectively.  There is little change between the medium/hard and soft/medium flux ($\\approx 20$\\%) for the low absorption sources.   The heavily absorbed sources, however, have average values of medium/hard and soft/medium flux of 0.08 and 0.02.  This is a 75\\% change in the values. Obviously, there is much less soft flux for the absorbed sources.  The complex sources with poorly defined spectral models, have intermediate values of the medium/hard and soft/medium colors of 0.17 and 0.11.  In terms of use as a diagnostic, we find that the plot of column density versus ratio of F$_{2 - 10\\,keV}$/F$_{14 - 195\\,keV}$ (medium/hard) is a good diagnostic of column density, for n$_H < 10^{24}$\\,cm$^{-2}$.  Sources with similar column densities occupy areas close to the regions of constant power law index plotted (for $\\Gamma = 1.5$ and 1.9).  This is not true for the plot of column density versus ratio of F$_{0.5 - 2\\,keV}$/F$_{2 - 10\\,keV}$ (soft/medium).  This appears to be a poor diagnostic, despite its wide spread use in deep X-ray surveys, with a large spread in the soft/medium color particularly seen in the sources with n$_H > 10^{23}$\\,cm$^{-2}$.  These results are not surprising, since the 2--10\\,keV flux is also affected by absorption.  From these color diagrams, we decided to construct a color-color diagram of the soft/medium flux ratio versus the hard/medium flux ratio in attempts to construct a better diagnostic diagram for sources with too few counts to measure column density.  In Figure~\\ref{fig-color}, we plot this diagram using the same symbols as in the previous diagram to indicate low absorption (blue), complex (red), and more heavily absorbed (black) sources.  In this figure, it is clear that the different types of sources are clearly separated.  The low absorption sources occupy the left hand upper corner, where the soft/medium and hard/medium colors are nearly equal.  The heavily absorbed sources are closer to the right bottom corner, where the hard/medium flux ratio is much higher than the soft/medium flux ratio.  Between these values, the complex sources as well as a source from each of the other two categories, reside.  All of these sources have measured column densities from $10^{22} - 10^{23}$\\,cm$^{-2}$, intermediary between the two classes.  This result is very nice in that it provides a good diagnostic tool for observations with few counts, but requires data above 15\\,keV.  In addition to the flux and column density measurements, we have measured power law indices, blackbody components (where present), and Fe K equivalent widths (with the physical width of the line fixed to 0.01\\,keV at 6.4\\,keV).  From our sample, we found no correlation between the hard band (2 -- 10\\,keV) luminosity and power law indices.  An important point to note is that the measured power law index for the complex spectra depends very much on the model used.  Comparing the results of the partial covering model with the double power law model (Tables~\\ref{tbl-7} and~\\ref{tbl-8}),  the average power law index for the partial covering model is significantly higher ($<\\Gamma> = 1.74$ compared to $<\\Gamma_2> = 1.36$) with smaller associated error bars. For the remaining 9 sources, the sources with spectra modeled by absorbed simple power law or power law and blackbody models, $<\\Gamma> = 1.75$, similar to the results from the partial covering model.  The values for the simple model/low absorption sources and those from the partial covering model are consistent with average photon indices for AGN ($\\approx 1.8$ from \\citet{mus82}).  Soft excesses, modeled with a blackbody component, were statistically significant in half of the spectra modeled by a simple power law model.  We find a significantly smaller fraction compared to the ROSAT sample of \\citet{gal07}, who find soft excesses in all of their sources, and a significantly larger fraction compared to the Lockman Hole \\xmm survey \\citep{mat05}, where only 11\\% of type 1 and 25\\% of type 2 sources show a soft excess.  Due to the low number of counts for our heavily absorbed sources, we can not quantify with certainty how many complex/heavily absorbed sources have this component, but at least one source (NGC 1142) has a statistically significant soft excess.  For the low absorption sources, $<kT> = 0.08$\\,keV, which is similar to but slightly lower than that seen for PG selected QSOs ($<kT_{BB}> = 0.14 \\pm 0.02$\\,keV) \\citep{por04, pic05}.  If the soft excess is the result of a thermal process, the lower kT values in our sample could be related to the lower luminosities of our sample, compared to the PG QSOs.  In fact, our average black body temperature is directly in the range of those found for type 1 AGN in the Lockman Hole sample, $<kT> = 0.09\\pm0.01$\\,keV \\citep{mat05}.  The final spectral component measured for our entire sample is the Fe K equivalent width (EW) at 6.4\\,keV. In Figure~\\ref{fig-iron}, we plot the Fe K EW versus the hard band (2 -- 10\\,keV) luminosity. We fit a line with the ordinary least squares bisector method to the upper limits of the EW measurements (see plot), yielding a fit of $\\log EW = (-0.697 \\pm 0.144) \\times \\log L_{2 - 10 kev} + (32.045 \\pm 6.164)$. The significance of this fit, indicated by R$^2 = 0.22$, where R$^2$ is the coefficient of determination, is very low.  Thus, our results show no indication of the X-ray Baldwin or IT effect \\citep{iwa93}, an anti-correlation of Fe K EW and hard band luminosity.  For the low absorption/simple model sources, we measured the significance of the \\ion{O}{7} and \\ion{O}{8} K edges to search for evidence of a warm absorber (see Table~\\ref{tbl-11}).  All of the simple model sources, with the possible exception of WKK 1263, are classified optically as Seyfert 1s and thus can be directly compared with the \\citet{rey97} sample.  In Figure~\\ref{fig-rey}, we plot the values of optical depth for each of the edges versus L$_{2-10 keV}$ for our sources as well as the \\citet{rey97} sources with luminosities in the same range.  As the figure shows, the optical depths we found for our sources are much lower than those from the Reynolds sample.  Only one source, ESO 490-G026, had a clear detection ($\\Delta\\chi^2 = 25$) with the optical depths of both edges having upper limits above 0.10.  Thus, where half of the Reynolds sample and 13/18 of the \\citet{geo98} sample showed evidence of a warm absorber in the line of sight, we find only 1/8 of our Seyfert 1 sources to show significant evidence of a warm absorber.  This result could be due to an incomplete or biased sample of Seyfert 1 sources, since these were among the low absorption sources with the lowest BAT flux in the 9-month sample.  Alternatively, our result could be representative of the entire BAT sample.  In this case, it is possible that the previous AGN samples showed more detections because they were from an optically selected/soft X-ray selected sample.  Thus, a possibility is that the emission that ionizes the gas, creates a region of warm, ionized gas, that also destroys dust.  An optical or softer X-ray survey could preferentially select these sources, missing more obscured sources.  Analysis of the remaining BAT sources will allow us to verify whether our result of few warm absorbers is consistent with  the properties of the entire 9-month catalog.   In addition to the spectral properties, we examined our sources for two types of variability: (1) during the \\xmm observations and (2) variability in spectral form and brightness between the \\xmm and XRT observations.  For the first type of variability, we created binned light curves for each object in our sample.  We found that 3/21 sources showed significant variability, with rates varying by 2.4 -- 2.8\\,ct\\,s$^{-1}$ over time scales of 0.8 -- 2.7\\,hr, in the PN observations.  The sources that varied the most were all Seyfert 1 sources with low absorption and X-ray spectra well-fit by simple power law or power law and a blackbody models.  These sources were among the brightest in our sample.  While only three sources showed short term variability, during the \\xmm observation which lasted $\\approx 10$\\,ks, nearly all of the sources (13/16) exhibited variability on longer time scales, of hours to months, from comparisons of the \\xmm and XRT observations.  This result agrees with earlier studies which found AGN more variable on long time scales than short time scales \\citep{bar86, nan97}.  From our comparison of the \\xmm and XRT spectra, we found 13/16 sources had varying fluxes, 6/16 had varying column densities, and 6/16 sources had varying power law indices.  Sources tended to vary more in the soft band than the hard band (see the Variability section).  Unfortunately, due to lower count rates, measuring variability for the heavily absorbed sources was more uncertain, particularly in the soft band (0.5 -- 2\\,keV).  In addition, we had less XRT observations with $> 100$\\,counts for these sources.  Based on the result that our sources varied more in the soft than hard band, it is likely that if we had more counts in the soft band for the heavily absorbed sources, as well as more observations for comparison, our results would agree with other AGN variability studies which found $ > 90$\\% of their sources to vary over the time scale of months to years (such as the AGN from the Lockman Hole \\citep{mat07} and Chandra Deep Fields \\citep{bau03, pao04} studies).  In Figure~\\ref{fig-varavg}, we plot column density versus flux and photon index versus flux for the sources that vary.  The plots show the observed parameter for each observation/average parameter for the source, where the observed values for each source are plotted with a different symbol.  In the column density figure, there is clearly no correlation seen between the column density and flux.  Similar results were found by \\citet{ris02} for a sample of Seyfert 2 galaxies, where they conclude that the variations in column density can not be caused by varying ionization states but by a clumpy absorber.  In the plot of spectral indices, however, we do find a correlation between the spectral index and the flux.  Therefore, we find that higher fluxes correspond to higher spectral indices.  This result has been seen for individual sources \\citep{mus93}.  Based on a variability study of the AGN sources in the Lockman Hole, \\citet{mat07} find no correlation between spectral variability and flux variability.  Further, they find flux variability much more prevalent in their sample than spectral variability, finding spectral variability in only $14 \\pm 8$\\% of Seyfert 1s and $34 \\pm 14$\\% of Seyfert 2s.   However, \\citet{mat07} note that the detection of spectral variability is related to the quality of the spectrum.  When they consider this factor, they predict a higher fraction of $\\approx 40$\\% to exhibit spectral variability.  We do not have a complete sample for Seyfert 2 sources, due to the low number of counts in the XRT observations for heavily absorbed sources.  However, we find 7/8 classified Seyfert 1 sources to exhibit a variation in either column density or power law index, much higher than the \\citet{mat07} estimated value.  Additionally, as we stated earlier, there is a clear correlation between changing flux and power law index for individual sources.  For the Compton-thick sources, variability or a lack thereof, gives clues to the size and location of the Compton-thick gas.  For our heavily obscured Compton-thick candidates, only MRK 417 had enough counts in an XRT observation ($> 100$\\,counts) to test for long term variability.  We found no statistically significant evidence of variability for this source between the \\xmm and XRT observations, taken 6 months apart.  This lack of variability in two observations does not give us much information.  For the low absorption Compton-thick candidates, however, we find a great deal of variability between the \\xmm and XRT observations.  In particular, for NGC 6860, the flux and spectral index are higher in the XRT observations, while the column density is lower.  Since the XRT observation was taken 4 months earlier, this puts a limit on the suggested change from Compton-thin to a reflection-dominated spectrum.  Similarly, significant changes are seen between the XRT and \\xmm observations of ESO 362-G018.  In this source, the most significant change is the disappearance of the strong Fe K line seen in the \\xmm observation. A smaller time constraint is placed on this source, 2 months between the last XRT and the \\xmm observation, for a change from a Compton-thin to a reflection-dominated spectrum.  Changes from Compton-thin to Compton-thick spectra have been noted before, particularly by \\citet{mat03}.  They discuss two possible scenarios to explain the changes in spectra, a change in column density of the absorber \\citep{ris02} and a ``switched-off'' source, that is, a state where the emission from the central source drastically decreases below our detection threshold.  For both ESO 362-G018 and NGC 6860, Seyfert 1.5 sources embedded in dusty host galaxies, a changing absorber is a more appealing explanation.  %\\subsection{Compton Thick Sources} %-- From this point, comment on Compton thick sources.  Compare with Bassani et al. 99 ?"
    },
    "0711/0711.0966_arXiv.txt": {
        "abstract": "Although at least one quarter of early-type barred galaxies host secondary stellar bars embedded in their large-scale primary counterparts, the dynamics of such double barred galaxies are still not well understood.  Recently we reported success at simulating such systems in a repeatable way in collisionless systems.  In order to further our understanding of double-barred galaxies, here we characterize the density and kinematics of the $N$-body simulations of these galaxies.  This will facilitate comparison with observations and lead to a better understanding of the observed double-barred galaxies. We find the shape and size of our simulated secondary bars are quite reasonable compared to the observed ones. We demonstrate that an authentic decoupled secondary bar may produce only a weak twist of the kinematic minor axis in the stellar velocity field, due to the relatively large random motion of stars in the central region. We also find that the edge-on nuclear bars are probably not related to boxy peanut-shaped bulges which are most likely to be edge-on primary large-scale bars. Finally we demonstrate that the non-rigid rotation of the secondary bar causes its pattern speed not to be derived with great accuracy using the Tremaine-Weinberg method.  We also compare with observations of NGC 2950, a prototypical double-barred early-type galaxy, which suggest that the nuclear bar may be rotating in the opposite sense as the primary. ",
        "introduction": "\\label{sec:intro}  Double barred (S2B) galaxies were first described over thirty years ago \\citep{devauc_75}.  The {\\it Hubble Space Telescope} has revealed secondary bars at the center of at least one quarter of early-type optically-barred galaxies \\citep{erw_spa_02}.  Dynamically decoupled\\footnote{In this context, by decoupled we mean only that $\\omg{s} \\neq \\omg{p}$, where \\omg{s} (\\omg{p}) is the pattern speed of the secondary (primary) bar.}  secondary bars in S2B galaxies have been hypothesized to be a mechanism for driving gas past the inner Lindblad resonance (ILR) of primary bars, to feed the supermassive black holes that power active galactic nuclei \\citep{shl_etal_89}.  The dynamics of secondary bars are still not well understood. The random apparent relative orientations of primary and secondary bars in nearly face-on galaxies points to dynamical decoupling \\citep{but_cro_93, fri_mar_93}.  But images alone cannot reveal much about how the two bars rotate through each other.  Kinematic evidence of decoupling, using either gas or stars, is harder to obtain \\citep{pet_wil_02, sch_etal_2002,moi_etal_04}.  Indirect evidence for decoupling was claimed by \\citet{ems_etal_01} based on rotation velocity peaks inside the secondary bars.  Conclusive direct kinematic evidence for a decoupled secondary bar was obtained for NGC 2950 by \\citet[][hereafter CDA03]{cor_deb_agu_03} who showed, using the method of \\citet{tre_wei_84}, that the primary and secondary bars cannot be rotating at the same pattern speed.  \\begin{figure*}[!ht] \\centerline{ \\includegraphics[angle=-90.,width=0.345\\hsize]{fig1aa.ps} \\includegraphics[angle=-90.,width=0.325\\hsize]{fig1ab.ps} \\includegraphics[angle=-90.,width=0.325\\hsize]{fig1ac.ps}} \\vspace{0.01\\hsize} \\centerline{ \\includegraphics[angle=-90.,width=0.345\\hsize]{fig1ba.ps} \\includegraphics[angle=-90.,width=0.325\\hsize]{fig1bb.ps} \\includegraphics[angle=-90.,width=0.325\\hsize]{fig1bc.ps}} \\vspace{0.01\\hsize} \\centerline{ \\includegraphics[angle=-90.,width=0.345\\hsize]{fig1ca.ps} \\includegraphics[angle=-90.,width=0.325\\hsize]{fig1cb.ps} \\includegraphics[angle=-90.,width=0.325\\hsize]{fig1cc.ps}}  \\caption{3-D surface density of 3 runs. In each row, from left to right, panels are for the large-scale total stellar distribution, central zoom-in of total stellar distribution, separate disk (dark) and bulge (red) surface density contours, respectively.  (a). {\\it top row}: the canonical run with a secondary bar formed (Run D); (b). {\\it middle row}: the run with an unrotating bulge, which formed only a single bar (run S); (c). {\\it bottom row}: the run without a bulge (run NB). The contours in $x$-$y$ plane are separated by half dex, while in vertical projections they are spaced more sparsely to avoid contour over-crowding.} \\label{fig:3dshape} \\end{figure*}  Simulations offer the best way to understand double barred systems. However, the decoupled nuclear bars that formed in early simulations did not last long.  For example, the most long-lived nuclear bar in \\citet{fri_mar_93} lasted for less than two turns of the primary bar, corresponding to about 0.4 Gyr, which is far too short to explain the observed abundance of nested bars.  Furthermore, their models usually require substantial amounts of gas to form and maintain these nuclear bars. \\citet{hel_etal_07, hel_etal_07_2} reported that nested bars form in a quasi-cosmological setting, but the amplitudes of the bars also seem to weaken rapidly after most of gas has formed stars \\citep[][Figure~2]{hel_etal_07}. \\citet{pet_wil_04} found that 4 out of 10 double-barred galaxies contain very little molecular gas in the nuclear region. These clues suggest that large amounts of molecular gas may not be necessary to maintain central nuclear bars. \\citet{rau_etal_02} reported that a secondary bar forms in a collisionless $N$-body simulation, although their secondary bar had a ``vaguely spiral shape.''  On the side of orbital studies, \\citet{mac_spa_97,mac_spa_00} discovered a family of loop orbits that may form building blocks of long-lived nuclear stellar bars (also \\citealt{mac_ath_07}). Their studies are very important for the understanding of double barred galaxies, but their models are not fully self-consistent, since nested bars in general cannot rotate rigidly through each other \\citep{lou_ger_88}. So fully self-consistent $N$-body simulations are still needed to check if their main results still hold when the non-rigid nature of the bars is taken into account.  Recently \\citet[hereafter DS07]{deb_she_07} demonstrated that long-lived secondary bars can form in purely collisionless $N$-body simulations, when a rotating pseudobulge is introduced in their model. The nuclear bars in their work are distinctly bars, and do not have a spiral shape. They showed that the behavior of their models were in good agreement with the loop orbit predictions of \\citet{mac_spa_00}.  In this report we analyze the photometrical and kinematical properties of high resolution models in detail. Our theoretical results here can also be compared to the observed 2-D kinematics of some double-barred galaxies, to achieve a better understanding of the dynamics of the secondary bars. ",
        "conclusions": "We have analyzed the photometrical and kinematical properties of our high resolution models, and contrasted them when with or without a secondary bar. This study also compared the simulated secondary bars with observations.   In general the shape of secondary bars in our models is reasonable compared to observed ones. The length ratio of two bars, determined by various methods, is in the range of 0.12 to 0.19, in good agreement with \\citet{erw_spa_02, erw_spa_03}. We also found the overall edge-on shape of boxy bulges is largely unaffected by the existence of a secondary bar. At lower inclinations, the central density profile along the primary bar major axis is lower than that along the minor axis, due to the secondary bar orienting to a different direction.  The primary extends roughly to its corotation radius, and therefore fits the definition of a fast bar (see for example \\citealt{agu_etal_03}). Although the secondary bar rotates more rapidly than the primary, its semi-major axis is much shorter than its corotation radius, even if we take the oscillation of the bar patterns speeds into account. We did not find evidence of CR-ILR coupling \\citep[e.g.][]{pfe_nor_90,fri_mar_93} in our models.  We find that the central twist of kinematic axes is quite weak even if a secondary bar is present, due to the relatively large velocity dispersion of stars in the central region. This is consistent with the 2-D stellar kinematics of secondary bars studied in \\citet{moi_etal_04}. There are no clear $h_4$ signatures associated with the presence of secondary bars. A $h_3$ reversal feature may appear in the nuclear region at some favorable orientations. We do not find a $\\sigma$ (velocity dispersion) drop for our secondary bar model. It is more likely that $\\sigma$-drops are just the signature of newly-formed stars, and it is not necessarily a unique feature of double-barred systems.  We showed that the Tremaine-Weinberg method is not very reliable even when the primary bar contribution is fully excluded.  The way in which the measurement fails is consistent with the proposal of CDA03, namely that the non-rigid rotation leads to internal motions that violate the stationary frame assumption of the method. Nonetheless, we find no example in our simulations where the behavior of the TW integrals mimics that observed in NGC 2950. Thus this galaxy may indeed have counter-rotating secondary and primary bars.  The general agreement between our simulations and observations of double barred galaxies gives us confidence that the simulations are capturing the same dynamics as in nature.  This is especially remarkable because secondary bars are not merely scaled down versions of primary bars, but have distinctly different kinematic properties. In the absence of self-consistent simulations, earlier orbit-based models could not directly confront the challenge from observations which found such differences.  This demonstrates the advantage of finally being able to simulate stellar double-barred galaxies, which had been puzzling for so long."
    },
    "0711/0711.2222_arXiv.txt": {
        "abstract": "We investigate the clustering of afterglow light curves observed at X-ray and  optical wavelengths. We have constructed a sample of 61 bursts with known distance and X-ray afterglow. This sample includes bursts observed by BeppoSAX, XMM-Newton, Chandra, and SWIFT. We correct the light curves for cosmological effects and compare the observed X-ray fluxes one day after the burst. We check for correlations between the observed flux and the burst spectral and temporal properties. We confirm the previous result of Bo\\\"er and Gendre (2000) that X-ray afterglow light curves cluster in luminosity, even when we consider the last SWIFT data. We observe this clustering only for the afterglow light curves; the inclusion of prompt-related data broaden the distribution. A similar clustering is observed for the optical light curves; GRB sources can be divided in three classes, namely optical and X-ray {\\it bright} afterglows, optical and X-ray {\\it dim} ones, and optically {\\it bright} -X-ray {\\it dim} ones. We argue that this clustering is related to the fireball total energy, the external medium density, the fraction of fireball energy going in relativistic electrons and magnetic fields. These parameters can be either  fixed to a standard value, or correlated.  We finally propose a method for the estimation of the GRB source redshift based on the observed X-ray flux one day after the burst and optical properties. Using this method, we compute a redshift of 1.4 $\\pm$ 0.2 for \\object{GRB 980519} and of $1.9 \\pm 0.3$ for \\object{GRB 040827}. We tested this method on three recently detected SWIFT GRBs with known redshift, and found it in good agreement with the reported distance from optical spectroscopy. ",
        "introduction": "Long Gamma-Ray Bursts (GRBs) are linked with the death of massive stars \\citep[for a review, see][]{mes06}. Their association with supernovae \\citep[e.g.][]{hjo03, sta03} and the fact that these events are at cosmological distance \\citep{met97} make them interesting for studies of cosmology in the redshift range 1-15. However, while the GRB features a prompt emission, usually seen in gamma-ray only, and an afterglow seen at all wavelengths, only the former emission has been considered for cosmological studies yet. One of the first attempt to do so, based on the $E_p-E_{iso}$ correlation found by \\citet{ama02}, was done by \\citet{ghi03}, who constrained the $\\Omega_m - \\Omega_\\Lambda$ parameters. While not very constraining, their findings were compatible with previous tests made with supernova samples. The use of GRB for cosmological studies, specially at high redshift where their detection at optical/IR wavelength is difficult, needs to build a robust indicator of their distance, whenever possible based on their intrinsic properties.  The afterglow emission is less studied from a cosmological point of view because of its diversity. However, there were hints of standardization of the X-ray afterglows. The first attempt was done by \\citet{boe00}, who found evidences for clustering in the X-ray light curves of BeppoSAX afterglows. This clustering was confirmed later by \\citet{gen05} who extended the sample to the XMM-Newton and Chandra data. In the following we will refer to these articles respectively as paper I and II. In addition, \\citet{kou04} showed that supernova and GRB light curves had similar behaviors and were converging with time toward a similar luminosity. In paper I we tried also to check if optical light curves were also   clustered: this attempt failed due to the poor knowledge of the intrinsic absorption in the burst host galaxy at that time. This study was completed by \\citet{nar06} and \\citet{lia06} who found independently that optical afterglows were also clustered in luminosity.  The high detection rate and throughput of SWIFT, which provide a large sample of X-ray observations and rapid, accurate localization, enabling a redshift estimation by optical/IR telescopes, allowed us to increase dramatically our sample. With this larger sample, we have tried to derive a method for estimating the burst redshift from the X-ray light curve \\citep{gen06}. However, before to apply this method it is necessary to understand the nature of these two groups, and to determine a consistent way to derive to which group belongs each burst. This is the purpose of this article. In Section \\ref{sec_ana} we present our sample and the data analysis we performed. We discuss the X-ray clustering in Section \\ref{sec_result}. We compare the X-ray and optical clusterings in Section \\ref{sec_discu}. In section \\ref{sec_redshift} we present our method of GRB source redshift estimation from the X-ray afterglow light curve. We finally test this method on several GRBs detected recently by SWIFT, and we propose an estimation of the redshift for two previously detected GRBs of unknown distance. ",
        "conclusions": "We have investigated the clustering of afterglow light curves observed previously (papers I and II) in X-ray and in optical from BeppoSAX, XMM-Newton and Chandra data, using the newly available, well-monitored SWIFT light curves. Adding SWIFT bursts to the previous sample reported in paper II, we still confirm our previous findings. On a sample of 61 events the X-ray light curves cluster in two groups, with a significance larger than 6 $\\sigma$. \\citet{wil07} describe the X-ray light curve as a sum of two components, the first one being related to the prompt emission: we found that it is the second component that clusters in luminosity. This finding is also supported by the fact that the clustering is not related to the properties of the prompt emission.  A similar clustering was observed also at optical wavelengths. We compared the classification within each group in X-ray and in optical, and found three classes: {\\it bright} optical and X-ray afterglows, {\\it dim} ones, and optically {\\it bright}-X-ray {\\it dim} ones. We propose that this clustering is related to the fireball properties and the surrounding medium density, which can be either constant or correlated. We used the mean X-ray and optical fluxes of each group to put constraints on the values of $\\epsilon_B$ and $\\epsilon_{\\bf e}$, and expressed them in terms of the total fireball energy and the surrounding medium density. We stress that the extension of this work at low frequency (radio, sub-millimeter and infrared) may help solving the exact origin of the clustering by strongly constraining the medium density and the position of the synchrotron self absorption and injection frequencies.  Using the observed X-ray afterglow properties of GRBs, we propose a new, simple, method for the determination of the source redshift based on X-ray data; optical photometry in U and V bands may help to clear the degeneracy between the two estimates found. This method has been established on a sample of bursts with known redshift, and we find an excellent correlation. We apply this method to a sample of GRB source of unknown redshift. We propose an estimation of the redshift for \\object{GRB 980519} (1.4 $\\pm$ 0.2) and for \\object{GRB 040827} ($1.9 \\pm 0.3$).  While we were building this distance indicator and writing this paper, several GRB were detected by SWIFT and their distance derived using optical spectroscopy. We thus have a sample of bursts which allows to test this correlation independently. To simulate actual conditions for an unknown event we used the SWIFT count rate light curve available from their web site \\citep{eva07}, and we apply a mean spectral index of 1.2 for the count-to-flux conversion. The results are indicated in Table \\ref{table_sup}, indicating an agreement between 10.3 and 24.5 \\% \\citep[in accordance with the error estimated in \\ref{toto}, we estimate far superior than other methods e.g.][]{att05}."
    },
    "0711/0711.2825_arXiv.txt": {
        "abstract": "In order to take full advantage of the long time series collected by the GONG and MDI helioseismic projects, we present here an adaptation of the rotation-corrected $m$-averaged spectrum technique in order to observe low radial-order solar p modes. Modeled profiles of the solar rotation demonstrated the potential advantage of such a technique \\cite{brown85, schou04, app00}. Here we develop a new analysis procedure which finds the best estimates of the shift of each $m$ of a given ($n,\\ell$) multiplet, commonly expressed as an expansion in a set of orthogonal polynomials, which yield the narrowest mode in the $m$-averaged spectrum. We apply the technique to the GONG data for modes with $1 \\leq \\ell \\leq 25$ and show that it allows us to measure lower-frequency modes than with classic peak-fitting analysis of the individual-$m$ spectra. ",
        "introduction": "In the search for long-lived, low-order p modes, the classic methods of peak-fitting the individual-$m$ spectra fail because of the decrease in mode amplitudes and increase in solar convective motions (solar noise) as the frequency decreases, reducing the signal-to-noise ratio (hereafter SNR) of those modes. Instead, various pattern-recognition techniques have been developed in an effort to reveal the presence of modes in the low-frequency range. In the case of spatially-resolved helioseismic data (for example observations collected by the space-based SOHO/MDI instrument and the ground-based multi-site GONG network), $m$-averaged spectra appeared to be a powerful tool. They were employed early in the development of helioseismology \\cite{brown85}, but have been replaced by fitting the $m$-spectra individually as the quality of the data improved. Years later, in order to fully take advantage of the long-duration helioseismic MDI and GONG instruments, $m$-averaged spectra corrected by the modeled solar rotation were used to detect new low radial-order p modes and to put some upper limits on the detectability of the g modes \\cite{schou04,app00}. These authors demonstrated the potential advantage of such rotation-corrected $m$-averaged spectra. We present here an adaptation of the $m$-averaged spectrum technique in which the $m$-dependent shift parameters are determined. ",
        "conclusions": "We developed a new method to measure low-frequency solar p modes in spatially-resolved helioseismic observations. The $m$-averaged spectrum technique, in which we determine the correction for the splitting coefficients (i.e., $a$-coefficients), allows us to measure lower frequencies in GONG data than with classic peak-fitting methods of the individual-$m$ spectra, which are limited by the low signal-to-noise ratio of these low-radial order modes. We are currently in the process of generalizing the method with a larger number of $a$-coefficients, both odd and even orders. We are also going to apply this technique to the MDI data. The lower level of background noise in the MDI data should allow us to measure even lower low-frequency modes than with GONG observations.     \\ack This work utilizes data obtained by the Global Oscillation Network Group (GONG) program, managed by the National Solar Observatory, which is operated by AURA, Inc. under a cooperative agreement with the National Science Foundation. The data were acquired by instruments operated by the Big Bear Solar Observatory, High Altitude Observatory, Learmonth Solar Observatory, Udaipur Solar Observatory, Instituto de Astrof\\'{\\i}sica de Canarias, and Cerro Tololo Interamerican Observatory. This work has been supported by the NASA SEC GIP grant NAG5-11703."
    },
    "0711/0711.1637_arXiv.txt": {
        "abstract": "We combine deep $UBVRIzJK$ photometry from the MUSYC survey with redshifts from the COMBO-17 survey to study the \\emph{rest-frame} ultraviolet ($UV$) properties of 674 high-redshift ($0.5<z<1$) early-type galaxies, drawn from the Extended Chandra Deep Field South (E-CDFS). Galaxy morphologies are determined through visual inspection of Hubble Space Telescope (HST) images taken from the GEMS survey. We harness the sensitivity of the $UV$ to young ($<1$ Gyrs old) stars to quantify the \\emph{recent} star formation history of the early-type population. We find compelling evidence that early-types of all luminosities form stars over the lifetime of the Universe, although the bulk of their star formation is already complete at high redshift. Luminous ($-23<M(V)<-20.5$) early-types form 10-15 percent of their mass after $z=1$, while their less luminous ($M(V)>-20.5$) counterparts form 30-60 percent of their mass in the same redshift range. ",
        "introduction": "The vast majority of work on early-type galaxies in the past has focussed on optical spectro-photometric data. A significant drawback of optical photometry is its lack of sensitivity to moderate amounts of recent star formation (RSF). While red optical colours imply a high-redshift formation epoch for the \\emph{bulk} of the stellar population in early-type galaxies, the optical spectrum remains largely unaffected by the minority of stellar mass that forms in these systems at low and intermediate redshift. As a result it is virtually impossible to quantify early-type star formation histories (SFHs) over the last 8 billion years ($0<z<1$).  RSF can be efficiently traced using the rest-frame ultraviolet ($UV$) spectrum, which is sensitive to young, massive main sequence stars with ages less than $\\sim$1 Gyr. Using a large sample of early-type galaxies detected by the GALEX $UV$ space telescope, Kaviraj et al. (2006) have recently shown that, contrary to the expectations of traditional `monolithic collapse' models \\citep[e.g.][]{BLE92}, local early-types show widespread recent star formation (RSF) - at least 30 percent show blue $UV$-optical colours indicative of unambiguous RSF in these systems. The work presented here extends these results to high redshift by exploiting deep optical photometry to trace the rest-frame $UV$ properties of early-type galaxies in the redshift range $0.5<z<1$. ",
        "conclusions": "\\begin{figure} $\\begin{array}{cc} \\includegraphics[width=3.3in]{Kaviraj1} & \\includegraphics[width=3.3in]{Kaviraj2} \\end{array}$ \\caption{LEFT-HAND PANEL: Rest-frame $(NUV-r)$ colour-magniutude relation of the E-CDFS galaxy population. RIGHT-HAND PANEL: Rest-frame $(NUV-r)$ colour plotted against the RSF in the E-CDFS galaxy population. The $NUV$ ($\\sim2300\\AA$) and $r$ band filters are taken from the GALEX and SDSS filtersets respectively.} \\label{fig:discussion} \\end{figure}  The early-type population as a whole exhibits a typical RSF between 5 and 13 percent in the redshift range $0.5<z<1$, while the early-types on the broad `red sequence', ($NUV-r>4$), typically show RSF values less than 5 percent. The reddest early-types (which are also the most luminous) are virtually quiescent with RSF values of $\\sim1$ percent. Since the timescale of this study is $\\sim2.5$ Gyrs, a simple extrapolation (from RSF values in Figure 1) indicates that luminous ($-23<M(V)<-20.5$) early-type galaxies typically form up to 10-15 percent of their mass after $z=1$ (with a tail to higher values), while less luminous early-types ($M(V)>-20.5$) form 30-60 percent of their mass after $z=1$.   \\nocite{Kaviraj2006b}"
    },
    "0711/0711.2772_arXiv.txt": {
        "abstract": "The shock-acceleration theory predicts a power-law energy spectrum in the test particle approximation, and there are two ways to calculate a power-law index, Peacock's approximation and Vietri's formulation. In Peacock's approximation, it is assumed that particles cross a shock front many times and energy-gains for each step are fully uncorrelated. On the other hand, correlation of the distribution of an energy-gain factor for a particle is considered in Vietri's formulation. We examine how Peacock's approximation differs from Vietri's formulation. It is useful to know when we can use Peacock's approximation because Peacock's approximation is simple to derive the power-law index. In addition, we focus on how the variance of the energy-gain factor has an influence on the difference between Vietri's formulation and Peacock's approximation. The effect of the variance has not been examined well until now. For demonstration, we consider two cases for the scattering in the upstream: the large-angle scattering (model A) and the regular deflection by large-scale magnetic fields (model B). Especially there is no correlation among the distribution of an energy-gain factor for every step in model A. In this model, we see the power-law index derived from Peacock's approximation differs from the one derived from Vietri's formulation when we consider the mildly-relativistic shock, and the variance of the energy-gain factor affects this difference. We can use Peacock's approximation for a non-relativistic shock and a highly-relativistic shock because the effect of the variance is hidden. In model B, we see the difference of the power-law converging along the shock velocity. ",
        "introduction": "Galactic cosmic rays with energies $E \\le 10^{15}$ eV, below the so called 'knee' in the cosmic ray spectrum, are thought to originate from shocks of supernova remnants. The shock-acceleration theory was proposed by various authors independently to explain the origin of Galactic cosmic rays. \\cite{b2}, \\citet{b14}, \\citet{b4} and \\citet{b5} studied the shock-acceleration for non-relativistic shocks. Their theory relies on multiple scattering of charged particles which results in round-trips between the upstream and the downstream of a shock. In the test particle approximation, shock-acceleration theory predicts that a shock produces energetic particles whose energy spectrum obeys a power-law which explains the observation of the high energy cosmic rays well. A power-law is obtained by two ways. 1) Solving the convection-diffusion transport equation with appropriate boundary conditions. 2) Considering an interaction of test particles with a shock front from the kinetic theory viewpoint. These independent two theories predict the same result in spite of different approaches.  Cosmic rays with energies $E \\sim 10^{18.5}-10^{20.5}$ eV are generally believed to come from extragalactic origins, based on their harder spectrum, isotropic arrival directions on the sky and the fact that they are not confined by the Galactic magnetic fields. Some of authors consider cosmic ray production for the relativistic shocks (\\citet{b21}; \\citet{b19}; \\citet{b15}). The relativistic shock-acceleration is first considered by \\citet{b17} considering relativistic kinematics based on the Bell's theory taking account of the beaming effect. Later, \\citet{b12},\\citet{b13} solved the transport equation with appropriate collision operators. They focussed on the diffusion approximation.  On the other hand, \\citet{b9} and \\citet{b20} formulated the shock-acceleration for the arbitrary shock speeds without the diffusion approximation. Their theories include the non-diffusive effect which is important for relativistic shocks and reproduce the same result in the non-relativistic limit. Both of them are constructed in two parts. First, they calculated $ P_u(\\mu_0,\\mu)$ ($ P_d(\\mu_0,\\mu)$ ) which denotes the conditional probability that a particle entering into the upstream (downstream) along a direction $\\mu_0$ will leave toward the downstream (upstream) along a direction $\\mu$. They used some phenominological scattering models to calculate $P_u$ and $P_d$. Second, they calculated a power-law index using different ways. \\citet{b17} calculated a power-law index assuming that particles cross a shock front many times and an energy-gain of each step is fully uncorrelated (Peacock's approximation). \\citet{b9} used Peacock's approximation. On the other hand, \\citet{b20} solved the transport equation exactly and calculated a power-law index (Vietri's formulation). He considered correlation among energy-gains of a particle. He used his original formulation to calculate a power-law index. Vietri's formulation gives the different result from Peacock's approximation in general but both approaches give the same result in the Newtonian limit (see also \\citet{b6}).  To calculate a power-law index, we have to calculate $P_u$ and $P_d$ using some scattering models. \\citet{b9} considered the large-angle scattering in both the upstream and the downstream to calculate $P_u$ and $P_d$ (Model A). The large-angle scattering model mimics the scattering in strongly turbulent fields. It is often supposed that the turbulence can be strong in astrophysical shock environments (see \\citet{b7}, and references therein). In this model, the initial information of the distribution function is lost when particles are scattered. So an energy-gain factor of each step is uncorrelated. Moreover, \\citet{b8} pointed out necessity to consider the regular deflection as far as relativistic shocks are concerned (Model B). It was understood that return of the particles to a shock surface from the upstream region can be warranted even in the absence of scattering, provided background magnetic fields are at an angle with the shock normal. This is due to the fact particles return to the shock surface from the upstream by the regular deflection before scattering occurs.  In this paper, we examine how the variance of the energy-gain factor influences the difference between Vietri's formulation and Peacock's approximation. It is useful to know when we can use Peacock's approximation because Peacock's approximation is simple to derive a power-law index. It is also important to understand Vietri's formulation and Peacock's approximation in detail because we can examine validity of the works in which Peacock's approximation is used. First, we calculate $P_u$ and $P_d$ in model A (where the large-angle scattering occurs at both sides of a shock) which is suitable for examining the effect of the variance of the energy-gain factor because there is no effect of the correlation in this model. Then we calculate a power-law index using both Peacock's approximation and Vietri's formulation. Next, we consider model B where the large-angle scattering occurs in the downstream and a particle is deflected by large-scale magnetic fields in the upstream. We calculate probability functions and a power-law index using Peacock's approximation and Vietri's formulation in the same way as the case of model A. Finally, we examine how the power-law index derived from Peacock's approximation differs from the one derived from Vietri's formulation to examine the effect of the variance. In model A, we see that the variance affects the difference between Peacock's approximation and Vietri's formulation, and explain what the effect of the variance means physically. To examine the above fact, we show the power-law index changing the shock velocity from a non-relativistic one to a highly-relativistic one. \\citet{b16} calculated a power-law index with $0.04\\leq\\Gamma_s\\beta_s\\leq10$ considering the regular deflection, where $\\beta_s$ is the shock velocity and $\\Gamma_s$ is the Lorentz factor of the shock velocity . We extend this calculation to the highly-relativistic range.  The plan of this paper is as follows. In section 2 we briefly summarize the theoretical framework introduced in \\citet{b9}. We also review Peacock's approximation and Vietri's formulation. In section 3, first, we consider the large angle scattering in both the upstream and the downstream. Next, we consider the regular deflection in the upstream and the large angle scattering in the downstream. We also discuss what makes Peacock's approximation inadequate. Discussion and conclusion are presented in section 4. ",
        "conclusions": "In this paper, we examine how the power-law index derived from Peacock's approximation is different from the one derived from Vietri's formulation. \\citet{b17} gives the approximate power-law index and \\citet{b20} gives the exact power-law index. \\citet{b20} claimed that Peacock's approximation and Vietri's formulation give the different result in general, but they did not explain the reason in detail (see also \\citet{b6}). The effect of the variance has not been studied well until now and we have examined the effect of the variance.  First, we conclude the difference of the power-law index derived from Peacock's approximation and Vietri's formulation. We considered two cases for the scattering in the upstream to model multiple shock crossing: the large-angle scattering (model A) and the regular deflection by large-scale magnetic fields (model B). We consider only the large-angle scattering in the downstream. In model A, the power-law index derived from Peacock's approximation is different from the one derived from Vietri's formulation for a mildly-relativistic shock. The difference decreases as the shock velocity approaches the highly-relativistic velocity. In model B, the difference becomes larger along the shock velocity at first. Finally, the difference converges when the shock velocity is mildly-relativistic. From the above result, we can use Peacock's approximation for a highly-relativistic shock when we consider model A and the difference between Peacock's approximation and Vietri's formulation is important when we consider model B for a relativistic shock. Moreover, we see the reason why Peacock's approximation and Vietri's formulation give the different results. We conclude the variance in the distribution of the energy gain factor affects the difference between Peacock's approximation and Vietri's formulation in model A.  In Fig.~3, we show that the power-law index calculated by Vietri's formulation is harder than the one calculated by Peacock's approximation in both cases of model A and model B. This tendency can be understood as follows: Let us consider the number of particles with $E_k$, which have experienced $k$ cycles in an average (see Eq.~(20)). If there is the variance of energy-gain factor's distribution, there should be contribution to the number from those that have experienced less than $k$-cycles and more than $k$-cycles. It is apparent that the contribution from the former is larger than the latter, because the energy spectrum obeys a power-law with index larger than 3. As a result, the spectrum becomes hard when the effect of the distribution width is taken into account. As above, we can understand why the power-law index becomes hard as the variance becomes large. However, this interpretation can not be suitable when the variance is larger than some value. The effect of the variance converges as we show in Fig.~5. What's more, we can interpret that the effect of the distribution's variance becomes weak as the power-law index approaches 3 because the fraction of particles that are accelerated by the mean energy-gain factor increases.  Next, let us explain the general characteristics of the power-law index when we change the scattering model. We found the power-law index converges as the shock velocity increases in model B. This is the same characteristic as \\citet{b3} who studied about the case of the small-angle scattering, although the spectrum derived from the large-angle scattering is steeper than that of the small-angle scattering. On the other hand, the index decreases monotonically in model A. For the limit of non-relativistic shocks, the result is consistent with the diffusive shock-acceleration theory.  The difference of the power-law index between model A and model B is large for a relativistic shock. The difference is small if one uses the small-angle scattering model instead of the large-angle scattering (e.g., \\citet{b1}). This is because the averaged energy-gain factor is very large in model A. The energy-gain factor can be written as $G(\\mu_{out},\\mu)= \\left( \\frac{1-V_r\\mu_{out}}{1-V_r\\mu} \\right)$. Thus, the energy-gain factor can be large for a highly-relativistic shock in the case of the large-angle scattering since the difference between $\\mu_{out}$ and $\\mu$ is large in the case of the large-angle scattering. On the other hand, the energy-gain factor remains the order of unity even for a highly-relativistic shock both in the case of the small-angle scattering and the regular deflection.  Peacock's approximation and Vietri's formulation give the different results in general. To see how Peacock's approximation is good, of course it is also important to investigate how many times scatterings should occur enough to use the central limit theorem and whether there is correlation among energy-gains for every step. Such consideration about these effects should be done not only for relativistic shocks but also for non-relativistic shocks. For example, we probably can not use a central limit theorem even for non-relativistic shocks in low energy range.This consideration might be useful to discuss the boundary between the supra-thermal particle and the non-thermal particle. The variance might also be considered to calculate the power-law index as we show in this paper.  We have considered the case of the large-angle scattering in this paper. We will consider the case of the small-angle scattering in the future work."
    },
    "0711/0711.2919_arXiv.txt": {
        "abstract": "{Structural parameters of model star clusters are measured in radial profiles built from number-density, mass-density and surface-brightness distributions, assuming as well different photometric conditions.} {Determine how the core, half-star count and tidal radii, as well as the concentration parameter, all of which are derived from number-density profiles, relate to the equivalent radii measured in near-infrared surface-brightness and mass-density profiles. We also quantify changes in the resulting structural parameters due to depth-limited photometry.} {Star clusters of different ages, structure and mass functions are modelled by assuming that the radial distribution of stars follows a pre-defined analytical form. Near-infrared surface brightness and mass-density profiles result from mass-luminosity relations taken from a set of isochrones. Core, tidal and half-light, half-mass and half-star count radii, together with the concentration parameter, are measured in the three types of profiles, which are built under different photometric depths.} {While surface-brightness profiles are almost insensitive to photometric depth, radii measured in number-density and mass-density profiles change significantly with it. Compared to radii derived with deep photometry, shallow profiles result in lower values. This effect increases for younger ages. Radial profiles of clusters with a spatially-uniform mass function produce radii that do not depend on depth. With deep photometry, number-density profiles yield radii systematically larger than those derived from surface-brightness ones.} {In general, low-noise surface-brightness profiles result in uniform structural parameters that are essentially independent of photometric depth. For less-populous star clusters, those projected against dense fields and/or distant ones, which result in noisy surface-brightness profiles, this work provides a quantitative way to estimate the intrinsic radii by means of number-density profiles built with depth-limited photometry. } ",
        "introduction": "\\label{intro}  Star clusters are a powerful tool in the investigation of Galaxy structure and dynamics, star formation and evolution processes, and as observational constraints to N-body codes. This applies especially to the long-lived and populous globular clusters (GCs) that, because of their relatively compact nature, can be observed in most regions of the Galaxy, from near the center to the remote halo outskirts.  In general terms, the structure of most star clusters can be described by a rather dense core and a sparse halo, but with a broad range in the concentration level. In this context, the standard picture of a GC assumes a isothermal central region and a tidally truncated outer region (e.g. \\citealt{Binney1998}). Old GCs, in particular, can be virtually considered as dynamically relaxed systems (e.g. \\citealt{NoGe06}). During their lives clusters are continually affected by internal processes such as mass loss by stellar evolution, mass segregation and low-mass star evaporation, and external ones such as tidal stress and dynamical friction e.g. from the Galactic bulge, disk and giant molecular clouds (e.g. \\citealt{Khalisi07}; \\citealt{Lamers05}; \\citealt{GnOs97}). Over a Hubble time, these processes tend to decrease cluster mass, which may accelerate the core collapse phase for some clusters (\\citealt{DjMey94}, and references therein). Consequently, these processes, combined with the presence of a central black hole (in some cases) and physical conditions associated to the initial collapse, can affect the spatial distribution of light (or mass) both in the central region and at large radii (e.g. \\citealt{GLO99}; \\citealt{NoGe06}).  It is clear from the above that crucial information related to the early stages of Galaxy formation, and to the cluster dynamical evolution, may be imprinted in the present-day internal structure and large-scale spatial distribution of GCs (e.g. \\citealt{MvdB05}; \\citealt{GCProp}). To some extent, this reasoning can be extended to the open clusters (OCs), especially the young, which are important to determine the spiral arm and disk structures and the rotation curve of the Galaxy (e.g. \\citealt{Friel95}; \\citealt{DiskProp}). Consequently, the derivation of reliable structural parameters of star clusters, GCs in particular, is fundamental to better define their parameter space. This, in turn, may result in a deeper understanding of the formation and evolution processes of the star clusters themselves and the Galaxy.  Three different approaches have been used to derive structural parameters of star clusters. The more traditional one is based on the surface-brightness profile (SBP), which considers the spatial distribution of the brightness of the component stars, usually measured in circular rings around the cluster center. The compilation of Harris (1996, and the 2003 update\\footnote{\\em http://physun.physics.mcmaster.ca/Globular.html}) presents a basically uniform set of parameters for 150 Galactic GCs. Among their structural parameters, the core (\\rc), half-light (\\rhL) and tidal (\\rt) radii, as well as the concentration parameter $c=\\log(\\rt/\\rc)$, were based mostly on the SBP database of \\citet{TKD95}. SBPs do not necessarily require cluster distances to be known, since the physically relevant information contained in them is essentially related to the relative brightness of the member stars. In principle, it is easy to measure integrated light. However, SBPs are more efficient near the cluster center than in the outer parts, where noise and background starlight may be a major contributor. Another potential source of noise is the random presence of bright stars, either from the field or cluster members, especially outside the central region in the less-populous GCs or most of the OCs. Structural parameters derived from such SBPs would certainly be affected. One way to minimise this effect is the use of wide rings throughout the whole radius range, but this would cause spatial resolution degradation on the profiles, especially near the center.  The obvious alternative to SBPs is to use star counts to build  radial density profiles (RDPs), in which only the projected number-density of stars is taken into account, regardless of the individual star brightness. This technique is particularly appropriate for the outer parts, provided a statistically significant, and reasonably uniform, comparison field is available to tackle the background contamination. On the other hand, contrary to SBPs, RDPs are less efficient in central regions of populous clusters where the density of stars (crowding) may become exceedingly large. In such cases it may not be possible to resolve individual stars with the available technology.  Finally, a more physically significant profile can be built by mapping the cluster's stellar mass distribution, which essentially determines the gravitational potential and drives most of the dynamical evolution. However, mass density profiles (MDPs) not only are affected by the same technical problems as the RDPs but, in addition, the cluster distance, age and a reliable mass-luminosity relation are necessary to build them.  In principle, the three kinds of profiles are expected to yield different values for the structural parameters under similar photometric conditions, since each profile is sensitive to different cluster parameters, especially the age and dynamical state. Qualitatively, the following effects, basically related to dynamical state, can be expected. Large-scale mass segregation drives preferentially low-mass stars towards large radii (while evaporation pushes part of these stars beyond the tidal radius, into the field), and high-mass stars towards the central parts of clusters. If the stellar mass distribution of an evolved cluster can be described by a spatially variable mass function (MF) flatter at the cluster center than in the halo, the resulting RDP (and MDP) radii should be larger than SBP ones. The differences should be more significant for the core than the half and tidal radii, since the core would contain, on average, stars more massive than the halo and especially near the tidal radius. Besides, the presence of bright stars preferentially in the central parts of young clusters (\\citealt{DetAnalOCs} and references therein) should as well lead to smaller SBP core and half-light radii than the respective RDP ones.  Another relevant issue is related to depth-limited photometry. When applied to the observation of objects at different distances, depth-limited photometry samples stars with different brightness (or mass), especially at the faint (or low-mass) end. Thus, it would be interesting to quantify the changes produced in the derived parameters when RDPs, MDPS and SBPs are built with depth-limited photometry, as well as to check how the structural parameters derived from one type of profile relate to the equivalent radii measured in the other profiles.  In the present work we face the above issues by deriving structural parameters of star clusters built under controlled conditions, in which the radial distribution of stars follows a pre-established analytical profile, and field stars are absent. Effects introduced by mass segregation (simulated by a spatially variable mass function), age and structure are also considered. This work focuses on profiles built in the near-infrared range. The main goal of the present work is to examine relations among structural parameters measured in the different radial profiles, built under ideal conditions, especially noise-free photometry and as small as possible statistical uncertainties (using a large number of stars). In this sense, the results should be taken as upper-limits.  \\begin{table*} \\caption[]{Model star cluster specifications} \\label{tab1} \\renewcommand{\\tabcolsep}{2.65mm} \\renewcommand{\\arraystretch}{1.2} \\begin{tabular}{ccccccccrcccc} \\hline\\hline Model&$R_t/R_c$&c&$\\chi_0$&$\\chi_t$&Age&\\feh&$m_i$&$m_s$&$\\langle m\\rangle$ &$\\rm M_J(TO)$&$\\rm M_J(bright)$&$\\rm M_J(faint)$\\\\ &    &    &     &    &(Myr)&&(\\ms)&(\\ms)&(\\ms)&(mag)&(mag)&(mag)\\\\ (1)  &(2)  &(3) &(4)  &(5)  &(6)  &(7) &(8)  &(9)  &(10)&(11) &(12) &(13)\\\\ \\hline GC-A &5&0.7&0.00&1.35&$10^4$&$-1.5$&0.15&1.02&0.43&$+2.86$&$-2.14$&$+9.12$\\\\ GC-B&20&1.3&0.00&1.35&$10^4$&$-1.5$&0.15&1.02&0.43&$+2.86$&$-2.14$&$+9.12$\\\\ GC-C&20&1.3&0.00&0.00&$10^4$&$-1.5$&0.15&1.02&0.46&$+2.86$&$-2.14$&$+9.12$\\\\ GC-D&40&1.6&0.00&1.35&$10^4$&$-1.5$&0.15&1.02&0.43&$+2.86$&$-2.14$&$+9.12$\\\\ OC-A&15&1.2&0.30&1.35&$10^3$&$~~0.0$&0.15&2.31&0.59&$+0.32$&$-2.68$&$+9.18$\\\\ OC-B&15&1.2&0.30&1.35&$100$&$~~0.0$&0.15&5.42&0.92&$-1.82$&$-4.82$&$+9.18$\\\\ OC-C&15&1.2&0.30&1.35&$10$&$~~0.0$&0.15&18.72&1.76&$-4.82$&$-8.82$&$+9.18$\\\\ \\hline \\end{tabular} \\begin{list}{Table Notes.} \\item Col.~3: concentration parameter $c=\\log(\\rt/\\rc)$. Cols.~4 and 5: mass function slopes at the cluster center and tidal radius. Cols.~8-10: lower, upper and average star mass. Col.~11: absolute J magnitude at the turnoff (TO). Cols.~12 and 13: absolute J magnitude at the bright and faint ends. \\end{list} \\end{table*}  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{fig1.eps}} \\caption{Model star cluster specifications. Panel (a): a random selection of $n$ in the range $0\\leq n\\leq1$ produces King-like RDPs in the range $0\\leq R\\leq\\rt$ (see Eq.~\\ref{eq2}). Panel (b): Radially-variable mass function slopes $\\left(\\frac{dN}{dm}\\propto m^{-(1+\\chi)}\\right)$ used in the models. Panel (c): Padova isochrones used to simulate the mass-luminosity relation of the star cluster models. The 10\\,Gyr, $\\feh=-1.5$ metallicity isochrone is adopted in the globular cluster models. Panel (d): distribution of concentration parameters of the GCs in H03 with peaks at $c\\approx1.6,~1.3,~{\\rm and}~0.7$. Panel (e): model fraction of stars brighter than $M_J=M_{J_{TO}}+\\Delta_{TO}$. In all cases, the fraction of stars brighter than the TO ($M_{J_{TO}}$) is below the $1\\%$ level.} \\label{fig1} \\end{figure}  This work is structured as follows. In Sect.~\\ref{ModelSCs} we present the star cluster models and build radial profiles with depth-limited photometry. In Sect.~\\ref{Struc} we derive structural parameters from each profile, discuss their dependence on depth, and compare similar radii derived from the different types of profiles. In Sect.~\\ref{N6397} we compare relations derived from model parameters with those of the nearby GC NGC\\,6397. Concluding remarks are given in Sect.~\\ref{Conclu}. ",
        "conclusions": "\\label{Conclu}  In this work we simulated star clusters of different ages, structure and mass functions, assuming that the spatial distribution of stars follows an analytical function, similar to \\citet{King62} profile. The mass and near-infrared luminosities of each star were assigned according to a mass function with a slope that may depend on distance to cluster center. They form the set of models from which we built number-density, mass-density and surface-brightness profiles, allowing for a variable photometric depth. The structural parameters core, half-light, half-mass and half-star count, and tidal radii, together with the concentration parameter, were measured in the resulting radial profiles. Next we examined relations among similar parameters measured in different profiles, and determined how each parameter depends on photometric depth.  We point out that the results should be taken as upper-limits, especially for open clusters, since we have considered noise-free photometry and a large number of stars, which produced small statistical uncertainties.  With respect to the adopted form of the radial distribution of stars, we note that \\citet{King62} isothermal sphere, single-mass profile has been superseded by more realistic models like those of \\citet{King66}, \\citet{Wilson75} and \\citet{EFF87}, which have been fit mostly to the SBPs of Galactic and extra-Galactic GCs (Sect.~\\ref{ModelSCs}). The analytical functions associated with these models are characterised by different scale radii (among other parameters) that are roughly related to \\citet{King62} radii. Thus, it is natural to extend the scaling with photometric depth undergone by \\citet{King62} radii to the equivalent ones in the other models.  The main results can be summarised as follows. \\begin{itemize} \\item {\\em (i)} Structural parameters derived from surface-brightness profiles are essentially insensitive to photometric depth, except perhaps the cluster radius in very young clusters.  \\item {\\em (ii)} Uniform mass functions also result in structural parameters insensitive to photometric depth.  \\item {\\em (iii)} Number-density and mass-density profiles built with shallow photometry result in underestimated radii, with respect to the values obtained with deep photometry. Tidal, half-star count and half-mass, together with the core radii are affected with increasing intensity.  \\item {\\em (iv)} Because of the presence of bright stars, radii underestimation increases for young ages.  \\item {\\em (v)} For clusters older than $\\sim1$\\,Gyr, number-density and mass-density radii present essentially the same values; for younger ages, RDP radii become increasingly larger than MDP ones, especially at the deepest profiles.  \\item {\\em (vi)} Irrespective of age, profiles deeper than the turnoff have RDP radii systematically larger than SBP ones, especially the core.  \\item {\\em (vii)} The concentration parameter also changes with photometric depth, reaching a maximum around the turnoff.  \\end{itemize}  Most of the above model predictions were qualitatively confirmed with radii measured in ground-based RDPs and SBPs of the nearby GC NGC\\,6397.  In principle, working with SBPs has the advantage of producing more uniform structural parameters, since they are almost insensitive to photometric depth. However, as discussed in Sect.~\\ref{intro}, SBPs usually present high levels of noise at large radii. Noise that is also present in SBPs of clusters projected against dense fields and/or the less populous ones.  A natural extension of this work would be to examine radial profiles built with photometry that includes observational uncertainties, differential absorption, metallicity gradients, binaries, and star cluster models with a number of stars compatible with those of open clusters.  As a consequence of the wide range of distances to the Galactic (and especially extra-Galactic) star clusters, interstellar absorption, and intrinsic instrumental limitations, the available photometric data for most clusters do not sample the low-mass stars. All sky surveys like 2MASS, usually are restricted to the giant branch, or the upper main sequence, for clusters more distant than a few kpc. In such cases, the structural parameters have to be derived from radial profiles built with photometry that does not reach low-mass stars. The present work provides a quantitative way to estimate the intrinsic (i.e. in the case of photometry including the lower main sequence) values of structural radii of star clusters observed with depth-limited photometry."
    },
    "0711/0711.2402_arXiv.txt": {
        "abstract": "Recent astrophysical observations of neutron stars and heavy-ion data are confronted with our present understanding of the equation of state of dense hadronic matter. Emphasis is put on the possible role of the presence of hyperons in the interior of compact stars. We argue that data from low-mass pulsars provide an important cross-check between high-density astrophysics and heavy-ion physics. ",
        "introduction": "The research areas of high-density astrophysics, as the physics of compact stars, and relativistic heavy-ion collisions are probing matter at extreme densities.  The properties of neutron stars are determined by the nuclear equation of state (EoS), as well as microphysical reactions in dense matter. The stiffness of the high-density matter controls the maximum mass of compact stars. New measurements of the global properties of pulsars, rotation-powered neutron stars, point towards large masses and correspondingly to a rather stiff equation of state (for a recent review on the equation of state for compact stars see \\cite{Lattimer:2006xb}). In a recent analysis of the x-ray burster EXO 0748--67 it was even claimed that soft nuclear equations of state are ruled out \\cite{Ozel:2006km}. Note that this analysis, if confirmed, would not rule out the presence of quark matter in the core of compact stars \\cite{Alford:2006vz}.  On the other side, strange particles (kaons) produced in relativistic heavy-ion collisions just below the threshold of the elementary reaction are sensitive to medium effects due to the created high-density matter (see e.g.\\ \\cite{Schaffner:1996kv}). Recent investigations conclude that the systematics of kaon production can only be explained by an extremely soft nuclear equation of state above normal nuclear matter density \\cite{Sturm:2000dm,Fuchs:2000kp,Hartnack:2005tr,Forster:2007qk}.  There seems to be conflict in determining the nuclear equation of state, which we will discuss in detail in the following. We investigate the impacts of the compression modulus and symmetry energy of nuclear matter on the maximum mass of neutron stars in view of the recent constraints from heavy-ion data on kaon production in dense matter. In particular, we delineate the different density regions probed in the mass-radius diagram of compact stars. We outline the importance of the Schr\\\"odinger equivalent potentials for subthreshold production of kaons. The possible effects from the presence of hyperons in dense neutron star matter are confronted with pulsar mass measurements. ",
        "conclusions": "The combined analysis of heavy-ion data on kaon production at subthreshold energies and neutron star mass measurements points towards a nuclear EoS that is soft at moderate densities and hard at high densities. A soft nuclear EoS as extracted from kaon production data is not in contradiction with heavy pulsars as mutually exclusive density regions are probed. The nuclear EoS above $n\\approx 2n_0$ determines the maximum mass of neutron stars, which is controlled by unknown high-density physics (as hyperons and quark matter). Compact star matter constrained by the heavy-ion data seems to result in rather low maximum masses for compact stars when hyperons are included. A measurement of a heavy pulsar will make it quite difficult for having hyperons inside a neutron star and could exhibit an emerging conflict between hypernuclear and pulsar data.  Properties of low-mass neutron stars ($M\\leq 1.2M_\\odot$), however, are likely to be entirely determined by the EoS of nucleons and leptons only up to $n\\approx 2n_0$, as hyperon and possibly quark matter could appear at larger densities. Thus, the measurement of the radii of low-mass pulsars provides the opportunity for a cross-check between heavy-ion and astrophysical data and possibly for the detection of an exotic phase in the interior of compact stars."
    },
    "0711/0711.2416_arXiv.txt": {
        "abstract": "{The physics of the pulsar magnetosphere near the neutron star surface remains poorly constrained by observations. Although about 2000~pulsars have been discovered to date, little is known about their emission mechanism, from radio to high-energy X-ray and gamma-rays.  Large vacuum gaps probably exist in the magnetosphere, and a non-neutral plasma partially fills the neutron star surroundings to form an electrosphere.} {In several previous works, we showed that the differentially rotating equatorial disk in the pulsar's electrosphere is diocotron unstable and that it tends to stabilise when relativistic effects are included.  However, when approaching the light cylinder, particle inertia becomes significant and the electric drift approximation is violated.  In this paper, we study the most general instability, i.e. by including particle inertia effects, as well as relativistic motions.  Electromagnetic perturbations are described in a fully self-consistent manner by solving the cold-fluid and Maxwell equations.  This general non-neutral plasma instability is called the magnetron instability by plasma physicists.} {We linearise the coupled relativistic cold-fluid and Maxwell equations. The non-linear eigenvalue problem for the perturbed azimuthal electric field component is solved numerically with standard techniques for boundary-value problems like the shooting method. The spectrum of the magnetron instability in a non-neutral plasma column confined between two cylindrically conducting walls is computed for several cylindrical configurations. For a pulsar electrosphere, no outer wall exists. In this case, we allow for electromagnetic wave emission propagating to infinity.} {First we checked our algorithm in the low-density limit. We recover the results of the relativistic diocotron instability. When the self-field induced by the plasma becomes significant, it can first increase the growth rate of the magnetron instability. However, equilibrium solutions are only possible when the self-electric field, measured by the parameter~$s_{\\rm e}$ and tending to disrupt the plasma configuration, is bounded to an upper limit, $s_{\\rm e,max}$.  For $s_{\\rm e}$ close to but smaller than this value $s_{\\rm e,max}$, the instability becomes weaker or can be suppressed as was the case in the diocotron regime.} {When approaching the light-cylinder, particle inertia becomes significant in the equatorial disk of the electrosphere. Indeed, the rest-mass energy density of the plasma becomes comparable to the magnetic energy density. The magnetron instability sets in and takes over the destabilisation of the stationary flow initiated by the diocotron instability close to the neutron star surface. As a consequence, the flow in the pulsar inner magnetosphere is highly unstable, leading to particle diffusion across the magnetic field line. Therefore, an electric current can circulate in the closed magnetosphere and feed the wind with charged particles.} ",
        "introduction": "This year, we celebrate the 40th~year of the discovery of the first pulsar. Nevertheless, the detailed structure of charge distribution and electric-current circulation in the closed magnetosphere of a pulsar remains poorly understood. Although it is often assumed that the plasma fills the space entirely and corotates with the neutron star, it is on the contrary very likely that it only partly fills it, leaving large vacuum gaps between plasma-filled regions. The existence of such gaps in aligned rotators has been very clearly established by \\cite{1985MNRAS.213P..43K, 1985A&A...144...72K}.  Since then, a number of different numerical approaches to the problem have confirmed their conclusions, including some work by \\cite{1989Ap&SS.158..297R}, \\cite{1989Ap&SS.161..187S}, \\cite{1993A&A...268..705Z}, \\cite{1993A&A...274..319N}, \\cite{1994ApJ...431..718T}, \\cite{2002ASPC..271...81S}, and ourselves \\citep{2002A&A...384..414P}. This conclusion about the existence of vacuum gaps has been reached from a self-consistent solution of the Maxwell equations in the case of the aligned rotator.  Moreover, \\cite{2001MNRAS.322..209S} have shown by numerical modelling that an initially filled magnetosphere like the Goldreich-Julian model evolves by opening up large gaps and stabilises to the partially filled and partially void solution found by \\cite{1985MNRAS.213P..43K} and also by \\cite{2002A&A...384..414P}. The status of models of the pulsar magnetospheres, or electrospheres, has recently been critically reviewed by \\cite{2005RMxAC..23...27M}. A solution with vacuum gaps has the peculiar property that those parts of the magnetosphere that are separated from the star's surface by a vacuum region are not corotating and so suffer differential rotation, an essential ingredient that will lead to non-neutral plasma instabilities in the closed magnetosphere, a process never addressed in detail.  This raises the question of the stability of such a charged plasma flow in the pulsar magnetosphere.  The differential rotation in the equatorial, non-neutral disk induces a non-neutral plasma instability that is well known to plasma physicists (\\citealt{1980PhFl...23.2216O, Davidson1990, 1992PhFlB...4.2720O}). Their good confinement properties (trapped particles can remain on an almost unperturbed trajectory for thousands of gyro-periods) makes them a valuable tool for studying plasmas in laboratory, by using for instance Penning traps.  In the magnetosphere of a pulsar, far from the light cylinder and close to the neutron star surface, the instability reduces to its non-relativistic and electrostatic form, the diocotron instability. The linear development of this instability for a differentially rotating charged disk was studied by \\cite{2002A&A...387..520P}, in the thin disk limit, and by \\cite{2007A&A...469..843P, 2007A&A...464..135P} in the thick disk limit. It both cases, the instability proceeds at a growth rate comparable to the star's rotation rate.  The non-linear development of this instability was studied by \\cite{2003A&A...411..203P}, in the framework of an infinitely thin disk model. They have shown that the instability causes a cross-field transport of these charges in the equatorial disk, evolving into a net out-flowing flux of charges. \\cite{2002ASPC..271...81S} have numerically studied the problem and concluded that this charge transport tends to fill the gaps with plasma. The appearance of a cross-field electric current as a result of the diocotron instability has been observed by \\cite{2002AIPC..606..453P} in laboratory experiments in which charged particles were continuously injected in the plasma column trapped in a Malmberg-Penning configuration.  A general overview of the equilibrium and stability properties of non-neutral plasmas in Cartesian and cylindrical geometry can be found in~\\cite{1991RvMP...63..341D}. \\cite{1986PhRvA..33.4284T} describe how to compute fully self-consistent general equilibria configuration for a cold-fluid plasma in a cylindrical diode. This is useful to investigate the stability properties in magnetron devices as presented in \\cite{1986PhFl...29.3832D}.  The aim of this work is to extend the previous work by \\cite{2007A&A...469..843P, 2007A&A...464..135P} on the diocotron instability by including particle inertia effects.  In this paper we present a numerical analysis of the linear growth of the relativistic magnetron instability for a non-neutral plasma column.  The paper is organised as follows.  In Sect.~\\ref{sec:Modele}, we describe the initial setup of the plasma column consisting of an axially symmetric equilibrium between two conducting walls. We give several equilibrium profiles useful for the study of the magnetron instability in different configurations. In Sect.~\\ref{sec:AnalLin}, the non-linear eigenvalue problem satisfied by the perturbed azimuthal electric field component is derived. The algorithm to solve the eigenvalue problem is checked against known analytical results in the low-density limit (diocotron instability), Sect.~\\ref{sec:Check}.  Then, applications to some typical equilibrium configurations are shown in Sect.~\\ref{sec:Results}. First we consider a plasma column with constant density. Next, we study the effect of the cylindrical geometry (curvature of the flow) and the transition to the planar diode limit. Finally, the consequences of the magnetron instability on the pulsar electrosphere is investigated.  The conclusions and the possible generalisation are presented in Sect.~\\ref{sec:Conclusion}. ",
        "conclusions": "\\label{sec:Conclusion}  We developed a numerical code to compute the eigenspectra and eigenfunctions of the magnetron instability in cylindrical geometry, including particle inertia, as well as electromagnetic and relativistic effects, in a fully self-consistent way. It is thus possible to study the behavior of the plasma in the vicinity of the light cylinder, i.e. where the particle kinetic energy becomes comparable to the magnetic field energy density.  Unstable modes are computed for a uniform external applied magnetic field and arbitrary velocity, density and electric field profiles. The resulting equilibrium configuration is computed self-consistently, according to the cold-fluid and Maxwell equations.  Application of the code to a plasma column as well as to the pulsar electrosphere have been shown. In both cases, the magnetron regime gives rise to instabilities that become less and less unstable when the self-field parameter of the flow, $s_{\\rm e}$, approaches its maximum allowed value, $s_{\\rm e} \\lesssim s_{\\rm e,max}$.  Whereas the growth rates can be comparable to the rotation period of the neutron star in the non-relativistic limit, it is found that for special rotation profiles, the magnetron instability is completely suppressed in the relativistic regime, as was already the case in the diocotron limit.  However, the magnetron instability is not restricted to the electrosphere. For instance, it is believed that the pulsed radio emission emanates from the magnetic poles close to the neutron star surface. Highly relativistic and dense leptonic flows are accelerated at these polar caps, forming non-neutral plasma beams able to radiate coherent electromagnetic waves by several processes like the cyclotron maser or the free electron laser instability. We will address this problem as well as the effect of finite temperature in a forthcoming paper."
    },
    "0711/0711.0353.txt": {
        "abstract": "{  We present a method to reconstruct the geometry and energy of high energy muon tracks in IceCube. Through a log-likelihood optimization procedure, an event hypothesis is obtained by maximizing the agreement of the expected amount of light (as function of time) in the optical modules with the shapes of the pulses recorded in the optical modules. This reconstruction method aims to use all information contained in the waveforms recorded in the IceCube digital optical modules (\\DOMs), by comparing those waveforms directly with the expected arrival time distribution of Cherenkov photons at the \\DOM\\ after emission from a hypothetical track, taking into account the optical properties of the South Pole ice. We expect that this method will be effective in particular for highly energetic events in which a significant fraction of the \\DOMs\\ records many photo-electrons.  Currently, for simulated events within an energy range of $100\\TeV$ to $32\\PeV$ which were reconstructed as throughgoing, we obtain an energy resolution of $0.34$ in $\\mathrm{Log}(E/\\mathrm{GeV})$ and an angular resolution of $0.62^\\circ$.  }  \\begin{document} ",
        "introduction": "%\\ramble{Some nice paragraph about high energy cosmic rays and GZK neutrinos.} % %Extremely high energy cosmic rays (EHECRs) ($>10^{20}\\eV$) %are observed by several experiments such as AGASA\\cite{agasa}, %HiRes\\cite{hires} and Auger\\cite{auger}.  There are several theories to %explain the EHECRs.  They are usually classified into two models, namely %``bottom up'' and ``top down''.  In the more conventional bottom up model, %the EHECRs are explained by statistical acceleration\\cite{fermiacc} %in sources such as AGNs and GRBs. In the more exotic top down models, %the EHECR are generated through topological defects or the decay %of supermassive particles. In many of these models high energy neutrinos %are produced. % %An EHECR loses its energy during its travel through interaction %with the cosmic microwave background, well known as the GZK %mechanism\\cite{GZK1966}.  Through this mechanism, extremely high energy (EHE) % neutrinos will %be generated.  A measurement of (or an upper limit on) the flux of EHE %neutrinos will help to constrain the models for EHECR generation.  %\\ramble{Some words about how icecube works, provide essential information so %that a new reader can have a rough idea of what we are talking about %when we use words like \"ATWD waveform\". E.g. mention ATWD is relatively %short and fine grained, but high dynamic range, while FADC is long with %low dynamic range.}  The IceCube telescope is being deployed in the Antarctic ice with its main goal to detect high energy neutrinos arriving from astrophysical sources. Nearly one third of the detector is installed and currently operational \\cite{ice3status}. When fully deployed, the instrumented volume will be approximately $1\\km^3$.  When a neutrino interacts in the ice in or near the detector, it produces a track or cascade signature. Some of the Cherenkov light emitted by the charged lepton and secondary charged particles triggers the \\DOMs. A \\DOM\\ digitizes the signal from a $10\\,\\mathrm{inch}$ photo-multiplier in two ways: with an analog transient waveform digitizer (ATWD) and with a fast analog to digital converter (fADC) \\cite{firstyear}.  The main purpose of the ATWD is to record precise timing information of photons arriving in \\DOMs\\ relatively close to the track or cascade. Therefore, it reads the same signal in 3 channels operating on different gains. Each channel has up to 128~bins with a bin witdh of $3.6\\ns$.  The main purpose of the fADC is to measure pulses with a wider time distribution from a further away track or cascade. It has 256~bins with a bin width of $25\\ns$, giving a total time window of $6.4\\us$.  Given that the IceCube neutrino observatory records the full waveform information, a new likelihood reconstruction technique to exploit the full waveform information is the goal of the research described in this paper.  Conventional reconstruction techniques~\\cite{nimreco} ported to IceCube from its predecessor, AMANDA, do not use the complete waveform. This is a reflection of the original AMANDA data acquisition system which recorded only the leading edge time of the pulse, the total charge of the pulse, and the total time over threshold of the pulse. These conventional reconstruction techniques in IceCube utilize this information by extracting pulse shapes from the ATWD or fADC waveforms and reconstruct a cascade or a muon hypothesis based on this information.  In this paper, we focus on the likelihood reconstruction of high energy muon tracks arising from extremely high energy (EHE) neutrinos with energies up to $10^{11}\\GeV$. %In section~\\ref{gzk} we show how this reconstruction can be %for example be used in the search for so-called GZK neutrinos, which %should be produced when EHE cosmic rays interact with cosmic microwave %background\\cite{GZK1966}. EHE neutrinos should be produced when EHE cosmic rays interact with the cosmic microwave background \\cite{GZK1966}. The significant background due to atmospheric muons presents a major challenge, however.  Since the zenith and energy distributions are different for signal and background, good geometry and energy reconstruction are vital for signal detection.  We hope that with the waveform-based event reconstruction method a significant improvement in sensitivity can be achieved for events at a wide energy range from $\\sim10\\TeV$ up to highest energies, $\\sim\\EeV$. At energies above $1\\PeV$ we expect to increase the sensitivity by effectively reconstruction the energy of non-contained events.  %The waveform-based event reconstruction method is a powerful tool that %can be applied to the GZK neutrino search. ",
        "conclusions": "%\\label{gzk} %\\myedit{To be written... See figures~\\ref{gzk1} and \\ref{gzk2} to understand why %it is vital to have good direction and energy reconstruction also at extremely high %energies. HOWEVER: with these figures, the paper will be too long...} % % %\\begin{figure} %\\begin{center} %%\\includegraphics[width=0.496\\textwidth,height=0.472\\textwidth]{EventRateVsZenithReco.gif} %\\end{center} %\\caption{Rate versus the cosine of the reconstructed zenith %angle. Background in red: atmospheric muon events (probably with %some strong cut on $N_\\mathit{ch}$. Signal (GZK neutrinos) in %black. \\myedit{For the proceedings we should probably avoid relying %on color. BTW: if Keiichi sent me EPS versions of these plots then I %cannot find them back in my mail folder...}}\\label{gzk1} %\\end{figure} % %\\begin{figure} %\\begin{center} %%\\includegraphics[type=gif,width=0.496\\textwidth,height=0.472\\textwidth]{EvetRateRecoWeightCorrect.gif} %\\end{center} %\\caption{Rate versus energy. Background in red: atmospheric muon events %(probably with some strong cut on $N_\\mathit{ch}$. Signal (GZK neutrinos) %in black. \\myedit{For the proceedings we should probably avoid relying %on color.}}\\label{gzk2} %\\end{figure}  %\\ramble{Show GZK log(E) plot. %Show GZK cos(zenith) plot. %Show GZK $A_\\mathit{eff}$ plot.}"
    },
    "0711/0711.2620_arXiv.txt": {
        "abstract": "Astronomers have constructed models of globular clusters for over 100 years.  These models mainly fall into two categories: (i) static models, such as King's model and its variants, and (ii) evolutionary models.  Most attention has been given to static models, which are used to estimate mass-to-light ratios and mass segregation, and to combine data from proper motions and radial velocities.  Evolutionary models have been developed for a few objects using the gaseous model, the Fokker-Planck model, Monte Carlo models and $N$-body models.  These models have had a significant role in the search for massive black holes in globular clusters, for example.  In this presentation the problems associated with these various techniques will be summarised, and then we shall describe new work with Giersz's Monte Carlo code, which has been enhanced recently to include the stellar evolution of single and binary stars.  We describe in particular recent attempts to model the nearby globular cluster M4, including predictions on the spatial distribution of binary stars and their semi-major axis distribution, to illustrate the effects of about 12 Gyr of dynamical evolution.  We also discuss work on an approximate way of predicting the ``initial\" conditions for such modelling. ",
        "introduction": "\\subsection{Some astrophysical questions} There are many reasons for constructing a dynamical model of a globular cluster, but they fall into two broad categories.  First there are problems that can be tackled by constructing static (equilibrium) models, such as \\begin{enumerate} \\item Inferring the mass from the surface brightness profile, radial velocities,\tproper motions and  mass functions.  In this way one can estimate the total mass of stars below the observational limit, such as faint white dwarfs \\citep[e.g.][]{Dr1988} \\item Inferring the global mass function from local mass functions \\citep[e.g.][]{Ri2004}: then one can address the question of whether this is the same for all clusters. \\item Measuring cluster distances by comparison of radial velocities and proper motions:  the model is used to correct for rotation, or to link the different locations and stellar components which are observed by the different techniques \\citep[e.g.][]{Ve2006}. \\end{enumerate}  Then there is a second range of questions which require dynamic {\\sl evolutionary} models, i.e. questions such as \\begin{enumerate} \\item Inferring the primordial mass function from local, present-day mass functions: the model is used to correct for preferential escape of low-mass stars \\citep[e.g.][]{BM2003}. \\item Inferring primordial parameters of the binaries from their present-day statistical properties: primordial abundance, period distribution, etc \\citep[e.g.][]{Kr2001} \\item  Determining the effect of dynamics on the estimation of mass through the virial theorem, which is affected by mass segregation \\citep{Fl2006} \\end{enumerate} Note that these last few references do not deal with specific objects (which is the focus of the rest of this review) but with general trends.  This paper begins with a review of the methods and observational constraints which have been used to construct models of individual globular clusters, often with a view to answering the above types of question.  Then we focus on one particular method, the Monte Carlo method, and its application to the nearby globular cluster M4.  \\subsection{Methods and constraints}  \\subsubsection{The methods} A number of techniques have been used to construct static models, to answer questions of the first type: \\begin{enumerate} \\item Plummer's model \\citep{Pl1911} \\item King's model \\citep{Ki1966,PK1975} \\item Anisotropic models \\citep{MB1963} \\item Multi-mass models \\citep{GG1979, Pr1986, Du1990} \\item Non-parametric models \\citep{GF1995} \\item Schwarzschild's method \\citep{Ve2006} \\item Jeans' equations \\citep{Le1992} \\end{enumerate} Even this list may not be exhaustive.  Much less work has been done on dynamical evolutionary models of {\\sl specific} globular clusters, but the methods include \\begin{enumerate} \\item Gas/fluid models (\\citealt{An1980} [M3]) \\item Fokker-Planck models (Cohn and co-workers: \\citealt{Gr1992} [N6624], \\citealt{Du1997} [M15]; \\citealt{Dr1993, Dr1995} [N6397]; Phinney 1993 [M15]) \\item Monte Carlo models (Giersz \\& Heggie 2003 [$\\omega$ Cen]) \\item $N$-body models \\end{enumerate}  The last of these should really not be on this list.  Though it should be the method of choice, it has not been used for the purposes which are the focus of this paper.  An example is the modelling of M15 by \\citet{Ba2003}, who constructed a small version which, when scaled up, corresponded approximately to the conditions expected for M15.  The difficulty is that unscaled $N$-body models are practically limited to $N$ of order $10^5$ at present \\citep[e.g.][]{BM2003,  Hu2007}, whereas the median for the globular clusters at the present time is about $5\\times 10^5$.  Though it might be hoped that the gap would be bridged by the next generation of computers, it must be recognised that all globular clusters at the present day have lost substantial numbers of stars, and the primordial median must have been higher.  Note that it has become possible only relatively recently to carry out full simulations of open star clusters.  The initial mass of M67, for instance, is estimated at about 19 000 $\\msun$, i.e. about 10 times its present mass, and this simulation, with a realistic complement of primordial binaries, took of order 1 month \\citep{Hu2005}.  \\subsubsection{The observational constraints}  Whichever method is chosen to model a globular cluster, there are a number of observational constraints to be satisfied.  In historical order of first use we have \\begin{enumerate} \\item Surface brightness and/or star counts, starting with \\citet{Ze1908}; \\item Radial velocities, whether central averaged values \\citep{Il1976} or radial velocities of individual stars \\citep{GG1979} \\item Pulsar accelerations \\citep{Ph1993, Gr1993} \\item {\\sl Deep} luminosity/mass functions, starting essentially with the advent of studies by several authors using HST (1995) \\item Accurate proper motions \\citep{Bo2006}.  (We refer here to models built to satisfy constraints imposed by observations of internal proper motions, and not to the use of proper motions to establish membership.) \\end{enumerate}   \\subsubsection{The Monte Carlo Model}\\label{sec:mcm}  This paper focuses on an application of the Monte Carlo code developed essentially by \\citet{Gi1998, Gi2001, Gi2006}.  In this approach we assume spherical symmetry and dynamic equilibrium, and characterise each star by its energy $E$ and angular momentum $J$. The code repeatedly alters $E$ and $J$ to mimic the effects of gravitational encounters, using the theory of relaxation.  The same theory underpins Fokker-Planck codes, and the basic Monte Carlo code provides essentially a Monte Carlo solution of the Fokker-Planck equation.  The Monte Carlo code is rather suitable for the addition of a number of other  process, chiefly: \\begin{enumerate} \\item The galactic tidal field, which is treated as a cutoff \\item Binaries, whose interactions are treated using cross sections (from \\citet{Sp1987} for interactions with single stars, and expressions based on \\citet{Mi1983,Mi1984a,Mi1984b} for interactions between binaries) \\item Stellar evolution of single stars \\citep{Hu2000} and binary stars \\citep{Hu2002}. \\end{enumerate} Each of these requires some comment, though further details are given by  Giersz \\& Heggie in this volume, and in Sec.4 below.  \\begin{enumerate} \\item There are significant differences between  a tidal field and a tidal cutoff, as these lead to somewhat different scalings of the dissolution time with $N$ \\citep{Ba2001}.  We have attempted to mimic this with a mass-dependent lowering of the escape energy. \\item The Monte Carlo code described here has no triples, and so hierarchical triples (which are a common product of binary-binary encounters) have to be bypassed.  Furthermore, the use of cross sections hinders the inclusion of star-star collisions during 3- and 4-body encounters.  Finally, the cross sections are not well known for unequal masses.% \\item Stellar evolution is implemented via the McScatter interface \\citep{He2006}.  Besides the stellar evolution packages of Hurley et al (referenced above) it also interfaces to the stellar evolution package SeBa in starlab \\citep{PZ1996}.  At present, however,   the latter is limited to solar metallicity, which is unsuitable for the study of old globular clusters. \\end{enumerate}  In view of the above approximations and uncertainties, the testing and calibration of the Monte Carlo code against the results of $N$-body models (in the regime of small enough $N$) is an essential safeguard. Such studies are described by Giersz \\& Heggie in this volume. ",
        "conclusions": "It is shown by Giersz \\& Heggie (this volume) that Monte Carlo models can provide similar results to $N$-body models, in the range where comparison is possible, with similar physics (binaries, stellar evolution, etc.), except for a number of restrictions:\\begin{enumerate} \\item Use of a tidal cutoff, instead of the tidal field.  Though Sec.\\ref{sec:mcm} summarises our current approach to this problem,  other treatments are possible, and worth trying. \\item Use of a {\\sl static} tide.  The effects of tidal shocks have been studied by a number of authors \\citep[e.g.][]{Ku1995}, and it would be possible to add the effects as another process altering the energies and angular momenta of the stars in the simulation. \\item Rotation: it has been shown \\citep{Ki2004} that, to the extent that rotating and non-rotating models can be compared, rotation somewhat accelerates the rate of core collapse.  Rotation is hard to implement in this Monte Carlo model, however. \\item Use of cross sections for triple/quad interactions: this limitation could be overcome by direct integration of the interactions, as is done by \\citet{FR2007} in their version of the Monte Carlo scheme.  It will be important to do so, as it would remove the dependence on cross sections which are not well known for unequal masses, and would also permit us to include the collisions between stars which commonly occur in long-lived few-body interactions \\citep{HI1985}. \\item Neglect of triples:  these are also commonly produced in binary-binary encounters \\citep{Mi1984a}, and it is desirable to include these as a third species (beyond single and binary stars). Their observable effects may be small, but of course there is one intriguing example in the very cluster we have focused on here \\citep{Th1993}. \\end{enumerate}  Despite these limitations, not all of which are easily curable, Monte Carlo models are feasible in reasonable time for globular clusters, which are too large for direct $N$-body models.  They yield predictions for mass segregation, luminosity functions, distributions of binary parameters, anisotropy, and many other kinds of data, which can hardly be obtained in any other way.  (The only comparable method of which we are aware is the hybrid code of \\citealt{GS2003}.) Even when $N$-body simulations eventually become possible, Monte Carlo models will remain as a quicker way of exploring the main issues, just as King models have continued to dominate the field of star cluster modelling even when more advanced methods (e.g. Fokker-Planck models) have become available.  In addition to some of the possible improvements mentioned above, it is our intention to extend the approach to a number of other objects, including a ``collapsed-core'' cluster such as NGC6624.  We welcome all suggestions for observational or theoretical comparisons, either on M4 or on other objects."
    },
    "0711/0711.0249_arXiv.txt": {
        "abstract": "{GCIRS~3 is the most prominent MIR-source in the central parsec {of the Galaxy}. NIR spectroscopy has failed to solve the enigma of its nature. The properties and peculiarities of extreme individual objects in the central stellar cluster contribute to our knowledge of star and dust formation close to a supermassive black hole.} {We initiated an unprecedented interferometric experiment to understand the nature of GCIRS~3, where we investigate its properties as a spectroscopic and interferometric reference star at 10~$\\mu$m.} {VLT/VISIR imaging separates a compact source from diffuse, surrounding emission. The VLTI/MIDI instrument was used to measure spectroscopically resolved visibility moduli at an angular resolution of $\\sim$ 10~mas of that compact 10~$\\mu$m source, still unresolved by a single VLT. Recent NIR/MIR photometry data were added to enable simple SED- and full radiative transfer-modeling of the data.} {The luminosity and size estimates show that IRS~3 is   probably a cool carbon star enshrouded by a complex dust distribution. Blackbody  temperatures were derived. The coinciding interpretation of single telescope and interferometric data confirm dust emission from several different spatial scales. The interferometric data resolve the inner rim of dust formation. Despite observed deep silicate absorption towards GCIRS~3, we favor a carbon-rich circumstellar dust shell. The silicate absorption most probably takes place in the outer diffuse dust, which is mostly ignored by MIDI measurements, but very observable in complementary VLT/VISIR data. This indicates physically and chemically distinct conditions of the local dust, changing with the distance to GCIRS~3.} {We have demonstrated that optical long baseline interferometry at infrared wavelengths is an indispensable tool for investigating sources at the Galactic center. Our findings suggest further studies of the composition of interstellar dust and the shape of the 10~$\\mu$m silicate feature in this extraordinary region. } ",
        "introduction": "\\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{6733fi01.eps} \\caption{Flux-calibrated Lucy-Richardson deconvolved image at 8.59~$\\mu$m after restoration with a 250~mas Gaussian beam. The pixel scale is 75~mas per pixel. The logarithmic contours levels are 1.6$^{n}\\cdot$7~mJy. The flux scale is given on top. IRS~3 and \\object{Sgr A*} are highlighted. Details of the data reduction are given in \\citet{2007A&A...462L...1S}.} \\label{fig:62} \\end{figure}   At a distance of $\\sim$7.6~kpc \\citep{2005ApJ...628..246E}, the center of the Milky Way is by far the closest center of a large spiral galaxy. Its astrophysical properties can be studied on a unique angular scale of $\\sim$ 40 mpc/arcsec, which is two orders of magnitude smaller than the angular scale at the nucleus of \\object{M31}, the next comparable galaxy \\citep{2005MNRAS.356..979M}. Star formation and the kinematics of the central stellar cluster can be studied in the region of direct influence of the supermassive black hole (SMBH) at the dynamic center of the Milky Way \\citep{1996Natur.383..415E,2000Natur.407..349G,2002Natur.419..694S}.  The active history of star formation, despite the tidal forces of the SMBH, is manifested in the existence of numerous massive, young stars in the central cluster \\citep{1995ApJ...447L..95K, 2000MNRAS.317..348G,2004ApJ...602..760E, 2004A&A...425..529M}. The 7 most luminous ($L\\,>\\,10^{5.75}\\,$L$_\\odot$), moderately hot blue supergiants ($T\\,<\\,10^{4.5}\\,$K) provide about 50\\% of the flux ionizing the region \\citep{1995ApJ...441..603B,1995ApJ...447L..95K,1997A&A...325..700N}. \\citet{2003A&A...398..229S} have shown that a few tens of carbon-rich supergiants can produce about 50\\% of the mass loss of a large stellar sample with solar neighborhood characteristics. Thus, similar to the ionizing flux, the mass loss and dust-formation properties of a stellar cluster can be dominated by a few prominent stars. These facts underline the importance of studying individual extreme objects like \\object{GCIRS~3}~\\footnote{in the following IRS~3} to understand properties of the entire surrounding stellar environment.  The recent advent of mid infrared\\footnote{3-20 $\\mu$m covering the atmospheric $L$, $M$, $N$, and $Q-$windows} (MIR) instruments on 8~m class telescopes enables us to study in detail the thermal dust at the Galactic center (GC) at unprecedented angular resolution. The investigation of the circum- and interstellar dust distribution at the GC uncovers stellar mass loss, zones of wind interaction, formation history, evolution, and kinematics.  Photometric and spectral properties of dusty stars at the GC have recently been published by \\citet{2004A&A...425..529M} and \\citet{2005A&A...433..117V,2006ApJ...642..861V}. Despite an average optical extinction of $A_{\\rm V}\\,\\sim\\,$25 \\citep{2003ApJ...594..294S,2005A&A...433..117V}, near-infrared spectroscopy and imaging reveal the nature of the underlying dust-embedded stars in most cases, since the dust extinction is wavelength-dependent and decreases  from the optical towards longer wavelengths.  This article focuses on the most prominent of the MIR bright dusty sources, \\object{IRS~3}, the embedded stellar source, which still eludes any spectral classification \\citep{2006ApJ...643.1011P,2006ApJ...641..891T}. It is located within the central 20''. A recent state-of-the-art narrow-band MIR image at 8.6~$\\mu$m with an angular resolution of only 250~mas is shown in Fig.~\\ref{fig:62}. The extended and diffuse dust emission surrounding IRS~3 is visible at unprecedented angular resolution in this image.  The NIR extinction studies reveal a spatial variation of only $\\sim$10\\% ($A_{\\rm K}\\,\\approx\\,3$) of the interstellar extinction over this region \\citep{2003ApJ...594..294S,2004A&A...425..529M,2007A&A...469..125S}.  In contrast, narrow-band, $N$-band photometry and spectroscopic observations are interpreted to indicate a significant amount of {\\em additional} 10~$\\mu$m silicate absorption along the line of sight towards IRS~3 with respect to other parts of the central 20'' \\citep{1978ApJ...219..121B,1985MNRAS.215..425R,2006ApJ...642..861V}. The published intrinsic optical depths of about $\\tau_{9.8}({\\rm IRS~3}) \\sim 1$, in addition to the average $\\tau_{9.8}({\\rm GC})\\,\\approx \\,3.5$ \\citep[e.g.][]{1985MNRAS.215..425R}, still underestimate the true value due to source confusion problems. While our new high-resolution VISIR data clearly indicate that more than 50\\% of the $N$-flux is diffuse, extended emission at 0.3'' resolution, one byproduct of the MIDI observations is the calibrated, low-resolution spectra of the compact emission, which show $\\tau_{9.8}$(IRS~3) to be much larger than the value given above.  \\citet{1978ApJ...220..556R} found the spectro-photometric MIR properties of IRS~3 to resemble either young stars or OH/IR stars. The latter interpretation is opposed by the lack of OH-maser emission. With an MIR color temperature of $\\sim$~400~K, the central emission of IRS~3 was found to be (together with the nearby GCIRS~7) the hottest and most compact of the sources dominating the thermal dust irradiation from the GC \\citep{1978ApJ...220..556R,1985ApJ...299.1007G,1990MNRAS.246....1S}. Extended dust emission around IRS~3 interacting with external stellar winds has been found in recent $L-$ and $M-$band observations \\citep{2005A&A...433..117V}. While hot star hypotheses are given by some authors \\citep{1995ApJ...447L..95K,2003ANS...324..597T}, the lack of ionizing gas leads \\citet{1985MNRAS.215..425R} to the assumption of IRS~3 being a cool dust-enshrouded star.  Within the past two years we collected a unique dataset of optical long-baseline interferometric (OLBI) data of IRS~3 at 10~$\\mu$m. These constitute the first successful OLBI observation of an object within the central parsec of our galaxy, opening the window to NIR/MIR GC observations at highest angular resolution \\citep{2005Msngr.119...43P}. In this article we show that the OLBI data strongly support the former hypothesis of a cool dust-enshrouded star. Furthermore our results shed light on the amount and spectral shape of the interstellar 10~$\\mu$m extinction towards IRS~3 at unprecedented angular resolution. Since MIDI is a relatively new instrument, the achievable precision of visibility measurements of such distant, challenging targets is not common knowledge yet. In Sect.~\\ref{sec:2} we therefore describe the observations, the extensive data-reduction and calibration efforts, and the evaluation of different data reduction techniques in some detail to ease and stimulate similar experiments and the comparison of their outcomes. Then the immediate observational results are given (Sect.~\\ref{sec:30}), followed by a detailed discussion of the results in the astrophysical context (Sect.~\\ref{sec:4}) and a summary of our conclusions in Sect.\\ref{sec:6}. ",
        "conclusions": "\\label{sec:6}  {We have presented new high-resolution 10~$\\mu$m data of the enigmatic IRS~3 object in the immediate vicinity of SgrA*. The analysis and interpretation of the data contribute to the understanding of both the nature of the embedded star and the location of the unusually strong interstellar silicate absorption towards IRS~3.  The new high-resolution VISIR imaging data clearly resolve IRS~3 into both a compact and a diffuse emission component for the first time. Most probably the deep 10~$\\mu$m silicate absorption feature towards IRS~3 takes place partially in the local interstellar dust, which is responsible for the substantial diffuse MIR emission.  The interferometric data bear convincing evidence that IRS~3 is a cool star within a carbon-rich dust-shell without significant circumstellar silicate absorption. The clear imprint in the circularly symmetric interferometric data of the existence of more extended circumstellar dust was shown.} A one-dimensional radiative transfer model can successfully explain the complete, available near and mid infrared data, spatially unresolved by single telescope measurements, if {outer dust emission on 40~mas scales and beyond and non-negligible additional optical depth is accounted for}. We estimate a radius of 35~AU and a temperature of $\\sim\\,1200~$K for the inner dust rim around IRS~3.  With a bolometric luminosity of 5~$\\cdot 10^4\\,L_\\odot$ and a stellar temperature of about 3000~K IRS~3 appears to be a cool luminous carbon star, {most probably in the helium-core burning phase. Such mass-losing stars are important for  enriching their environment. {Only because of the high angular resolution of the interferometer could we demonstrate that a large fraction of the dust is produced by IRS~3 itself, whereas most other MIR bright sources in the central parsec have been shown to heat the surrounding interstellar dust that has already formed \\citep{2005ApJ...624..742T,2006JPhCS..54..273P}. Thus IRS~3 appears to be the prime candidate for observing ongoing dust formation in the immediate vicinity of a supermassive black hole.}}  In addition to the interstellar silicate absorption, we found indications of SiC absorption. {For the first time the new generation of MIR spectrometers at 8m class telescopes offer the possibility of studying in detail the spectral properties of the interstellar absorption towards GC stars at angular resolution high enough to account for source confusion and substantial diffuse emission in this outstandingly dense and dusty region. Since a C-rich circumstellar dust shell is free of any silicate feature and IRS~3 is embedded in an exceptionally large amount of ISM, the resulting deep absorption of the smooth background continuum of IRS~3 appears well-suited to studying the shape of the interstellar absorption towards the GC and the extraordinary strength of the silicate absorption feature in more detail. The absorption profile depends on the chemical composition and is needed to accurately correct for the ISM extinction, yielding the intrinsic spectra of all GC sources.}  On the basis of our results we propose to grant IRS~3 a central role in such new studies of the interstellar silicate absorption towards the GC. Furthermore, a longterm MIR variability study is proposed to investigate periodical dust formation around IRS~3, suggested by this complex dust morphology."
    },
    "0711/0711.1553_arXiv.txt": {
        "abstract": "We present deep mid-infrared spectroscopy with the {\\it Spitzer Space Telescope} of 13 submillimeter galaxies (SMGs) in the Great Observatories Origins Deep Survey North (GOODS-N) field. We find strong polycyclic aromatic hydrocarbon (PAH) emission in all of our targets which allows us to measure mid-IR spectroscopic redshifts and place constraints on the contribution from star formation and active galactic nuclei (AGN) activity to the mid-IR emission. In the high signal-to-noise ratio composite spectrum of SMGs, we find that the hot dust continuum from an AGN contributes at most 30\\% of the mid-IR luminosity. Individually, only $2/13$ SMGs have continuum emission dominating the mid-infrared luminosity; one of these SMGs, C1, remains undetected in the deep X-ray images but shows a steeply rising continuum in the mid-IR indicative of a Compton-thick AGN. We find that the mid-IR properties of SMGs are distinct from those of 24$\\mu$m-selected ULIRGs at $z\\sim2$; the former are predominantly dominated by star formation while the latter are a more heterogeneous sample with many showing significant AGN activity. We fit the IRS spectrum and the mid-IR to radio photometry of SMGs with template spectral energy distributions to determine the best estimate of the total infrared luminosity from star formation. While many SMGs contain an AGN as evinced by their X-ray properties, our multi-wavelength analysis shows that the total infrared luminosity, $L_{IR}$, in SMGs, is dominated by star formation and not AGN activity. We find that high redshift SMGs lie on the relation between $L_{IR}$ and $L_{\\rm{PAH,}6.2}$ (or $L_{\\rm{PAH,}7.7}$ or $L_{\\rm{PAH,}11.3}$) that has been established for local starburst galaxies. This suggests that PAH luminosity can be used as a proxy for the star formation rate in SMGs. SMGs are consistent with being a short lived cool phase in a massive merger. Although there appears to be an AGN present in many SMGs, it does not appear to have become strong enough to heat the dust and dominate the mid- or far-infrared emission. ",
        "introduction": "\\label{sec:intro} The {\\it COsmic Background Explorer} ({\\it COBE}) revealed that the extragalactic background light at infrared (IR) wavelengths was comparable to that at optical wavelengths indicating that half of the star formation and active galactic nuclei (AGN) emission, averaged over cosmic time, is obscured by dust (e.g.~Puget et al.~1996; Hauser \\& Dwek 2001). In the local Universe, the volume density of dusty, infrared luminous galaxies is insufficient to account for this background. Therefore the density of such galaxies must evolve strongly with redshift (e.g.~see Lagache, Puget, \\& Dole 2005 for a review). In fact, imaging observations with the {\\it Infrared Space Observatory} ({\\it ISO}) and the {\\it Spitzer Space Telescope} (hereafter {\\it Spitzer}) have shown that about 70\\% of the comoving star formation rate density (SFRD) at $0.5<z<3$ is obscured by dust (Chary \\& Elbaz 2001; LeFloc'h et al.~2005; P{\\'e}rez-Gonz{\\'a}lez et al.~2005).  In the local Universe, we refer to these dusty galaxies as luminous infrared galaxies (LIRGs, $11\\le \\rm{log}\\,(L_{\\rm{IR}}\\,[L_{\\odot}])<12$) and ultraluminous infrared galaxies (ULIRGs, $\\rm{log}\\,(L_{\\rm{IR}}\\,[L_{\\odot}]) \\ge 12$) where $L_{\\rm{IR}}$ is the total IR luminosity from 8--1000$\\,\\mu$m (see Sanders \\& Mirabel 1996 for a review). It is now becoming clear that IR luminous galaxies are a heterogenous population with selection functions at different mid- and far-infrared wavelengths revealing different subsets of this population (Daddi et al.~2005; Lutz et al.~2005b; Yan et al.~2005; Pope et al.~2006).  Perhaps the most enigmatic members of the high redshift ULIRG population are the submillimeter (submm) galaxies (SMGs). These galaxies were initially discovered in the late 1990s with the Submillimetre Common User Bolometer Array (SCUBA, Holland et al.~1999) on the James Clerk Maxwell Telescope (JCMT) and then later with MAx-Planck Millimetre BOlometer (MAMBO) camera on the Institut de Radio Astronomie Millimetrique (IRAM) 30-m telescope (see Blain et al.~2002 for a review). They have proven to be difficult to study, due to the large beam size of submm observations and the difficulty in finding counterparts at other wavelengths (e.g.~Ivison et al.~2002; Pope et al.~2005). Nevertheless, it has been found that they are very massive systems (Swinbank et al.~2004; Borys et al.~2005; Greve et al.~2005; Tacconi et al.~2006) at $z\\sim2$ (Chapman et al.~2005; Pope et al.~2006; Aretxaga et al.~2007) which are thought to be connected to the most massive present day galaxies via an evolutionary sequence (Lilly et al.~1999). A crucial aspect that needs to be explored in SMGs is what powers their extreme luminosities: AGN activity or star formation?  The mid-IR spectral regime is a particularly good probe of what is producing the total infrared luminosity since it has contributions from both star-formation and AGN each with a unique signature at these wavelengths. The 5--15$\\,\\mu$m mid-IR spectral energy distribution (SED) can be decomposed into 3 main components (e.g.~Sajina et al.~2007): (1) emission features from what are thought to be polycyclic aromatic hydrocarbon (PAH) molecules (Puget \\& Leger 1989; Allamandola et al.~1999); (2) power-law (or warm blackbody) emission; and (3) extinction characterized by prominent silicate absorption features (e.g.~Draine 2003). The first is understood to be powered entirely by star formation while the second is predominantly emission from the AGN (Genzel et al.~1998). The second component can also contain contributions from hot dust (very small grains) present in the most energetic HII regions (e.g.~Tran et al.~2001). The third component can be found in both starbursts and AGN since it just requires a dust screen around hot dust emission.  Measuring the contribution from ULIRGs to the global star formation rate requires dissecting their IR emission into contributions from starbursts and AGN. Mid-IR spectroscopy with {\\it ISO} made it possible to decompose the various contributors in local ULIRGs with the result that they are predominantly starburst powered, although they also show signs of AGN activity (e.g.~Genzel et al.~1998; Lutz et al.~1999; Rigopoulou et al.~1999; Sturm et al.~2000; Tran et al.~2001). In particular, {\\it ISO} studies of local ULIRGs showed that at low luminosities, the ULIRGs are powered mainly by a starburst but at the very highest luminosities ($L_{\\rm{IR}}\\gtrsim3\\times10^{12}\\rm{L}_{\\odot}$), the dominant emission source is the AGN (Lutz et al.~1998; Tran et al.~2001). From this one might expect that SMGs, with their extreme IR luminosities ($L_{\\rm{IR}}\\sim6\\times10^{12}\\rm{L}_{\\odot}$, Pope et al.~2006), are primarily AGN-dominated.  The {\\it Spitzer} Infrared Spectrograph (IRS; Houck et al.~2004), with its sensitivity and spectral resolution, can extend this effort outside the local Universe, allowing much larger samples to be observed to much fainter flux limits. The IRS can detect emission features and continuum in ULIRGs out to $z\\sim4$ (Valiante et al.~2007). The 3 mid-IR spectral components listed above have been fit to the {\\it Spitzer} IRS spectra of ULIRGs at various redshifts to determine the AGN contribution to the mid-IR luminosity (e.g.~Sajina et al.~2007). Although X-ray imaging is very effective at identifying the presence of an AGN in a galaxy, it is subject to  obscuration in extreme, Compton-thick sources. The X-ray luminosity provides a direct measurement of the power of the AGN at X-ray energies but it is necessary to make assumptions about the geometry of the obscuring dust to determine the bolometric luminosity of the AGN (e.g.~Alexander et al.~2005). Furthermore, AGN indicators at UV and optical wavelengths can be very obscured in dusty SMGs (Swinbank et al.~2004; Chapman et al.~2005). Mid-IR spectroscopy, by directly detecting the thermally re-radiated emission from the obscuring dust, provides one of the most effective ways to identify AGN, including Compton-thick sources, and measure their contribution to the infrared luminosity (e.g.~Lutz et al.~2004).  Two previous studies have presented IRS spectroscopy of SMGs. Menendez-Delmestre et al.~(2007) presented IRS spectra of 5 radio-detected, spectroscopically confirmed SMGs. They found that four SMGs with lower redshifts ($z<1.5$) showed PAH emission and a composite spectrum of these four sources fits well to a scaled M82 spectrum plus a power-law component. They conclude that these systems host both star formation and AGN activity. Their fifth source at $z=2.4$ was found to have a more AGN-type spectrum. Valiante et al.~(2007) discussed IRS spectra of 13 bright SMGs, 9 of which are detected with the IRS (the first two spectra from this sample were presented in Lutz et al.~2005a). This sample was more representative of the redshift distribution of SMGs. They also found that SMGs are likely to be mainly starburst powered. Sajina et al.~(2007) used IRS spectra to study a sample of {\\it Spitzer} 24$\\,\\mu$m-selected galaxies at $z\\sim2$. This sample was selected to have $S_{24\\,\\mu m}>900\\,\\mu$Jy and very red 24-to-8$\\,\\mu$m and 24-to-$R$ colors (Yan et al.~2007). 75\\% of this sample consists of continuum-dominated sources, with some weak PAH emission indicative of star formation activity. They conclude that these sources display mid-IR spectra similar to those of the most luminous local ULIRGs. Follow-up studies of these galaxies have shown that they are not generally (sub)mm detected sources ($S_{1.2\\rm{mm}}\\lesssim 1.5\\,$mJy for $33/40$ sources, Lutz et al.~2005b).  In this paper we present {\\it Spitzer} IRS spectra of a sample of SMGs in GOODS-N which have extensive multi-wavelength information from X-ray to radio wavelengths. This allows us to produce full spectral energy distributions and to calculate the bolometric luminosity. Our {\\it Spitzer} IRS observations go substantially deeper ($S_{24\\,\\mu m}>200\\,\\mu Jy$) than previous work. This enables a study of the typical blank-field submm galaxy population which have $S_{850\\,\\mu m}>2\\,$mJy. The main goal of our program is to decompose the mid-IR spectra into star formation and AGN components, in order to infer how much of the total infrared luminosity is coming from each. From the mid-IR spectroscopy we are also able to obtain spectroscopic redshifts and measure the luminosities and equivalent widths of individual PAH lines to investigate various line diagnostics. We compare our measurements with line strengths in other star-forming galaxies to put constraints on the evolution of SMGs and the relationship between star formation and black hole growth.  This paper is formatted as follows. Our {\\it Spitzer} observations are described in Section 2 and the data analysis is explained in Section 3. Section 4 presents the main results of the IRS spectroscopy including redshifts, spectral decomposition and AGN classification. In Section 5 we fit the full SEDs of these SMGs to determine their total IR luminosities. We explore correlations between the PAH and IR luminosities in Section 6. We discuss our results in the context of the role that SMGs play in galaxy evolution in Section 7. Finally, our conclusions are listed in Section 8. We include an appendix which discusses the noise properties of our {\\it Spitzer} IRS spectra and includes notes on individual sources.  All magnitudes in this paper use the AB system. We use a standard cosmology with $H_{0}=73\\,\\rm{km}\\,\\rm{s}^{-1}\\,\\rm{Mpc}^{-1}$, $\\Omega_{\\rm{M}}=0.3$ and $\\Omega_{\\Lambda}=0.7$. ",
        "conclusions": "{\\it Spitzer} IRS spectroscopy has been obtained for a sample of 13 SMGs brighter than 200$\\,\\mu$Jy at 24$\\,\\mu$m. This has effectively confirmed the identification of these SMG counterparts, and in some cases has provided a new redshift estimate. The SMGs show strong PAH emission and, on average, only a small continuum contribution from hot dust.  We have explored several diagnostics from the mid-IR spectra of SMGs to determine the level of AGN contributing to the luminosity at these wavelengths. All of them seem to converge on a picture in which SMGs are starburst-dominated systems with at most a 30\\% contribution from an AGN at mid-IR wavelengths. Their mid-IR spectra are similar to a scaled spectrum of local starburst galaxies like M82.  The classification of an SMG as AGN or starburst-dominated from the mid-IR often disagrees with the classification from the X-rays. While X-ray observations are better suited for determining the presence of an AGN in an SMG, the mid-IR spectra can determine how important the AGN is to the total infrared luminosity since it is directly detecting the hot dust emission and not subject to details about the geometry of the obscuring material.  Full IR SED fits to the mid-IR data, far-infrared and radio photometry show that SMGs are best-fit by scaled up versions of local IR-luminous galaxy templates with additional extinction from the Draine (2003) extinction curves. These models have cool dust temperatures ($T=32\\,$K) and high $L_{\\rm{IR}}$s which imply SFRs of $\\sim700\\,\\rm{M_{\\odot}yr^{-1}}$ assuming a Salpeter initial mass function.  SMGs lie on the local relation between $L_{\\rm{IR}}$ and $L_{\\rm{PAH,}6.2}$ (or $L_{\\rm{PAH,}7.7}$ or $L_{\\rm{PAH,}11.3}$), which means that the PAH line flux can be used to estimate the SFRs in these systems, albeit with large uncertainties. Equivalent widths, which have been the focus of some other studies, are much less useful in this regard.  SMGs appear to be an earlier phase in the evolution of massive galaxies than other high redshift ULIRGs. The average high redshift ULIRG shows warmer dust and less PAH emission. It is therefore likely that the AGN in SMGs has not yet grown strong enough to heat the dust substantially and destroy the PAH molecules.  We have been able to put constraints on the contribution from AGN activity to the restframe 5--12$\\,\\mu$m mid-IR emission, and have extrapolated this to far-IR wavelengths. Without more data from 30--100$\\,\\mu$m rest frame it is difficult to directly probe the emission at far-IR wavelengths. We will be able to model this part of the spectrum in more detail once data from the {\\it Herschel Space Observatory}\\footnote{\\tt{http://herschel.esac.esa.int/}} and SCUBA-2 (Holland et al.~2006) become available. Details of the SMG phase will be understood only after the spatially resolved spectroscopy capabilities of ALMA\\footnote{\\tt{http://www.eso.org/projects/alma/}} become available."
    },
    "0711/0711.3556_arXiv.txt": {
        "abstract": "In high-mass microquasars (HMMQ), strong interactions between jets and stellar winds at binary system scales could occur. In order to explore this possibility, we have performed numerical 2-dimensional simulations of jets crossing the dense stellar material to study how the jet will be affected by these interactions. We find that the jet head generates strong shocks in the wind. These shocks reduce the jet advance speed, and compress and heat up jet and wind material. In addition, strong recollimation shocks can occur where pressure balance between the jet side and the surrounding medium is reached. All this, altogether with jet bending, could lead to the destruction of jets with power $<10^{36}\\,\\rm{erg/s}$. The conditions around the outflow shocks would be convenient for accelerating particles up to $\\sim\\,$TeV energies. These accelerated particles could emit via synchrotron and inverse Compton (IC) scattering if they were leptons, and via hadronic processes in case they were hadrons. ",
        "introduction": "\\label{aba:sec1} The interaction between a jet and its environment allows the study of both the properties of the jet and its ambient medium\\cite{mar97,sax05}. Moreover, this interaction is important for the evolution of jets\\cite{per05}. In the case of extragalactic jets, it is known to be different in case of FRI and FRII sources; the collision and interaction of jets with galactic clouds and collective stellar winds is thought to be partly responsible for the jet mass loading and deceleration in the former\\cite{lai96}, whereas the latter can propagate and interact with the medium up to Mpc scales. In the case of microquasars, not much is known about the interaction of jets with their environments\\cite{hei02,bor08}, and only several sources show evidence of such an interaction (e.g. SS~433\\cite{zealey80}; XTE~J1550$-$564\\cite{corbel02}; Cygnus~X-1\\cite{gallo05}).  In HMMQ, the primary star suffers severe mass loss in the form of a supersonic wind\\cite{lam93}, which can embed the jet rendering their mutual interaction likely. Here we report on a work studying this interaction and its effects on the jet evolution and radiation. ",
        "conclusions": "We show that, with the reasonable overpressure and physical parameters in the jets and stellar winds of HMMQ, assuming negligible wind porosities, the jets form strong standing shocks due to interaction with the cocoons formed first, and the stellar wind from the primary later, on top of the transient reverse/bow shocks. We estimate that standing shocks must form within the system region for typical jet temperatures ($>10^{9-10} {\\rm K}$). Moreover, we have shown that jets with power $< 10^{36}\\,\\rm{erg/s}$ could be deflected and disrupted by the primary star wind. The influence of hydrodynamical instabilities such as Kelvin-Helmholtz is to be neglected on the basis of our results, as the evolution and dynamics of these jets are nonlinear from the start. In this scenario, the inertia of the jet compared to that of the wind is responsible for the fate of the former. Finally, the interaction of the jet head with the wind, and the jet sides with the cocoon and the stellar wind, in HMMQ, could lead to efficient non-thermal multiwavelength radiation, showing the importance of the role of the primary star for high energy emission in these sources, which is not just a provider of seed photons for IC scattering and photon photon absorption, or target nuclei for hadronic processes."
    },
    "0711/0711.4140_arXiv.txt": {
        "abstract": "We present a review of observed parameters of giant radio pulses, based on the observations conducted by our group during recent years. The observations cover a broad frequency range of about 3  octaves, concentrating between 600 and 4850 MHz. Giant pulses of both the Crab pulsar and the millisecond pulsar B1937+21 were studied with the 70-m Tidbinbilla, the 100-m GBT, 64-m Kalyazin and Westerbork radio telescopes.  We discuss pulse energy distribution, dependence of peak flux density from the pulse width, peculiarities of radio spectra, and polarization properties of giant radio pulses. ",
        "introduction": "Giant pulses (GPs) are abnormally short  (down to nanoseconds), bright (up to MJy), polarized (up to 100\\%) bursts of pulsar radio emission. They occur in the very narrow windows and are known only for a handful of pulsars.  Here we present the detailed study of two brightest sources of giant pulses -- the Crab pulsar and millisecond pulsar B1937+21.  Both the Crab pulsar and the PSR B1937+21 have many similar characteristics that can be a clue for understanding the giant pulse phenomenon. They both have an interpulse in their regular  profiles and  highest values of magnetic field at the light cylinder (about $10^6$~G). Popov et al.\\cite{MV2006} have recently proposed that the Crab regular profile consists of only GPs whereas the precursor observed at the frequencies below 1 GHz represents the regular profile. Hence, for both the Crab pulsar and the pulsar B1937+21 GPs occur at the very trailing edge of the regular emission profile.  Polarization properties are also similar. In general, GPs from the Crab pulsar are wider (up to $100~\\mu$s), mainly with some degree of linear polarization. However, Hankins et al. \\cite{Hankins}  showed that these pulses consist of very narrow (<2 ns) bursts with 100\\% circular polarization, either left or right. The number of these fundamental bursts in any individual broad GP was estimated by  Popov et al.\\cite{MV2006} to be about 100. Giant pulses from the B1937+21 reveal the broad spectrum of polarization properties: they can be almost 100\\% circular polarized, and they may have high linear polarization as well. This means that probably they also consist of narrower pulses.  Many similar properties are undoubtedly manifestations of the same physical process responsible for GP emission. However, there are still some differences, that represent either the individual character of these two pulsars or the influence of the ISM, or both. ",
        "conclusions": "It seems rather unlikely that any plasma mechanism is responsible for GP emission. Extermely short GP duration implies very small GP emission region, $10^2 - 10^3$ cm. For strongest GPs, volume density of the emission exceeds by many orders density of plasma energy, it is comparable with the density  of the energy of the magnetic field. Huge peak flux densities for the strongest GPs from the Crab pulsar detected at the Earth means an ultimate field strength of EM wave inside pulsar's magnetosphere. Interaction of such strong wave with plasma is very specific: the wave accelerates charged particles up to relativistic speed over one wave cycle, and the particles emit secondary waves at different frequency in different directions than the incident wave. The process may explain unusual HFC1 and HFC2 components in the profile of the Crab pulsar."
    },
    "0711/0711.4230_arXiv.txt": {
        "abstract": "{  Ring-shaped morphologies of nuclear star-forming regions within the central $40-200$~pc of disk galaxies have been barely resolved so far in three composite Sy~2 nuclei, the Sy~2 Circinus galaxy and in three non-AGN galaxies. Such morphologies resemble those of the standard $\\sim 1$\\,kpc-size nuclear rings that lie in the inner Lindblad resonance regions of disk galaxies and, if they have a similar origin, represent recent radial gas inflows tantalisingly close to the central supermassive black holes. We aim to identify the population of such ultra-compact nuclear rings (UCNRs) and study their properties in relation to those of the host galaxies. From archival {\\it Hubble Space Telescope} UV and \\ha\\ images and from dust structure maps of the circumnuclear regions in nearby galaxies, we analyse the morphology of the star formation and dust, specifically searching for ring structures on the smallest observable scales. In a sample of 38 galaxies studied, we have detected a total of four new UCNRs, $30-130$~pc in radius, in three different galaxies.  Including our confirmation of a previous UCNR detection, this yields a UCNR fraction of roughly 10\\%, although our sample is neither complete nor unbiased. For the first time we resolve UCNRs in two LINERs. Overall the UCNR phenomenon appears widespread and limited neither to late-type galaxies nor exclusively to AGN hosts. ",
        "introduction": "Star-forming nuclear rings, almost exclusively found in barred galaxies (Buta \\& Combes 1996), are attributed to shock focusing of gas near the location of one or more inner Lindblad resonances (ILRs; Schwarz 1984; Athanassoula 1992; Heller \\& Shlosman 1996; see also Shlosman 1999; Knapen 2005 for reviews).  Their sizes, of typically $\\sim 1$\\,kpc, allow them to be easily resolved in ground- and space-based observations. On the other hand, high resolution imaging with the {\\it Hubble Space Telescope (HST)} has led to the detection of three unusually compact nuclear star-forming regions, where unambiguous signatures of recent episodes of star formation (SF) coexist with the active galactic nucleus (AGN)---so-called {\\it composite} Sy~2 nuclei (Gonz{\\'a}lez Delgado et al. 1998). These barely resolved star-forming regions exhibit morphologies which are not incompatible with those of ``standard'' nuclear rings---an issue that was never addressed in the literature. The prevalence of such ``ultra-compact nuclear rings'' (UCNRs) and their origin are currently unknown. While conducting a first comprehensive survey of UCNRs, supplemented by statistical, phenomenological and dynamical studies of their origin and evolution, we report here the results of the initial search of {\\it HST} archival imaging, which has revealed a number of UCNRs in NGC~2985, NGC~4579 and NGC~4800.  Besides the composite Sy~2 nuclei mentioned above, star-forming UCNRs of less than 200\\,pc in radius are known in the Sy~2 Circinus galaxy (Wilson et al.~2000) and in galaxies without AGN: NGC~5248 (Laine et al. 2001; Maoz et al. 2001; Jogee et al. 2002) and possibly in NGC~3245 (Barth et al.~2001a) and IC~342 (B\\\"oker, F\\\"orster-Schreiber, \\& Genzel 1997).  The smallest of these UCNRs are around $\\sim40$\\,pc in radius. Even more tantalising is the detection of molecular rings or disks of $\\sim 1-100$\\,pc in low-luminosity AGN, such as M51 (Kohno et al. 1996) or in NGC~1068 (Galliano et al. 2003). The Milky Way, a ``normal'' galaxy, hosts a massive molecular ring within $\\sim 1.6 - 7$\\,pc from the Galactic Centre.  Presently, it is not clear that all these rings have a similar origin---but this is a possibility. A major open question is whether the UCNR population is the small-radius tail of the general population of nuclear rings (e.g., shown in Laine et al. 2002) or whether it characterises a different physical phenomenon.  Nuclear rings in general are taken to be by-products of gas inflow in response to gravitational torques induced by bars or galaxy interactions, so it is significant that rings are found so close to the region dominated by the supermassive black hole (SMBH). UCNR hosts are nuclear starburst galaxies (by definition) and potentially related to the triggering and fuelling of AGN.  There are two closely related scenarios that could account for the tiny UCNRs compared to their ``normal'' counterparts, namely, a shallower radial mass distribution in the central kpc of the host galaxy, implying a smaller bulge, or ring  formation in response to a non-axisymmetric mass distribution in the inner few hundred pc, that tumbles exceedingly fast. In both cases this implies that the ILR region is located close to the centre. The first scenario would indicate that one might expect a deficiency of UCNRs in early-type spirals. The second scenario favours the presence of nuclear gaseous or stellar bars (Shlosman 1999). These alternatives can be distinguished using a number of observational tests. ",
        "conclusions": "In this paper, we present the discovery of UCNRs in three spiral galaxies, identified from {\\it HST} archival UV and \\ha\\ imaging, which, for the first time, include LINERs. Overall, four galaxies among our sample of 38 display this phenomenon. Although the sample is neither complete nor unbiased, this rough estimate puts the frequency of UCNRs at $\\sim 10\\%$. Moreover, they are found now in Sy~2s, LINERs and non-AGN galaxies---hence their origin lies in galactic dynamical processes and not in radiative processes related to the AGN activity. The UCNRs found range in radius from 30 to 130\\,pc, and lie in galaxies of types Sab and Sb.  Two of our galaxies host UCNRs that look very much like classical nuclear rings and are probably resonance rings, formed as a consequence of gas inflow slowing down near one or more ILRs. One, NGC~2985, has close and influencing companions, while the other, NGC~4800, is weakly barred.  The UCNR in the third galaxy, NGC~4579, is off-centre which can be explained as a result of a superposition of $m=1$ and $m=2$ perturbations. The inner arcs detected in this object appear to have a natural explanation in the combined emission from the jet-induced SF and the postshock gas.  The small number of UCNRs discovered does not yet allow us to reach any conclusions on their formation or evolution as a class. However, these results suggest  that there is no qualitative difference between these rings and the nuclear rings seen in disk galaxies, although they may constitute a separate population from the kpc-size rings. More UCNRs are expected to be found in our ongoing search, which will allow us to establish the fraction of galaxies with UCNRs, and to study in detail the relations to their host galaxies. Given the proximity of UCNRs to central SMBHs, these findings will be of great interest to the issue of fuelling of AGN activity, as well as to that of galactic dynamics on scales of tens to a hundred pc."
    },
    "0711/0711.1309_arXiv.txt": {
        "abstract": "This paper revisits the classical Kennicutt method for inferring the stellar initial mass function (IMF) from the integrated light properties of galaxies.  The large size, uniform high quality data set from the Sloan Digital Sky Survey DR4 is combined with more in depth modeling and quantitative statistical analysis to search for systematic IMF variations as a function of galaxy luminosity. Galaxy H$\\alpha$ equivalent widths are compared to a broadband color index to constrain the IMF.  This parameter space is useful for breaking degeneracies which are traditionally problematic.  Age and dust corrections are largely orthogonal to IMF variations.  In addition the effects of metallicity and smooth star formation history e-folding times are small compared to IMF variations.  We find that for the sample as a whole the best fitting IMF slope above 0.5 ${\\rm M}_\\odot$ is $\\Gamma =1.4535$ with a negligible random error of $\\pm 0.0004$ and a systematic error of $\\pm 0.1$.  Galaxies brighter than around ${\\rm M}_{r,0.1}=-20$ (including galaxies like the Milky Way which has ${\\rm M}_{r,0.1}\\sim -21$) are well fit by a universal $\\Gamma \\sim 1.4$ IMF, similar to Salpeter, and smooth, exponential star formation histories (SFH). Fainter galaxies prefer steeper IMFs and the quality of the fits reveal that for these galaxies a universal IMF with smooth SFHs is actually a poor assumption.  Several sources of sample bias are ruled out as the cause of these luminosity dependent IMF variations.  Analysis of bursting SFH models shows that an implausible coordination of burst times is required to fit a universal IMF to the ${\\rm  M}_{r,0.1}=-17$ galaxies.  This leads to the conclusions that the IMF in low luminosity galaxies has fewer massive stars, either by steeper slope or lower upper mass cutoff, and is not universal. ",
        "introduction": "A precise measurement of the stellar initial mass function (IMF) and its functional dependence on environmental conditions would impact astronomy over a wide range of physical scales.  It would be of great help to theorists in untangling the mysteries of star formation and it is a key input in spectral synthesis models used to interpret the observed properties of galaxies both nearby and in the early universe.  The current question is whether the IMF is \\textit{universal}-- the same regardless of time and environmental conditions.  \\citet{Ken98} concisely states the current understanding of IMF universality.  It is difficult to believe that the IMF is universal given the diversity of galaxy types, environments, star formation rates, and populations within galaxies over the range of observable lookback times.  On the other hand, while IMF measurements do vary they are all consistent with a universal IMF within measurement errors and sampling statistics.  The only way to proceed then is to strive for smaller measurement errors and improving sample sizes.  A definitive theoretical derivation of the IMF does not yet exist. Theoretical approaches to the IMF usually center around the Jeans mass, $M_J$, the mass at which a homogeneous gas cloud becomes unstable.  At first the collapse of a cloud is isothermal and the Jeans mass decreases which leads to fragmentation of the cloud \\citep{Hoyle53}.  Both \\citet{Rees76} and \\citet{Low76} suggested that at some point during the cloud collapse the line cooling opacity becomes high enough that the collapse is no longer isothermal. At this point the Jeans mass increases and fragmentation stops.  The minimum Jeans mass is the smallest fragment size at this point and it provides a lower limit to size of the stars formed.  Authors have calculated Jeans masses and minimum Jeans masses using a variety of methods.  In the classical derivation of the Jeans mass $M_J \\propto T^{3/2} \\rho^{-1/2}$.  \\citet{Low76} finds that $M_{J,min} \\propto m^{-16/7} \\kappa^{-1/7}$ where $m$ is the mass of gas atoms or molecules and $\\kappa$ is the opacity at the final fragmentation.  More recently turbulence in clouds has been studied.  \\citet{Padoan97} found $M_{J,min} \\propto n^{-1/2} T^2 \\sigma_v^{-1}$ where $n$ is the number density and $\\sigma_v$ is the velocity dispersion of the gas.  Other investigators have looked at the hierarchical fractal geometry of molecular clouds, thought to arise from turbulence, as a generator of the IMF (e.g. \\citet{Elmegreen97}).  On a related note \\citet{Adams96} point out that molecular clouds exhibit structure on all resolvable spacial scales suggesting that as no characteristic density exists for the clouds neither does a single Jeans mass.  They develop a semi-empirical model for determining the final masses of stars from the initial conditions of molecular clouds without invoking Jeans mass arguments and use it to construct IMF models.  The key components of their model are sound speed and rotation rate of cloud cores and the idea that stars help determine their final masses through winds and outflows.  But from the beginning the study of the IMF has been driven by measurements. In 1955 Salpeter was the first to make a measurement of the IMF inferring it from his observed stellar luminosity function \\citep{S55}.  We parameterize the IMF by: \\begin{equation} \\frac{{\\rm d}n}{{\\rm d}\\log m} \\propto \\left \\{ \\begin{array}{ccc} -0.5 & \\mbox{for}& 0.1 <m/{\\rm M}_\\odot < 0.5\\\\ -\\Gamma & \\mbox{for}& 0.5 < m/{\\rm M}_\\odot < 120 \\end{array} \\right. \\label{eq:imfbaldry} \\end{equation} following \\citet{BG03}.  Salpeter found that $\\Gamma = 1.35$.  It is often overlooked that his measurement only covered masses for which $0.4 \\lesssim m/\\rm{M}_\\sun \\lesssim 10$.  Nonetheless his original measurement is surprisingly consistent with modern values over a wide range of masses.  The IMF in equation \\ref{eq:imfbaldry} is similar to the Salpeter IMF for $\\Gamma = 1.35$.  The difference is that there are fewer stars with masses less than 0.5 ${\\rm M}_\\odot$.  We adopt a two part power law as there is agreement amongst several authors that there is a change in the IMF slope near 0.5 ${\\rm M}_\\odot$ \\citep{Kroupa01}.  The technique we will use is not sensitive to the IMF at low masses so we assume a constant value in that regime.  Salpeter's idea continues to be used today in IMF measurements of resolved stellar populations.  The technique can be applied to field stars as well as clusters.  However Salpeter's method has several inherent limitations.  The nature of stars presents a challenge.  On the high mass end stars are very luminous, but live only a few million years, while on the low end stars are faint but have lifetimes many times longer than the current age of the universe.  There are very few star clusters which are both young and close enough to allow us access to the IMF over the full mass range.  In addition, the main sequence mass-luminosity relationship is a function of age, metallicity, and speed of rotation in addition to mass.  It is not yet well-known at the low and high mass extremes.  Unresolved binaries can also affect the measured IMF \\citep{Kroupa01}.  The light from unresolved binaries is dominated by the more massive of the pair.  As a result the less massive star is typically not detected which leads to a systematic under-counting of low mass stars.  $42\\% \\pm 9\\%$ of main sequence M stars \\citep{Fischer92} and 43\\% of main sequence G stars \\citep{Duquennoy91} are primary stars in multiple star systems.  These are both lower limits as some companions may have eluded detection.  As roughly half of stars are in multiple systems it has a potentially large effect on the observed luminosity function.  Except for at the high mass end field stars in the solar neighborhood offer the best statistics for IMF measurements.  However the solar neighborhood IMF is found to be deficient in massive stars when compared to other galaxies.  For example the \\citet{MS79} and \\citet{Scalo86} solar neighborhood IMFs are rejected by integrated light approaches, i.e.  \\citet{K83}, hereafter K83, \\citet{K94} (KTC94), and \\citet{BG03}.  Analysis of individual star clusters can be used to detect IMF variations.  As methods and data quality can vary between authors comparisons between individual clusters are difficult.  However \\citet{Phelps93} studied eight young open clusters with the same technique.  On the extremes they measured $\\Gamma =1.06 \\pm 0.05$ for NGC 663 and $\\Gamma = 1.78 \\pm 0.05$ for NGC 581 over masses from around 1 to 12 M$_\\sun$.  \\citet{Kroupa01} considered the case in which one could have perfect knowledge of the masses of all stars in a cluster.  Clusters have a finite size so even with no measurement errors uncertainty arises from sampling the underlying IMF.  He shows that the observed scatter in IMF power law values above 1 M$_\\sun$ can be accounted for by sampling bias for clusters with between $10^2$ and $10^3$ members (the \\citet{Phelps93} clusters have memberships in this range).  Furthermore, dynamical evolution of clusters can affect measurements of the IMF.  This can happen by preferentially expelling lower mass stars from the cluster and by breaking up binary systems within the cluster.  In total \\citet{Kroupa01} finds that for stars with $m \\gtrsim 1{\\rm M}_\\sun$ the spread in the observed IMF power-law slopes in clusters is around 1 when both binary stars and sampling bias are considered even when the underlying IMFs are identical.  Stochastic processes can also influence the ability to detect systematic IMF variations.  O and B stars produce ionizing photons and cosmic rays which affect the surrounding nebula.  However the probability of creating one of these massive stars is comparatively small.  If one of these massive stars happens by chance to form first it may drive up the nebular temperature and depress the formation of less massive stars compared to regions without massive stars \\citep{Robberto04}.  An alternative to IMF measurements based on the stellar luminosity functions of resolved stellar populations is to infer the characteristics of stellar populations from the integrated light of galaxies using spectral synthesis models.  The advantage of using integrated light techniques is that many of the problems plaguing IMF investigations of resolved stellar populations are avoided.  Stochastic effects are washed out over a whole galaxy.  Unresolved multiple star systems are irrelevant.  The number of observable galaxies is large and their environments span a much larger range than those of Milky Way clusters.  Integrated light techniques can be applied to the high redshift universe.  This creates a strong motivation to develop IMF techniques and test them for galaxies in the local universe which can later be used to probe earlier stages of galaxy formation.  The assumption of a universal IMF has a huge influence on the interpretation of the high redshift universe.  The reionization of the universe and the Madau plot-- the global star formation rate of the universe over cosmic history-- are two areas where IMF variations could impact the current picture of galaxy formation and evolution.  On the downside conclusions are dependent on the stellar evolution models used, which are not well constrained at high masses or with horizontal branch stars at low metallicities.  The biggest problem is that changes in the IMF, star formation history (SFH) or age of a galaxy model can have similar effects in the resulting spectral energy distribution (SED).  Any integrated light technique needs to address these degeneracies.  It is also difficult to probe the IMF at sub-solar masses using integrated light.  On the level of galaxies the concept of a universal IMF has recently become more complicated.  A number of recent studies have investigated the effect of summing the IMF in individual clusters over a galaxy in the presence of power law star cluster mass functions.  \\citet{WK05} argue that the integrated galaxial IMF will appear to vary as function of galactic stellar mass even if the stellar IMF is universal.  However \\citet{Elmegreen06} argues that the galaxy wide IMF should not differ from the IMF in individual clusters based on analytical arguments and Monte Carlo simulations.  Our approach can only measure the IMF averaged over whole galaxies and cannot address this distinction.  Even so systematic variations of any kind from the Salpeter slope have not been measured outside of some evidence for non-standard IMFs in low surface brightness galaxies (LSB) and galaxies experiencing powerful bursts of star formation \\citep{Elmegreen06}.  Either way, observational evidence for systematic variations of the IMF in galaxies would be highly valuable.  The plan of this paper is as follows:  \\S 2 explains our method for constraining the IMF.  In \\S 3 we describe the SDSS data and our sample selection.  \\S 4 details our modeling scheme.  In \\S 5 we discuss our statistical techniques.  \\S 6 reports our results.  In \\S 7 we check our results against the H$\\delta_{\\rm A}$ distribution of the data and \\S 8 presents our conclusions. ",
        "conclusions": "The goal of this paper was to revisit the K83 method for inferring the IMF from integrated stellar populations and to harness the richness of the SDSS data, improved spectral synthesis models and greater computational power available today to make a state-of-the-art measurement of the IMF.  The quality of the SDSS spectroscopy allowed us to address several of the limitations of K83 and KTC94-- we resolve the [\\ion{N}{2}] lines (a significant improvement in the accuracy of the H$\\alpha$ EWs of individual galaxies), eliminate contamination from AGN, make extinction corrections for individual galaxies and fit underlying stellar absorption of H$\\alpha$.  We succeeded in achieving more accurate EWs for individual galaxies.  The median total EW error for our sample is 17\\% compared to a 10\\% uncertainty in EWs combined with a 20-30\\% uncertainty in the extinction correction for K83.  We expanded the grid of models to allow for a range of ages and metallicities. We used $\\chi^2$ minimization to go beyond differentiating between two or three IMF models to actual fitting for the best IMF slope.  The vast size of the SDSS sample allowed us to both drive down random errors and to cut the data into narrow parameter ranges which were still statistically viable.  The size of the DR4 sample yielded $\\Delta \\Gamma=0.0011$ 95\\% confidence region due to random error for the sample as a whole.  Even the volume limited ${\\rm M}_{r,0.1}=-17$ luminosity bin with only 329 objects has a random error of $\\Delta \\Gamma=0.0086$.  Only in bins with fewer than 10 objects do the random errors become significant.  Our IMF fitting is therefore dominated by systematics.  Originally we believed that our systematics would be dominated by the effects of SFH discontinuities.  However we conducted several experiments where we selected populations of galaxies from models with bursting or gasping SFHs and gave them measurement errors consistent with those in the ${\\rm M}_{r,0.1}=-17$ luminosity bin.  To our surprise our $\\chi^2$ minimization revealed the true IMFs with $\\Delta \\Gamma=0.01$.  The main effect was to reduce the quality of the fits.  This is due to the fact that H$\\alpha$ EWs return to nominal levels in a relatively short time after SFH discontinuities.  Another way to estimate the size of the systematic errors is to look at the trends of the ${\\rm M}_{r,0.1}=-21$ and $-22$ luminosity bins, because they have the largest membership, in Figures \\ref{fig:redshift}, \\ref{fig:aperture} and \\ref{fig:extinction}.  Assuming that the IMF is universal and that our method is perfect we should get the same answer for any subset of the data we might choose.  The largest ranges are $\\Delta \\Gamma \\sim 0.12$ for redshift binning, $\\Delta \\Gamma \\sim 0.19$ by aperture and $\\Delta \\Gamma \\sim 0.19$ by extinction.  Conservatively then the systematic error is $\\pm 0.1$.  There are two points to be kept in mind about this estimate of the systematic error.  For one it is the systematic error in the exact value of $\\Gamma$. Even in Figure \\ref{fig:avap} where more narrow bands of measured extinction and aperture fraction are considered the same trends with luminosity are seen as with the sample as a whole.  The relative systematics between luminosity bins in these narrow slices is much smaller.  The second thing to remember is the way in which we empirically defined our systematic error discounts the possibility of IMF variations.  What we have called systematics could actually be science.  If galaxies have radial IMF gradients or if dust content plays a strong role in star formation the systematic error could be much smaller.  The main area in which we were unable to improve upon the K83 and KTC94 studies is that they were able to match the aperture size to the galaxies which avoids the issue of aperture effects.  In spite of a more quantitative approach, like K83 and KTC94 the results are mostly qualitative.  However there are four key results from our investigation.  First, for galaxies brighter than ${\\rm M}_{r,0.1}\\sim-20$ the best fitting IMFs are Salpeter-like ($\\Gamma \\sim 1.4$).  In addition the assumption of a universal IMF and smoothly varying SFHs is a good fit.  This is reassuring as it follows the conventional wisdom and provides confidence that the method works.  Secondly, galaxies fainter than ${\\rm M}_{r,0.1}\\sim-20$ are best fit by steeper IMFs with larger fractions of low mass stars.  For these galaxies a universal IMF and smooth SFH is a poor assumption.  This result is in qualitative agreement with evidence that LSBs have bottom-heavy IMFs \\citep{Lee04}.  Thirdly, while breaking the IMF-SFH degeneracy for individual galaxies using the H$\\alpha$ EW and color is hopeless, for a statistical sample of galaxies the degeneracy can be broken.  Lastly, given our analysis of discontinuous SFHs it appears that the IMF is not universal in low luminosity galaxies and fewer massive stars are being created in these galaxies.  It is worth mentioning the main caveat of our $\\Gamma$ values again.  As illustrated in Figure \\ref{fig:oddimf} IMF parameterizations are themselves degenerate in our parameter space.  Increasing the IMF slope has a similar effect to lowering the highest mass stars that are formed or increasing the fraction of intermediate mass stars.  This method cannot explicitly determine if two populations have the same underlying IMF.  Figure \\ref{fig:all-like} shows that for the sample as a whole the \\citet{Scalo98} three part power law yields nearly the same result as our two part power law.  However our method is sensitive in many cases if the IMFs are different.  In terms of star forming cloud temperatures the harsher ambient radiation and larger number of sources of cosmic rays present in more luminous galaxies agree qualitatively with our results.  With the extra energy hitting the star forming clouds larger masses may be needed for contraction and fractionization may end sooner, suppressing the formation of less massive stars \\citep{Larson98}.  \\citet{CCT05} find that while the \\ion{H}{2} regions of the luminous grand design spiral NGC 5457 (M 31) can be reproduced by a single Salpeter IMF, for the low luminosity flocculent galaxy NGC 4395 a blend of two IMFs is required.  However, such trends are not seen in studies of well-resolved stellar populations \\citep{Kroupa02}.  Another explanation for the absence of massive stars is that the massive stars are there, but are not visible.  Extinction to the center of star forming regions, where massive stars preferentially exist, can reach ${\\rm A_V}\\sim 20$ \\citep{Engelbracht98}.  However the low luminosity galaxies have the lowest observed extinctions (see Table \\ref{tab:lumbins}) which is the opposite of what would be expected given our IMF results.  It is also possible that the IMF is in fact universal, but the way in which it is sampled in embedded star clusters leads to an integrated galaxial IMF which varies from the true IMF.  \\citet{WK05} use a universal IMF with the assumption that stars are born in clusters where the maximum cluster mass is related to the star formation rate.  For a range of models this leads to a narrow range for the apparent IMF in high mass galaxies.  For low mass galaxies the IMF is steeper with a wider range of slopes.  The results here are in qualitative agreement for some of the integrated galactic IMF scenarios in  \\citet{WK05} given that there should be a rough correlation between galaxy luminosity and mass.  Once again, \\citet{Elmegreen06} argues that the galaxy wide IMF should be the same as the IMF in clusters regardless.  In light of the theory of \\citet{WK05}, whether the results of this paper speak to a relationship between environment and the formation of individual stars is open to interpretation.  However the impact on the modeling and interpretation of the properties of galaxies is clear. \\citet{KWK07} note that the integrated galaxial IMF is the correct IMF to use when studying global properties of galaxies.  Even if the IMF of stars is in truth universal it may currently be misused in the modeling of galaxies.  Furthermore a varying integrated galaxial IMF could open the door to new insights in galaxy evolution.  For instance, \\citet{KWK07} suggest that the observed mass-metallicity relationship in galaxies naturally arises from a variable integrated galaxial IMF similar to the results of this paper.  Future work will expand in several directions.  The success constraining the IMF with only two observed parameters (albeit carefully chosen to be orthogonal to systematic errors) motivates a more expansive analysis with more parameters.  Information from wavelength regimes beyond the SDSS can be used. For instance the absorption strength and P-Cygni profile shape of the \\ion{C}{4} $\\lambda 1550$ line due to massive stars is sensitive to the IMF slope and upper mass cutoff \\citep{LRH95}.  However it could be contaminated by absorption from the interstellar medium and would need to be disentangled \\citep{Shapley03}.  A full Markov Chain Monte Carlo analysis could be implemented fitting to multiple spectral features and marginalizing over a range of SFHs.  In addition to luminosity, surface brightness and gas phase metallicity can be tested for systematic IMF variations.  While the results for luminous galaxies are already dominated by systematics, the continued progress of the SDSS can provide better statistics for a more detailed analysis of what physical processes are behind the IMF variations in faint galaxies."
    },
    "0711/0711.2831_arXiv.txt": {
        "abstract": "Several SU~UMa-type dwarf novae, in particular, WZ~Sge-type stars tend to exhibit rebrightenings after superoutbursts.  The rebrightening phenomenon is problematic for the disk instability theory of dwarf novae since it requires a large amount of remnant matter in the disk even after superoutbursts.  Here, we report our optical and infrared observations during the first-ever outburst of a new dwarf nova, SDSS~J102146.44$+$234926.3.  During the outburst, we detected superhumps with a period of $0.056281\\pm 0.000015\\,{\\rm d}$, which is typical for superhump periods in WZ~Sge stars.  In conjunction with the appearance of a long-lived rebrightening, we conclude that the object is a new member of WZ~Sge stars.  Our observations, furthermore, revealed infrared behaviors for the first time in the rebrightening phase of WZ~Sge stars.  We discovered prominent infrared superhumps.  We calculate the color temperature of the infrared superhump source to be 4600---6400~K.  These temperatures are too low to be explained with a fully-ionized disk appearing during dwarf nova outbursts.  We also found a $K_{\\rm s}$-band excess over the hot disk component.  These unprecedented infrared activities provide evidence for the presence of mass reservoir at the outermost part of the accretion disk.  We propose that a moderately high mass-accretion rate at this infrared active region leads to the long-lived rebrightening observed in SDSS~J102146.44$+$234926.3. ",
        "introduction": "Dwarf novae are a sub-class of cataclysmic variables which contain a Roche-lobe filling secondary star and an accreting white dwarf (for a review, see \\cite{war95book}).  Dwarf nova outbursts are characterized with amplitudes of 2---5~mag and durations of a few days.  SU~UMa-type stars are dwarf novae showing two types of outbursts, that is, normal outbursts and superoutbursts which have longer durations ($\\sim 10$---$20$~d) and amplitudes $\\sim 1\\,{\\rm mag}$ larger than those of normal outbursts.  Superoutbursts are characterized by the appearance of periodic short-term modulations, called ``superhumps'', whose period is slightly longer than the orbital period (\\cite{war95suuma}).  Dwarf nova outbursts are caused by a sudden increase of mass-accretion rate in an accretion disk.  The most widely accepted theory for SU~UMa stars is the thermal---tidal instability model for the accretion disk (\\cite{osa89suuma}).  The thermal instability model considers two kinds of thermally stable states of the disk, that is, a cool disk of neutral hydrogen gas and a hot disk of ionized hydrogen gas, corresponding to the quiescent and outburst states, respectively. According to the disk instability model, the cool disk becomes thermally unstable when the density of the accumulated gas reaches a critical value.  The disk, then, experiences a state transition to the hot disk, observed as an outburst (\\cite{osa74DNmodel}; \\cite{hos79diskinst}; \\cite{mey81DNoutburst}; \\cite{sma82alphadisk}; \\cite{can82DNburst}).  The outburst leads to an expansion of the disk.  In the case that the disk has an enough angular momentum, its outer edge reaches the 3:1 resonance radius, which triggers the tidal instability (\\cite{whi88tidal}).  The disk becomes eccentric, as a result, the strong tidal torque yields an additional dissipation of the angular momentum in the disk, which leads to superoutbursts. Because of the eccentricity of the disk, periodic modulations associated with the orbital motion of the secondary star are expected. They are actually observed as superhumps.  The superhump period is longer than orbital periods due to the precession of the eccentricity wave (\\cite{pat05epsilon}).  The disk instability model can explain several basic behaviors of ordinary dwarf novae, for example, their frequency of outbursts and two kinds of outbursts in SU~UMa stars (for a review, see \\cite{osa96review}).  This standard model is, however, required to be modified in order to reproduce rebrightenings after superoutbursts which are observed in several SU~UMa stars.  A peculiar subclass, WZ~Sge-type dwarf novae, in particular, tend to exhibit various types of rebrightenings (\\cite{how95TOAD}; \\cite{kat01hvvir}).  WZ~Sge-type stars have shortest orbital periods ($P_{\\rm orb}\\sim 85\\,{\\rm min}$) and longest outburst cycles ($\\gtrsim 10\\,{\\rm yr}$) among dwarf novae.  Based on characteristics of light curves, rebrightenings can be divided into three types; i)~a single short rebrightening (RZ~Leo, \\cite{ish01rzleo}; ASAS~J002511$+$1217.2, \\cite{tem06J0025}), ii)~repetitive short rebrightenings, or sometimes called ``echo outbursts'' (WZ~Sge, \\cite{pat02wzsge}; EG~Cnc, \\cite{pat98egcnc}, \\cite{kat04egcnc}; SDSS~J080434.20$+$520349.2, \\cite{pav07j0804}), and iii) a long-lived plateau (AL~Com, \\cite{nog97alcom}, \\cite{ish02wzsgeletter}; CG~CMa, \\cite{kat99cgcma}; TSS~J022216.4$+$412259.9, \\cite{ima06j0222}; V2176~Cyg, \\cite{nov01v2176cyg}).  Additionally to WZ~Sge stars, several short-period SU~UMa stars also exhibit single or a few short rebrightenings (\\cite{bab00v1028cyg}; \\cite{uem02j2329}; \\cite{ima06j0137}).  On the other hand, several WZ~Sge stars show no sign of rebrightenings (HV Vir, \\cite{kat01hvvir}).  The mechanism of rebrightenings and the origin of their diversity are poorly understood.  In a simple picture based on the disk instability model, most of the mass and angular momentum are lost from the disk by the end of superoutbursts (\\cite{osa96review}).  The long-lived rebrightening phase is problematic because it apparently requires an unknown delayed mass-accretion even after the termination of superoutbursts when the disk should nearly be empty (\\cite{kat04egcnc}).  This discrepancy can be reconciled if a mass-transfer rate from the secondary is enhanced due to the irradiation by the disk (\\cite{war98CVreviewWyo}; \\cite{ham00DNirradiation}; \\cite{pat02wzsge}).  This scenario, however, has a difficulty to reproduce sudden onset and cessation of the rebrightening phase (\\cite{pat02wzsge}; \\cite{osa03DNoutburst}; \\cite{osa04emt}).  \\citet{kat98super} propose an alternative scenario that a considerable amount of matter left over the 3:1 resonance radius is responsible for the delayed accretion.  The remnants at the outermost region is expected to work as a mass reservoir, leading to the additional mass accretion during the rebrightening phase (\\cite{hel01eruma}; \\cite{osa01egcnc}; \\cite{kat04egcnc}).  The behavior of the outermost region of the accretion disk, hence, definitely provides crucial clues for the rebrightening phenomenon. Infrared (IR) observations can be a strong tool to detect variations from a low temperature region at the outermost accretion disk.  Such IR observations during the rebrightening phase have, however, been difficult because of low frequencies of outbursts of WZ~Sge stars.  Here, we report results of our observation of a new variable star which was a promising candidate for WZ~Sge stars.  The discovery of its outburst was reported by E.~J.~Christensen in Nov. 2006 (\\cite{chr06cbet746}).  There is a quiescent counterpart, SDSS J102146.44$+$234926.3 (hereafter, J1021) in the Sloan Digital Sky Survey (SDSS) DR5 database with magnitudes of $g=20.74$ and $r=20.63$ (\\cite{SDSSDR5}).  We conducted an international campaign for J1021, in which we succeeded in performing IR time-series observations during a rebrightening phase for the first time.  In the next section, we describe our photometric observations.  In section~3, results of our observations are presented.  In section~4, we discuss the results and the nature of the rebrightening phenomenon.  In the final section, we summarize our findings. ",
        "conclusions": "\\subsection{Infrared behavior during dwarf nova outbursts in general}  In the previous section, we reported the IR activity during the rebrightening phase of J1021.  The uniqueness of its IR behavior among SU~UMa stars are, however, unclear due to the lack of enough samples for IR observations during dwarf nova outbursts.  \\citet{nay87oycar} reported optical and IR observations during a superoutburst of an ordinary SU~UMa-type dwarf nova, OY~Car.  In this paper, optical and IR superhumps have similar amplitudes, that is, $\\sim 0.3$~mag both in $B$- and $J$-bands.  They are consistent with the classical picture of the accretion disk in dwarf nova outbursts, in which a superhump source having 8000---10000~K lies at the outermost part of the disk having $\\gtrsim 10000$~K.  The IR behavior of OY~Car is totally different from that observed in the rebrightening phase of J1021.  To date, this observation of OY~Car is all information about IR behavior during superoutbursts.  About J1021, furthermore, we have no IR data during the main outburst which should be compared with those during the rebrightening.  In order to reveal the nature of the IR active region, future IR observations are required during superoutbursts and rebrightenings of ordinary SU~UMa and other WZ~Sge stars.  \\subsection{Validity of the color temperature}  In the previous section, we calculated the color temperatures of superhump sources from observed colors.  As a result, we obtained low color temperatures, which were rather close to those observed at quiescence.  The color temperature provides a good approximation if the source is optically thick.  In dwarf novae, however, the optical light from the accretion disk at quiescence is, in general, dominated by the optically thin emission (\\cite{how04wzsge}, and references therein).  Even in outbursts, the continuum radiation at the outermost region is predicted to be optically thin (\\cite{wil80ADemissionline}).  In the case of J1021 during the rebrightening, the temperature of the IR superhumps was rather close to the quiescent one, although the entire disk luminosity was still high.  The radiation mechanism of the IR superhump is, hence, less obvious.  If the IR emission is optically thin, the color temperatures of the low temperature region in table~\\ref{tab:bbsim} may have no complete validity as an indicator of real temperatures. We, then, need to solve the radiation transfer in the disk in order to reproduce the observed light curves, which is, however, beyond the scope of this paper.  Results of such a full calculation, on the other hand, have only little influence on our order-of-magnitude discussion. IR spectroscopic observations are required to conclude whether the outer low temperature region is optically thin or thick during the rebrightening phase.  \\subsection{Delayed maximum of rebrightenings in WZ~Sge stars}  Among 5 systems showing long plateau-type rebrightenings, 3 systems exhibit a precursor just before their rebrightening maximum, as mentioned in \\textsection~3.2.  In ordinary SU~UMa stars, precursors are occasionally observed just before supermaxima.  They are interpreted as a normal outburst triggering a subsequent superoutburst (\\cite{osa03DNoutburst}).  After the precursor maximum, the object starts a rapid fading since it takes a few days for the full evolution of the tidal instability.  When the tidal dissipation begins working effectively, the rapid fading is terminated by a rising trend toward the supermaxima.  The similar mechanism may also work in the structure of the long plateau-type rebrightening.  Light curves between precursors and rebrightening maxima are, however, peculiar in WZ~Sge stars compared with the precursor---supermaximum structure in ordinary SU~UMa stars.  Several WZ~Sge stars keep a gradually rising trend after initial rapid risings of rebrightenings. AL~Com and V2176~Cyg, for example, reached the maximum of rebrightenings $\\sim 5$ and $\\sim 4\\,{\\rm d}$ after the precursors, respectively (\\cite{nog97alcom}; \\cite{nov01v2176cyg}). TSS~J022216.4$+$412259.9 is an extreme case, in which a gradual rising apparently continued for $\\sim 8\\,{\\rm d}$ until the rebrightening maximum (\\cite{ima06j0222}).  This characteristic is clearly different from that observed in ordinary superoutbursts in which objects reach their supermaximum within a day after the onset of outbursts (\\cite{war95suuma}).  While J1021 shows no delay of the rebrightening maximum, we discuss the delayed maximum in terms of the mass reservoir at the outermost part of the disk.  According to the disk instability model, the rising speed corresponds to the propagation speed of the heating wave in the hot disk (\\cite{min85DNDI}).  The delayed rebrightening maximum is possibly attributed to a late propagation of the heating wave into the outer mass reservoir, namely the IR active region.  If the IR superhump observed in J1021 is a characteristic feature for the rebrightening phase, however, this scenario is unlikely because it would yield an entire hot disk and prominent optical superhumps.  Another possible scenario for the delayed rebrightening maximum is that the inner hot disk continuously receives a sufficient amount of gas from the outer mass reservoir.  In the case of J1021, the tidal dissipation effectively works at the outer low temperature region as demonstrated by the IR superhumps.  We can, thus, naturally expect a significant mass accretion from the outer cool to the inner hot part of the disk.  \\subsection{Origin of the $K_{\\rm s}$-band excess}  During the rebrightening phase, the $K_{\\rm s}$-band flux exhibited a clear excess over the blackbody component which dominated between the $V$ and $J$ ranges, as mentioned in \\textsection~3.3.  A noteworthy feature of the $K_{\\rm s}$-band light curve is a relatively rapid evolution compared with $V$- and $J$-band ones.  The $V$- and $J$-band emission is presumably dominated by the optically thick emission from the disk.  The evolution time-scale of them is a viscous diffusion time-scale of the hot optically thick disk.  The time-scale becomes longer in an outer, in other words, lower temperature region of the disk.  The rapid evolution of the $K_{\\rm s}$-band emission is, thereby, inconsistent with the theory of optically thick disk.  It implies that the $K_{\\rm s}$-band emission is not from an optically thick disk, but an optically thin region which is located outside the optically thick disk.  The source of the $K_{\\rm s}$-band emission might be expelled from the optically thick disk at the onset of the rebrightening as a result of a rapid expansion of the hot disk.  \\subsection{On the nature of the periodicity in the early fading tail}  The periodic modulations during the early fading tail has a slightly shorter period than the superhumps, as can be seen from figures~\\ref{fig:pdm} and \\ref{fig:o-c}.  The nature of the shorter period modulations is unclear.  The period is possibly the orbital period of J1021.  If it is the case, the superhump excess $\\varepsilon = (P_{\\rm SH}-P_{\\rm orb})/P_{\\rm orb}$ is calculated to be $0.0052\\pm 0.0003$.  The empirical relationship between the $\\varepsilon$ and the mass ratio of binary systems reported in \\citet{pat01SH} yields the mass ratio of $M_2/M_1=0.024$ for J1021. The $\\varepsilon$ and the mass ratio are the smallest one among previously known dwarf novae (\\cite{pat01SH}).  J1021 is possibly the most evolved dwarf novae if the period of the modulations in the fading tail indicate the orbital period.  An alternative possibility is that they are late superhumps.  Late superhumps are periodic modulations appearing just after superoutbursts (\\cite{woe88vwhyi}; \\cite{rol01iyumatomography}).  Their period is in agreement with a superhump period, while their profile is $\\sim 0.5$ phase shifted from that of superhumps.  In the case of J1021, we cannot confirm a clear phase shift in figure~\\ref{fig:o-c}.  We, however, consider that the late superhump scenario is more plausible because the long fading tail indicates an effective tidal dissipation, in which we can expect a long-lived eccentric disk and the superhump periodicity.  As can be seen in figure~\\ref{fig:o-c}, $O-C$s of humps are apparently changed smoothly from the main outburst to the rebrightening and the early fading tail.  This also supports the late superhump scenario."
    },
    "0711/0711.0005_arXiv.txt": {
        "abstract": "To probe the distribution and physical characteristics of interstellar gas at temperatures $T\\:\\approx \\:3\\times 10^{5}$~K in the disk of the Milky Way, we have used the {\\it Far Ultraviolet Spectroscopic Explorer (FUSE)\\/} to observe absorption lines of O~VI~$\\lambda$1032 toward 148 early-type stars situated at distances $>\\:1$ kpc.  After subtracting off a mild excess of O~VI arising from the Local Bubble, combining our new results with earlier surveys of O~VI, and eliminating stars that show conspicuous localized X-ray emission, we find an average O~VI mid-plane density $n_0\\:=\\:1.3\\times 10^{-8}{\\rm\\: cm}^{-3}$.  The density decreases away from the plane of the Galaxy in a way that is consistent with an exponential scale height of 3.2~kpc at negative latitudes or 4.6~kpc at positive latitudes.  Average volume densities of O~VI along different sight lines exhibit a dispersion of about 0.26~dex, irrespective of the distances to the target stars. This indicates that O~VI does not arise in randomly situated clouds of a fixed size and density, but instead is distributed in regions that have a very broad range of column densities, with the more strongly absorbing clouds having a lower space density.  Line widths and centroid velocities are much larger than those expected from differential Galactic rotation, but they are nevertheless correlated with distance and $N$(O~VI), which reinforces our picture of a diverse population of hot plasma regions that are ubiquitous over the entire Galactic disk.  The velocity extremes of the O~VI profiles show a loose correlation with those of very strong lines of less ionized species, supporting a picture of a turbulent, multiphase medium churned by shock-heated gas from multiple supernova explosions. ",
        "introduction": "}  Very hot plasmas are recognized as important constituents in a broad range of astrophysical contexts, ranging from the solar and stellar transition layers and coronae up to the largest arrangements of matter in the universe that form web-like structures that bridge groups and clusters of galaxies.  In all cases, the processes responsible for heating the gas give important insights on evolutionary and dynamical phenomena in the systems under study.  On the largest scales, there is a recognition that hot gas in the intracluster media and even the intergalactic media at low redshifts probably dominate the baryonic content of the universe \\citep[e.g.][]{CenOstr, tripp00a,danforth05,bregman07}. In this paper, we focus on the presence of hot gas within the interstellar medium (ISM) of the Milky Way, a system intermediate in size between the extremes just mentioned.  This is a topic that has been studied over the last 3 decades, and shortly (\\S\\ref{sect_intro_early}) we will look back upon some of the key developments in this field to help define the motivations for our present study.  \\subsection{Basic Principles of Detecting Hot Gas \\label{sect_intro_basic} }  Atoms within plasmas at temperatures $T>10^5\\,$K have many of their outer electrons stripped by collisions.  As a consequence, there is an upward shift in the energy of their resonance transitions, pushing important spectroscopic information into the ultraviolet and X-ray bands.  It is for this reason that progress in understanding the distribution and properties of hot gases in space have had to await the development of space-borne instruments that could observe radiation at energies above those in the visible band.\\footnote{Forbidden transitions in the visible from highly ionized atoms are generally too weak to observe \\citep{graney90}, except as seen in emission for exceptionally bright sources such as supernova remnants \\citep{woodgate74,woodgate75,teske90,sauv95,sauv99,szen00}.}  Of all the transitions from highly ionized atoms at wavelengths longward of the Lyman limit, the $\\lambda\\lambda 1032,~1037$ doublet of O~VI is the most suitable for study.  The advantages of observing O~VI are threefold: (1) the ionization potential of O$^{+4}$ is high enough (114$\\,$eV) to insure that there is no appreciable contamination by cooler material photoionized by starlight,\\footnote{Except for small contributions from very hot, He-poor white dwarfs \\citep{dupree83}.} as is the case with other partly stripped, Li-like atoms such as Si~IV and C~IV \\citep{cowie81}, (2) the transitions are strong [$\\log f\\lambda = 2.138$ for the 1032~\\AA\\ transition \\citep{morton03}], and (3) the cosmic abundance of oxygen is high.  One mildly negative factor is that the peak abundance of the five-times ionized form of oxygen (22\\% at $T=3\\times 10^5\\,$K under collisional ionization equilibrium conditions) is lower than representative fractions that can be reached by other ionization stages just below or above that of O~VI \\citep{shapiro76b, shull82,sutherland93}. Also, the temperature range over which O~VI is the dominant species is smaller than that of many other ions, such as O~VII~$\\lambda$~21.6 seen in the X-ray band.  The easiest way to observe the O~VI transitions is to record them in absorption in the spectrum of a UV-bright source.  In the disk of the Galaxy, the most suitable stars are those hotter than spectral type B3 that have large projected rotation velocities $v \\sin i$. Difficulties with poorly behaved stellar continua sometimes arise to create ambiguities in the absorption strengths; this issue is one that we address later in \\S\\ref{sect_cont} for our current study.  For probing the Galactic halo, AGNs and quasars make excellent background sources.  Energetic collisions in a hot gas not only ionize the atoms, but they also excite them to electronic states at energies well above the ground state.  As a consequence, the atoms emit radiation in their strong resonance lines.  O~VI can be seen in emission \\citep[see, e.g., Table~1 of][]{dixon06,otte06} although the flux levels are very faint ($\\sim1-4\\times 10^{-18}$ ergs~cm$^{-2}$~s$^{-1}$~arcsec$^{-2}$) and require long integration times to achieve a satisfactory detection.  A principal difference between the absorption and emission measurements is that the former is proportional to the line integral of the electron density, while the latter senses the square of the electron density [but with some sensitivity to temperature beyond just the variation from the ion fraction curve \\citep{shull94}].  This difference can be exploited to measure the filling factor of the gas, $\\langle n_e\\rangle ^2/\\langle n_e^2\\rangle$.  One drawback in measuring the O~VI emission is that self absorption and resonant scattering can complicate the interpretation if the optical depths are of order or greater than unity \\citep{shelton01}.  Also, the emission can be attenuated by the absorption caused by dust in the foreground.  Line emission from atoms other than O~VI also radiate from hot gases. As a rule, the characteristic energy bands where the emissions are strongest and most plentiful increase toward higher energies as the temperature rises.  For gas at temperatures of 282,000$\\,$K where O~VI has its maximum ion fraction in collisional equilibrium, the characteristic energy of the strongest emission-line features is below 300$\\,$eV \\citep{kato76}.  Unfortunately, there are a number of drawbacks that limit our ability to make good use of the X-ray background measurements as we interpret the O~VI absorption data. First, at energies below 300$\\,$eV the absorption by foreground neutral hydrogen and helium is strong \\citep{morrison83}.  Second, it is often difficult to know how much emission arises from beyond the target used for the O~VI measurement.  Third, over recent years there has been an increasing awareness that the soft X-ray background is contaminated by emission from charge exchange between heavy solar wind ions and the local interstellar gas \\citep{cravens00, cravens01}, a process identical to that which produces X-ray emission from comets \\citep{kharchenko00, kras02, kharchenko03}.  This contamination may be as much as half of the observed diffuse background level \\citep{lallement04}.  Finally, with the exception of the lowest energy X-ray measurements, such as those that operate below the 111$\\,$eV absorption cutoff of a beryllium filter \\citep{bloch86}, the observations have much stronger responses to gases that are too hot ($T\\gtrsim 2\\times 10^6\\,$K) to yield much O~VI.  All but the second reason listed above may help to explain why \\citet{savage03} found a poor correlation between soft X-ray emission and the strength of O~VI absorption toward extragalactic targets at high Galactic latitudes.   \\subsection{Early Developments \\label{sect_intro_early}}  Prior to the advent of space astronomy, an ingenious insight on the probable existence of a large volume of space that contains hot gases was provided by \\citet{spitzer56}, who noted that cold clouds found in the Galactic halo had to have some means of confinement.  He speculated that such clouds were stabilized by the external thermal pressure of a hot Galactic corona.  It was not until some 18 years later that firm evidence on the pervasiveness of this type of material within the Galactic disk was demonstrated by direct observations of absorption by the O~VI doublet at 1032 and 1038~\\AA\\ in the spectra of 32 stars \\citep{ebj74, york74}, accompanied by a recognition by \\citet{will74} that the diffuse soft X-ray emission observed by instruments on sounding rockets \\citep{bowyer68, bunner71, davidson72, yentis72} probably came from the same type of material.  Following the initial discoveries, \\citet{ebj78a} (hereafter J78) conducted a survey of O~VI absorption which increased the accumulation of cases to 72 Milky Way stars.  From this body of data, he offered some conclusions on the nature of the hot gas and how it is distributed in space \\citep{ebj78b, ebj78c}.  In particular, the statistical fluctuations of absorption strengths and velocities from one sight line to the next indicated that approximately 6 hot gas regions, each with a column density $N({\\rm O~VI})\\approx 10^{13}~{\\rm cm}^{-2}$, are distributed across a representative distance of 1$\\,$kpc.  Later, inspired by evidence from the soft X-ray surveys that the Sun is surrounded by a volume of hot gas out to a distance of around 100$\\,$pc \\citep{sanders77, mccammon90, snowden90, snowden98}, a structure now called the Local Bubble (hereafter, `LB'), \\citet{shelton94} conceived of a slightly more elaborate model that accounted for a possible constant foreground contribution of O~VI, with $N({\\rm O~VI})$ averaging about $1.6\\times 10^{13}{\\rm cm}^{-2}$, which reduced the apparent fluctuations of a much more patchy distribution at greater distances.  Their re-interpretation of the O~VI data suggested that outside the Local Bubble the Galactic disk is populated by regions that have column densities in the range $2-7\\times 10^{13}{\\rm cm}^{-2}$ with separations 450 $-$ 1300$\\,$pc.  With the advent of the {\\it Far Ultraviolet Spectroscopic Explorer} (FUSE), \\citet{bill_local} looked at weak O~VI lines towards nearby white dwarfs (WDs), and determined the average volume density to be $2.4\\times10^{-8}$~cm$^{-3}$.  They concluded that the local O~VI contribution is much less than that estimated by \\citet{shelton94} unless it is confined to a thin shell at the edge of the LB, i.e., just beyond the white dwarfs. A much larger survey of nearby WDs was conducted by \\citet{savage06}, who measured the average O~VI density to be $3.6\\times10^{-8}$~cm$^{-3}$.   A challenge that was presented to the early O~VI investigators was a proposal by \\citet{castor75} that all of the O~VI arose from bubbles that were created by stellar mass-loss winds from the target stars. \\citet{ebj78c} addressed this issue and concluded that some O~VI could arise from the bubbles, but that most of the hot gas was truly interstellar in nature.  We return to this topic once again in this paper (\\S\\ref{sect_bubbles}).  After the end of the \\Copernicus\\ era, there was a long period where only the transitions of N~V, Si~IV and C~IV could be seen in absorption using the {\\it International Ultraviolet Explorer\\/} ({\\it IUE\\/}) and the {\\it Hubble Space Telescope\\/} ({\\it HST\\/}). Because these species can be created at lower energies (ionization potentials of the next lower stages are only 33, 48 and 77$\\,$eV, respectively), there was a persistent concern that photoionized gases could be responsible for these species, especially in low density environments.  Even so, important conclusions emerged about the distributions of highly ionized material in the Galactic halo \\citep{savage87, ken91, sembach92}, the Magellanic Clouds \\citep{deboer80, fitzpatrick85}, and within the Galactic disk \\citep{cowie81, walborn84, edgar92, spitzer92, ken93, ken94, tripp93, savage94, ken94_zetoph, huang95}, often using distant UV sources fainter than the sensitivity limit of {\\it Copernicus}.  Further studies of O~VI, however, were confined to results that were more limited in scope \\citep{ferguson95, hurwitz95, dixon96, hurwitz96, widmann98, ken99} because they came from short-lived missions, such as the {\\it Hopkins Ultraviolet Telescope\\/} ({\\it HUT\\/}) \\citep{davidsen93} and the {\\it Orbiting Retrievable Far and Extreme Ultraviolet Spectrograph\\/} ({\\it ORFEUS-SPAS\\/}) \\citep{hurwitz98}. With the availibility of \\FUSE , we can now return to large-scale surveys of O~VI absorption.  Already, interstellar O~VI has been studied in the Galactic halo \\citep{bill00, richter01, howk02a, fox04, savage03, sembach_hvcO6, wakker03}, Large Magellanic Cloud \\citep{howk02b}, Small Magellanic Cloud \\citep{hoopes02}, and the Local Bubble \\citep{bill_local,savage06}.  Likewise, intergalactic O~VI has been found for absorption systems at redshifts that are too low for coverage with {\\it HST\\/} \\citep{bill00, sembach01_3c273, tripp01, savage02, phl1811_1, prochaska04, richter04, sembach04, danforth05}.  With the current paper, we cover yet another important realm, the Galactic disk.  Through the years, theoretical investigations kept pace with the observations, starting with initial assessments by \\citet{shapiro76a} on the nature of a possible Galactic fountain, and an early proposal by \\citet{cox74} on the origin of a pervasive network (``tunnels'') of hot gas that is kept alive by the channeling of shock waves (from supernovae) within it.  This picture was investigated in more detail by \\citet{smith77}.  \\citet{mckee77b} presented a seminal treatise on the different phases of the medium, accounting for mass and energy balance between a hot interstellar medium created by supernovae and the more conventional phases at lower temperatures.  This theory was a very important advance over an earlier model for ISM phases developed by \\citet{field69} that overlooked the hot phase and its secondary effects.  Some noteworthy consequences that might arise from the hot gas production include mechanical energy transport to the embedded cool clouds \\citep{cox79} and a feedback mechanism that may regulate star formation in disk galaxies \\citep{cox81}.  Enhanced accumulations of O~VI-bearing gas will undoubtedly arise from hot gas at the edges of, and within, the interiors of supernova remnants, which is discussed in the models presented by \\citet{chevalier74}, \\citet{cowie81_snr}, \\citet{slavin92}, \\citet{edgar93}, \\citet{smith01}, and \\citet{shelton06}.  More recently, exciting realizations of how the ISM may be configured from supernova explosions have been constucted by \\citet{avillez00}, \\citet{avillez04} and \\citet{avillez05_mag} using hydrodynamic and magnetohydrodynamic simulations. Their models predict global equilibrium in the disk-halo circulation when the explosions in the ISM are traced over several hundred Myr, and demonstrate how effective turbulent diffusion can be in mixing hot and cold gas. \\citet{avillez05} predict specifically the expected O~VI column densities along random lines of sight over distances of up to 1~kpc, and find good agreement with the \\FUSE\\ WD and \\Copernicus\\ data. Their models predict a clumpy turbulent ISM in which O~VI vortices and filaments are interecepted every $\\sim 100$~pc, each with $N$(O~VI)$\\:\\simeq\\:1-2\\times 10^{12}$~\\pcm .      \\subsection{Current Topics \\label{sect_current} }  A persistently debated unknown quantity is the volume filling factor of the hot gas in the Galactic disk \\citep{ferriere95}.  Equally important is the topology of the gas: is it confined to isolated, compact bubbles or does it comprise a matrix of multiply connected regions within which the rest of the ISM is confined?  For observers, it is easy to measure the average amount of O~VI in the Galactic disk, and with reasonable assumptions we may relate this quantity to the total amount of gas in the temperature regime $10^5-10^6\\,$K. However, a much more difficult task is to understand the structure of the gas along the viewing direction, i.e., information that is lost by the projection of everything along the sight line, which is compounded by the restriction that not all directions in the sky can be sensed. \\citet{ebj78b} placed an approximate upper limit of 20\\% for the filling factor from the apparent lack of an anticorrelation between the amount of O~VI and denser (mostly neutral) hydrogen, as revealed by the color excesses of the target stars (however, this argument can equally well be used to say the filling factor is greater than 80\\%). Models that propose that the hot gas has a very large filling factor [e.g. that of \\citet{mckee77b}] must be restrained by observations by \\citet{heiles80} that indicate that the neutral, warm medium (gas that emits 21-cm radiation but has a high spin temperature and thus negligible absorption) has its own significant filling factor.  One feature of the McKee \\& Ostriker model is that it assumes that supernovae occur at random locations and time.  Actually, supernovae are highly correlated \\citep{mccray79, heiles90}, and this should influence theories on the pervasiveness of O~VI-bearing material, leading to fewer, but larger volumes whose interiors hold the hot gas. The existence of such large, coherent structures seems to be validated by observations of enormous shells of gas seen in 21-cm line emission \\citep{heiles79, heiles84, mcclure02}.   The kinematics of O~VI and its relationship with other forms of gas, those that are highly ionized and those that are not, can provide important clues on the nature of the hot gas on a more microscopic level.  One of the first conclusions to emerge from the velocity profiles of the O~VI features is that they were broader than those produced by less ionized species, such as Si~II and Si~III \\citep{york74}, indicating that the O~VI was not simply a more ionized counterpart of the lower excitation gas.  The absorptions exhibited by York were of very good quality, which enabled him to place upper limits to the Doppler motions that corresponded to a range $T=4\\times 10^5$ to $2\\times 10^6$ for different cases.  The most striking conclusion from the survey by \\citet{ebj78b} was that the scatter of velocity centroids, amounting to only 26~\\kms , is far below the shock speeds needed to create the post-shock temperatures that can strip 5 electrons from the oxygen atoms, i.e., at least $130-160$~\\kms . \\citep{ray79, shull79} A similarly small dispersion of $21\\,{\\rm km~s}^{-1}$ was reported recently by \\citet{savage03} for sight lines through the Galactic halo.  The assertion that O~VI does not have extraordinarily large velocities must be viewed with caution however.  We can not rule out the presence of some additional O~VI that is moving very rapidly, for if the absorption profiles of such high velocity material were very broad and shallow, they would be very difficult to detect against the undulating stellar continua.  Simply put, O~VI is not created in zones that immediately follow the shocks that heated the gas to high temperatures.  Instead, the low velocities indicate that what we see arises either from some interaction with the normal, low-ionization material or, alternatively, perhaps ions whose velocities are regulated by magnetic fields in the disk \\citep{cox88} that might have been compressed and amplified by the motions of expanding supernova remnants \\citep{spitzer90}. For the former possibility, we may consider two alternatives.  One is that the O~VI originates in transition layers where heat flows by conduction across interfaces established between regions that hold very hot ($T>10^6\\,$K) gas and denser, cooler material \\citep{cowie77, mckee77a, begelman90_2, bertoldi90, mckee90}. Material is transported between the regions: at early times cool gas evaporates into the hot region, whereas later, after radiative cooling becomes important, hot gas condenses onto the cool gas \\citep{bork90}. If there is a substantial difference in velocity between the cool and hot material, an additional means of transporting heat and matter operates: instabilities in the transition zone create turbulence which acts as a source of mechanical mixing of the cool and hot phases \\citep{begelman90, slavin93}.   Observational support for links between O~VI-bearing gas and the cooler material was provided by \\citet{cowie_ebj79}, who noted correlations in velocity between O~VI and the ions Si~III and N~II. Additional support was provided by the good velocity correspondence between O~VI and C~II absorption along short path lengths through the local ISM \\citep{savage06}.  One way to gain insights on the relative importance of the different processes discussed above, as well as to see if the gas is simply radiatively cooling \\citep{edgar86} instead of being cooled by contact with low temperature material, is to test whether or not various other highly ionized atoms (N~V, Si~IV and C~IV) agree with theoretical predictions for the different regimes \\citep{spitzer96, fox03}.    In addition to the possibilities listed above, we must be watchful for another physical condition, that, under certain circumstances, may apply to the regions that are responsible for O~VI absorption.  In a discussion of hot gas within the Local Bubble, \\citet{breit94} considered the possibility that an explosion in a dense region followed by expansion in a lower density medium would lead to adiabatic cooling of the hot gas.  The lag in recombination as the gas cools would be even more severe than for a radiatively cooling gas. They envision the possibility that we are surrounded by gas at $T\\sim 50,000\\,$K that still has an appreciable concentration of ions that normally appear in equilibrium at much higher temperatures.  It is important to note that we cannot depend on identifying this gas on the basis that the lines have a velocity dispersion characteristic of the Doppler broadening at the lower temperatures, since the expansion process itself creates a broadening that replaces the thermal broadening.  In the discussion presented up to now, we have considered a number of contemporary problems in the study of gases responsible for O~VI absorption in the disk of the Galaxy.  More comprehensive discussions are provided by review papers written by \\citet{cox87_araa}, \\citet{savage87_tetons}, \\citet{ebj87_iue}, and \\citet{spitzer90}.  \\subsection{Outline of This Paper---A Summary}  The capability of \\FUSE\\ to record spectra of stars much fainter than the magnitude limit of \\Copernicus\\ ($V\\leq 7.5$), allows us to probe regions well beyond the original O~VI survey of J78.  In this paper, we present a new \\FUSE\\ sample of disk stars which has about twice as many stars as the older survey.  These two considerations allow us to overcome some shortcomings of the earlier work and address such issues as the possible presence of very large scale structures containing either more or less than the usual amount of O~VI, and whether or not the apparent decrease of O~VI at moderate distances from the Galactic plane agrees with an exponential scale height of $\\sim 2.3\\,$kpc as determined from sources well outside the Galaxy \\citep{savage03}. Also, with the larger column densities that arise from the greater lengths of the sight lines, we should be less sensitive to distortions of the results arising from {\\it i)} the presence of bubbles around the target stars and {\\it ii)} an over-representation of the region surrounding the Sun that arises from the constraint that all sight lines must emanate from a single point.  This paper focuses on three principal topics. First, we describe in detail the data collection and reduction of the \\FUSE\\ spectra, and the analysis of the O~VI absorption line profiles. Second, we present a re-evaluation of the distances to the targeted stars, which is an important parameter for understanding how O~VI is distributed in the Galactic disk. Third, we discuss the global properties of O~VI absorption in the disk, which, of course, is the main science goal of this survey.   As far as the first of these topics is concerned, the data reduction of \\FUSE\\ spectra is not unlike that of any other data taken with modern photon-counting devices. We were concerned, however, that we should be able to precisely calibrate the wavelength scale of the spectra, since \\FUSE\\ does not take calibration spectra contemporaneously with an object. This led us to compare the data from many of our sight lines with high quality, high resolution {\\it HST} data, from which we found that the standard \\FUSE\\ wavelength scale for the pipeline calibration (CalFUSE 2.0.5) required a correction of 10~\\kms. This analysis is introduced in \\S\\ref{sect_wavecal} and described in detail in Appendix~\\ref{sect_cl1}.  In analyzing the O~VI lines, we faced two main obstacles. First, J78 realized that any \\sixa\\ line is contaminated to some degree by the HD $6-0$ R(0) line at 1031.9~\\AA . Before measuring the physical parameters of an O~VI line, the contribution from this HD line must be removed.  Fortunately, in the \\FUSE\\ data, there exist other HD lines that are free of the dense H$_2$ `forest' and which can be used to model the contaminating HD line.  We found that removing the model using Voigt profile line-fitting techniques was highly successful and actually presented few problems for the subsequent analysis of the \\sixa\\ line. Details are given in \\S\\ref{sect_contaminants} and \\ref{sect_hdremoval}; \\S\\ref{sect_epoch} also discusses a consistency check on the efficacy of the HD removal using multi-epoch data.  A more problematic obstacle in analyzing the O~VI lines arose in determining the shape of a star's continuum. With the data reduced, we quickly discovered that the gradient of the continuum was probably varying on scales less than the width of an O~VI line. The most extreme variations were prevelant among the O9-B1 type stars.  It became clear that we needed to track three possible continuum fits: the adopted best fit, and two possible variations above and below this fit --- `upper' and `lower' continuum fits, as we will refer to them throughout this paper.  Although measuring continuum errors in absorption-line studies is far from unusual, in many cases, the interstellar lines under study are quite narrow, and continuum errors are largely irrelvant. For our \\FUSE\\ data, however, we found that continuum errors actually dominated the O~VI column density errors, and not the usual Poisson counting statistical errors.  We discuss how we chose our continua in \\S\\ref{sect_cont}.  With the spectra normalized, we were able to derive column densities, Doppler parameters and line velocities using both Voigt profile fitting (\\S\\ref{sect_voigt}) and the Apparent Optical Depth (AOD) method (\\S\\ref{sect_colsAOD}), both of which gave very similar results, as expected. For the subsequent analysis, we assumed that the absorbing gas had no small-scale velocity structures hidden within the broad profile that might invalidate either analysis method (\\S\\ref{sect_saturation}). We favored the results from the Voigt profile fitting procedure, but used AOD measurements when no line was clearly present and an upper limit to the O~VI column density was required (\\S\\ref{sect_colsAOD} and Appendix~\\ref{sect_upper_limits}). The spectra themselves are presented in \\S\\ref{sect_explain_spectra}.  The second principal goal of this paper is to re-evaluate the distances to the stars used in our survey.  There are several databases now available (the \\Hipparcos , \\Tycho, and 2MASS catalogs) which provide a more reliable estimate of a distance than was possible a decade ago.  We introduce the problem first in \\S\\ref{sect_the_disk}, where we show the distribution of stars in the Galactic plane. We expound a full discussion in Appendix~\\ref{sect_distances}, since the methods used to derive distances are technical and outside the main scope of our discussion on O~VI in the Galactic disk.  A reader disinterested in the mechanisms of deriving stellar distances can certainly skip this Appendix, and instead simply note that the most egregious error in determing a distance arises from the difficulty in determing the absolute magnitude of a star, $M_V$. This is not because of a problem in knowing $M_V$ for the particular spectral type or luminosity class of a star, but because of the problems inherent in defining the type and class correctly in the first place. The error in $M_V$ is not only a function of the type and class, but depends on the quality of the stellar spectra used to define these parameters, and on the skill and experience of the observer making the measurements.  Despite the significant improvements in measuring various stellar parameters with modern techniques, we believe that errors in stellar distances range from $10-30$\\%, and that even larger deviations are possible. A full explanantion is given in Appendix~\\ref{sect_absmags}.  The final topic of this paper is, of course, the global properties of O~VI absorption in the disk of the Milky Way. Our motivation for the study has already been introduced in the previous sections, but we note the following. With the huge amounts of data available from satellites that cover many different regions of the electromagnetic spectrum, it is now possible to compare the results from absorption-line surveys with results from other types of studies. In this paper, we concentrate on X-ray emission data taken with the {\\it R\\\"{o}ntgen Satellite\\/} (\\ROSAT ). Although X-ray emitting gas is nominally at a higher temperature than we would expect for O~VI absorbing gas, one can imagine a natural confluence of hot gas in structures that might be detected in both O~VI absorption and X-ray emission. In \\S\\ref{sect_bubbles} and \\ref{sect_SNR}, we describe the types of X-ray emission in \\ROSAT\\ maps seen towards our \\FUSE\\ stars. In many cases, these maps indicate regions of hot gas along the line of sight from bubbles of gas around stellar associations or supernova remnants. These observations lead us to classify the morphology of X-ray emitting patches (or absence thereof) around or near our target stars.   Using these classifications, we show in \\S\\ref{sect_rho_rosat} that the average line of sight density $N$(O~VI)/$d$ is somewhat higher for sight lines which pass through X-ray enhanced regions, leading us to conclude that a small fraction of $N$(O~VI) towards some stars does indeed arise in circumstellar material or wind-blown bubbles.  This information is important, because by selecting sight lines with no clear X-ray emission, and by correcting O~VI column densities for absorption by the Local Bubble using recently published data, we can derive a less biased value of the mid-plane density of O~VI, $n_0$. We discuss this value at some length in \\S\\ref{sect_density}, concluding with our best determination of $n_0$ and its likely errors in \\S\\ref{sect_finally}. In calculating these values, we quantify the significance in the difference in $N$(O~VI) towards stars in the north and south Galactic hemispheres originally found by \\citet{savage03} (\\S\\ref{sect_NS}). The difference has no effect on the value of $n_0$ we derive, but does produce two different values of the scale height of the O~VI gas layer in the plane of the Milky Way.  In \\S\\ref{sect_Nwithde} we show that $N$(O~VI) is directly correlated with the distance to a star in the plane of the disk. This correlation can most easily be understood as the effect of longer sight lines passing through more absorbing clouds.  This shows that O~VI absorption cannot arise primarily from the circumstellar environment of the observed stars, but must be a truly interstellar phenomenon. The data also reveal a discontinuity in $N$(O~VI) with distance beyond a few hundred pc, in that column densities appear to be too high for a given distance. As we show, however, this arises from excess contributions to the measured $N$(O~VI) from the Local Bubble.  However, the data also suggest that the hot O~VI-absorbing ISM is neither a smoothly distributed homogeneous gas layer, nor a population of randomly distributed clouds with similar properties. First, there is no acceptable fit to the decrease in O~VI density with height above the plane of the Galaxy unless an additional `error' is introduced, a term which represents an intrinsic `clumpiness' in the interstellar medium (\\S\\ref{sect_clump}). Second, this clumpiness does not decrease with distance, as would be expected if a sight line simply passed through more a greater number of generic clouds (\\S\\ref{sect_Nwithde}). To explain this, we propose a picture of the O~VI absorbing ISM that contains a mixture of cloud sizes, with small, randomly distributed small clouds and sparsely distributed large clouds. We also note an alternative explanation, that we may be seeing the clustering of many small clouds.  There are other lines of evidence that suggest that the O~VI absorbing `clouds' are actually one component in a complex, evolving, turbulent ISM. The velocities of the bulk of the O~VI absorption do not follow the general pattern of differential co-rotation in the disk of the Galaxy (\\S\\ref{sect_vel_with_long}). Moreover, if we compare the extremes of the O~VI absorption line profiles with the extremes of other lower ionization species (O~I, C~II, C~III and Si~III) we find that they match each other quite closely (\\S\\ref{sect_highVlow}). This apparent `coupling' of at least some of the components in the absorption line profiles (at least those with the highest peculiar velocities along a sight line) suggests that the structures giving rise to absorption from hot gas are multi-ionized.  A natural explanation for all these phenomena has been discussed above. The numerical simulations of \\citet[ and references\\/ therein]{avillez05} vividly realize how the ISM may look when driven by supernovae explosions. Their two dimensional density maps (for example) in the plane of the Galaxy show how hot gas arises in bubbles around supernovae, which is then sheared through turbulent diffusion, destroying the bubble and stretching the hot absorbing gas into filaments and vortices that dissipate with time.  In \\S\\ref{sect_Nwithb} we confirm the relationship between $N$(O~VI) and Doppler parameter found by other authors, and determine that the O~VI lines are too wide to arise from differential Galactic rotation alone. Clearly, the velocity fields that govern the O~VI components along a sight line are much larger than can be reproduced by the simple motions of the Galactic disk, and of course, expanding supernova shells naturally explain these large velocities.  When considering an ensemble of such regions along long sight lines, it seems plausible that the correlation of $b$ with $N$(O~VI) could arise from the existence of exceptionally large internal motions inside large clouds, or that smaller clouds within clusters are moving more rapidly than the gases within individual clouds.   \\citet{avillez05} found good apparent agreement between their models and the \\FUSE\\ WD data and the results from \\Copernicus\\ over a path length of $\\sim 1$~kpc.  Our results are also likely to be consistent with these models, although a detailed comparison is beyond the scope of this paper.  The data presented herein should provide sufficient information to test these, and any other detailed simulations, over larger scales. ",
        "conclusions": "\\begin{enumerate}   \\item We have used \\FUSE\\ to observe early-type (mainly O2$-$B3) stars in the Galactic disk in order to characterize \\sixa\\ absorption lines arising from interstellar gas in the plane of the Milky Way. A total of 120 stars was observed as part of \\FUSE\\ PI team programs P102 and P122, of which nine were eventually rejected from our survey; the remaining 111 stars constitute a total investment of 754 ksec of \\FUSE\\ time.  We supplemented these sight lines with archival data available April 2003 or earlier; our final sample consists of 148 stars, most of which are at Galactic latitudes $|b|<10\\degr$ and distances of $>\\: 1$~kpc.  \\item In order to derive accurate spectroscopic distances to the stars, we have used several modern databases to improve upon published distances (Appendix~\\ref{sect_distances}). \\Hipparcos\\ and \\Tycho\\ catalogs have provided precise optical magnitudes, as well as identification of multiple systems and variable stars. 2MASS infrared magnitudes have enabled us to estimate the interstellar $V-$band extinction along a sight line, and the absolute magnitudes of stars --- given a spectral type and luminosity class --- have been constructed by collating available surveys from the literature. Despite these improvements, imprecise knowledge of stellar types and classes continues to have the greatest uncertainty in deriving a spectroscopic distance, limiting the precision in a distance measurement to $\\sim\\:10-30$\\%.  \\item Where appropriate, we have added the O~VI measurements from several other O~VI studies to our sample, in order to probe distances not covered by our \\FUSE\\ stars. These include nearby ($\\sim\\: 20 -100$ pc) white dwarf stars (WDs) \\citep{savage06}, intermediate distance ($\\sim 100-1000$~pc) stars observed by the \\Copernicus\\ satellite \\citep{ebj78a}, distant halo stars \\citep{zsargo03}, and extragalactic sight lines \\citep{savage03,wakker03}. For many of the stars (but not including the WDs), we have rederived distances using the same set of tools used to derive distances towards the \\FUSE\\ sample of stars.  \\item \\ROSAT\\ X-ray maps show that many stars lie towards regions of enhanced X-ray emission. This emission likely arises from hot bubbles blown out from the star or its stellar association. Other sight lines, however, are found at positions with very little X-ray emission, which we label as $R\\:=\\:0$ sight lines.  In contrast, lines of sight intercepting X-ray emitting regions have designations of $R\\:>\\:0$ (\\S\\ref{sect_bubbles} and \\S\\ref{sect_SNR}).  The average volume density towards stars in the Galactic disk, $n\\:=\\:N$(O~VI)$/d$, depends on whether we use $R\\:=\\:0$ or $R\\:>\\:0$ sight lines, and whether we subtract a contribution to the O~VI column density from the Local Bubble (LB; \\S\\ref{sect_rho0}). When we subtract a value of $N_{\\rm{LB}}$(O~VI)$=1.11\\times10^{13}$~\\pcm\\ arising from the LB, we find little difference in $n$ between stars with distances $0.2\\:<\\:d\\:<\\:2.0$~ kpc and those with $d\\:>\\:2$~kpc (\\S\\ref{sect_rho_lb}); we also find that for stars $d\\:>\\:2$~kpc, the $R\\:>\\:0$ sight lines have slightly higher densities than $R\\:=\\:0$ sight lines, even after excluding stars in the direction of the Vela SNR and Carina Nebula, which have the highest densities in our sample (\\S\\ref{sect_rho_rosat}).  This is probably because the circumstellar environment of the star contributes a small excess to the O~VI column density.   \\item Our data well fit the relationship $N$(O~VI)$\\sin |b|\\:=\\:n_0 h (1-\\exp^{-|z|/h})$, which naturally arises from assuming that the density $n$ of O~VI absorbing gas falls with height $|z|$ above the Galactic plane as $n\\:=\\:n_0 \\exp^{-|z|/h}$. Here, $n_0$ is the mid-plane density and $h$ is the characteristic scale height. We find $n_0\\:=\\:1.3\\times 10^{-8}$~\\pcmV\\ [after subtracting a contribution to $N$(O~VI) from the LB and using the $R\\:=\\:0$ sight lines; \\S\\ref{sect_finally}], but that $h$ depends on whether we use extragalactic lines of sight in the northern (N) or southern (S) Galactic hemispheres \\citep[as first noted by][ ---see \\S\\ref{sect_NS}]{savage03}. We derive values of $h\\:=\\:[4.6,3.2]$~kpc when using the [N,S] sample.  We also find it neccessary to include a clumpiness factor $\\sigma_{\\rm{CL}}$ to represent the fact that the dispersion in $\\log [N$(O~VI)$\\,\\sin |b|$] around the theoretical fit is much larger than can be accounted for by experimental errors alone (\\S\\ref{sect_clump}). Our best estimate is that $\\sigma_{\\rm{CL}}\\:\\simeq\\:0.26\\pm0.02$~dex.   \\item $N$(O~VI) and the distance to a star are correlated (\\S\\ref{sect_Nwithde}). This shows that the processes which give rise to O~VI absorption are ubiquitous over the entire Galactic disk, and certainly do not arise primarily from circumstellar environments. The only inhomogeneity of the OV~I volume density $n$ can be found in a concentration of $N$(O~VI) upper limits (10 out of 23 sight lines) towards $d\\:>\\:1$~kpc \\FUSE\\ sight lines in the region $l\\:=\\:80^\\circ-140^\\circ$, the outer Perseus arm of the Galaxy. There is no evidence that the volume density of O~VI is influenced by the presence of Galactic spiral arms along the stellar sight lines.  \\item Although O~VI is distributed smoothly enough for $N$(O~VI) to correlate with distance, the relative {\\it dispersion} on either side of this correlation does not decrease with distance, as would be expected if more distant sight lines intercepted more clouds. Quantitively, we find no difference in the dispersion ($\\sigma_{\\rm{CL}}$) between nearby ($0.2-2$~kpc) and distant ($>\\:2$~kpc) stars. This indicates that in addition to a random placement of small clouds, there may also be sparsely distributed large clouds.  Or, alternatively, some of the small clouds may simply be clustered over very large scales.  \\label{point_Nwithde}   \\item The velocity edges of the O~VI lines and those of the strong transitions of C~III~$\\lambda 977.02$, O~I~$\\lambda 1039.23$, Si~III~$\\lambda 1206$ and C~II~$\\lambda 1335$, are all correlated. This suggests that, at least to first order, the processes that move the hot gas to high velocity also affect the velocities of other species (\\S\\ref{sect_highVlow}).    \\item The smallest reliable O~VI Doppler parameter we find in our sample is $b\\:\\simeq\\:16-17$~\\kms, which is consistent with that expected for a thermally broadened cloud at a temperature close to the peak in the ionization fraction curves, $\\log T_{\\rm{max}}\\:\\simeq\\: 5.45$ (\\S\\ref{sect_bees}).  We also confirm earlier research which found that $b$ and $N$(O~VI) are correlated. We show, however, that the width of O~VI lines cannot be accounted for by differential Galactic rotation; the lines are far wider than expected if absorption came from a smoothly distributed ISM corotating with the disk of the Milky Way (\\S\\ref{sect_Nwithb}). The correlation suggests that the large clouds mentioned in Point~\\ref{point_Nwithde} above have large internal motions, or, alternatively, that the small clouds within clusters are moving more rapidly than the gases within individual clouds.    \\end{enumerate}  \\bigskip"
    },
    "0711/0711.3695_arXiv.txt": {
        "abstract": "Using techniques from singular perturbation theory, we explicitly calculate the cosmological evolution in a class of modified gravity models. By considering the (m)CDTT model, which aims to explain the current acceleration of the universe with a modification of gravity, we show that Einstein evolution can be recovered for most of cosmic history in at least one $f(R)$ model. We show that a standard epoch of matter domination can be obtained in the mCDTT model, providing a sufficiently long epoch to satisfy observations. We note that the additional inverse term will not significantly alter standard evolution until today and that the solution lies well within present constraints from Big Bang Nucleosynthesis. For the CDTT model, we analyse the ``recent radiation epoch'' behaviour ($a\\propto t^{1/2}$) found by previous authors. We finally generalise our findings to the class of inverse power-law models. Even in this class of models, we expect a standard cosmological evolution, with a sufficient matter domination era, although the sign of the additional term is crucial. ",
        "introduction": "The observation that the universe appears to be accelerating at present times has caused one of the greatest problems to modern cosmology. High precision data from Type Ia supernovae, the Cosmic Microwave Background and large scale structure seem to hint that the universe is presently dominated by an unknown form of energy, dubbed {\\it dark energy} \\cite{Riess:1998cb,Perlmutter:1998np,Tonry:2003zg,Bennett:2003bz,Netterfield:2001yq,Halverson:2001yy}. The problem resides in the lack of good physical motivation for this energy density. One obvious contender for the role of dark energy is Einstein's cosmological constant, but particle physics fails to predict the correct density. Other candidates include scalar fields or extra dimensions.  We refer the reader to \\cite{Copeland:2006wr} for a review.  Recently, a modification of General Relativity itself was suggested to explain this accelerating universe \\cite{Deffayet:2001pu,Capozziello:2002rd,Freese:2002sq,Arkani-Hamed:2002fu,Dvali:2003rk,Nojiri:2003ni,Arkani-Hamed:2003uy,Carroll:2003wy,Capozziello:2003tk, Nojiri:2003ft,Carroll:2004de,Abdalla:2004sw,Capozziello:2006dj,Appleby:2007vb}. For reviews see e.g. \\cite{Easson:2004fq,Nojiri:2006ri,Trodden:2006qk}. The postulation is that in which the Ricci scalar of the Einstein-Hilbert lagrangian is replaced by an arbitrary function of the Ricci scalar, $f(R)$. The most simple example includes inverse powers of $R$: $f(R)=R-a/R^n$~\\cite{Carroll:2003wy}. While this modified action was shown to give an accelerated attractor, due to the equivalence to scalar-tensor theories with a vanishing Brans-Dicke parameter\\cite{Chiba:2003ir,Chiba:2006jp,Erickcek:2006vf}, solar system experiments rule these models out ($\\omega_{BD}>40,000$~\\cite{Bertotti:2003rm}).  Attempts have been made to overcome this problem~\\cite{Nojiri:2003ft}, but solar system gravity still raises doubt as to these models~\\cite{Brookfield:2006mq}.  These models have raised much recent interest, due to their perhaps simple nature.  The presence of ghosts and stabilities have been studied \\cite{Dick:2003dw,DeFelice:2006pg,Faraoni:2005vk}, as well as the implications to the CMB, large scale structure, solar system gravity and gravitational waves ~\\cite{Amarzguioui:2005zq,Cembranos:2005fi,Song:2006ej,Navarro:2006mw,Bean:2006up, Faulkner:2006ub,Li:2007xn,Fairbairn:2007qp, Hu:2007nk,DeFelice:2007zq,Starobinsky:2007hu, DeFelice:2007ez,Song:2007da, Capozziello:2007ms,Pogosian:2007sw, Carloni:2007yv,Ananda:2007xh}.  Recently, it has been claimed that, in all theories which behave as a power of $R$ at large or small $R$, standard cosmological evolution can not be obtained~\\cite{Amendola:2006kh,Amendola:2006eh}. Further, it is claimed that no $f(R)$ will lead to an acceptable matter era\\cite{Amendola:2006we}. Specifically, during the matter phase, the authors find that in these theories the scale factor evolves as $a\\propto t^{1/2}$ and not as $t^{2/3}$ as required. In order for a modified theory of gravity to be convincing, we require a close-to-Einstein evolution in the past, with a standard, sufficiently-long period of matter domination.  We note that several models have been proposed which lead to matter dominance before acceleration~\\cite{Nojiri:2006gh,Nojiri:2006be,Hu:2007nk} and that this issue is a matter of debate in the literature.  If the modification to gravity exists as a perturbative term in addition to the Einstein-Hilbert terms, then it is the belief of the authors of this paper that Einstein cosmology should be recovered in the limit that the perturbation term disappears. This system has great similarity with perturbation theory~\\cite{Hinch}. This paper studies the asymptotic expansion of the (m)CDTT model \\cite{Carroll:2003wy}, to show that standard cosmological evolution can be produced in at least one model. Specifically, we aim to recover an acceptable matter era and show that there is at least a class of $f(R)$ for which this is possible. For the purposes of this paper, we do not concern ourselves with the constraints from solar system experiment.  The paper is laid out as follows.  In Section \\ref{sec:theory}, we define the theoretical framework and give some useful formulae.  We state the simplest example in Section \\ref{sec:CDTT} and derive the asymptotic solution in Section \\ref{sec:asymp}. The relevance of this solution to cosmology is highlighted in Section \\ref{sec:relevance}, where we comment on the era of matter domination, as well as Big Bang Nucleosynthesis.  After generalising our specific example to the class of inverse power-law models in Section \\ref{sec:gen}, we conclude in Section \\ref{sec:conc}.  We discuss the eigenmode analysis in the Appendix. ",
        "conclusions": "\\label{sec:conc}  A modification of gravity has been suggested in the form of $f(R)$ theories.  However, until now, it has been believed that these cannot produce a standard era of matter domination. By studying the asymptotic behaviour of a class of $f(R)$ models, we show that a sufficiently-long period of matter domination can be achieved.  We consider one specific example, the (m)CDTT model, $R+\\sigma\\mu^4/R$ with $\\sigma=\\pm1$, and explicitly calculate the asymptotic behaviour. In terms of the CDDETT variables, where $x=-H$, we find that the system gives a singular perturbation problem, comprising a two region structure: an outer region $-x=O(1)$ and an inner region $x=O(\\mu)$. We show that the mCDTT solution recovers Einstein cosmology when $-x=O(1)$ (i.e. in the past when $H\\gg\\mu$) and matches an accelerated cosmology when $-x=O(\\mu)$ (i.e. today when $H\\sim\\mu\\sim H_0$). We directly relate this solution to cosmological variables and find the condition for which we obtain sufficient matter domination to satisfy large scale structure data.  We also show that the perturbation causes no significant effect at the time of Big Bang Nucleosynthesis.  We generalise our results to include the class of inverse power-law functions.  Although these are also ruled out by solar system tests, we aim to exhibit the generality of our previous results, which could be valid for other functions. We find that the original CDTT model cannot match to the outer $v\\sim2/3$ (matter dominated) solution and the ``recent radiation'' behaviour of \\cite{Amendola:2006kh,Fairbairn:2007qp} is found. Generally, we find that for models with $R\\pm \\epsilon AR^{-n}$, standard cosmological evolution is expected until $-x=O(\\epsilon^{1/4n})$, which can be coincident with the present time.  After this point, an accelerated solution is found.  This solution depends strongly on the sign of $A$.  Our main result is that, for a model with lagrangian $R+f(R)$, we expect the leading order outer (past-time) solution to be given by Eqn.~{\\ref{eq:rhor0}}, the Einstein solution. The next order terms will be proportional to the perturbative parameter in $f(R)$. The leading order inner (future) solution can be found using the system in Eqn.~(\\ref{eq:gendV}).  We therefore expect the solutions to match across the boundary and the universe should find the accelerating attractor, given by Eqn.~(\\ref{eq:Vatt}).  Whether the ``recent radiation epoch'' is found (instead of a matter epoch) seems to depend on the sign of the corrective term.  It is the authors' belief that this perturbation method can be applied to even more general functions of higher curvature invariants, for example $f(R,P,Q)$, where $P=R_{\\mu\\nu}R^{\\mu\\nu}$ and $Q=R_{\\alpha\\beta\\gamma\\delta}R^{\\alpha\\beta\\gamma\\delta}$, such as the CDDETT model.  This is the focus of future work.  With solutions such as we have shown here, observational constraints can be applied to a much larger class of modified gravity models than is possible at present.   \\appendix"
    },
    "0711/0711.4760_arXiv.txt": {
        "abstract": "We used MIDI, the mid-infrared interferometric instrument of the VLTI, to observe the massive protostellar candidate IRS 9A, located at a distance of about 7 kpc at the periphery of the NGC 3603 star cluster. Our ongoing analysis shows that MIDI almost fully resolves the object on all observed baselines, yet below 9\\,$\\mu$m we detect a steep rise of the visibility. This feature is modelled as a combination of a compact hot component and a resolved warm envelope which lowers the correlated flux at longer wavelengths. The extended envelope can already be seen in both MIDI's acquisition images and in complementary data from aperture masking observations at the Gemini South telescope. Its shape is asymmetric, which could indicate a circumstellar disk inclined against the line of sight. The compact component is possibly related to the inner edge of this (accretion) disk. The uncorrelated mid-infrared spectrum appears featureless and could be caused by optically thick emission without a significant contribution from the disk atmosphere. ",
        "introduction": "The giant H\\,II region NGC 3603 is powered by one of the densest clusters of high mass stars known in our galaxy. Due to its vicinity to this cluster, IRS 9A has been liberated from most of the gas and dust of its natal molecular cloud by the strong stellar winds and ionising radiation arising from the nearby O-type main-sequence stars. This offers the unique possibility to observe a high mass star at infrared wavelengths during its relatively early evolutionary phase. ",
        "conclusions": "Our ongoing analysis of the MIDI data shows IRS 9A to be almost completely resolved on baselines of the order of 30\\,m. The mid-infrared spectrum is featureless, and most likely caused by emission from a warm envelope. The models are so far unable to account for the steep rise of the visibility towards the short wavelength end, so that the nature of this compact component remains unclear and further refinements of the models are needed."
    },
    "0711/0711.3140_arXiv.txt": {
        "abstract": "The brightest supernova remnant in the Magellanic Clouds, N132D, belongs to the rare class of oxygen-rich remnants, about a dozen objects that show optical emission from pure heavy-element ejecta. They originate in explosions of massive stars that produce large amounts of O, although only a tiny fraction of that O is found to emit at optical wavelengths. We report the detection of substantial amounts of O at X-ray wavelengths in a recent 100 ks {\\it Chandra} ACIS observation of N132D. A comparison between subarcsecond-resolution {\\it Chandra} and {\\it Hubble} images reveals a good match between clumpy X-ray and optically emitting ejecta on large (but not small) scales. Ejecta spectra are dominated by strong lines of He- and H-like O; they exhibit substantial spatial variations partially caused by patchy absorption within the LMC. Because optical ejecta are concentrated in a 5 pc radius elliptical expanding shell, the detected ejecta X-ray emission also originates in this shell. ",
        "introduction": "\\label{intro}  The heavy-element ejecta of core-collapse (CC) supernovae (SNe) are dominated by oxygen. By studying the properties of O and other freshly synthesized elements in supernova remnants (SNRs), we can infer the total amount of ejected O, relative abundances, and the spatial distribution of heavy elements. With this information, we can then constrain the progenitor main-sequence mass and learn more about how massive stars explode. Oxygen is most readily seen in X-rays, as evidenced by a growing number of O-rich SNRs detected by modern X-ray satellites such as {\\it Chandra} and {\\it XMM-Newton}. X-ray detection of plasma enriched in O but deficient in Fe allows us to identify CC SNRs in the first place. High-resolution studies of SNRs containing both optical and X-ray--emitting O ejecta hold particular promise for advancing our understanding of ejecta in CC SNRs.   N132D, an X-ray--bright SNR in the Large Magellanic Cloud (LMC), contains optically emitting O-rich ejecta \\citep{sutherland95,morse95,morse96}.  \\citet{morse95} interpreted their observations in terms of an expanding (1650 km s$^{-1}$) 10 pc ellipsoidal ejecta shell.  Assuming undecelerated expansion, they estimated the remnant's age at 3150 yr, 10 times older than Cas A. Spectroscopy with {\\it ASCA} \\citep{hughes98} and {\\it XMM-Newton} \\citep{behar01} showed that the X-ray emission is dominated by the shocked ambient medium.  X-ray emission from ejecta has not yet been convincingly demonstrated, although it might be expected based on the high (30-35 M$_{\\odot}$) mass of the SN progenitor inferred from optical and UV spectroscopy \\citep{blair00}.  We report here the discovery of clumpy X-ray--emitting O ejecta in N132D with new {\\it Chandra} X-ray observations, matching (on large scales) the optical ejecta morphology as seen by {\\it Hubble}. ",
        "conclusions": "\\label{summ}  High spatial resolution imaging and spectroscopy of N132D with {\\it Chandra} reveal the presence of clumpy O-rich ejecta in its center. The spatial distribution of X-ray--emitting ejecta correlates well with optically emitting ejecta on large (but not small) scales. This shows that optical and X-ray emission can be closely linked in dynamically advanced remnants such as N132D, where most of the O might have already been shocked. In N132D, ejecta are located in an expanding ellipsoidal shell. No such shell is present in either G292.0+1.8 or Pup A, two O-rich Galactic SNRs with ages similar to N132D. Optically emitting ejecta in G292.0+1.8 are distributed throughout much of the remnant, with little detailed correlation with X-rays \\citep{winkler06}. Asymmetrically distributed optical ejecta filaments in Pup A are moving in the opposite direction to its neutron star \\citep{winkler07}. This demonstrates that the innermost metal-rich ejecta are strongly affected by poorly understood processes leading to CC explosions. No correlation of optical emission with more evenly distributed X-ray--emitting ejecta has been reported. X-ray expansion velocities in N132D are poorly known; \\citet{hwang93} found (marginal) evidence for an X-ray expansion consistent with the optical expansion in spatially integrated spectra obtained by {\\it Einstein}, but archival {\\it Chandra} and {\\it XMM-Newton} grating spectra are more suitable for studying the ejecta kinematics. The origin of the X-ray and optically emitting ejecta shell is unknown at this time. It may be the location of a reverse shock propagating into ejecta; more likely, it may be an ejecta shell created by a Ni bubble effect shortly after the SN explosion, similar to what is seen in the O-rich SNR B0049-73.6 \\citep{hendrick05}. More detailed analysis of existing {\\it Chandra} and {\\it XMM-Newton} observations is warranted."
    },
    "0711/0711.4626_arXiv.txt": {
        "abstract": "We revisit theoretical and observational constraints on geometrically-thin disk accretion in Sagittarius\\,A$^\\star$ (\\SgrA). We show that the combined effects of mass outflows and electron energization in the hot part of the accretion flow can deflate the inflowing gas from a geometrically-thick structure. This allows the gas to cool and even thermalize on an inflow timescale. As a result, a compact, relatively cool disk may form at small radii. We show that magnetic coupling between the relativistic disk and a steady-state jet results in a disk that is less luminous than a standard relativistic disk accreting at the same rate. This relaxes the observational constraints on thin-disk accretion in \\SgrA (and by implication, other Low-Luminosity Active Galactic Nulcei, LLAGN). We find typical cold gas accretion rates of $\\mbox{a few} \\times 10^{-9} M_\\odot \\, {\\rm yr}^{-1}$. We also find that the predicted modified disk emission is compatible with existing near-infrared (NIR) observations of \\SgrA in its quiescent state provided that the disk inclination angle is $\\gtapprox 87^\\circ$ and that the jet extracts more than $75$\\% of the accretion power. ",
        "introduction": "The exceptionally compact radio source  \\SgrA  is spatially coincident with the dynamical centre of the Galaxy, which contains a mass $M \\approx 3.7 \\times 10^6 M_\\odot$, deduced from measurements of stellar orbital motions \\citep[e.g.][]{b2,Schodel02,b1}. High-resolution VLBA imaging of \\SgrA at millimetre (mm)  wavelengths provides the most compelling observational evidence yet for the existence of supermassive black holes (SMBHs), with $M$ now constrained to lie within a radius $\\approx 1$\\,AU of \\SgrA \\citep{b23}.  The observed bolometric luminosity of \\SgrA, $L \\approx 10^{36}\\,{\\rm erg \\, s}^{-1} \\approx 2 \\times 10^{-9} \\LEdd$ (where $\\LEdd$ is the Eddington luminosity), is unusually low, even compared to that of other LLAGN \\citep{b3}. This can be attributed to a very low mass accretion rate, $\\mdot$, or a low radiative efficiency, $\\epsilon = L/\\mdot c^2$, or to a combination of both. The most popular accretion models for \\SgrA are Radiatively Inefficient Accretion Flows (RIAFs), which have attributed the low $L$ to a low $\\epsilon$ and which are based on the hot ion torus \\citep*{SLE76,Ichimaru77,Rees82} and Advection Dominated Accretion Flow (ADAF; \\citealt{NarYi94}) models.  Applications of RIAFs to \\SgrA have evolved considerably in recent years in response to increasing observational constraints (see \\citealt{Yuan06} for a review). In the original models, the inferred mass accretion rate was comparable to the Bondi rate for spherical accretion, $\\simeq 10^{-5} \\, M_\\odot \\, {\\rm yr}^{-1}$, and the radiative efficiency was very low, $\\epsilon \\simeq \\mbox{a few} \\times 10^{-6}$, as a result of preferential ion viscous heating and negligible electron-ion coupling. The most recent RIAF model \\citep*{YQN03,YQN04} now takes into consideration a mass outflow component as well as energization of the electrons. As a result, the revised RIAF model now attributes the low luminosity of \\SgrA to a combination of low \\mdot \\, and moderately low $\\epsilon$. The outflow reduces the rate at which hot ionized gas is fed onto the central SMBH to levels that are consistent with mm and sub-mm polarization measurements, which require $\\mdot \\ltapprox 4 \\times 10^{-8} \\, M_\\odot \\, {\\rm yr}^{-1}$ \\citep*{b16,b17,b18,Bower03,Marrone06,b26}. Electron heating and acceleration are then required to produce the observed levels of radio and sub-mm emission \\citep{QuatNar99}, resulting in a radiative efficiency $\\epsilon \\simeq 10^{-2}$.    The chanelling of binding energy away from the ions as a result of outflows and electron heating (as well as electron-ion coupling) necessarily implies that the ion energy is subvirial. There may then be insufficient ion pressure to support a geometrically-thick structure and consequently, some of the gas can radiatively cool on an inflow time \\citep{Rees82}. This raises the possibility that a compact, cool accretion disk may form at small radii. To test this possibility, the effects of a mass outflow on the continuity, angular momentum and energy equations need to be self-consistently considered in the exisiting RIAF model for \\SgrA  \\citep{YQN03}. Numerical approaches \\citep[e.g.][]{HawBalb02,SQHS07} have so far been limited by the nonconservative numerical scheme and the neglect of radiative cooling.  Observationally, increasing evidence is emerging to suggest that geometrically-thin disk accretion may persist at or near the last marginally stable orbit in low-luminosity sources \\citep{Miller06a,Miller06b,Maoz07,Rykoff07,Liu07}. Young massive stars seen orbiting the Galactic Center (GC) and believed to have formed \\textit{in situ} \\citep{NaySun05,Nay06,Pau06,Nayakshin06,Lev06,Bel06} suggest that at least some of the hot gas detected by \\textit{Chandra} has condensed into a cool phase. Additionally, new models for stellar wind accretion in \\SgrA indicate that slow winds can radiatively cool within a dynamical timescale to produce a cold phase of accreting gas \\citep{Cuadra06}. However, a strong constraint on the presence of a cold, optically-thick disk is the absence of eclipses in the orbit of S2, the best studied star in the central S-cluster. This limits the size of a putative accretion disk to $\\ltapprox 10^{16} \\, {\\rm cm} \\simeq 3 \\times 10^5 \\rg$ \\citep*{Cuadra03}, where $\\rg = GM/c^2 \\approx 5 \\times 10^{11} (M/3.7\\times 10^6 M_\\odot) \\, {\\rm cm}$ is the gravitational radius.  If a compact, cool disk forms near \\SgrA, much of its emission would suffer strong interstellar extinction. Even so, the mass accretion rate would have to be extremely low and the spin axes of the disk and galactic plane would have to be closely aligned for the emission from a standard disk to remain below observational limits in the NIR \\citep{Cuadra03,Trippe07}. However, it has been suggested \\citep*[e.g.][]{NagWilFal01,Falcke01,FKM04,Gallo07,MH07} that the spectral properties of LLAGN and galactic X-ray binaries in their low/hard X-ray state can be attributed to a jet-dominated mode of accretion \\citep[see also][]{MerlFab02}. Although jets have not been observed in \\SgrA, various arguments have been made for their presence \\citep*[see e.g.][for a comprehensive discussion]{MarkBowFalck07}. If jets are present, an accretion disk can no longer be described by the standard model \\citep{b9,b10} because the radial disk structure is modified by the magnetic torque responsible for jet formation \\citep{KB04,KB07b}. Indeed, extraction of accretion power by a magnetized jet results in an accretion disk that is less luminous than a standard disk accreting at the same rate \\citep{KB07a}. That is, the radiative efficiency is lower than that predicted by standard accretion disk theory.  In this paper, we revisit theoretical constraints on the formation of a cool accretion disk  and we investigate observational constraints on jet-driven disk accretion in \\SgrA. Unlike the jet-RIAF model for \\SgrA \\citep*{FalckMark00,YMF02} and the disk-corona model for LLAGN \\citep{MerlFab02}, we explicitly model the magnetic coupling between the accretion flow and the outflow. This model has been successfully applied to M87 \\citep{JK07}. The organization of this paper is as follows. In Section \\ref{sec2}, we show that a RIAF becomes geometrically-thin as a result of diversion of binding energy from the inflowing ions. This allows the electrons to radiatively cool and to thermalize with the ions on an inflow time, thus forming a cool disk at small radii. We also present the relevant equations for relativistic disk accretion modified by MHD stresses and we calculate the modified disk flux radial profile and corresponding modified disk emission spectrum using parameters appropriate for \\SgrA. In Section \\ref{sec3}, we calculate the steady-state synchrotron spectrum resulting from a jet magnetically coupled to the underlying accretion flow. We compare the predicted spectra with the observed quiescent spectrum for \\SgrA. A discussion and concluding remarks are given in Sections \\ref{sec4}  and \\ref{Conclusions}, respectively. ",
        "conclusions": "We have revisited theoretical constraints on thin-disk accretion in \\SgrA, showing that a geometrically-thick, hot, two-temperature accretion flow cannot be sustained in the presence of outflows and electron heating. We have revised observational constraints, taking into consideration modified disk emission as a result of magnetic coupling to a jet. The magnetic torque which drives the jet efficiently extracts accretion power from the disk at small radii. This results in an accretion disk that is dimmer than a standard relativistic disk accreting at the same rate, so the constraints on disk luminosity and inclination are less stringent. For \\SgrA, we infer a mass accretion rate of $\\mdot \\simeq \\mbox{a few} \\times 10^{-9} M_{\\odot}\\, {\\rm yr}^{-1}$. We find that a disk luminosity of $\\Ld \\approx 5 \\times 10^{36} \\, {\\rm erg \\, s}^{-1} $ and an inclination angle $\\gtapprox 87^\\circ$ are compatible with existing observational constraints. However, more NIR observations of \\SgrA in quiesence are needed to provide tighter constraints on the presence of a cool accretion disk. Numerical simulations of radiative, low-$\\mdot$ accretion flows are also warranted."
    },
    "0711/0711.4410_arXiv.txt": {
        "abstract": "As dust settles in a protoplanetary disk, a vertical shear develops because the dust-rich gas in the midplane orbits at a rate closer to true Keplerian than the slower-moving dust-depleted gas above and below.  A classical analysis (neglecting the Coriolis force and differential rotation) predicts that Kelvin-Helmholtz instability occurs when the Richardson number of the stratified shear flow is below roughly one-quarter.  However, earlier numerical studies showed that the Coriolis force makes layers more unstable, whereas horizontal shear may stabilize the layers.  Simulations with a 3D spectral code were used to investigate these opposing influences on the instability in order to resolve whether such layers can ever reach the dense enough conditions for the onset of gravitational instability.  I confirm that the Coriolis force, in the absence of radial shear, does indeed make dust layers more unstable, however the instability sets in at high spatial wavenumber for thicker layers.  When radial shear is introduced, the onset of instability depends on the amplitude of perturbations: small amplitude perturbations are sheared to high wavenumber where further growth is damped; whereas larger amplitude perturbations grow to magnitudes that disrupt the dust layer.  However, this critical amplitude decreases sharply for thinner, more unstable layers.  In 3D simulations of unstable layers, turbulence mixes the dust and gas, creating thicker, more stable layers.  I find that layers with minimum Richardson numbers in the approximate range 0.2 -- 0.4 are stable in simulations with horizontal shear. ",
        "introduction": "It is a remarkable fact that planets start out as microscopic grains within the protoplanetary disks of gas and dust in orbit around newly-formed protostars, somehow growing roughly $10^{40}$ orders of magnitude in mass in a period no more than $10^7$ years corresponding to disk lifetimes \\citep{lissauer93}.  There is no one physical process that can explain growth over this enormous range of sizes: the very smallest grains (micron to millimeter sizes) can grow via collisional agglomeration in which the sticking mechanism is electrostatic in nature; whereas, on the other end of the size spectrum, objects in the kilometer to tens of kilometers regime can grow via gravity-enhanced collisions \\citep{beckwith00}.  The least understood stage of growth is how millimeter-size particles grow to kilometer-size; grains in this regime are too large for sticking via electrostatic forces, yet far too small to have any significant self-gravity.  Even more problematic is the fact that particles in this intermediate-size regime are strongly affected by aerodynamic drag of the surrounding gas: meter-size objects, for example, have radial drift speeds on the order of $10^4$~cm/s at 1~AU and thus spiral inward onto the protostar on the timescale of a few hundred years \\citep{weidenschilling77}.  Whatever process is responsible for grain growth through this range of sizes must act on timescales faster than this inspiral time if any raw materials are to be available to build protoplanets.  \\citet{goldreich73} and \\citet{safronov69} proposed that a very thin, very dense sheet of settled dust in the midplane of the protoplanetary disk might be be gravitationally unstable; the nonlinear evolution would result in the layer clumping-up directly into gravitationally-bound kilometer-size planetesimals on a timescale of order the orbital period.  According to this scenario, there is a direct jump from small particles to kilometer-size planetesimals without growing slowly through the intermediate sizes which have very short orbital decay timescales.   As attractive as this mechanism is for planetesimal formation, a significant obstacle is turbulence which can stir and mix the dust with the gas and prevent the dust layer from settling into a thin enough, dense enough sheet for the gravitational instability to operate \\citep{weidenschilling80}.  Even in the absence of any mechanism to drive turbulence, the settling of the dust particles into a thin sublayer in an initially laminar midplane would create a vertical shear that might be unstable to Kelvin-Helmholtz instability.   Because of a relatively weak outward radial pressure gradient, pure gas in a protoplanetary disk orbits the protostar at a rate slightly slower than true Keplerian: $V_{gas} = V_K(1-\\eta)$, where $\\eta\\sim 10^{-3}$ for typical conditions at 1~AU \\citep{adachi76,weidenschilling77}.  Pure dust, in the absence of any gas, would orbit exactly at the Keplerian rate.   One can show that in the limit of perfect dust-gas coupling, the orbital velocity of a mixture of gas and dust is a function of the local dust-to-gas ratio \\citep{adachi76}; as dust settles into the midplane, the dust-rich gas in the midplane orbits faster than the dust-depleted gas above and below the midplane.  Two-fluid (gas and dust) numerical simulations by \\citet{cuzzi93}, \\citet{champney95} and \\citet{dobrovolskis99} suggested that turbulent diffusion would prevent settling of grains into thin enough sheets for gravitational instability.  The onset of Kelvin-Helmholtz instability is determined by a competition between the stabilizing effects of stratification and the destabilizing effects of vertical shear.  If there is no rotation and no horizontal shear, then a necessary (but not sufficient) condition for instability is the Richardson number criterion \\citep{chandra60,drazin81}: \\begin{equation}\\label{E:Richardson_criterion} \\textrm{instability: } \\quad Ri(z) \\equiv \\frac{N^2}{(dU/dz)^2} < Ri_{crit} \\sim 0.25 \\quad \\textrm{for some z,} \\end{equation} where $N$ is the Brunt-V\\\"{a}is\\\"{a}l\\\"{a} frequency (the frequency of buoyant oscillations of a stably-stratified medium) and $dU/dz$ is the vertical shear.  The critical Richardson number of one-quarter corresponds to a state in which the kinetic energy of the vertical shear is sufficient to lift the heavier gas out of the gravitationally potential well and remix it with the lighter overlying fluid.  Assuming that dust would settle into thinner and thinner layers down to the limit set by the classic Richardson criterion,  \\citet{sekiya98}, \\citet{sekiya00}, and \\citet{youdin02} determined critical vertical quasi-equilibrium dust profiles and investigated the conditions necessary for such layers to be gravitationally unstable.  \\citet{garaud04} pursued two-fluid linear calculations (without rotation and radial shear) of dust sedimentation into sheets and found that the Kelvin-Helmholtz instability was excited before gravitational instability unless the global dust-to-gas ratio was greatly enhanced over solar abundance. \\citet{youdin04} suggested that such enhancements might be attainable because the inward drift speed decreases as particles migrate inward, resulting in ``particle pile-ups.''  It is not at all clear whether the classic Richardson number criterion is appropriate for the Kelvin-Helmholtz instability in protoplanetary disks in which both rotation and radial shear could significantly affect the onset  of instability as well as alter the nonlinear evolution of any instability that develops. \\citet{ishitsu03} pursued a purely linear analysis to study the effect of radial shear on the time evolution of unstable Kelvin-Helmholtz modes and showed that the modes were sheared to higher spatial wavenumber and could eventually be stabilized. On the other hand, \\citet{gomez05} investigated the effects of the Coriolis force on the onset of instability, and found that settled layers were unstable at much higher Richardson numbers (corresponding to thicker dust layers) than in the classic case.  However, these simulations were two-dimensional and did not include the radial shear.  The motivation for this work is to investigate the apparently opposing influences of rotation (destabilizing) and radial shear (stabilizing) on the Kelvin-Helmholtz instability of settled dust layers and resolve whether or not such layers can ever reach the thin enough, dense enough conditions for the onset of gravitational instability.   Apart from the impact of Kelvin-Helmholtz instability on the formation of planetesimals, it is important to determine the vertical distribution of dust in order to correctly interpret observations \\citep{brittain05, rettig06, dullemond07}.  The approach we take here is that the dust is perfectly coupled to the gas; in terms of timescales, the friction or stopping time (see \\eqref{E:epstein} below) is much shorter than the evolution of Kelvin-Helmholtz instabilities if they are present.  We will not simulate the formation of the dust layer from a well-mixed state, but assume that the layers have already formed with a given profile of the dust-to-gas ratio.  If the layer turns out to be unstable on a fast timescale, it implies that such a layer would never have formed in the first place.  In \\S2, we present equations for the hydrodynamic evolution of settled dust layers in the limit of perfect dust-gas coupling.  In \\S3, we revisit the cases of  (i) no rotation, no shear, and (ii) rotation, no shear.  In \\S4, we present new three-dimensional simulations of the evolution of settled dust layers with both rotation and shear.  Finally, in \\S5, we discuss the impact of these results on the planetesimal formation via gravitational instability. ",
        "conclusions": "We revisited the case of no horizontal shear, both with and without the Coriolis force.  The case without the Coriolis force is the most similar to classic Kelvin-Helmholtz instability, and the critical Richardson number for the onset of instability is close to the expected value of one-quarter, consistent with there being sufficient kinetic energy in the shear to lift the heavier fluid out of the gravitational well and mix it with the overlaying lighter fluid \\citep{garaud04}.  As first noted by \\citet{gomez05}, the case with the Coriolis force is surprisingly different, with instability occurring at much higher Richardson numbers.  However, the wavelengths of the most unstable eigenmodes for these thicker layers is smaller than the thickness of the layer itself.  \\citet{ishitsu03} showed that the horizontal shear is able to eventually stabilize unstable eigenmodes by shearing them out to high wavenumber.  However, unstable eigenmodes are able to grow for a period of time; the question is whether they can grow to a large enough amplitude to trigger nonlinear effects and disrupt the dust layer before they are damped as they are sheared.  In order to investigate the competing influences of the Coriolis force (destabilizing) and horizontal shear (stabilizing), we used a 3D spectral, anelastic, shearing-box code \\citep{barranco06} to simulate settled dust layers in the limit of perfect dust-gas coupling.  We find that the stability of dust layers depends on the amplitude of initial perturbations: small perturbations grow for a period of time, but are damped before they reach sufficient magnitude to trigger nonlinear effects; whereas larger amplitude perturbations are able to grow to magnitudes that can disrupt the dust layer, resulting in turbulence and mixing.  Kelvin-Helmholtz instability in thicker layers has slower growth rates, implying that the magnitude of initial amplitudes would have to be larger.  This was seen in Runs 06 and 07 in which the critical threshold for perturbations was an order of magnitude larger for a layer that was 50\\% thicker.  Thick dust layers that are stable according to the classic Richardson number criterion, but are unstable with the addition of the Coriolis force exhibit Kelvin-Helmholtz instability with very slow growth rates and at high spatial wavenumber (see Figure~\\ref{F:growth_coriolis}).  Three-dimensional simulations of these layers indicate that the horizontal shear is able to damp the instability no matter how large the initial perturbations for these layers (see Runs 04 and 05).  Thinner layers, on the other hand, have very fast growth rates, and the threshold amplitude is so low as to be practically irrelevant in real protoplanetary disks where there is no doubt fluctuations of such low magnitude.  In two-dimensional simulations of unstable dust layers, the turbulence that develops fills the computational domain and leads to large-scale mixing.  Characteristic of ``inverse cascades'' in two-dimensional turbulence, small eddies merge with other eddies to form larger coherent vortices which chaotically advect the dust, homogenizing the dust-to-gas ratio.  In three-dimensional simulations, the turbulence is locally confined to a region right around the original unstable dust layer and does not propagate through the rest of the computational domain.  The unstable thin layers evolve to thicker, more stable layers with minimum Richardson numbers tantalizingly close to the value of one-quarter for the onset of instability in the absence of the Coriolis force and horizontal shear.  These results hold when the global dust to gas ratio $\\Sigma_d/\\Sigma_g$ and the global gas radial pressure gradient $\\eta_0V_{K0}/c_{s0}$ are varied.  There is no reason to expect that the unstable layers in these simulations would evolve to a final state that just happens to be at the critical thickness for stability.  One could imagine that the turbulence is so efficient at mixing that it results in layers whose widths are significantly thicker than thinnest stable layer.  Because the simulations presented here are in the limit of perfect dust-to-gas coupling, no further sedimentation of the dust is allowed so that the thicker final states cannot settle to the critical state.  Future work will involve two significant improvements.  First, we will relax the perfect dust-to-gas coupling assumption and allow there to be a finite value for the stopping time.  Dust will be treated as a second fluid with its own velocity field \\citep{cuzzi93}.  Layers that are unstable will develop turbulence and re-mix the dust with the gas, but then would be allowed to further settle.  \\citet{johansen06} performed 2D simulations with two fluids and showed that layers evolved toward a self-regulated state in which further settling was inhibited by turbulence generated by Kelvin-Helmholtz instability, maintaining the layer in a dynamic equilibrium right at the critical thickness for instability.  However, because their simulations were 2D, they found the dust layers had very high Richardson number similar to the layers investigated by \\citet{gomez05}.  Second, we will add the effects of non-ideal magnetohydrodynamics with finite resistance.   \\citet{turner07} investigated the magneto-rotational instability in protoplanetary disks with ``dead zones'' in the midplane where the ionization is too low to couple the fluid to the magnetic fields.  However, they find that turbulence originating in the cosmic-ray-ionized surface layers can mix free charges into the interior and weakly couple the midplane gas to the magnetic fields, effectively eliminating the dead zone.  Their analysis does not include the role that dust grains play in removing free charges and reducing the ionization.  If turbulence lofts the particles throughout the disk, ionization can be suppressed, which will have the tendency to decouple the interior gas from the magnetic fields, allowing the dead zone to reform.  Dust particle could then re-settle, ionization could increase, re-coupling the gas to the fields and generating a new phase of turbulence.  We plan to investigate if there is indeed a limit cycle to the formation and destruction of the dead zone in protoplanetary disks.  The author would like to thank the National Science Foundation for support via the Astronomy \\& Astrophysics Postdoctoral Fellowship program.  Computations for this project were done at the Institute for Theory \\& Computation at the Harvard-Smithsonian Center for Astrophysics, and at the San Diego Supercomputing Center.  The author would like to thank Philip Marcus, Ramesh Narayan, and Niyash Ashfordi for fruitful conversations related to this project."
    },
    "0711/0711.2065_arXiv.txt": {
        "abstract": "The chemical composition of the ultra-high-energy (UHE) cosmic rays serves as an important clue for their origin. Recent measurements of the elongation rates by the Pierre Auger Observatory  hint at the possible presence of heavy or intermediate mass nuclei in the UHE cosmic rays. Gamma-ray bursts (GRBs) and hypernovae have been suggested as possible sources of the UHE cosmic rays. Here we derive the constraints on the physical conditions under which UHE heavy nuclei,  if they are accelerated in these sources, can survive in their intense photon fields.  We find that in the GRB external shock and in the hypernova scenarios, UHE nuclei can easily survive photo-disintegration. In the GRB internal shock scenario, UHE nuclei can also survive, provided the dissipation radius and/or the bulk Lorentz factor of the relativistic outflow are relatively large, or if the low-energy self-absorption break in the photon spectrum of the prompt emission occurs above several KeV. In internal shocks and in the other scenarios, intermediate-mass UHE nuclei have a higher probability of survival against photo-disintegration than UHE heavy nuclei such as Fe. ",
        "introduction": "Ultra-high-energy (UHE) cosmic rays above $\\sim10^{18}$ eV are thought to be of extra-galactic origin, since charged particles with such high energies cannot be confined by the magnetic field of our Galaxy. Galactic cosmic ray accelerators, such as supernova remnants, are expected to reach their maximum energy below $\\sim10^{18}$ eV. The transition between galactic and extra-galactic cosmic rays is therefore believed to happen either at the 'second knee' around $10^{18}$ eV, where the chemical composition changes significantly as measured by HiRes, followed by a 'dip' in the spectrum (Berezinsky et al. 2006), or at the 'ankle' around $10^{19}$ eV where the cosmic ray spectrum becomes flatter.  At these ultra-high energies, the chemical composition is a subject of debate. It has been speculated that these cosmic rays are made of pure protons, up to the highest energies. On the other hand, there are also reasons for favoring a cosmic ray spectrum dominated by heavy or intermediate mass nuclei at the highest energies, since according to the Hillas criterion (Hillas 1984) astrophysical sources are able to accelerate particles up to a maximum energy proportional to their nuclear charge. Recently, a mixed composition scenario has been invoked to fit the UHECR spectrum above $\\sim10^{18.5}{\\rm eV}$ (Allard et al. 2005). By studying the details of the development of the resulting air showers, one can in principle infer the species of the primary UHECRs, since at a given energy, showers initiated by heavy nuclei develop higher up in the atmosphere than proton-induced showers. Both AGASA and HiRes data favor a dominance of light hadrons, consistent with being pure protons, in the composition of UHECRs above $10^{19}$ eV (Hughes et al. 2007), which is consistent with models assuming UHECRs above $10^{18}$ eV are due to extra-galactic protons. {On the other hand, a recently report of an analysis of the muon component of UHE air showers in the Yakutsk data showed a heavy nuclei component in the UHECR spectrum (Glushkov et al. 2007). }  One of the aims of Pierre Auger Observatory is to study the composition of UHECRs. It provides more precise measurements of the depth of the UHECR-initiated shower maximum $X_{\\rm max}$ at extremely high energies ($>10^{19}$ eV), with uncertainties which are about a factor 4 smaller than those of the best measurement from the HiRes experiment (Unger et al. 2007). The elongation rate data presented by the Pierre Auger Observatory team is satisfactorily represented by a fit containing a break point in the slope at $10^{18.35}$ eV (Unger et al. 2007). Due to the uncertainties in the hadronic interactions at the highest energies, the interpretation of these elongation rate depends on the hadronic physics models used in the analysis, and is therefore rather ambiguous at present.  However, regardless of which interaction models are used, the data appears to require the presence of a substantial fraction of heavy or intermediate-mass nuclei in the range of GZK cutoff energies.  A possible conflict in the chemical composition between the above Auger elongation results and recent results by the Pierre Auger collaboration (2007) on large scale spatial correlations is an issue which remains to be resolved.  The proposed astrophysical source models of UHECRs include AGN jets (e.g.  Biermann 1987; Berezinsky et al. 2006), intergalactic accretion shocks (e.g. Inoue, Aharonian \\& Sugiyama, 2005), and gamma-ray bursts (e.g.  Waxman 1995, 2004a; Vietri 1995; Wick et al. 2004; Dermer \\& Atoyan 2006; Murase et al. 2006). Recently, we have proposed that extra-galactic hypernovae associated with sub-energetic GRBs are also a probable source for UHECRs, whose energetics are sufficient large to account for UHECRs from the second knee and above (Wang et al. 2007). These UHECR accelerators also accelerate electrons, which produce optical, x-ray or gamma-ray photons through synchrotron and inverse Compton emission. Since a sufficiently high density of these low-energy photons can potentially lead to photo-disintegration of the cosmic ray heavy nuclei, a natural question is under what conditions do these sources allow the survival of UHE heavy nuclei from the acceleration site. In this paper we explore this problem for GRBs and hypernovae sources{\\footnote{The photo-disintegration of UHE nuclei in AGN scenario has been discussed in Anchordoqui et al. (2007b)  and Dermer (2007a).}}, since they have the most compact source sizes and might therefore pose the greatest hurdle for UHE heavy nuclei survival.   Long-duration GRBs are generally believed to result from the core collapse of massive stars, direct evidence for which comes from the spectroscopic identification of bright supernovae in association with these GRBs (see Woosley \\& Bloom 2006 for a review). The collapse leads to a black hole (or magnetar) plus accretion disk system. The annihilation of neutrinos or Poynting flows arising from the inner hot accretion disk produce a high entropy fireball outflow, which expands and converts its internal energy into the kinetic energy of a small amount of baryon material, while propagating through the star along the rotation axis of the collapsing core. The outflow gets collimated and finally two highly relativistic jets break out of the stellar envelope. In the GRB scenario for UHECRs, both internal shocks (Waxman 1995) and external shocks (Vietri 1995; Wick et al.  2004; Dermer \\& Atoyan 2006) in the jets are suggested to be able to accelerate baryons to ultra-high energies. Internal shocks develop due to internal collisions between the shells in the fireball jet ejecta, while the external shocks occur when the fireball ejecta interacts with the surrounding interstellar medium (see \\Meszaros 2006 and Zhang 2007 for recent reviews). In the internal shock scenario, viewed as a source of the cosmic ray nuclei, the jets themselves must contain heavy nuclei, whereas in external shocks the heavy nuclei could come from the swept-up circumburst material that the jet is running into, presumably the stellar wind or interstellar medium. Since the two shocks have quite different dissipation radii and emission properties, the disintegration problem of the UHE comic ray nuclei are different, as discussed in \\S\\S \\ref{sec:int} and \\ref{sec:ext} respectively.  Hypernovae are a peculiar type of supernovae with higher ejection velocities of the remnant shell and generally larger explosion energies than typical supernovae (Paczy{\\'n}ski 1998). The prototype of hypernovae is SN1998bw, a type Ic supernova associated with an under-energetic GRB, GRB980425 (Galama et al. 1998). Mildly relativistic ejecta components are inferred to be present in all three of the well-identified hypernova/under-energetic GRB systems so far, namely SN1998bw/GRB980425, SN2003lw/GRB031203 and SN2006aj/GRB060218. Based on this mildly relativistic ejecta component and the event rates of these objects, we have shown that hypernovae can accelerate particles to $\\sim Z\\times10^{19}$ eV, and the energetics and occurrence rate is sufficient to account for the flux of UHECRs above the second knee ($\\sim 6\\times10^{17}$) eV, where $Z$ is the nuclear charge of the accelerated particles (Wang et al. 2007). The particles are accelerated in the hypernova blast wave formed by the interaction between the hypernova ejecta with the surrounding stellar wind medium. In \\S \\ref{sec:hn} we study whether UHE heavy nuclei accelerated in such hypernovae can survive their photon environment.  {The goal of this paper is therefore to answer the question, if UHE heavy nuclei are accelerated in GRBs or hypernovae, what are the possible origins  of these nuclei, and under what conditions would they survive their source environment; or conversely, under what conditions would one expect these sources to accelerate mainly protons to UHE energies.} ",
        "conclusions": "\\label{sec:disc}  The possible presence of heavy nuclei in UHECRs raises interesting questions: what is the origin of these nuclei? can these nuclei survive in the sources where they get accelerated? In this paper, we have endeavored to address these questions for two proposed UHECR sources discussed in the literature, namely GRBs and hypernovae. For GRBs, both internal shocks and external shocks have been suggested to be able to accelerate particles to UHE energies, and in this paper we have considered the role of both of these in the context of UHECR heavy nuclei. We have sketched out some possible { mechanisms for the presence of} heavy nuclei in these scenarios. For GRB internal shocks, we suggest that the nuclei are entrained from the progenitor stellar core, e.g. the Fe core, the O core, etc., during the stage when the jets arising from accretion in the innermost collapsing core, are making their way out through the star.  {Since instabilities favoring entrainment are predominant mainly during the initial stellar crossing phase of the jet, a jet composition with a substantial fraction of heavy nuclei may only be expected in long bursts with observed $\\gamma$-ray durations $t_\\gamma\\lesssim 10-15(1+z)$ s, while longer burst jets would be expected to consist mainly of protons. }  For the GRB external shock and hypernova source scenarios, the nuclei may be the heavy elements present in the stellar wind of Wolf-Rayet stars, which are thought to be the progenitors of GRBs and hypernovae. The stellar wind of Wolf-Rayet stars, especially the WO, WC sub-types stars, is heavily enriched with intermediate mass nuclei, such as O, C, etc.  {After they escape from their sources (see e. g. Dermer 2007b), UHECR nuclei can be subject to photo-disintegration also in intergalactic space, before arriving at Earth.  An UHE heavy nucleus with an energy $\\ga10^{19}{\\rm eV}$ will mainly collide with cosmic infrared background (CIB) photons. Using new constraints on CIB data, it has been found that UHE iron nuclei with energy $\\la10^{20}{\\rm eV}$ have a mean free path of $\\ga 300{\\rm Mpc}$ and that this mean free path increases  as the energy of the nucleus decreases (Hooper et al. 2007; Stecker \\& Salamon 1999). Lighter nuclei have relatively shorter mean free path (Hooper et al. 2007). Thus, UHE nuclei can in principle originate from sources at cosmological distance, such as GRBs and hypernovae. The main question appears to be whether they survive the environment of their original sources. For this reason, most of the present paper has been devoted to a quantitative  study of the survival of UHE nuclei in the above mentioned sources.  We find that: }  i) In GRB internal shocks, heavy nuclei can survive in the sources if the internal shock radius and/or the Lorentz factor of the relativistic jets are relatively large, as given by Eq.(11).  Thus one might expect acceleration of Fe nuclei to be more favored in bursts with smoother, longer variability timescale light-curves. For a smaller internal shock radius, the photodisintegration process due to prompt X-ray photons may become optically thick. However, if a synchrotron self-absorption break is present above several keV in the photon spectrum, the reduced number density of X-ray photons will lower the photodisintegration optical depth accordingly. In general, compared to heavy nuclei, the UHE intermediate-mass nuclei find it easier to survive photo-disintegration in internal (and external) shocks.  ii) In GRB external shocks, due to  the much larger dissipation radii compared to internal shocks, UHE nuclei can easily survive in the sources. We have also calculated the evolution of the photo-disintegration optical depth during the afterglow phase for both an ISM external medium and a stellar wind medium. The results show that UHE nuclei can survive in the afterglow shock for a relatively long time  in both cases.  iii) In the hypernova remnant acceleration scenario, UHE nuclei can survive in the sources, except in the early short period of time ($\\la 100 {\\rm days}$ for typical parameters, see Eq.(31)). In this early short period, however, only a very small amount of ejecta energy has been converted to UHECRs, so their contribution to the UHECR flux is negligible. Most of the UHECRs originate from the blast wave ejecta with $\\Gamma\\beta\\la0.5$, which  decelerate after $\\sim 1000 {\\rm days}$, typically. UHE nuclei accelerated during this time are safe from photo-disintegration.  { Earlier  calculations on the photo-disintegration of UHE nuclei in GRB internal shocks by Anchordoqui et al. (2007b) consider only one specific set of typical parameters for the internal shock radius and the relativistic Lorentz factor. We have improved this by considering the whole parameter space and working out the constraints  on the physical conditions under which UHE heavy nuclei  can survive in internal shocks. We have also considered the effect of a self-absorption break in the prompt emission on the photo-disintegration problem. Besides internal shocks, we have explored the photo-disintegration problem of UHE nuclei in GRB external shocks and hypernova remnants. }  In summary, we have suggested possible scenarios for the injection of heavy nuclei into the acceleration zones of extragalactic sources such as GRB and hypernovae. We found that, if  heavy nuclei are accelerated in these sources, they will survive the threat of photo-disintegration under fairly general conditions for the case of GRB external shocks, and for hypernovae. They could survive also in GRB internal shocks, if the latter occur at relatively large radii and/or the bulk Lorentz factors are large. On the other hand, for small shock radii and/or smaller bulk Lorentz factors, a pure proton UHECR composition would be favored if the self-absorption break in the photon spectrum is not high. Since the instability-induced entrainment process of heavy nuclei in internal shocks is currently not well-known, we stress that the expected fraction of heavy nuclei injected into the internal shock acceleration process is uncertain. A significant fraction is plausible, but if the entrainment is inefficient, we would in any case expect a proton-dominated composition in GRB internal shock models. In GRB external shocks and hypernova shocks, on the other hand, the abundance of heavy nuclei is dependent on the external medium or the stellar outer envelope and wind enrichment fraction.  { As we were completing our work, the Pierre Auger Collaboration (2007) reported a plausible correlation between UHECR at energies above $6\\times10^{19}$ eV, assuming that they are protons, and AGNs within 75 Mpc selected from a particular catalog. Although this result is statistically significant at the 2.8$\\sigma$ level, as the authors themselves stress, it does not rule out a possible origin in sources, e.g. galaxies, which are distributed in a similar manner as these AGNs.  There is an unresolved tension between the above Auger spatial correlation analysis suggesting protons, and previous Auger results on maximum shower elongations $X_{max}$ suggesting a significant heavy element component for UHECR in the same energy range. { The AGN hypothesis for the UHECRs has been questioned in the analysis of Gorbunov et al. (2007).} In our present paper, we have investigated the conditions under which UHECR originating from GRB or hypernovae could contain a significant fraction of heavy nuclei, as well as the conditions under which these UHECR would be expected to be mainly protons. If any of the observed UHECR were heavy nuclei, these could reach the Earth from distances beyond 100 Mpc, and the number of galaxies which can host a GRB or hypernova increases significantly with distance. Within the proton-inspired 3.2 degree circle containing an AGN in the Auger correlation analysis, there would be many more galaxies which could host a GRB or a hypernova. In the case of heavy nuclei, the deflection angles are larger (e.g.  for Oxygen with energy $\\ga 6\\times10^{19}$ eV it is $\\la 16$ degree, Sommers 2007), which would contain even more galaxies. While a study of the angular correlations is beyond the scope of this paper, our results are compatible with and relevant for both the Auger elongation studies and spatial correlation studies, providing constraints which can be used in future analyses.}  {"
    },
    "0711/0711.0126_arXiv.txt": {
        "abstract": "We report the discovery of WASP-3b, the third transiting exoplanet to be discovered by the WASP and SOPHIE collaboration. WASP-3b transits its  host star USNO-B1.0 1256-0285133 every $ 1.846834 \\pm 0.000002  $ days. Our high precision radial-velocity measurements present a variation with amplitude characteristic of a planetary-mass companion and in-phase with the light-curve. Adaptive optics imaging shows no evidence for nearby stellar companions, and line-bisector analysis excludes faint, unresolved binarity and stellar activity as the cause of the radial-velocity variations. We make a preliminary spectroscopic analysis of the host star finding it to have $T_{\\rm eff} = 6400 \\pm 100$\\,K and $\\log g = 4.25 \\pm 0.05$ which suggests it is most likely an unevolved main sequence star of spectral type F7-8V. Our simultaneous modelling of the transit photometry and reflex motion of the host leads us to derive a mass of $1.76 ^{+ 0.08 }_{- 0.14 } M_J$  and radius $1.31 ^{+ 0.07 }_{- 0.14 } R_J$ for WASP-3b. The proximity and relative temperature of the host star suggests that WASP-3b is one of the hottest exoplanets known, and thus has the potential to place stringent constraints on exoplanet atmospheric models. ",
        "introduction": "Since the discovery by \\citet{henry2000} and \\citet{c5} of the first exoplanet found to transit its host star, HD209458b, a further 22 transiting systems have been announced (see http://obswww.unige.ch/\\~{}pont/TRANSITS.htm). Transiting exoplanets are highly prized because the transit geometry constrains the orbital inclination, and this in turn allows their masses and radii to be determined directly. The mass-radius relation for exoplanets allows us to probe their internal structure, since it is these parameters which are compared with models of planetary structure and evolution \\citep{sato2005, guillot2006}. The limited numbers of transiting exoplanets studied so far show remarkable diversity in their physical parameters. For example, planets with masses $M\\sim 1M_J$ range in size from 0.8--1.5$R_J$ for reasons that still elude us, although several plausible explanations  have been proposed \\citep{arras2006, burrows2007, fortney2007, gu2004}. The discovery of transiting planets in greater numbers will allow us to further explore the mass-radius plane, and thereby constrain theories of planetary formation, migration and evolution.  The discovery of the first exoplanet, 51\\,Peg\\,b, was by the radial-velocity  method \\citep{mq1}, and this technique is responsible for the discovery of the vast  majority of the known exoplanetary systems, including HD209458b. However, following the discovery of the transits of HD209458b \\citep{henry2000, c5} it was  widely believed that the multiplex advantage of wide-field photometric imaging could lead to this technique becoming the dominant method for detecting exoplanets.  While it is true to say that, at least initially, photometric surveys have been slow to realise their expected detection rates \\citep{h1}, recently this situation has begun to change with 14 new systems published in 2006-07 alone. This improved detection rate is largely due to the development of a better understanding of noise characteristics, especially the correlated noise inherent in such photometric surveys  \\citep{p2}. Transiting planets now comprise approximately 10\\% of the known exoplanets.  Successful exoplanet photometric surveys include Wide Angle Search for Planets (WASP) Project  \\citep{p1}, the Hungarian Automatic Telescope  (HAT) Network \\citep{b2}, OGLE \\citep{ogle}, the Transatlantic Exoplanet (TrES) Survey \\citep{d1,o1}  and the XO group \\citep{m1}. The WASP project has published two new systems in the last year WASP-1b and WASP-2b \\citep{c4} and  WASP-1b, in particular, has proved to be especially interesting, being well oversized compared to other planets for its mass.  In this paper the WASP and SOPHIE collaboration announce the discovery of a new, relatively high mass, strongly-irradiated gas-giant exoplanet, WASP-3b. ",
        "conclusions": "In this study we have found WASP-3b to be a transiting gas-giant exoplanet  with mass $ 1.76 ^{+ 0.08 }_{- 0.14 } M_J$  and radius $ 1.31 ^{+ 0.07 }_{- 0.14 } R_J$.  Its host star, WASP-3, has a photospheric temperature of $6400 \\pm 100$\\,K and $\\log g = 4.25 \\pm 0.05$, consistent with its F7-8V spectral type derived from 2MASS photometry. This places WASP-3b amongst the most massive of known transiting exoplanets (Figure~\\ref{fig:mr}). Given the hosts relatively large rotational velocity and the large radius of the planet we would expect a significant amplitude for Rossiter-McLaughlin  effect (the model is included in Figure~\\ref{fig:sol_rvlc}).  \\begin{figure} \\begin{center} \\psfig{figure=fig_mass_rad.eps,width=7.50cm,angle=0} \\caption[]{ The known confirmed transiting exoplanets plotted in the Mass-Radius plane.  Iso-density contours are plotted in cgs units. For clarity of scale we have not plotted  HD17156 (mass 3.12\\,$M_J$, radius 1.15\\,$R_J$) or the extremely high density object HD147506b (mass 8.04\\,$M_J$, radius 0.98\\,$R_J$).  WASP-3b is marked as a filled circle (data from  http://obswww.unige.ch/\\~{}pont/TRANSITS.htm and references therein).} \\label{fig:mr} \\end{center} \\end{figure}  \\citet{sozz07} demonstrate a correlation of planet radius with host  mass for 14 confirmed transiting exoplanets, and a correlation of planet mass  with orbital period for the same sample (first noted by \\citet{mazeh2005}). Figure~\\ref{fig:planetparas} shows these relationships updated while in Figure~\\ref{fig:planetgravity} we also update the apparent correlation of surface gravity with orbital period noted by \\citet{southworth2007}.  Despite the additional objects the $R_P$ v's $M_*$ correlation remains week (even ignoring the two most massive objects HD17156b and HD147506b). For both the $M_P$ v's $P$ and $g$ v's $P$ we contend that these relationships arise partly through observational selection and partly through the effects of the intense radiation fields that these planets are experiencing. We believe the absence of high gravity/mass bodies at  longer periods is primarily a detection effect, while the absence of low gravity/mass planets at short period could indeed be caused by rapid evaporation.   \\begin{figure} \\begin{center} \\psfig{figure=fig_planetparas.eps,width=8.5cm,angle=0} \\caption[]{ Modified version of Figure 5 from \\citet{sozz07}, incorporating a further seven newly discovered systems (data from  http://obswww.unige.ch/\\~{}pont/TRANSITS.htm and references therein, but again excluding HD17156b and HD147506b). In each case WASP-3b is marked by the filled circle. } \\label{fig:planetparas} \\end{center} \\end{figure}  \\begin{figure} \\begin{center} \\psfig{figure=fig_planetgravity.eps,width=8.5cm,angle=0} \\caption[]{ The correlation of planetary surface gravity with orbital period for the 22 shortest period transiting planets(updated version from Figure~2 from \\citet{southworth2007}). } \\label{fig:planetgravity} \\end{center} \\end{figure}  The closeness of the orbit and the large radius and high effective temperature of the star combine to make WASP-3b one of the most strongly-irradiated, and hence one of the hottest of the known exoplanets, second only to  OGLE-TR-56b and comparable to OGLE-TR-132. This raises the possibility that the atmosphere may be hot enough for TiO and VO to remain in the gas phase above the temperature minimum, creating a hot, strongly-absorbing stratosphere \\citep{fortney2006hd149026b}. This would give an anomalously high infrared brightness temperature, as \\citet{harrington2007hd149026b} inferred from {\\em SPITZER/IRAC} secondary-eclipse photometry of  HD 149026b at 8$\\mu$m. Being much closer and brighter than any of the OGLE host stars, WASP-3b is thus an excellent candidate for future observational tests of the hot-stratosphere hypothesis."
    },
    "0711/0711.3125_arXiv.txt": {
        "abstract": "A high-speed halo-type coronal mass ejection (CME), associated with a GOES M4.6 soft X-ray flare in NOAA AR 0180 at S12W29 and an EIT wave and dimming, occurred on 9 November 2002. A complex radio event was observed during the same period. It included narrow-band fluctuations and frequency-drifting features in the metric wavelength range, type III burst groups at metric--hectometric wavelengths, and an interplanetary type II radio burst, which was visible in the dynamic radio spectrum below 14 MHz. To study the association of the recorded solar energetic particle (SEP) populations with the propagating CME and flaring, we perform a multi-wavelength analysis using radio spectral and imaging observations combined with white-light, EUV, hard X-ray, and magnetogram data. Velocity dispersion analysis of the particle distributions (SOHO and {\\it Wind} {\\it in situ} observations) provides estimates for the release times of electrons and protons. Our analysis indicates that proton acceleration was delayed compared to the electrons. The dynamics of the interplanetary type II burst identify the burst source as a bow shock created by the fast CME. The type III burst groups, with start times close to the estimated electron release times, trace electron beams travelling along open field lines into the interplanetary space. The type III bursts seem to encounter a steep density gradient as they overtake the type II shock front, resulting in an abrupt change in the frequency drift rate of the type III burst emission. Our study presents evidence in support of a scenario in which electrons are accelerated low in the corona behind the CME shock front, while protons are accelerated later, possibly at the CME bow shock high in the corona. ",
        "introduction": "Coronal mass ejections (CMEs) are large-scale structures carrying plasma and magnetic field from the Sun. CMEs are associated with flares, filament eruptions, shocks, radio bursts, solar energetic particle (SEP) events, and have been found to be the primary cause of major geomagnetic disturbances and changes in the solar wind flow (see, {\\it e.g.}, Schwenn {\\it et al.}, \\citeyear{schwenn05}; Vilmer {\\it et al.}, \\citeyear{vilmer03}; Mann {\\it et al.}, \\citeyear{mann03}; Gopalswamy {\\it et al.}, \\citeyear{gopalswamy02}; Leblanc {\\it et al.}, \\citeyear{leblanc01}; Gosling, \\citeyear{gosling97}).  Radio type III bursts are fast frequency-drifting features that can be used as tracers of energetic electrons as they are caused by subrelativistic ($\\approx$2.5 to $\\approx$80 keV) electron streams travelling outward from the Sun at speeds of about 0.1\\,c to 0.5\\,c (see, {\\it e.g.}, Wild, \\citeyear{wild50}; Dulk, \\citeyear{dulk85}). Near-relativistic electrons (30\\,--\\,315 keV) can be observed {\\it in situ} during SEP events, in association with type III bursts that continue into the IP space (Simnett \\citeyear{simnett05}). Radio type II bursts, on the other hand, are thought to be caused by MHD shock waves propagating through the corona and interplanetary medium ({\\it e.g.}, Nelson and Melrose, \\citeyear{nelson85}). Some studies suggest that coronal metric type II bursts and interplanetary (IP) decameter--hectometer (DH) type II bursts do not have the same driver ({\\it e.g.}, Cane and Erickson, \\citeyear{cane05}), and that the latter are exclusively created by the CME bow shocks ({\\it e.g.}, Cliver, Webb, and Howard, \\citeyear{cliver99}).  SEPs are believed to be accelerated in the lift-off phase of CMEs, which could include either acceleration in flaring processes ({\\it e.g.}, de Jager, \\citeyear{jager86}), acceleration in coronal/interplanetary shocks ({\\it e.g.}, Reames, \\citeyear{reames99} and references therein), or in both ({\\it e.g.}, Klein and Trottet, \\citeyear{klein01}; Klein {\\it et al.}, \\citeyear{klein99}; Kocharov and Torsti, \\citeyear{kocharov02}). Recently, the role of CME interaction in SEP production has also been discussed ({\\it e.g.}, Gopalswamy {\\it et al.}, \\citeyear{gopalswamy02}; Kahler, \\citeyear{kahler03}; Wang {\\it et al.}, \\citeyear{wang05}, and references therein). One of the difficulties in determining the relevant processes lies in the uncertainties of timing. SEPs are typically released no earlier than the maxima of the flare impulsive phases (Kahler, \\citeyear{kahler94}), yet at the time of SEP release the height of the associated CME can vary from one to over ten solar radii (Kahler, \\citeyear{kahler94}; Krucker and Lin, \\citeyear{krucker00}; Huttunen-Heikinmaa, Valtonen, and Laitinen, \\citeyear{huttunen05}). It has also been observed that different particle species can be released non-simultaneously (Huttunen-Heikinmaa, Valtonen, and Laitinen, \\citeyear{huttunen05}; Krucker and Lin, \\citeyear{krucker00}; Mewaldt {\\it et al.}, \\citeyear{mewaldt03}).  The {\\it in-situ} observations of SEPs are strongly influenced by the magnetic connection between the spacecraft and the flare site ({\\it e.g.}, Cane, Reames, and von Rosenvinge, \\citeyear{cane88}). For example, the nominally well-connected area for the spacecraft in the L1-point is around 60$^{\\circ}$ West on the solar hemisphere. \\inlinecite{kallenrode93} found that at larger angular distances between the observer's magnetic footpoint and the flare location, the onsets of energetic particles in relation to the microwave maximum are more delayed, and the proton onsets are relatively more delayed than the electron onsets.  Assuming that SEPs originated from the flaring processes, \\inlinecite{kallenrode93} concluded that the propagation of $\\approx$20 MeV protons from the flare site to the connecting magnetic field line is systematically slower than of $\\approx$0.5 MeV electrons, and the electron-to-proton ratio tends to increase with angular distance from the flare site. Of course, the modern view of particle acceleration in the solar corona does not rely on flares as the only source of energetic particles during SEP events. Thus, instead of reflecting the azimuthal particle propagation conditions in the corona, the findings of Kallenrode can be re-interpreted as reflecting the azimuthal and/or temporal evolution of the source, usually a coronal shock wave driven/launched by a CME; a delay between the plasma eruption and the observed onset of the SEP event reflects the establishment of the magnetic connection between the source and the observer and/or the efficiency of the acceleration and escape of particles in different parts (bow/flanks) of the global coronal shock wave.  Consistent with the current views and the conclusions of \\inlinecite{kallenrode93} are the results of \\inlinecite{krucker99}, who found that in some events the injection time of the electrons is consistent with the timing of the type III radio emission, while in other cases the electrons were apparently released up to half an hour later than the start of the type III burst. The non-delayed electron events were concentrated in the well-connected area, while the delayed events were more spread out. Similar delays were also observed between protons and helium nuclei by \\inlinecite{huttunen05}. The near-simultaneous proton and helium events originated from the well-connected area, and the events with delayed helium nuclei were more spread out.  In this paper we study the possible SEP accelerators on 9 November 2002. In Section 2 we analyse the multiwavelength observations of this event: in 2.1 we describe how SEP onset times were estimated, in 2.2 we analyse the flaring activity and CMEs near the SEP onset times, and in 2.3 we identify accelerated particles via radio emission. In Section 3 we calculate heights for the radio features and compare them with other structures. Results are presented in Section 4 and they are further discussed in Section 5. ",
        "conclusions": "The release times for solar energetic particles, electrons and protons, can be obtained from, {\\it e.g.}, velocity dispersion analysis. The feasibility of this method can be argued, as the path lengths can vary and there can be unknown propagation effects in the IP space. In our analysis protons in the 14\\,--\\,80 MeV energy range yield path lengths of 1.0\\,--\\,1.4 AU, while the analysed electrons suggest much higher path lengths of 2.4\\,--\\,2.7 AU. The path lengths of course have a strong effect on the estimated release times: with the low path lengths, electrons would be released just before the protons, and with adopting the high path lengths just for the electrons, electron release would have preceded proton release by more than half an hour.  Very high path lengths are usually taken as an indication of not-so-accurate timing results (Lintunen and Vainio, \\citeyear{lintunen04}; Saiz {\\it et al.}, \\citeyear{saiz05}). High path lengths can be the result of delayed onset observation on low-energy channels due to high pre-event backgrounds, high contrast between the spectral indices of pre-event background and the SEP event, prolonged injection of the particles, or interplanetary scattering of the particles. Therefore the release times yielded by the lower path lengths are more plausible.  We do not speculate on the idea of a ``seed population'' for the released particles, although there is the possibility of injecting (flare) particles into the corona where the particles get accelerated later. The timing and the high start frequencies of the metric type III bursts that continue to the IP space suggest, however, that at least the energetic electrons are accelerated low in the corona and after the flare maximum.  In summary, from the estimated release times for the electrons and protons compared with the multi-wavelength signatures of flare and coronal mass ejection processes, we conclude that electrons were most probably accelerated at low coronal heights in the wake of a fast halo-type CME -- the source region could be reconnecting structures either on the side of the CME or near its legs -- and that the best candidate for proton accelerator is the CME bow shock, at a time when the CME was propagating high in the corona, above 9 R$_{\\odot}$.  The tilts and brightenings in the type III bursts at the times when they pass the IP type II burst is a new observation. We have presented a possible geometry leading to the observed dynamics, and point out that the tilted type III bursts can reveal propagating shock fronts in cases when the type II burst lanes are not visible in the spectrum."
    },
    "0711/0711.4369_arXiv.txt": {
        "abstract": "% In this contribution, I review the applications and potential limitations of the spectral energy distribution fitting tool that I have developed, with a strong emphasis on the limits to which this tool can be used to improve our understanding of massive star formation. I discuss why our current grid of models cannot be used to distinguish between the several competing theories of massive star formation. I also discuss stellar mass determinations, artificial correlations between parameters in the grid of models, multiplicity, confusion, dust assumptions, and unique fits. I briefly review the improvements we intend to carry out for our next grid of models, which will eliminate many of these limitations. Finally, I show examples of applications of this tool to massive young stars. ",
        "introduction": "In Robitaille et al. (2006, hereafter R06), we presented a large grid of model spectral energy distributions (SEDs) for young stellar objects (YSOs) for a range of evolutionary stages (from embedded protostars to optically thin disks) and stellar masses (from 0.1M$_\\odot$ to 50M$_\\odot$), and we used this grid of models to explore the relation between the physical conditions in young stars and their colours at near- and mid-infrared wavelengths. We then presented a linear regression SED fitting technique (Robitaille et al., 2007, hereafter R07) which we used to analyse the SEDs of young stars in the Taurus-Auriga star-formation region, as a proof of concept that this technique can correctly identify the physical conditions in YSOs. The tool has subsequently been used for a number of studies, including an analysis of the YSO population in M16 (Indebetouw et al., 2007), the IRAS 18507+0121 star formation region (Shepherd et al., 2007), the Rosette Molecular Cloud (Poulton et al., 2007), and star formation regions in the SMC (Simon et al., 2007) and LMC (Whitney et al., 2008). The models and the fitting tool are available to the community through a web interface\\footnote{http://www.astro.wisc.edu/protostars}. The web-based tool has already been used by various groups to model massive YSOs, and some of their results were presented at this meeting.\\\\  Since this tool is being widely used to analyse the SEDs of low, intermediate, and high-mass YSOs, I use this contribution to re-iterate the main caveats and limitations that apply when using this tool (\\S\\ref{limit}), to discuss planned improvements in the future grid of models (\\S\\ref{future}), and to show applications of this tool to massive young stellar objects (\\S\\ref{discuss}). ",
        "conclusions": "\\label{discuss}  \\begin{figure}[t] \\begin{center} \\plotone{g34.4mm.eps} \\caption{\\textit{Left} - broadband SED for the embedded protostellar object G34.4+0.23MM (filled circles) shown with several model SEDs providing a good fit (solid lines). \\textit{Right} - \\textit{Spitzer} IRAC 8.0$\\mu$m image of the source (shown circled), which shows two diffuse sources which could be tracing scattered light from a bipolar cavity carved by an outflow.\\label{g34.4mm}} \\label{default} \\end{center} \\end{figure}  In this contribution, I have placed a strong emphasis on the limitations of the YSO SED modeling technique we have developed. However, with these limitations in mind, the models and SED fitting tool can be used to determine useful information about massive YSOs.\\\\  An example of useful quantitative application of the fitting tool to the study of massive YSOs is the determination of bolometric luminosities of massive YSOs by the Red MSX Source (RMS) survey team (Hoare et al., 2004). As discussed in Section \\ref{tracks}, parameters such as the stellar mass, age, or the envelope accretion rate are not parameters which relate directly to the SEDs, as opposed to parameters such as the bolometric luminosity and stellar temperature. While luminosities could be determined simply by integrating SEDs and assuming isotropic emission, the strength of the SED fitting technique is that it implicitly takes into account that the emission is not necessarily isotropic (since the models can have non-isotropic emission depending on the dust geometry), and therefore gives a handle on the uncertainties on the bolometric luminosity.\\\\  Another example, qualitative this time, is that of the embedded protostellar object G34.4+0.23MM (Shepherd et al., 2007). This source is very bright as seen by \\textit{Spitzer} at 24$\\mu$m, whereas IRAC observations show two faint diffuse sources, shown in Figure \\ref{g34.4mm}. By using the SED fitting tool, we found that we can fit the observations well with models of an embedded massive star seen edge-on with bipolar cavities (which tend to scatter light). The two diffuse sources seen at IRAC wavelengths might then correspond to the scattered light from the bipolar cavities. Further support to this interpretation are that a CO outflow is seen emanating from this source, and is aligned with the potential bipolar cavities.\\\\  In summary, I have provided in this contribution a review of some of the main limitations in our models and SED fitting technique, in particular when applied to massive YSOs, as well as examples of use of the fitting tool to derive quantitative or qualitative information about massive YSOs. The work carried out so far is only the first step towards improving our ability to use SEDs to learn about massive star formation, and we will address a number of limitations with the next grid of models that will be computed at the start of 2008."
    },
    "0711/0711.1857_arXiv.txt": {
        "abstract": "We present a catalog of 535 planetary nebulae discovered in the flattened elliptical galaxy NGC 4697, using the FORS1 Cassegrain spectrograph of the Very Large Telescope of the European Southern Observatory at Cerro Paranal, Chile. The catalog provides positions ($x, y$ coordinates relative to the center of light of NGC 4697, as well as $\\alpha, \\delta$), and, for almost all PNs, the magnitude $m$(5007) and the heliocentric radial velocity in km s$^{-1}$. ",
        "introduction": "We present a catalog of 535 planetary nebulae (PNs) discovered in the flattened elliptical galaxy NGC 4697, using the FORS1 Cassegrain spectrograph of the Very Large Telescope of the European Southern Observatory at Cerro Paranal, Chile. The observations and scientific results have been described in M\\'endez et al. (2001, in what follows Paper I). The interpretation of the most important result, which can be described as a Keplerian decline of line-of-sight velocity dispersion as a function of angular distance from the center of the galaxy, has been quite difficult. Several theoretical groups are working on the development of better numerical models, including three-integral models better adapted to a flattened dynamical system like NGC 4697. We have decided to publish the catalog now, so as to make it available to all interested groups. We have recalculated the radial velocities in one of the observed fields, using a different procedure than the one described in Paper I. After confirming that no important changes were necessary, we have combined the two sets of measurements into one final catalog. The catalog provides positions ($x, y$ coordinates relative to the center of light of NGC 4697, as well as $\\alpha, \\delta$), and, for almost all PNs, the magnitude $m$(5007) and the heliocentric radial velocity in km s$^{-1}$.  We also provide a list of peculiar, mostly emission-line objects that are not planetary nebulae, for eventual future studies. In section 2 we describe the radial velocity recalculations mentioned above, and compare the new versus the old velocities. Section 3 contains the catalog, and Section 4 is devoted to the peculiar objects we detected. ",
        "conclusions": ""
    },
    "0711/0711.1585_arXiv.txt": {
        "abstract": "\\small{A model is developed in which the inflaton potential experiences a sudden small change in its second derivative (the effective mass of the inflaton). An exact treatment demonstrates that the resulting density perturbation has a quasi-flat power spectrum with a break in its slope (a step in $n_s$). The step in the spectral index is modulated by characteristic oscillations and results in large running of the spectral index localised over a few e-folds of scales. A field-theoretic model giving rise to such behaviour of the inflationary potential is based on a fast phase transition experienced by a second scalar field weakly coupled to the inflaton. Such a transition is similar to that which terminates inflation in the hybrid inflationary scenario. This scenario suggests that the observed running of the spectral index in the WMAP data may be caused by a fast second order phase transition which occurred during inflation.} ",
        "introduction": "\\label{sec:intro} The present decade appears to have ushered in a golden age for precise cosmological observations. Observations of the CMB, most recently by the WMAP satellite, combined with measurements of the matter power spectrum from large scale structure, weak lensing surveys and Ly-$\\alpha$ absorption, permit parameters of the `standard cosmological model' to be determined to great accuracy. A consensus appears to be emerging that the late time behaviour of this model is `close to' LCDM with an approximately scale invariant primordial spectrum for density perturbations such as those predicted by the simplest inflationary models. Indeed, the inflationary paradigm has been remarkably successful in providing explanations for some well known properties of the universe including its spatial flatness. It also provides a mechanism for seeding galaxies by generating an approximately flat (scale invariant) primordial perturbation spectrum. Both these predictions of the inflationary scenario have received considerable observational support from measurements of anisotropies in the cosmic microwave background (CMB) as detected by WMAP and other CMB experiments \\cite{wmap,bond_boom,archeops}.  However, it is well known that, although primordial fluctuations spectra expected from inflation are likely to be approximately flat, or scale-invariant ($n_s(k) \\equiv d\\ln P_{\\cal R}(k)/d\\ln k \\simeq 1$), exact scale invariance ($n_s=1$) is achievable only for a very specific class of models \\cite{st05}, while a slightly red spectrum ($n_s \\lleq 1$) appears to be a generic prediction \\cite{pert} of the simplest viable one-parameter family of inflationary models including, in particular, $R+R^2$ as well as new and chaotic inflation \\cite{models}. More sophisticated inflationary models belonging to the slow-roll class may also have $n_s$ slightly exceeding unity, a notable example being hybrid inflation \\cite{hybrid}. Still for all these models $|n_s(k)-1|\\ll 1$, and the running of the slope $\\tilde\\alpha(k)\\equiv d\\,n_s(k)/ d\\ln k$ is also expected to be small : $|\\tilde\\alpha(k)|\\sim |n_s(k)-1|^2\\ll 1$ (tilde is used here to avoid confusion with the Bogoliubov parameter below). Existing CMB and other observational data are just approaching the level of accuracy necessary to detect deviations from exact scale invariance and to distinguish between different inflationary models.  However, while on the one hand the recent WMAP data provide the first evidence for $n_s$ being close to (but slightly less than) unity \\cite{wmap} (this result is at present inconclusive, see e.g. \\cite{PE06}), on the other, WMAP data \\cite{wmap} also suggest a rather large value of the running $|\\tilde\\alpha(k)|\\sim |n_s(k)-1|$, and the existence of local spikes like the `Archeops feature' at $l\\sim 40$ (first found in the Archeops data \\cite{archeops} and subsequently confirmed by WMAP \\cite{wmap}) which may indicate that inflation is altogether more complex than the simplest paradigms presented above. Therefore, though it is not yet clear if these features really exist in the primordial perturbation spectrum and are not foreground effects or statistical flukes, it is important to have a list of possible local 'features' such as bumps, wells, wiggles or spikes, superimposed on an approximately scale invariant smooth spectrum, which are expected in more complicated inflationary models. In particular, the large value of the running, if confirmed, should also be a local feature around the present Hubble scale since its persistence until the very end of inflation is incompatible with the requirement that the inflationary epoch be of sufficient duration (\\ie the number of e-folds $N\\sim 50$) \\cite{EP06}.  Such features in the primordial spectrum of fluctuations are likely to be measured to great accuracy with next generation CMB satellites such as Planck \\cite{planck}. The precise nature of the primordial spectrum is of great importance to cosmology not only because it would lead to an in depth understanding of inflation but also because of its bearing on the values of the remaining cosmological parameters. (Cosmological parameters are usually estimated by means of a procedure such as a {\\em Markov Chain Monte Carlo} (MCMC) scheme which assigns probabilities to cosmological parameters in a multi-dimensional parameter space. Better knowledge of one set of parameters values  can therefore significantly influence the probabilities for the remaining.)  A generic way to obtain local deviations from the approximately flat spectrum in the inflationary scenario is to add additional scalar fields to the inflaton, either explicitly -- which leads to multiple inflation \\cite{KLS85,KL87,SBB89} or implicitly, through a rapid change in the effective potential of the inflaton \\cite{S92}. In both cases, at least one of these additional scalar fields should {\\em not} be in the slow-rolling regime, otherwise the spectrum remains approximately flat \\cite{S85}. In this paper we further investigate the second case with the aim of finding new characteristic features being grafted on the primordial inflationary spectrum which are less peculiar than those investigated previously (in view of the observational fact that there is not too much room for such features).  So, let some scalar field which is weakly coupled to the inflaton experience a {\\em fast} phase transition during inflation, by fast is meant that its characteristic time is much less than the Hubble time $H^{-1}(t)$. This phase transition may be induced by the coupling of the scalar either to the inflaton directly, or to slowly changing space-time curvature. This requirement usually implies that the second field be much heavier than the inflaton as well as $H$. Then, assuming that this transition occurs adiabatically and the second field is always in the state of local thermodynamical equilibrium for fixed external values of the inflaton and space-time curvature, it can be integrated out leaving only some correction to the inflaton effective potential $V(\\varphi)$ that may be non-analytic (see e.g. \\cite{YY06} for detailed derivation).  Following the classification in \\cite{S98}, let us consider possible local non-analytic features of $V(\\varphi)$ arising at the point $\\varphi = \\varphi_0$ when the fast phase transition discussed above occurs during inflation. These features can be discontinuities (jumps) either in $V(\\varphi)$ itself or in one of its derivatives. We denote by $[\\,]$ a jump in the relevant quantity, so that $[A] \\equiv A(\\varphi_0+0) - A(\\varphi_0-0)$. The first three most interesting cases placed in the order of their decreasing peculiarity are:   \\begin{enumerate}  \\item $[V] \\neq 0$ -- a step in the effective potential $V(\\varphi)$. This feature is the most peculiar one. Note that it may not arise from a fast quasi-equilibrium phase transition since that would contradict energy (more precisely, free energy) conservation. The only way to obtain such feature is to assume that duration of the phase transition is comparable to $H^{-1}$, so it is not too fast actually.  Not unexpectedly, this results in a non-universal form of the corresponding local feature in the adiabatic perturbation spectrum which depends on detailed dynamics of the phase transition. In the most generic case, a bump modulated by strong oscillations appears in the power spectrum \\cite{S92,G05}. Significant bumps in the observable part of the spectrum are excluded though one may introduce them at very large $k$, close to the end of inflation, to obtain a significant number of primordial black holes (PBH), see e.g. \\cite{PNN94,GLW96}). If the size of the step is very small, the bump disappears and only a burst of oscillations remains. The latter case was studied in \\cite{step_pot}. However, the most recent analysis shows that there is no definite evidence for such a feature in the observable part of the spectrum \\cite{HCMS07}.  \\item $[V] = 0,~[V'] \\neq 0$. A kink in the potential $V(\\varphi)$ leads to a step in its slope $V'(\\varphi)$. The resulting feature in the power spectrum has a {\\em universal} form which does not depend on the details of the fast phase transition, namely a step with superimposed oscillations which are not as pronounced as in the previous case \\cite{S92}. As pointed in \\cite{S98}, to produce such a discontinuity in $V(\\varphi)$, the phase transition should be of first order.  Though initially such a feature was used to describe an apparent excess in the rich cluster power spectrum at $k\\sim 0.05\\,h^{-1}$ Mpc$^{-1}$ \\cite{LPS98} ($h$ is the present Hubble constant in terms of 100 km s$^{-1}$Mpc${^-1}$ and the present value of the scale factor is taken as unity), it seems now that the only place where such feature with a significant value of the step in the primordial spectrum may still exist is in the vicinity of the present Hubble scale \\cite{step_deriv}. Also, as in the previous case, such a feature may be introduced at scales close to the end of inflation to produce PBH \\cite{BBKP03}.  \\item $[V] = [V'] = 0, [V''] \\neq 0$ and $|[V'']|\\ll H^2$. A sudden small change in the slope of the potential (a kink) leads to a step in its second derivative $V''(\\varphi)$. This, mildest of all discontinuities, can be caused by a fast second order phase transition during inflation \\cite{S98}. The last inequality guarantees that slow-roll inflation continues during and after the phase transition, in contrast to the case of the hybrid inflation, or the case $|[V'']|\\sim H^2$ considered in \\cite{SBB89} where a second phase transition with the opposite sign of $V''$ had to be introduced to restore inflation after a short break (that required much fine tuning, of course). As a result, in contrast to previous cases, corrections to the primordial power spectrum $P(k)$ arise in the next to leading order only (at the same order as the Stewart-Lyth correction \\cite{SL93} in case of a smooth inflaton potential). However, due to the feature in $V(\\phi)$, corrections to $(n_s(k)-1)$ appear to be of the same order as the standard leading ones while corrections to the running of $n_s(k)$ dominate the smooth part of $\\tilde\\alpha(k)$ over a few e-fold interval of scales around the feature. As in the previous case, they have a {\\em universal} form, too$^1$ {\\footnotetext[1] {The second-order phase transition considered in [\\cite{subir}] occurs during a time period $\\delta t \\sim H^{-1}$ that results in the temporal breaking of slow-roll during the transition, in a more peculiar behaviour of the effective mass in the equation (3.1) for $\\xi_k$ and in a step-like or even bump-like behaviour of the perturbation spectrum similar to those occurred in the cases 1 and 2.} }.  \\end{enumerate}  Most recent observational constraints on the inflaton potential $V(\\phi)$ and the Hubble function $H(\\phi)$ obtained using only the assumption that these functions may be Taylor--expanded up to the third order in the range of $\\phi$ corresponding to the observable cosmological window of scales ($1-10^4$ Mpc) show that the first two slow-roll parameters $\\epsilon$ and $\\delta$ (defined in Sec. 2 below) are really small but the validity of the slow-roll expansion beyond them is not established, see \\cite{LV07} and references to previous papers therein. The last, third type of peculiarity is just the one satisfying these conditions. That is why in this paper we study it in detail and find the universal form of the corresponding feature. We shall show that a small step in $V''$ leads to a small step in the primordial spectral index $n_s$ accompanied by oscillations with a decreasing amplitude.  The plan of the rest of the paper is as follows. In Sec. 2 small corrections to the background behaviour due to a jump in $V''(\\phi)$ are calculated and their contribution to the Sasaki--Mukhanov equation for scalar perturbations is found. In Sec. 3 an exact solution for the resulting feature in $P_{\\cal R}(k)$ is obtained. A microscopic model that can produce such a feature in the effective potential is considered in Sec.4, and the required values of its parameters are found. Sec. 5 contains conclusions and discussion. ",
        "conclusions": "The rapid advance and sophistication of cosmology experiments, most notably those associated with measuring fluctuations in the cosmic microwave background and which purport to throw light on the physics of the very early universe, necessitate a close examination of different possibilities for the generation of primordial fluctuations responsible for the CMB signal. In this paper we have demonstrated the possibility of a new kind of perturbation spectrum generated during inflation. We have shown that, if during inflation, the effective mass of the inflaton changes rapidly, then this change results in a universal local feature being imprinted onto the primordial spectrum of density perturbations. Namely, a sudden change in $m_{\\rm eff}^2$ satisfying the condition $|[m_{\\rm eff}^2]|\\ll H^2$ leads to a break in the spectral slope -- equivalently -- to a step in the value of the primordial spectral index $n_s$. This break is accompanied by rapid oscillations decaying both in $P_{\\cal R}(k)$ and $n_s(k)$ away from the transition point. These oscillations are rather small in magnitude as compared to all exact solutions for perturbation spectra with features considered before. The amplitude of the running of the spectral index is rather large, $\\sim (n_s-1)$, at the transition point but decays away from it, too. The precise location of the step in $n_s$ and its amplitude are free parameters of this model whose values must be set after comparing the predictions of this scenario with observations. Note also we have not specified the form of the background power spectrum $\\mathcal{P_R}_0(k)$ on which our feature is superimposed. The form of $\\mathcal{P_R}_0(k)$ must clearly be derived from a concrete physical model. (One might even conjecture, for arguments sake, that $\\mathcal{P_R}_0(k)$ could contain additional features generated by other physical effects such as the existence of a radiative epoch prior to inflation, etc.) We have also demonstrated that a field--theoretic model which can give rise to a step in $V''(\\phi)$ is similar to that used to end inflation in the hybrid inflationary model, though different values of its parameters are required in our case which should satisfy the inequalities (\\ref{k1}), (\\ref{k2}) and (\\ref{k3}). It describes a fast second--order phase transition during inflation that occurs in some other scalar field weakly coupled to the inflaton. Some estimates of the values of the spectral index in this model are given and analytical formulae relating $n_s$ to the fundamental parameters in the inflationary potential are derived; see (\\ref{eq:n1}) \\& (\\ref{eq:n2}). The reader may also like to note that while the treatment in section 3 is quite general and allows the running of the spectral index to be positive (${\\tilde\\alpha} > 0$) as well as negative (${\\tilde\\alpha} < 0$), the microphysical model in section 4 favours  ${\\tilde\\alpha} < 0$ which is suggested by (\\ref{eq:delta_n}). To conclude, the scenario discussed in this paper suggests that the observed running of the spectral index in the WMAP data may be caused by a fast second order phase transition which occurred during inflation. A detailed comparison of this model with observations remains an important problem for further study."
    },
    "0711/0711.4858.txt": {
        "abstract": "{We present the results of optical wavelength  observations of the unusual SMC eclipsing binary system HD 5980 obtained in 1999 and 2004--2005.  Radial velocity curves for the  erupting LBV/WR  object ({\\it star A}) and its close WR-like companion ({\\it star B}) are obtained by deblending the variable emission-line profiles of NIV and NV lines.  The derived masses  $M_A=$58--79 M$_\\odot$ and $M_B=$51--67 M$_\\odot$, are more consistent with the the stars' location near the top of the HRD than previous estimates.  The presence of a  wind-wind interaction region  is inferred from the orbital phase-dependent behavior of He I P Cygni absorption components. The emission-line intensities  continued with the declining trend previously seen in UV spectra. The behavior of the photospheric absorption lines is consistent with the results of Schweickhardt (2002) who concludes that the third object in the combined spectrum, {\\it star C}, is also a binary system with $P_{starC}\\sim$96.5 days, $e=$0.83. % it is also a curious fact that its apparent orbital period is nearly a factor of 5 larger that %the orbital period in the{\\it star A+B} pair. } \\addkeyword{Binaries: close} \\addkeyword{Stars: mass loss} \\addkeyword{Stars: individual (HD5980)}  \\begin{document} \\RescaleTitleLengths{0.95} ",
        "introduction": "The Small Magellanic Cloud system HD 5980 consists of two very luminous  components referred to as  {\\it star A}  and  {\\it star B}  in a relatively close, eclipsing and eccentric orbit  and a third equally luminous source, referred to as {\\it star C}, that may merely be a line-of-sight coincidence.  The eclipsing nature of HD 5980 was discovered by Hoffmann, Stift \\& Moffat (1978), although the actual 19.3-day orbital period was found by Breysacher \\& Perrier (1980; henceforth BP80), and more recently refined by  Sterken \\& Breysacher (1997).  The two eclipses were not separated equally in time, indicating an eccentric  orbit ($e=0.32$, Breysacher \\& Perrier 1991; $e=0.27$ Moffat et al. 1998; $e=0.30$ Kaufer et al. 2002). The presence of strong and broad emission lines as well as  photospheric absorption lines led to its early classification as a WN+OB binary system (Azzopardi \\& Vigneau 1975; Walborn 1977; Breysacher \\& Westerlund  1978). However, Niemela (1988) showed that the Doppler shifts in the N IV 4057 and N V 4603 emission lines moved in anti-phase, concluding that {\\em both} stars in the 19.3-day binary orbit were Wolf-Rayet stars of the nitrogen sequence (WN) and that the absorption-line spectrum had to arise in a third object.  With this result, Niemela et al. (1997)  estimated  the masses of the two stars, M$_A\\sim$50 M$_\\odot$ and M$_B\\sim$28 M$_\\odot$.  From the eclipse light curve, Breysacher \\& Perrier (1991, henceforth BP91) derived stellar radii, R$_A\\sim$21 R$_\\odot$ and R$_B\\sim$15 R$_\\odot$. Its location in the SMC implied a visual absolute magnitude M$_V^{sys}=-$7.3 mag (BP91) and the  BP91 light curve analysis provided good constraints on the  visual absolute magnitude for each of the three stars within the system, from which M$_V^A=$-6.3 mag, M$_V^B=$-5.8 mag, and M$_V^C=$-6.1 mag. These and other parameters for HD 5980 are summarized in the first column of Table 1, which contains information derived from data obtained prior to  1981.  Results from data of other epochs is listed in column 2 (1994; near the time of eruption) and in column 3 (1999-2005; declining phase of the eruption).  The spectral characteristics and visual brightness of HD 5980  gradually changed between  the late 1970's and 1993, when it entered an eruptive state that lasted $\\sim$1 year (Bateson \\& Jones 1994; Barb\\'a et al. 1995; Koenigsberger et al. 1995).  The  activity  involved an increase in visual brightness and mass-loss rate, and a decrease in  wind velocity  and  effective temperature, similar to the eruptive phenomena observed in LBVs.  The radial velocity (RV) variations  in the  post-eruption emission-line spectrum  led Barb\\'a et al. (1996, 1997)  to conclude that the instability producing the outbursts originated in {\\it star A}. This was also confirmed by independent observations of Moffat et al. (1998).   The  hydrogen abundance in {\\it star A} that was derived  near the time of maximum eruption, when its emission-lines are believed to have completely dominated the combined spectrum,  was found to be depleted  (Koenigsberger et al. 1998b). Thus, {\\it star A} is a very massive star that has left  the Main Sequence.  Another fascinating aspect of the system is that in the process leading up to the eruption, {\\it star A} seems to have traversed the entire nitrogen-sequence (WN) of the WR classification system, starting out as a WN3 (Niemela 1988) and reaching WN11 (Drissen et al. 2001) during maximum eruption.  This behavior is unprecedented.  It is important to note that the eruptive activity in HD 5980 has been recorded only once. Thus, {\\it star A} may  still be basically in the pre-LBV evolutionary phase, resembling one of the H-rich (or at least not yet H-poor), luminous WR stars of type WNh, or probably even WNha, like most, if not all of the 10 other WN stars in the SMC (Foellmi et al. 2003).  These are not really {\\it bona fide} WR stars in the classical sense, i.e. in the He-burning stage.  {\\it Star B} is also inferred to be a WR star based on pre-eruption spectroscopic observations that indicated emission-line RV variations consistent with this notion (Breysacher, Moffat \\& Niemela 1982). If both stars are, indeed, WR's, the time-variable wind of {\\it star A} presents  a unique opportunity for studying the effects that a changing wind momentum ratio  produces on the wind-wind interaction (WWI).  Curiously, however, the change in the WWI region  during declining phases of the eruptive event appears to have had little effect on the observed X-ray emission (Naz\\'e et al. 2007).  %The wind momentum ratio determines the geometry of the WWI shock cone. %Moffat et al. (1998) discussed the expected effects due to the different wind momentum %ratios prior to, during and after the eruptions in HD 5980.   Koenigsberger (2004) reviewed %the UV observational data and its implications for the WWI scenario, reaching the conclusion %that {\\it star A}'s wind still dominated the wind momentum ratio in the early 2000's. From %a series of FUSE observations, Star A's wind velocity law  was deduced to be typical of %hot massive stars  except in the direction of its close companion %(Koenigsberger et al. 2006).  This raises the possibility   that the WWI region %may reach  far ``down\" into the wind of {\\it star A} (i.e., near its  photosphere), even though %this star's wind dominates the  overall momentum ratio.  Such a phenomenon could result from %``sudden radiative braking\" (Gayley et al. 1997) due to the strong UV radiation field of {\\it star B}. %Alternatively, this result could be indicating that {\\it star A}'s wind structure on the %hemisphere facing its companion is {\\em intrinsically} different  from the wind structure %emerging in other directions.  Koenigsberger \\& Moreno (2007) have recently suggested %that the tidal flows on the surface of {\\it star A} could lead to a shearing action.  The %associated energy dissipation  rates,  largest on the side facing the companion, could %alter the wind acceleration process. % %Central to resolving these issues is  the line profile variability and the degree to which it %is produced by the WWI. The time-dependent shape of the emission lines is determined by a combination of %mechanisms  that include wind and physical eclipses, wind-wind interactions, sudden radiative braking %and X-ray heating.  Of these, the effects on the line profiles due to eclipses can be easily %modeled (Auer \\& Koenigsberger 1994; Stevens et al. 1992).  However, the effects due to wind-wind %collisions are more difficult to assess. %If we assume that the WWI zone merely represents a boundary between the stellar winds of %the two stars, the computed line profile variability  (Stevens 1993;  Georgiev \\& %Koenigsberger 2004a) consists of {\\em the absence} of emission  at particular wavelengths, %or the weakening of P Cygni absorption components at particular orbital phases. On the %other hand, if the WWI zone is assumed to be a thin conical layer of enhanced density %where additional emission (with respect to the stellar winds) is produced, the line %profile variability can be modeled in terms of variable emission peaks that are superposed %upon the underlying broad wind emission (L\\\"uhrs et al. 1997; St.-Louis et al. 2006). %Very little guidance is provided from theoretical models regarding the emissivity of %the WWI region at optical and UV wavelengths. % The masses of {\\it star A} and {\\it star B} are the central issue for establishing their evolutionary state.  But a vexing problem arises when radial velocity curves of the emission lines are used to derive the stellar masses.  The  centroid of an emission line is ill-defined when the line is asymmetrical and undergoes strong variability over the orbital cycle, as is the case for HD 5980.  Without  a clear understanding of the line profile variations (lpvs), the orbital motion cannot be reliably ascertained and, hence, one cannot establish the primary cause of the variability.  The objectives of this paper are to use new and archival observations to  establish constraints on the sources of the line-profile variability in HD 5980 and thus  obtain  RV curves that may be more representative of the orbital motion.  Section 2 summarizes the  data; in Section 3 we describe the line profile variability, leading to a new set of RV curves (Section 4) from which the masses of {\\it star A} and {\\it star B} are estimated;  Section 5 summarizes the evidence suggesting that the wind density of {\\it star A} is diminishing over long timescales; Section 6 describes the stationary photospheric absorptions; Section 7 presents the results of the photometric monitoring; in Section 8 we discuss the wind-wind interaction region; and in Section 9 we present the conclusions. ",
        "conclusions": "In this paper we  present  results from  spectroscopic and photometric observations of HD 5980 obtained in  1999 and 2003-2005.  We first analyze the line-profile variability in the 2005 spectra, concluding that in lines that are formed in the dense, inner wind regions, line-profile-variability may be interpreted mostly in terms of the presence of two emission components, the first belonging to {\\it star A} (the eruptor) and the second associated with {\\it star B} (the companion).   Assuming that the emission-line profiles can be effectively separated with a two-function fit,  we find that the centroids of the two functions are consistent with the orbital motion of the two stars, leading to a new estimate of the masses, M$_A\\sim$58--79 M$_\\odot$ and M$_B\\sim$51--67 M$_\\odot$. These values are significantly larger than previously derived (Niemela et al. 1997), but still imply that the stars must have lost considerable amounts of mass before reaching their current evolutionary state, given their location on the HRD near 120 M$_\\odot$ ZAMS evolutionary tracks (Koenigsberger 2004).  The mechanisms by which such large amounts of mass may have been stripped require detailed examination.  In addition, it would be very interesting to get more theoretical insight into their current internal stellar structure and hence, be able to establish their current evolutionary state.  As pointed out by Koenigsberger et al. (1998b), if {\\it star B} is a {\\it bona fide} WNE remnant \\footnote{that is, the bare  He- or C-burning core of a massive progenitor that has been stripped of its outer layers, as opposed to a massive star in an earlier evolutionary stage undergoing strong wind mass-loss}  of the originally more massive member of the binary,  its  interior  may have already reached a very advanced stage of evolution and, thus, it may be  nearing the supernova stage.  On the other hand, the significantly larger mass that we have now derived for {\\it star B} (M$_B\\sim$56--67 M$_\\odot$ {\\it vs.} $\\sim$28 M$_\\odot$ as believed previously) suggests that it may not  yet have reached the core He burning  evolutionary phase.  Since {\\it star B} is believed to be the source of  the highly He-enriched WR-spectrum observed in the 1970's (Breysacher, Moffat \\& Niemela 1982), the implications for stellar evolution theory of such a massive, He-rich object are potentially very interesting. %If this is the case, the  presence in this star of prominent He II, N V and N IV emission lines %is intriguing.  Line strengths in the 2005 spectra are significantly weaker than those observed in 1999, consistent with a similar decreasing trend observed in UV lines, and suggesting that the eruptor is reverting to a more quiescent state.   The question of whether another activity cycle will start in the near future is an important one, since the timescale for the eruptive events provides an estimate of the total amount of mass that may be lost through this mechanism.  In addition, the activity timescale provides a clue to understanding the underlying mechanism responsible for the instability.  Continued observations are required to determine  the long-term variability pattern.  Visual photometric observations in 2005 provide a light curve covering the entire orbital cycle, as well as a detailed description of the system's brightness variations during selected orbital phases. In particular,  relatively large-amplitude ($\\sim$0.035 mag) fluctuations appear around secondary minimum, when {\\it star B}, the non-erupting star, is ``in front\".  Continued systematic observational monitoring around this orbital phase would be very valuable because of the  strong and rapid variations  observed in visual photometry, polarimetry and UV P Cygni line profiles. In addition, the X-ray appears to be greater here than around  the opposite eclipse.  The temporal behavior of the P Cygni absorption components of the  He I lines is qualitatively consistent with the presence of a denser wind structure folding around {\\it star B} which we associate with the WWI region.  However, the observed speeds are not consistent with the simple models. In particular, the flow speeds are significantly larger than predicted in the radiative shock limit and thus suggest that at least one of the shocks is adiabatic.  It is interesting to note that the hemisphere of {\\it star A} that faces the companion is predicted to be particularly active due to tidal interaction effects (Koenigsberger \\& Moreno 2007), raising the question of how this activity may affect the structure of the WWI region.  The photospheric absorption lines remain stationary over the 19-day orbital  cycle.  However, we find a systematic shift of $\\sim$20 km/s in  their radial velocities observed in 1999 with respect to those observed in 2005. This is consistent with the conclusion of Schweickhardt (2000) that {\\it star C}, the third source in the combined HD 5980 spectrum, is itself a binary system, with P$_{starC} \\sim$96 days. It is unclear, however,  whether the {\\it star C} system is gravitationally bound to the {\\it star A+star B} pair with a much longer orbital period.  It is  a curious fact that the {\\it star C} orbital period is nearly a factor of 5 longer than the orbital period in the {\\it star A+ star B} pair, which may indicate that the two systems are physically bound.  If this were the case, then the eruptive behavior in HD 5980 could be linked to the gravitational perturbations produced during periastron passages. In this case the orbital period of the {\\it star C} system around the {\\it star A+ star B} system would be $>$25 years, since only one large eruption has been seen over the $\\sim$40 years that the system has been studied.  In the Appendix we list the SWOPE and SMARTS photometry averages for each day of observation (Tables 7--10). The data presented in this paper will be made publicly available (please contact CF or GK)."
    },
    "0711/0711.4255_arXiv.txt": {
        "abstract": "The World Space Observatory Ultraviolet (WSO/UV) is a multi-national project grown out of the needs  of the astronomical community to have future access to the UV range. WSO/UV consists of a single UV  telescope with a primary mirror of 1.7~m diameter feeding the UV spectrometer  and  UV imagers. The spectrometer comprises three different spectrographs, two high-resolution echelle spectrographs (the High-Resolution Double-Echelle Spectrograph, HIRDES) and a low-dispersion long-slit instrument. Within HIRDES the 102--310~nm spectral band is split to feed two echelle spectrographs covering the UV range 174--310~nm and the vacuum-UV range 102--176~nm with high spectral  resolution ($R>50\\,000$). The technical concept is based on the heritage of two previous ORFEUS SPAS missions. The phase-B1 development activities are described in this paper considering performance aspects, design drivers, related trade-offs (mechanical concepts, material  selection etc.)  and a critical functional and environmental test verification approach. The current state of other WSO/UV scientific instruments (imagers) is also described. ",
        "introduction": "The World Space Observatory (WSO/UV) will provide future access to high-resolution far-UV  spectroscopy. WSO/UV is an international collaboration led by Russia (Roscosmos) to build a UV (102--310~nm)  mission with capabilities which are presently and in the near and long-term future unavailable to the world-wide  astronomical community. The mission is scheduled for launch in 2010. The planned instrument sensitivity will exceed that of HST/STIS by a factor of 5--10  and all observing time will be available for UV astronomy. The present mission design comprises a 1.7~m telescope. The focal-plane (FP) instruments consist of two high-resolution spectrographs ($R\\sim 55\\,000$) covering  the 102--310~nm range and a long-slit (1''~$\\times$~75'') low-resolution ($R\\sim 500-5000$) spectrograph. Although the primary science of the WSO/UV mission is spectroscopy, high spatial-resolution UV imaging instruments are foreseen. Additionally, a direct imager which samples the best diffraction-limited resolution of the optical system is implemented. WSO/UV will be operated like a ground-based telescope, i.e., it will perform  ``real-time'' operations in an orbit with reduced visibility constraints (high-Earth orbit).  Overviews of the WSO/UV mission and its science case were given by \\citet{Ba03} and \\citet{Go06}, respectively. ",
        "conclusions": "The WSO/UV mission will provide high-resolution UV spectroscopy and imaging. We have presented parts of the phase-B1 study of the main instrument of the WSO/UV mission, the High-Resolution Double-Echelle Spectrograph HIRDES. With the choice of a ceramic material (CeSiC) for the structure of the optical benches of the individual standalone spectrographs, the high spectral resolution can be achieved without active control of optical elements. Furthermore it is not necessary to install complex focusing mechanisms.  The introduced mirror, detector and spectrometer suspensions are already successfully qualified by other space missions.  {\\bf Acknowledgement} The phase-B1 study of HIRDES was supported and financed by the Bundesministerium f\\\"ur Wirtschaft und  Technologie, Deutsches Zentrum f\\\"ur Luft und Raumfahrt e.V. (DLR), through grant FKZ 50\\,QV\\,0503."
    },
    "0711/0711.3196_arXiv.txt": {
        "abstract": "{} {Very high energy (VHE; E$>$100 GeV) $\\gamma$-ray studies were performed for 18 active galactic nuclei (AGN) from a variety of AGN classes.} {VHE observations of a sample of 14 AGN, considered candidate VHE emitters, were made with the High Energy Stereoscopic System (HESS) between January 2005 and July 2007. Large-zenith-angle observations of three northern AGN (Mkn\\,421, Mkn\\,501, 1ES\\,1218+304), known to emit VHE $\\gamma$-rays, were also performed in order to sample their spectral energy distributions (SEDs) above 1 TeV.  In addition, the VHE flux from 1ES\\,1101$-$232, previously detected by HESS in 2004-2005, was monitored during 2006 and 2007.} {As significant detections from the HESS observation program are reported elsewhere, the results reported here are primarily integral flux upper limits. The average exposure for each of the 14 VHE-candidate AGN is $\\sim$7 h live time, and the observations have an average energy threshold between 230 GeV and 590 GeV. Upper limits for these 14 AGN range from $<$0.9\\% to $<$4.9\\% of the Crab Nebula flux, and eight of these are the most constraining ever reported for the object. The brief ($<$2.2 h each) large-zenith-angle observations yield upper limits for Mkn\\,501 ($<$20\\% Crab above 2.5 TeV) and 1ES\\,1218+304 ($<$17\\% Crab above 1.0 TeV), and a marginal detection ($3.5\\sigma$) of Mkn\\,421 (50\\% Crab above 2.1 TeV). 1ES\\,1101$-$232 was marginally detected ($3.6\\sigma$, 1.7\\% Crab above 260 GeV) during the 2006 (13.7 h live time) observations, but not in the 2007 (4.6 h live time) data. The upper limit in 2007 ($<$1.9\\% Crab above 260 GeV) is below the average flux measured by HESS from 2004-2006.} {} ",
        "introduction": "Active galactic nuclei (AGN) represent the only class of extragalactic objects known to emit VHE $\\gamma$-rays. These objects are found in the core of at least 5\\% of all galaxies.  Generally, they are characterized by very bright, highly variable, non-thermal emission spanning the entire electromagnetic spectrum from radio waves to TeV $\\gamma$-rays.  They are believed to be powered by accretion of matter onto a super-massive ($10^6$-$10^9$ solar mass) black hole. In the unified description of AGN (as reviewed in \\cite{AGN_model}), this black hole is surrounded in the inner regions by an accretion disk, and in the outer regions by a thick torus of gas and dust. Viewing these objects at various orientation angles with respect to the torus plane is believed to be the underlying cause of the wide variations in their observed properties, and hence numerous AGN classifications.  In about 10\\% of all AGN, collimated relativistic outflows of particles (known as relativistic jets) exist, presumably along the magnetic field in the vicinity of the black hole, and approximately perpendicular to the accretion disk and torus plane. Of particular interest to high-energy astronomy is a class of radio-loud AGN known as blazars which includes both  BL\\,Lacertae (BL\\,Lac) type objects and Flat Spectrum Radio Quasars (FSRQ). For blazars, one relativistic jet is pointed close the observer's line-of-sight (i.e. towards Earth), causing the observed emission to be relativistically beamed. Blazars are typically characterized by a double-peaked broad-band SED, and for different objects the peak energy of the lower or higher-energy components can differ by several orders of magnitude. BL Lac objects are commonly categorized \\cite{BLLacs} into three groups depending on the position of the first peak: low (LBL), intermediate (IBL), and high-frequency-peaked (HBL). Essentially all\\footnote{Out of the $\\sim$20 VHE AGN, the only published exceptions are the radio galaxy M\\,87 (\\cite{HEGRA_M87,HESS_M87}) and BL\\,Lacertae \\cite{MAGIC_BlLac} (an LBL). See Section~\\ref{VHE_exception} for more details.} AGN detected at VHE energies are HBL.  The HESS array (\\cite{HESS_jim, HESS_crab}) of four imaging atmospheric-Cherenkov telescopes located in Namibia is used to search for VHE $\\gamma$-ray emission from various classes of astrophysical objects. Approximately 30\\% of the HESS observation program is dedicated to studies of AGN, primarily blazars. Since 2005, the $\\sim$300 hours per year of AGN observations are divided between monitoring known VHE-bright AGN for bright flaring episodes, and searching for new VHE sources. For the discovery part of the AGN program, a candidate from a large, diverse sample of relatively nearby AGN is typically observed for $\\sim$10 hours.  If any of these observations show an indication for a signal (e.g., an excess with significance more than $\\sim$3 standard deviations), a deeper exposure is promptly scheduled to increase the overall significance of the detection and to allow for a spectral measurement. The detections resulting from the HESS AGN observation program are reported elsewhere (see Table~\\ref{HESS_detections} for references). Results of HESS AGN observations taken from January 2005 through July 2007, where VHE emission is not significantly detected, are presented here.  \\begin{table} \\begin{minipage}[t]{\\columnwidth} \\caption{Ten AGN detected by HESS in order of redshift ($z$).} \\label{HESS_detections} \\centering \\renewcommand{\\footnoterule}{} \\begin{tabular}{c c c} \\hline\\hline \\noalign{\\smallskip} AGN\\footnote{The asterisk denotes the seven AGN which were discovered to emit VHE $\\gamma$-rays by HESS. The dagger marks the objects with more than one HESS publication.} & $z$ & Reference\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} M\\,87 & 0.004 & \\cite{HESS_M87}\\\\ Mkn\\,421 & 0.030 & \\cite{HESS_421}\\\\ PKS\\,0548-322$^{*}$ & 0.069 & \\cite{HESS_0548}\\\\ PKS\\,2005-489$^{*}$ & 0.071 & \\cite{HESS_2005}\\\\ PKS\\,2155$-$304$^{\\dagger}$ & 0.116 & \\cite{HESS_2155}\\\\ 1ES\\,0229+200$^{*}$ & 0.139 & \\cite{HESS_0229}\\\\ H\\,2356$-$309$^{* \\dagger}$ & 0.165 & \\cite{HESS_2356}\\\\ 1ES\\,1101$-$232$^{* \\dagger}$ & 0.186 & \\cite{Gerd_1101}\\\\ 1ES\\,0347$-$121$^{*}$ & 0.188 & \\cite{HESS_0347}\\\\ PG\\,1553+113$^{*}$ & ? & \\cite{HESS_1553}\\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{minipage} \\end{table} ",
        "conclusions": "HESS observed a large sample of AGN between January 2005 and July 2007 as part of a campaign to identify new VHE-bright AGN. Results presented here detail the observations of 14 candidates for which no significant excess was found. The corresponding upper limits on the VHE flux are the most stringent to date for eight of the candidates, and are only surpassed by those from earlier HESS observations \\cite{HESS_AGN_UL} for the other six candidates. In addition the results from HESS observations of four AGN known to emit VHE $\\gamma$-rays are presented.  Although only a marginal excess is measured for two of these objects, and no excess is observed from the other two, the presence of simultaneous X-ray data for three of these objects enables more accurate SED modeling than possible with archival data.  With the detection of ten VHE AGN, including the discovery of seven, the HESS AGN observation program has been highly successful. However, despite more than five years of operation, the observation program is not complete as many proposed candidates have either not yet been observed or only have a fraction of their intended exposure. Therefore, the prospects of finding additional VHE-bright AGN with HESS are still excellent."
    },
    "0711/0711.3461_arXiv.txt": {
        "abstract": "We present the first results from a new 250, 350, and 500\\,\\micron\\ Galactic Plane survey taken with the Balloon-borne Large-Aperture Submillimeter Telescope (BLAST) in 2005. This survey's primary goal is to identify and characterize high-mass proto-stellar objects (HMPOs).  The region studied here covers 4\\,deg$^2$ near the open cluster NGC~6823 in the constellation Vulpecula ($\\ell=59^\\circ$). We find 60 compact sources ($<60''$ diameter) detected simultaneously in all three bands.  Their spectral energy distributions (SEDs) are constrained through BLAST, {\\it IRAS}, {\\it Spitzer} MIPS, and {\\it MSX} photometry, with inferred dust temperatures spanning $\\sim 12$--40\\,K assuming a dust emissivity index $\\beta=1.5$.  The luminosity-to-mass ratio, a distance-independent quantity, spans $\\sim0.2$--130\\,L$_\\odot$\\,M$_\\odot^{-1}$. Distances are estimated from coincident $^{13}$CO$(1 \\rightarrow 0)$ velocities combined with a variety of other velocity and morphological data in the literature. In total, 49 sources are associated with a molecular cloud complex encompassing NGC~6823 (distance $\\sim 2.3$\\,kpc), 10 objects with the Perseus Arm ($\\sim 8.5$\\,kpc) and one object is probably in the outer Galaxy ($\\sim 14$\\,kpc).  Near NGC~6823, the inferred luminosities and masses of BLAST sources span $\\sim 40$--$10^4$\\,L$_\\odot$, and $\\sim15$--$700$\\,M$_\\odot,$ respectively.  The mass spectrum is compatible with molecular gas masses in other high-mass star forming regions.  Several luminous sources appear to be Ultra Compact \\ion{H}{2} regions powered by early B stars. However, many of the objects are cool, massive gravitationally-bound clumps with no obvious internal radiation from a protostar, and hence excellent HMPO candidates. ",
        "introduction": "\\label{sec:intro}  The Balloon-borne Large Aperture Submillimeter Telescope (BLAST) is a 2-m stratospheric balloon telescope that observes simultaneously at 250, 350, and 500\\,\\micron\\ using bolometric imaging arrays \\citep{pascale2007}.  During the first BLAST science flight, a 100~hr Arctic flight from Sweden to Canada in June 2005 (BLAST05), BLAST conducted the first sensitive large-scale Galactic Plane surveys at these wavelengths.  The focus of these surveys is the earliest stages of massive star formation. As noted in the recent review by \\citet{zinn2007} our understanding of this evolutionary phase is still limited.  Over the past 15 years submillimeter observations at longer wavelengths (350--1200\\,\\micron), e.g. with SCUBA on the 15-m James Clerk Maxwell Telescope \\citep{holland1999}, or MAMBO on the IRAM 30-m telescope \\citep{kreysa1998,motte2007}, have opened up studies of the earliest evolutionary stages of molecular core collapse and proto-stellar formation.  Such cores lack an internal source of radiation and are very cold ($\\la 25$\\,K), so that they emit the bulk of their radiation at submillimeter wavelengths. Since this emission is optically thin in the submillimeter band, observed flux densities are proportional to column density and mass.  There is a variety of terminology used for these elusive early stages.  For example, \\citet{zinn2007} distinguish cores (size 0.1\\,pc) embedded within clumps (size 0.5\\,pc).  \\citet{motte2007} refer to molecular cloud fragments which could be high-luminosity infra-red protostars, infra-red quiet protostars, or high-mass pre-stellar cores. Along the lines of the latter, here we refer to high-mass proto-stellar objects (HMPOs) as compact sources residing in dense molecular clouds, that have the potential to form (one or more) massive OB stars, having luminosities in the range $\\sim10^2$--$10^5$\\,L$_{\\odot}$, but without associated radio continuum emission.  This latter qualification distinguishes them from massive young stellar objects, a later stage in which a hot star has formed (perhaps still accreting), providing the ionizing radiation necessary to form a high emission measure Ultra Compact (UC) \\ion{H}{2} region (and hence radio emission).  Note that at the distances of high-mass star forming regions surveying instruments detect quite massive clumps, possibly harboring groups of stars or their precursors, and not resolving the substructure associated with individual nodes of collapse; however, once star formation is underway the most massive objects dominate the luminosity and ionization.  BLAST is presently unique in its ability to detect and characterize cold dust emission from a range of pre- and proto-stellar sources, constraining the temperatures of objects with $T \\la 25$\\,K ($\\beta=1.5$) using its three-band photometry near the peak of the spectrum.  An earlier, less sensitive balloon-borne submillimeter telescope called ProNaOS \\citep{dupac2001} had similar goals. Previous studies which do not sample the spectral peak of the thermal emission \\citep[e.g.,][]{johnstone2000, pierceprice2000, kirk2005, reid2005, thompson2005, enoch2006, hill2006, schneider2006, thompson2006, young2006, enoch2007, moore2007} have been limited by their relative inability to measure the temperature, producing large uncertainties in the derived luminosities and masses.  Recent surveys with {\\it Spitzer} MIPS can constrain the temperatures of warmer objects \\citep{car05}, but the youngest and coldest objects are potentially not detected even in the long-wavelength {\\it Spitzer} bands.  BLAST is also very efficient at mapping, and so conducted a series of surveys spread across the entire portion of the Galactic Plane that was available during the flight.  These maps encompass (primarily) known high-mass star forming regions: a 10\\,deg$^2$ map of Cygnus X; a 4\\,deg$^2$ map in Vulpecula (described in detail in this paper); a 3\\,deg$^2$ map in Sagitta; and finally a 6\\,deg$^2$ map in Aquila towards the Galactic Ring Survey molecular cloud GRSMC~45.60+0.30 \\citep{rathborne2004}. The locations and sizes of these maps are shown in Figure~\\ref{fig:coverage}. The Vulpecula map was studied first, as it is the deepest of the four regions, and consists of 6.5 hours of data. The analysis of the other three fields, and those from the 2006 Antarctic flight, including a diffraction-limited 50\\,deg$^2$ map of the Vela molecular ridge, will be presented in forthcoming papers.  \\begin{figure*} \\centering \\includegraphics[width=\\linewidth]{f1.eps} \\caption{Locations of fields for the BLAST05 Galactic Plane surveys. Results from the Vulpecula field are discussed in detail in this paper. The background intensity image is 100\\,\\micron\\ emission from \\citet{schlegel1998}.} \\label{fig:coverage} \\end{figure*}  To put these surveys in context, \\citet{motte2007} state that ``large fast-mapping images with MAMBO-2 are currently the best tool to study density structure of molecular clouds with high spatial resolution.'' In about 33 hours of observations (and 30 more with the less efficient MAMBO) they mapped 3\\,deg$^2$ in Cygnus~X; not the entire region but targeted areas of high extinction.  They catalogued 129 compact sources (down to a scale of 15\\arcsec) to a peak flux level of 80\\,mJy ($5\\,\\sigma$) and proposed an additional 40 somewhat more extended structures (2\\arcmin).  Larger structures (10\\arcmin) were filtered out. In the survey reported here, we mapped a contiguous (no threshold for extinction) 4\\,deg$^2$ in 6.5 hours.  Although the very largest scales (including the DC level) are absent from our maps, our images maintain power well beyond 10\\arcmin; the images reveal interesting large-scale structures crossing the entire region.  As reported below, we catalog 60 sources to a flux level of 10\\,Jy at 250\\,\\micron.  Extrapolating this detection limit using a typical spectral shape ($T = 20$\\,K and $\\beta = 1.5$) this corresponds to 170\\,mJy at 1.2\\,mm.  We also find fewer sources per square degree compared to the \\citet{motte2007} survey of Cygnus~X because this region is slightly more distant, not as active, and our survey did not target areas of high extinction. There are exciting large sensitive surveys to look forward to at these same wavelengths with the SPIRE instrument for {\\it Herschel} \\citep{griffin2003}.  The region chosen for this study was centered near NGC~6823 in the constellation Vulpecula ($\\ell=59^\\circ$), and covers approximately 4\\,deg$^2$. This region of the Galactic Plane is prominent in images of thermal dust emission (e.g., {\\it IRAS} in Figure~\\ref{fig:coverage}), as well as in the radio and the optical. In the radio, the identifiable extended \\ion{H}{2} region excited by the massive stars is called Sh2-86 (nominal position indicated in Figure~\\ref{fig:blastmaps}).  The stellar HR diagram for NGC~6823 has been examined by \\citet{massey1995}.  They find an age of 5--7\\,Myr for the bulk of the stars.  There are several evolved OB supergiants of mass $\\sim 25$\\,M$_\\odot$.  The most massive star, O7 V ((f)), with mass $\\sim40$\\,M$_\\odot$, appears younger (2\\,Myr) than the rest.  In the optical this massive star is responsible for the illumination of many ``elephant trunks'' or ``pillars'' in the eastern portion of the \\ion{H}{2} region.  Very high-mass stars evolve quickly into supernovae.  We note that a supernova remnant (SNR) discovered by \\citet{taylor1992} is within 20\\arcmin\\ of the O7 star in projection ($< 12$\\,pc laterally, assuming the same distance), and even closer to one B supergiant --- nominal positions are indicated in Figure~\\ref{fig:blastmaps}. \\citet{massey1995} find several pre-main sequence stars of mass 5--7\\,M$_\\odot$, which pass through this stage in less than 2\\,Myr. Earlier stages of evolution are revealed by the bright Infra Red Astronomy Satellite ({\\it IRAS}) sources clustered in this region.  At least seven of these objects have been observed in molecular line studies \\citep{beu02,zha05,bel06}  This paper describes a technique developed for deconvolving the out-of-focus beams common to all data from BLAST05 \\citep[see][]{truch2007} and for detecting compact objects in the BLAST maps. We report a robust list of 60 submillimeter sources with sizes $\\la60''$ diameter that are detected simultaneously in the three BLAST bands (\\S\\ref{sec:obs}).  Infra-red photometry for BLAST sources is obtained from comparisons with {\\it IRAS} and {\\it Spitzer} MIPS maps, and the Midcourse Space Experiment ({\\it MSX}) point source catalog (\\S\\ref{sec:firsed}).  The BLAST and infra-red data are combined to constrain dust temperatures and integrated far infra-red (FIR) fluxes using isothermal modified blackbody SED fits (\\S\\ref{sec:coldsed}).  Our analysis has benefited from a $^{13}$CO$(1 \\rightarrow 0)$ data cube obtained as a part of a Galactic Plane survey at FCRAO \\citep{brunt2007}. Noting coincidences between compact structures in the $^{13}$CO emission, and comparing velocities and morphologies with the VLA Galactic Plane Survey, {\\it Spitzer} GLIMPSE \\citep{whitney2005} and Digitized Sky Survey images, distance estimates are obtained for all of the sources in \\S\\ref{sec:dist}. Of 60 BLAST sources, 49 objects are likely associated with the molecular clouds surrounding NGC~6823 at $\\sim2.3$\\,kpc, 10 appear to be in the more distant Perseus arm ($\\sim8.5$\\,kpc), and a single object is thought to lie in the outer Galaxy ($\\sim14$\\,kpc). These distances are used to calculate source luminosities and masses in \\S\\ref{sec:mass}.  In this section we also present an estimate of the mass function for sources associated with NGC~6823.  In this survey high-mass clumps are seen in a range of evolutionary stages: those with embedded UC \\ion{H}{2} regions; luminous objects without substantial radio emission (HMPOs); and cold gravitationally-bound clumps (or cores) with low luminosity-to-mass ratio and no evidence of star formation (\\S\\ref{sec:disc}). ",
        "conclusions": "\\label{sec:concl}  In this paper we present the first maps of the Galactic Plane observed with BLAST during its 2005 LDB flight from Sweden to Canada, specifically the survey in the direction of star forming clouds in Vulpecula.  Fourier deconvolution is used to improve the angular resolution of the maps from $\\sim3\\farcm5$ to $\\sim1$\\arcmin\\ full-width half-power (the theoretical 500\\,\\micron\\ diffraction limit). In these maps, 60 compact submillimeter sources are detected simultaneously at 250, 350, and 500\\,\\micron.  Complementary {\\it IRAS}, MIPSGAL, and {\\it MSX} photometry are used to constrain the submillimeter-MIR SEDs of these objects, and hence infer their cold dust temperatures and bolometric fluxes. Our sample has convincingly revealed objects with a range of dust temperatures, from $\\sim12$--40\\,K under the assumption of an isothermal modified blackbody with dust emissivity index $\\beta=1.5$. With these SED fits, we derive luminosity-to-mass ratios (independent of distance) in the range $0.2$--130\\,L$_\\odot$\\,M$_\\odot^{-1}$. Those with low values are highly shielded and cool, with no evidence for star formation, while those with high values have embedded high-mass star formation.  Using a $^{13}$CO$(1 \\rightarrow 0)$ data cube, the VLA Galactic Plane Survey \\ion{H}{1} cube, {\\it Spitzer} IRAC and optical DSS images, we argue that 49 of the 60 sources lie in a molecular cloud complex associated with the open cluster NGC~6823 at $\\sim2.3$\\,kpc, 10 objects are associated with the Perseus arm at $\\sim8.5$\\,kpc and 1 object is in the outer Galaxy at $\\sim14$\\,kpc.  With these distance estimates we calculate bolometric luminosities and cloud masses associated with the thermal emission from cold dust. The most luminous object ($7\\times10^4$\\,L$_\\odot$) is in the Perseus arm.  Near NGC~6823, the ranges are $\\sim 40$--$10^4$\\,L$_\\odot$, and $\\sim 15$--$700$\\,M$_\\odot$.  A mass function is constructed for the 49 objects associated with NGC~6823. It is compatible with the spectrum of molecular gas masses in other high-mass star forming regions, with a power-law index of $-1.7$ for the $\\ga100$\\,M$_\\odot$ sources. A flattening at lower masses might be present, but is affected by detection completeness of cool sources in this mass range.  The luminosity-mass distribution we find is broadly consistent with an evolutionary sequence, from cool high-mass, low-luminosity clumps (most are not detected with {\\it IRAS} or {\\it MSX} and there is no evidence for any star formation), to more evolved HMPOs and UC \\ion{H}{2} regions.  But interestingly, among the embedded objects in this molecular complex near NGC~6823 -- the next generation -- there are none quite as massive as the 40\\,M$_\\odot$ stars currently powering the nebula."
    },
    "0711/0711.3182_arXiv.txt": {
        "abstract": "{} {We present a study of the diffuse X-ray emission in the halo and the disc of the starburst galaxy \\n253. } {After removing point\\bold{\\colblue{-like}} sources, we analysed \\xmm\\ images, hardness ratio maps and spectra from several regions in the halo and the disc. We introduce a method to produce vignetting corrected images from the EPIC~pn data, and we developed a procedure that allows a correct background treatment for low surface brightness spectra, using a local background, together with closed filter observations. } {Most of the emission from the halo is at energies below 1~keV. In the disc, also emission at higher energies is present. The extent of the diffuse emission along the major axis of the disc is 13.6~kpc. The halo resembles a horn structure and reaches out to $\\sim$9~kpc perpendicular to the disc. Disc regions that cover star forming regions, like spiral arms, show harder spectra than regions with lower star forming activity. Models for spectral fits of the disc regions need at least three components: two thermal plasmas with solar abundances plus a power law and galactic foreground absorption. Temperatures are between 0.1 and 0.3~keV and between 0.3 and \\bold{0.9}~keV for the soft and the hard component, respectively. The power law component may indicate an unresolved contribution from X-ray binaries in the disc. The halo emission is not uniform, neither spatially nor spectrally. The southeastern halo is softer than the northwestern halo. To model the spectra in the halo, we needed two thermal plasmas with solar abundances plus galactic foreground absorption. Temperatures are around 0.1 and 0.3~keV. A comparison between X-ray and UV emission shows that both originate from the same regions. The UV emission is more extended in the southeastern halo, where it seems to form a shell around the X-ray emission. } {} ",
        "introduction": "The diffuse X-ray emission of starburst galaxies can be quite prominent. Especially in galaxies that we see edge-on, we can find very complex emission from galactic halos. One famous example is the starburst galaxy \\n253\\ in the Sculptor Group. It is close enough \\citep[2.58~Mpc, 1\\arcmin=750~pc,][]{PC1991} to resolve structures in the disc and halo, and to separate the detected point\\bold{\\colblue{-like}} sources from the diffuse emission. Also, it is seen almost edge-on \\citep[78.5\\deg,][]{P1980}, so an unobscured analysis of the halo emission is possible. \\n253 has been observed in X-rays many times. There are observations with \\ein\\ \\citep[e.g.][]{FT1984}, \\ros\\ \\citep[e.g.][]{P1992,RP1997,DW1998,VP1999, PV2000}, \\asca\\ \\citep[e.g.][]{PS1997}, \\sax\\ \\citep[e.g.][]{CP1999}, \\xmm\\ \\citep[e.g.][]{PR2001, BP2007}, and \\chandra\\ \\citep[e.g.][]{WH2002, SH2002, SH2004A, SH2004B}. While with some instruments one was not able to separate emission from point sources and diffuse emission, other instruments, especially \\ros, \\xmm, and \\chandra, do have a narrow enough point spread function to do so. We here report on the first extensive analysis of the diffuse emission in \\n253\\ with \\xmm. ",
        "conclusions": "\\subsection{The extent of the diffuse emission of \\n253} Extended emission from the soft northwestern halo was first reported from \\ein\\ observations \\citep[][]{F1988}. Later, observations with \\ros\\ also discovered the southeastern halo in X-rays \\citep[\\eg][]{PV2000}. The \\ros\\ images in the soft band trace the emission in the outer halo to projected distances of up to 9~kpc, both in the northwest and the southeast direction. With \\xmm, the emission is detected out to 9.0~kpc to the northwest and 6.3~kpc to the southeast. This difference in the southeastern halo can be explained by the high \\ros\\ sensitivity extending down to 0.1~keV. The useful \\xmm\\ EPIC~pn range is limited to 0.2~keV. This makes a big difference as there are many strong lines from \\ion{O}{IV}, \\ion{Ne}{VIII}, \\ion{Mg}{IX}, \\ion{Mg}{X}, \\ion{Si}{IX}, and \\ion{Si}{X} in the energy band between 0.1 and 0.2~keV. For a thermal plasma at a temperature of $\\sim$0.1~keV, these lines are even stronger than the \\ion{O}{VII} and \\ion{O}{VIII} lines, and about 60\\% of the total flux in the energy band from 0.1 to 2.0~keV originates from lines below 0.2~keV. The southeastern halo shows softer emission than the northwestern halo, therefore, the effect is strongest in the southeastern halo.  Also in the disc the extent of the emission is different. The \\ros\\ images trace the soft emission $\\sim$6.8~kpc towards the northeast and $\\sim$5.3~kpc towards the southwest. With \\xmm, the disc emission has an extent of $\\sim$7.2~kpc and $\\sim$6.3~kpc to the northeast and southwest, respectively. The disc spectra are harder than the halo spectra, and therefore the higher \\xmm\\ sensitivity at energies $>$0.4~keV comes into play.  \\subsection{Is the diffuse emission in the disc really from hot interstellar gas?}  \\begin{figure*} \\centering \\fig{\\includegraphics[width=\\textwidth]{8935f11.ps}} \\caption{Comparison between the K-band \\citep[left, 2MASS,][]{JC2003} and X-ray (right) morphology. The white contours overplotted on the X-ray image represent the K-band brightness levels. Both images are on a linear colour scale, and on the same spatial scale as indicated in the K-band image.} \\label{fig:XrayVsKband} \\end{figure*}  \\bold{In this section we investigate the possibility that the observed diffuse emission originates in fact not from a truly diffuse hot interstellar gas component, but from a large population of weak stellar-type X-ray sources that shows the same characteristics.}  \\bold{This effect was first discovered in the Milky Way's ridge X-ray emission \\citep[\\eg][]{RS2006}, where evidence was found that the bulk of the Galactic ridge X-ray emission is composed of weak X-ray sources, mostly cataclysmic variables and coronally active binary systems, with a luminosity of most of these sources of less than \\oergs{31}. Also in other galaxies, \\cite{RC2007} found that the apparently diffuse emission is consistent with the emission from an old stellar population like in the Milky Way.}  Following the method by \\cite{RC2007}, we used K-band observations, to infer the emissivity of the diffuse X-ray component per unit stellar mass. We derived the near-infrared luminosity and stellar mass of \\n253, using the total K-band magnitude of 3.772 \\citep{JC2003}, the distance modulus of 27.06, corrected for interstellar extinction of 0.007 \\citep{SF1998}, and the colour-dependent K-band mass-to-light ratio from \\cite{BJ2001}, $\\log(M_*/L_{\\rm{K}})=-0.692+0.652\\times(B-V)$, with $(B-V)=-0.16$ \\citep{CG2003}. This yielded a total K-band luminosity $L_{\\rm{K}}=1.6\\times10^{11}$~$L$\\subsun\\ and a total stellar mass $M_*=2.6\\times10^{10}$~$M$\\subsun. With a X-ray luminosity of \\n253\\ of 2.0\\ergs{39} (0.5--10~keV), the emissivity of the diffuse X-ray component per unit stellar mass then resulted in  \\begin{equation*} \\frac{L_{0.5-10\\,\\,\\mathrm{keV}}}{M_*}=7.6^{+0.9}_{-0.3} \\,(\\pm 2.3)\\times10^{28}\\,\\,\\mathrm{erg}\\,\\, \\mathrm{s}^{-1}\\,\\, M^{-1}_{\\odot}. \\end{equation*}  The errors are statistical errors on the measured X-ray flux. Additionally, we assumed an uncertainty of $\\sim$30\\% (given in parentheses), which might be associated with the $L_{\\mathrm{K}}$ to $M_*$ conversion \\citep{BM2003}. The emissivity of \\n253 should only be considered a lower limit. We cut out a quite large region in the centre of \\n253 and corrected for this by filling the hole with the average flux of the disc. Therefore, the obtained X-ray luminosity as well as the emissivity are probably \\bold{underestimated}.  From the luminosity and other properties of the Galactic ridge X-ray emission \\citep[\\eg][]{RS2006} and from direct measurements of the luminosity function of sources in the solar neighbourhood \\citep{SR2006}, the combined 0.5--10~keV emissivity of cataclysmic variables and coronally active stars has been estimated as $L_{\\mathrm{X}}/M_*\\sim1.2\\pm0.3\\times10^{28}\\,\\,\\mathrm{erg}\\,\\,\\mathrm{s}^{-1}\\,\\,M^{-1}_{\\odot}$. The value derived for \\n253\\ is larger than the value for the Milky Way, indicating the presence of a hot gaseous component.  An even stronger argument is the following: if the diffuse X-ray emission is produced by an old stellar population, then their morphologies should be similar. A comparison of the X-ray emission with the 2MASS K-band image \\citep{JC2003} is shown in Fig.~\\ref{fig:XrayVsKband}. We found that the X-ray morphology does not match the K-band morphology, therefore the diffuse emission is indeed not simply due to an old stellar population, but has to have a truly diffuse component.  \\subsection{Spectral fits and variations in the halo} As it was mentioned already in earlier publications, there is an ambiguity in the spectral fits between a pure multi-temperature thermal plasma model and a combination of thermal plasmas plus a power law component \\citep[e.g.][]{DP2000,SH2002}. This ambiguity in the halo emission still exists with the \\xmm\\ data. Fits to the halo spectra with a thermal plasma plus a power law model (see Table~\\ref{tab:HaloFits}) resulted in similar $\\chi^2_{\\nu}$, as for a multi-temperature thermal plasma model. A power law component from point sources could be excluded, since we were careful to remove any point source contribution. Another source for non-thermal emission could be synchrotron emission from cosmic ray electrons that are advected with the superwind or are accelerated locally in internal wind shocks. \\bold{This is possibly what we observe in \\colblue{the spectrum of }region 6 (and spectra containing region 6). \\colblue{This region }lies along the northwestern outflow axis. Here the spectral fit improves significantly by adding a power law component to the two thermal plasmas}. A comparison of the X-ray emission to the 330~MHz and 1.4~GHz radio emission \\citep{CH1992} showed that the radio emission is more extended, and does not show the horn structure that we see in X-rays. Because of this inequality, we prefer the multi-temperature thermal plasma \\bold{over the single temperature model plus power law model for most of} the X-ray halo emission at the moment.  \\bold{The appearance of a ``multi-temperature'' halo is a natural consequence of the plasma driven out of collisional ionisation equilibrium (CIE) towards a non-equilibrium ionisation (NEI) state in a galactic outflow. It has to be borne in mind that the processes of ionisation, recombination and adiabatic cooling occur all  on different time scales. Whereas the former two are controlled by microphysics the latter is due to the hydrodynamics of the flow. As has been described by \\eg \\cite{BS1994} adiabatic cooling can become the shortest in an accelerating outflow, like the one occurring in a starburst driven wind, leading to delayed recombination of the plasma. This means essentially the presence of highly ionised species in a plasma flow with low kinetic electron temperature, similar to the ions in the solar wind. Releasing the energy stored in these highly ionised species results in diffuse soft X-ray emission and some spectral hardening. Observing this with the limited spectral resolution of \\chandra\\ or \\xmm\\ CCDs, reveals a multi-temperature halo, although physically the plasma does by no means have ``different temperatures'', but is characterised by a non-equilibrium distribution of ions, which by itself is the result of the thermodynamic path the system has taken and reflects directly the hydrodynamic flow characteristics}. A currently ongoing analysis with non-equilibrium models \\citep[\\eg][]{BS1999} might also be able to explain the observations \\bold{and is the subject of a forthcoming paper}.  The northwestern halo shows significant hardness variations in HR1, as opposed to the findings by \\cite{SH2002}. We checked if this can be caused by a different energy band selection, but the result is independent whether we use the bands from \\cite{SH2002} (0.3--0.6~keV and 0.6--1.0~keV) or our own. These hardness variations might also be a sign of NEI X-ray emission.  \\subsection{Temperatures, abundances and column densities} \\label{sec:abund} The X-ray emission from \\n253 has been observed before with several other X-ray observatories. Especially the early observatories did not allow to separate the point sources from the diffuse emission since the point spread function was quite large. Hence, only a combined fit of the emission from the halo, the disc, and the nuclear region was possible. Temperatures of multi-temperature models ranged between 0.1 and 0.3~keV for the low, and between 0.6 and 0.7~keV for the high temperature component \\citep{DW1998, WH2000, DP2000}. Reported abundances were mostly highly subsolar and therefore unphysical for a supposedly metal enriched starburst galaxy plasma.  Only the X-ray observatories like \\xmm, \\chandra, and to some degree \\ros\\ allow us to separate the halo from the disc emission and to remove contribution by point sources via a spatial selection. From \\ros\\ data, \\cite{PV2000} inferred a foreground absorbed two-temperature thermal model with temperatures of 0.13 and 0.62~keV for the northwestern halo emission. No highly subsolar abundances were required. The disc emission could be explained by a 0.7~keV thermal plasma and an additional thermal plasma (kT=0.2~keV) in front of the disc coming from a coronal component.  The first \\xmm\\ results by \\cite{PR2001} required a two-component model for the disc emission with temperatures of 0.13 and 0.4~keV plus residual harder emission, possibly from unresolved point sources. For the nuclear region three temperatures were needed (0.6, 0.9, and 6~keV). Both models used solar abundances. No analysis of the halo emission was presented in their paper.  The best spatial resolution is provided by the \\chandra\\ observatory. Results on the diffuse disc and halo emission were first published by \\cite{SH2002}. For the halo emission they needed a multi-temperature model (apec) with at least two temperatures of 0.24 and 0.71~keV (the latter with quite large errors) and with a foreground absorption of $5.3\\times 10^{20}$~\\cm-2. A power law ($\\Gamma=3.3$) plus a thermal plasma (kT=0.24~keV) gave a similarly good fit. A combination of other thermal models (mekal) or non-equilibrium models (vnei) did not result in better fits. The diffuse emission from the disc was fitted with the same models, however, the temperatures were lower than in the halo, with 0.17 and 0.56~keV, respectively. The foreground absorption yields $4.7\\times 10^{20}$~\\cm-2. In all cases unphysically sub-solar abundances had to be assumed.  The temperature values for the halo emission, as found by our analysis, are lower than the ones from previous observations. Our soft component is about 0.10~keV, which is still compatible with the \\ros\\ results. However, the hard component is only $\\sim$0.31~keV for the northwestern halo and $\\sim$0.29~keV for the southeastern halo. A higher temperature was not necessary in any of our fits. A possible explanation could be the way the spectra were background subtracted. We used a sophisticated method (see App.~\\ref{app:bkgspec}) that uses the local background at the border of the field of view, where no emission from \\n253\\ is expected, while other authors used \\eg blank-sky observations \\citep{SH2002}. Using a background from different times and different fields on the sky can lead to systematic effects in the background substraction. A background region with a higher contribution of the local bubble could, for example, lead to an over-correction, especially at very soft energies ($<$0.5~keV).  In the disc we found temperatures between 0.1 and 0.3~keV and between 0.3 and \\bold{0.9}~keV, for the soft and the hard component, respectively. This is consistent with earlier results.  We also tried to constrain the metal abundances in our fits. However the errors on the obtained values are so large, that we are not able to give well constrained abundances (northwestern halo: $Z=0.3^{+4.7}_{-0.2}$~$Z$\\subsun, southeastern halo: $Z=0.4^{+4.6}_{-0.3}$~$Z$\\subsun, disc: $Z=1.0^{+0.9}_{-0.7}$~$Z$\\subsun). Since we do not expect highly subsolar abundances in an environment which is enriched with metals from the starburst via the superwind and galactic fountains, we fixed the abundances in our analysis to solar. This is entirely consistent with the above values. A reason for the low abundances, found with different instruments, could be that due to a lower spatial resolution and or/and sensitivity more point sources contribute to the final spectrum, increasing the continuum flux. The ratio of line emission to continuum flux is therefore decreased, which mimics the spectral shape of a plasma with low metal abundances. A similar effect can be achieved when a NEI spectrum is fitted with CIE models. Also, a too simplistic model could be the reason, combining regions with different temperatures.  In the disc, \\bold{\\colblue{two}} of the regions required an extra absorption component in the spectral model. The additional column densities range between \\bold{0.76\\hcm{21} and 5.8\\hcm{21}}. Direct radio measurements of the \\ion{H}{I} column density showed lower values than we derived from the X-ray data. An interpolation of the \\ion{H}{I} maps by \\cite{PC1991} and \\cite{KW1995} resulted in \\bold{($\\sim$2.4 and 3.4)\\comment{, and 4.8}\\hcm{21} for regions 9 and 16,\\comment{ and 17} respectively}. However, the \\ion{H}{I} value for the region including the nucleus of \\n253 (region 16) is affected by \\ion{H}{I} absorption, so the resulting column density can only be considered as a lower limit. Additional absorption is expected from molecular hydrogen. \\cite{MH1995} derived the H$_2$ column density in the direction of the nucleus of \\n253 to 3.7\\hcm{23}. Taking this value as an upper limit for the column density in the disc regions around the nucleus, the column densities derived from X-ray spectra are within the limits from radio observations.  \\subsection{X-ray versus UV morphology}  \\begin{figure} \\centering \\fig{\\includegraphics[width=88mm]{8935f12.ps}} \\caption{Two-colour UV image of \\n253 with NUV (175--280~nm) in yellow and FUV (135--175~nm) in blue. The intensity was rescaled to emphasise the faint, diffuse emission. Overplotted are the 0.2--0.5~keV X-ray contours.} \\label{fig:UV} \\end{figure}  Fig.~\\ref{fig:UV} shows the X-ray contours from the energy band 0.2--0.5~keV overplotted on a two-colour UV image, taken with the GALEX observatory (Galaxy Evolution Explorer, a UV space telescope) on 2003 October 13. For the northwestern halo there is quite a good agreement between the FUV and X-ray emission regions. The FUV emission traces the western horn to a distance of $\\sim$7.5~kpc above the disc, as well as the broad base emission in soft X-rays quite well. In the southeastern halo, again, the UV and the soft X-ray emission show the western horn structure, where the UV horn extends to about 7~kpc away from the disc. However, the FUV horn is slightly offset by $\\sim$700~pc to the northeast with respect to the X-ray horn. Images obtained with the Optical Monitor onboard \\xmm\\ are not sensitive enough to show the extraplanar UV emission.  \\cite{HH2005} proposed the following model for the origin of the UV emission: Since the UV luminosities are too high to to be produced by continuum and line emission from photoionised or shock-heated gas, the UV emission could be explained by dust in the outflow that scatters the stellar continuum from the starburst into our line of sight. They also found that the UV halo emission, as seen with GALEX, correlates with the H$\\alpha$ emission, which could originate from gas that is photoionised by UV photons from the starburst. The UV and H$\\alpha$ emission would originate in the same cold regions in the halo.  How does the warm gas that is responsible for the \\ha\\ and UV emission get out into the halo? There are two possibilities: either it has already been there from the beginning in the form of a cold and maybe clumpy halo component, or it was transported by the superwind and galactic fountains from the disc out into the halo.  There are models, where it is possible to drag up clouds of cold gas into the halo \\citep[\\eg model 3 of][]{SH2002}. In a sheet surrounding these clouds, X-ray emission could be produced by shocks or in conductive or turbulent mixing interfaces on the cloud surface. This model would also account for the non-uniformity of the X-ray emission as seen in the \\xmm\\ images (Fig.~\\ref{fig:PN_images} and \\ref{fig:PN_RGB_image}). However, the model cannot explain the displacement of the UV emission in the southeastern halo, since in the model the clouds are located within or at the inner border of the superwind.  Could the dust even survive this transport from the disc into the halo embedded in a hot plasma environment? \\cite{DS1979} give the sputtering time for a spherical dust grain of radius $a$ embedded in a plasma of hydrogen with temperatures between \\oexpo{6} and \\oexpo{9}~K and the density $n_{\\mathrm{H}}$ as \\begin{equation} t_{\\mathrm{sput}}\\sim 10^6 \\left(\\frac{a}{\\mu\\mathrm{m}}\\right)\\left(\\frac{n_{\\mathrm{H}}}{\\mathrm{cm}^{-3}}\\right)^{-1}\\,\\mathrm{yr}. \\end{equation} For $n_{\\mathrm{H}}$ between $2.5\\times10^{-2}$~cm$^{-3}$  in the outflow close to the centre \\citep{BP2007} and $3.2\\times10^{-3}$~cm$^{-3}$ out in the northwestern halo, and a grain size of $a$=0.1~$\\mu$m, $t_{\\mathrm{sput}}$ varies between 4.0 and 31~Myr. So to reach a height above the disc of 7.5~kpc in less than 31~Myr, an average velocity of at least 240~km/s is required. This is well compatible with measurements of outflow velocities in different wavelengths, that range from 260~km/s \\citep[Na D absorption,][]{HL2000} to about 400--600~km/s \\citep[\\ha, \\ion{N}{II}, \\ion{S}{II}, and \\ion{O}{II} emission,][]{U1978,DB1970}. Therefore it is quite possible that the dust survives the transport from the disc out into the halo.  Another model to explain the UV and X-ray morphology \\citep[\\eg model 5 of][]{SH2002} requires a thick disc component, through which the superwind emerges into the halo. On the contact surface between the hot superwind fluid and the cold thick disc material we get a heated layer through shocks and turbulent mixing where the X-rays are produced, surrounded on the outside by a colder layer where the UV emission originates. The thick disc component was originally created by lifting material up from the disc through the star formation activity \\citep[simulations by][]{RB1995}. This model would easily explain the UV displacement from the X-rays, however we would only get a hollow cone with X-ray emission. The latter is not what we see in the \\xmm\\ observations. Though a mix of both models would be able to explain the observed morphology.  \\bold{For completeness, the three other models that \\cite{SH2002} present are:}  \\bold{(i) Shocked clouds in the halo (model 2). In this model preexisting neutral clouds are overrun by the superwind, which can explain the often asymmetrical patchiness as a result of the spatial distribution of the clouds. However, no \\ion{H}{I} structures outside the disc have been found by \\cite{PC1991} and \\cite{KW1995}.}  \\bold{(ii) Hot swept-up shell of halo gas (model 4). In this model a superbubble bounded by a shock expands into the low-density halo medium. The X-ray radiation originates from the shocked ambient gas following the morphology of the surface of the superbubble. Possible UV/\\ha\\ emission could originate from dense cool disc gas entrained in the wind. Especially this last aspect is not supported by the observations.}  \\bold{(iii) Cooling radiation from the wind (model 1). As presented by \\cite{SH2002} this model cannot explain the UV or \\ha\\ emission, and also does not explain the horn shaped morphology of both the X-ray and UV emission.}  The magnetohydrodynamics ISM simulations of \\cite{AB2005} also shows a clumpy halo structure, characterised by turbulent mixing layers, which could explain the UV and X-ray filamentary structure. In some regions, the magnetic field forms loops surrounded by shells which may exhibit enhanced UV emission."
    },
    "0711/0711.4077_arXiv.txt": {
        "abstract": "{Central galaxies in galaxy clusters may be key discriminants in the competition between the cold dark matter (CDM) paradigm and modified Newtonian dynamics (MOND). } {We investigate the dark  halo of NGC\\,1399, the central galaxy of the Fornax cluster, out to a galactocentric distance of 80\\,kpc.} {The data base consists of  656 radial velocities of globular clusters obtained with MXU/VLT and GMOS/Gemini, which is the largest sample so far for any  galaxy. We performed a Jeans analysis for a non-rotating isotropic model.}{An NFW halo with the parameters $r_s = 50\\,\\rm{kpc}$ and $\\varrho_s = 0.0065 M_\\odot/\\rm{pc}^3$ provides a good description of our data, fitting well to the X-ray mass.  More massive halos are also permitted that agree with the mass of the Fornax cluster as derived from galaxy velocities. We compare this halo with the expected MOND models under isotropy and find that additional dark matter on the order of the stellar mass is needed to get agreement. A fully radial infinite globular cluster system would be needed to change this conclusion. } {Regarding CDM, we cannot draw firm conclusions. To really constrain a cluster wide halo, more data covering a larger radius are necessary. The MOND result appears as a small-scale variant of the finding that MOND in galaxy clusters still needs dark matter.} ",
        "introduction": "The dark matter halos of early-type galaxies are  investigated less well than those of spiral galaxies (for a recent review, see Romanowsky \\cite{romanowsky06}). Globular clusters (GCs) and planetary nebulae are  the main stellar dynamical tracers, while X-rays can be used at even larger galactocentric distances for a gas dynamical approach. The reliability of a deduced mass profile depends strongly on the number of dynamical probes available. Regarding GCs, it is therefore natural that the rich GC systems of bright elliptical galaxies have so far been  the preferred targets for such studies, for example M87 (C{\\^o}t{\\'e} et al.~\\cite{cote01}), NGC\\,4472 (C{\\^o}t{\\'e} et al.~\\cite{cote03}), and NGC\\,4636 (Schuberth et al.~\\cite{schuberth06}). The largest sample until today has been measured for NGC\\,1399, the central galaxy of the Fornax galaxy cluster. Richtler et al.~(\\cite{richtler04}, Paper\\,I) analyzed the velocities of about 450 GCs out to a galactocentric distance of 40\\,kpc and found an approximately constant velocity dispersion. The corresponding dark matter halo was found to be consistent with an NFW-halo (see section 5 for the numerical values). The relatively poor knowledge that we have until now regarding the dark halos of early-type galaxies, does not indicate a common behavior such as  the constant rotation curves found in spiral galaxies. While the inferred circular velocities of NGC\\,1399 or NGC\\,4636 appear to be flat, they seem to rise for M87, and most strikingly so for NGC\\,6166 (Kelson et al.~\\cite{kelson02}), the central galaxy in Abell 2199. In elliptical galaxies of lower luminosity, the projected velocity dispersion can even be falling (Romanowsky et al.~\\cite{romanowsky03}), perhaps including the circular velocity. One may speculate that these relate to the environment. The dark halos of central galaxies might be determined by the dark halos of their host galaxy clusters rather than being halos belonging to  individual galaxies. However, a larger sample of well-studied galaxies is necessary to arrive at firm conclusions.\\\\ We have now augmented the number of GC velocities around NGC\\,1399 to nearly 660 reaching a galactocentric distance of 80\\,kpc, which is an unprecedently large and extended sample. In this \\emph{Letter}, we present the implications of this sample for the dark halo, updating the conclusions of Paper\\,I. Moreover, we discuss the dark halo in the context of Modified Newtonian Dynamics (MOND) (e.g. Milgrom \\cite{milgrom83}, Sanders \\& McGaugh \\cite{sanders02}). The detailed description of the observations, the reduction, and the analysis will be presented in a forthcoming paper (Schuberth et al., in prep), taking into account the population substructure of the GC system as well as more anisotropic models.\\\\ In the following, we adopt a distance of 19\\,Mpc for NGC\\,1399. ",
        "conclusions": "We have used our new sample of about 660 velocities of GCs to constrain the dark halo of NGC\\,1399 within 80 kpc. The resulting halo, assumed to be of the NFW type, is more concentrated than the halos from cosmological simulations, but marginally in agreement. A more massive alternative halo, which would account for the total mass of the Fornax cluster as estimated by Drinkwater et al.~(\\cite{drinkwater01}), may also be consistent with our data. We are thus not able to conclusively constrain  the cluster-wide halo from probes within 80 kpc. We also consider the MOND case and find that MOND is not able to reproduce the observed GC kinematics under isotropy without invoking an additional dark matter component, as already found for other galaxy clusters. A realistic radial anisotropy does not change this conclusion. Moreover, only the isotropic model reproduces the X-ray mass well. The hypothetical dark halo must have a small core, and thus neutrinos do not seem to be  viable candidates. It may be in gaseous form.  Under our assumptions (particularly the interpolation between the MONDian and the Newtonian regime), MOND makes a strong prediction, which perhaps will be falsifiable in the future."
    },
    "0711/0711.1187_arXiv.txt": {
        "abstract": "{We re-investigate UZ Libr\\ae~spectra obtained at KPNO in 1998 and 2000. From the 1998 data we compose 11 consecutive Doppler images using the Ca\\,{\\sc i}-6439, Fe\\,{\\sc i}-6393  and Fe\\,{\\sc i}-6411 lines. Applying the method of average cross-correlation of contiguous Doppler images we find anti-solar differential rotation with a surface shear of $\\alpha\\approx-0.03$. The pilot application of the local correlation tracking technique for the same data qualitatively confirms this result and indicates complex flow pattern on the stellar surface. From the cross-correlation of the two available Doppler images in 2000 we also get anti-solar differential rotation but with a much weaker shear of $\\alpha\\approx-0.004$.} ",
        "introduction": "In our time, the Sun is the only star for which the surface differential rotation and other surface flows can be measured directly. The detection of differential rotation on active stars, however, is a challenging observational task. Starspots are regarded as manifestations of solar-like active regions and as such, are suitable for tracking and recovering surface flows. In this paper we make one of the first attempts at recovering surface flows on spotted stars (see also the paper by K\\H{o}v\\'ari et al. \\cite{sgemposter}).  Our target is the RS\\,CVn-type active giant star UZ\\,Libr\\ae. The star has a very fast rotation rate of $P_\\mathrm{rot}=4.76$ days, which is equivalent to the orbital period. Doppler images of the star were presented by Strassmeier (\\cite{1996A&A...314..558S}) and Ol\\'ah, Strassmeier \\& Weber (\\cite{2002A&A...389..202O}). In the latter paper the signature of a weak anti-solar differential rotation (DR) was found but the authors regarded the detection inconclusive. The possible DR was also examined by Ol\\'ah Jurcsik and Strassmeier (\\cite{2003A&A...410..685O}) using photometric observations. They found two permanent, distinct periods in the Fourier-spectrum of the long-term data, which were assigned to the different rotation rate of spots at high and low latitudes, resulting in $\\alpha=\\Delta\\Omega/\\Omega$ of $-0.0026$. Adopted stellar parameters of UZ\\,Lib are summarised in Table \\ref{tab:params}, the values which are not marked were taken from Ol\\'ah, Strassmeier \\& Weber (\\cite{2002A&A...389..202O}).    \\begin{table}[] \\caption{Stellar parameters of UZ Libr\\ae}\\label{tab:params} \\leavevmode \\footnotesize \\begin{tabular}{ll} \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} spectral type$~^1$ &K0 III \\\\ $P_{rot}=P_{orb}~^{2,3}$ & 4.768241 days \\\\ $v\\sin i~^{2,3}$ & 67 $\\pm$ 1 km\\,s$^{-1}$\\\\ $i~^4$ & $50^\\circ\\pm$10\\degr\\\\ T$_\\mathrm{eff}~^3$ & 4800 K\\\\ $\\log g~^2$ & 2.5\\\\ Microturbulence $\\xi$& 0.8 km s$^{-1}$\\\\ Macroturbulence $\\zeta_\\mathrm{R}=\\zeta_\\mathrm{T} $ & 4 km s$^{-1}$\\\\ Chemical abundance & solar\\\\ \\noalign{\\smallskip} \\hline\\\\[-5pt] \\end{tabular}  \\noindent$^1$ Bopp \\& Stencel (\\cite{1981ApJ...247L.131B})\\\\ \\noindent$^2$ Strassmeier (\\cite{1996A&A...314..558S})\\\\ \\noindent$^3$ Fekel et al. (\\cite{1999A&AS..137..369F}) \\end{table} ",
        "conclusions": "\\begin{itemize} \\item Applying our cross-correlation technique ACCORD for the time-series Dopper images of UZ\\,Lib in 1998 and in 2000 we derived weak surface DR of anti-solar type, in agreement with former DR detections. However, no sign of coherent meridional flow was found. \\item LCT technique is applied first time for stellar data (see also the paper by K\\H{o}v\\'ari et al. \\cite{sgemposter}) and delineated a complex surface flow network for UZ\\,Lib. \\item Averaging the zonal components of the LCT flow map along latitude resulted in an uncertain, weak anti-solar DR. Meridional flow components show no evidence for any coherent pattern, similarly to the ACCORD non-detection.   \\end{itemize}  \\begin{figure} \\includegraphics[angle=-0,width=1.00\\columnwidth,height=0.3\\columnwidth]{vida_krisztian_fig5.eps}  \\hspace*{-0.22cm} \\includegraphics[angle=-0,width=0.88\\columnwidth,height=0.3\\columnwidth]{vida_krisztian_fig6.ps} \\caption{Best correlated latitudinal shifts from ACCORD (top) and meridional flow components from LCT (bottom). No sign of coherent meridional flow can be seen.} \\label{fig:merflows} \\end{figure}"
    },
    "0711/0711.4288_arXiv.txt": {
        "abstract": "We revisit the paradigm of unified dark energy discussing in detail the averaging problem in this type of scenarios, highlighting the need for a full non-linear treatment. We also address the question of if and how models with one or several dark fluids can be observationally distinguished. Simpler and physically clearer derivations of some key results, most notably on the relation between the generalized Chaplygin gas and the standard ($\\Lambda$CDM) `concordance' model and on a Jeans-type small-scale instability of some coupled dark energy/dark matter models are presented. ",
        "introduction": "Introduction}  Observations have been steadily piling over the years that can only be `explained', in the context of General Relativity, by the presence of so-called \\emph{dark} forms of energy: ordinary baryonic matter and radiation are simply not enough \\cite{Riess,Perlmutter,Percival,Dodelson,Spergel}. In broad terms, a dark energy component (violating the strong energy condition) is required to explain the recent acceleration of the Universe and a (cold) dark matter component to account for the amount of structure observed.  Unfortunately, these components have never been observed directly. In fact, they may not \\emph{exist} at all. It is certainly possible to engineer dark energy and dark matter out of the picture, by modifying General Relativity directly; often this process involves extra-dimensions with non-trivial topologies and/or fiddling with TeVeS gravity sources \\cite{Teves} (see also \\cite{Freese,Alternatives}). At the moment, however, most modifications seem far too \\emph{ad hoc} to be particularly pleasing. In this paper, we will mostly work within the confines of General Relativity, i.e., we will assume that dark energy and dark matter are real entities. On the other hand, if this assumption is correct, it is truly mind boggling to find out that roughly $96\\% $ of the Universe should be in this dark form. Cosmology is thus hard-pressed to explain what these components are.  Certainly the realization that dark energy and dark matter may not be \\emph{independent} entities should play a significant role. Although we frequently perceive them as being different, this does not have to be the case. The Chaplygin gas \\cite{Chap1,Chap2}, for instance, experienced a short burst of popularity precisely for  being able to mimic both dark energy and dark matter (depending on the \\emph{local} density), and still be just one form of energy---an exotic one, granted, but the same can be said of quintessence, $k$-essence, and the rest of the phenomenological zoo. Thus, we should be open to the possibility that dark energy and dark matter are actually just  different manifestations of the same underlying field, not different entities \\emph{per se}.  On the theoretical front, the idea of a unified description of dark energy and dark matter has incredible heuristic potential and should be properly investigated. In case of failure, we would at least come back justified in treating dark energy and dark matter as an independent pair. Either way, something of value could be gained.  While the background and linear properties of unified dark energy models have been extensively studied by the community, the non-linear aspects of the problem have been largely ignored. In the following section, we will discuss in detail the averaging problem in the context of these models and show that non-linearities are a crucial aspect  that cannot be neglected. ",
        "conclusions": ""
    },
    "0711/0711.4131_arXiv.txt": {
        "abstract": "We present \\cha\\ and \\hst\\ observations of the ultraluminous X-ray source (ULX) IC 342 X-1. The \\cha\\ and \\hst\\ images are aligned using two  X-ray emitting foreground stars. The astrometry corrected position for X-1 is $\\rm R.A. = 03^h45^m55\\fs61$, $\\rm Decl. = +68\\arcdeg04\\arcmin55\\farcs3$ (J2000) with an error circle of $0.2\\arcsec$.  One extended optical source is found in the error circle, which could be the optical counterpart of X-1.  The source shows an extended feature in \\hst\\ images at long wavelengths, which is likely to be a superposition of two point sources, although it is possible that the dimmer one could be a jet.  Both sources are much redder than typical for ULX optical counterparts.  The brighter one has an absolute magnitude $M_V=-5.2\\pm0.2$ and $(B-V)_0 = 0.66 \\pm 0.13$ and the dimmer star is not detected in $B$ and has $(B-V)_0 > 2.1$.  Their colors are consistent with an F8 to G0 Ib supergiant or a carbon star, respectively.  However, it is likely that part or most of the optical emission may be due to X-rays reprocessed by the companion star or the accretion disk.  The stellar neighborhood of IC 342 X-1 lacks O stars and has a minimum age of $\\sim 10$~Myr.  This excludes the possibility that the surrounding nebula is powered by an energetic explosion of a single massive star that formed a black hole. We suggest that the nebula is most likely powered by an outflow from the X-ray source. ",
        "introduction": "Ultraluminous X-ray sources, the brightest nonnuclear X-ray point sources in external galaxies, still hide their nature from us.  Their luminosity of $3\\times10^{39}$ to $10^{41}$~\\ergs, inferred assuming isotropic emission, is too high to be powered by accretion onto stellar mass black holes that are formed in the collapse of normal stars with normal metallicity, and thus may suggest the presence of intermediate mass black holes \\citep[IMBHs;][]{col99,kaa01}.  However, if the emission is beamed \\citep{kin01,kor02} or truly super-Eddington \\citep{wat01,beg02}, IMBHs are not required.  Multiwavelength observations have provided powerful means to determine the nature of ULXs.  \\citet{pak02} carried out an optical survey of bright ULXs and discovered that several ULXs are associated with optical nebulae.  These nebulae have very large spatial scales of a few hundred parsecs, and present line emission. In particular, He~{\\sc ii} $\\lambda$4686 emission line was detected in the nebula surrounding Holmberg II X-1, and was used to determine that the flux of reprocessed X-rays from the central ULX was at least $4\\times10^{39}$~\\ergs\\ \\citep{kaa04}.  Identification of an optical counterpart of a ULX has been accomplished in several cases. However, it is hard to determine whether or not the optical light from the companion star is contaminated by the outer region of the accretion disk, thus stellar classification based on optical colors may not be reliable.  The far-ultraviolet spectrum of NGC 5204 X-1 suggests a B0 Ib companion, but contains some peculiar emission lines \\citep{liu04}.  Identification of a giant branch companion has been made for the ULX X41.4+60 in M82 based on the detection of a 62~day X-ray modulation interpreted as the orbital period \\citep{kaa06a,kaa06b,kaa07}.  Identification of the companion stars for more ULXs would enable us to place constraints on the binary evolution of these systems.  IC 432 is a nearly face-on starburst spiral galaxy.  IC~342 X-1 was discovered with {\\it Einstein} \\citep{fab87} and confirmed as a point source with {\\it ROSAT} observations \\citep{bre93}.  Using \\asca\\ observations in 1993, \\citet{oka98} found significant variability on time scales of hours from X-1, suggesting it was a compact object.  The source changed its spectrum dramatically in a 2000 \\asca\\ observation \\citep{kub01}, similar to the spectral state transitions found in Galactic black hole binaries.  A possible X-ray periodicity of 31 or 41 hours from X-1 was also detected in the same \\asca\\ observation in 2000 \\citep{sug01}.  X-1 is associated with a ``tooth''-shaped optical nebula \\citep{pak02} with emission lines similar to a supernova remnant \\citep[SNR;][]{pak02,rob03,abo07}.  \\citet{rob03} estimated that the initial energy required to power the nebula was at least $2\\times10^{51}$~ergs, greater than a canonical SNR.  Also, \\citet{rob03} found two regions of [O~{\\sc iii}] emission located oppositely on the shell of the nebula, which could be X-ray photoionized and thus indicate anisotropic emission from X-1. Instead, \\citet{gri06} found that the [O~{\\sc iii}] emission closely follows the H$\\alpha$ emission and is smoothly distributed in the nebula, implying it is shock excited. \\citet{abo07} found evidence for photoionization in the central regions of the nebula.  To search for an optical counterpart of IC~342 X-1, we obtained \\cha\\ and \\hst\\ observations. The observations and data are described in \\S~\\ref{sec:obs}.  The \\cha\\ and \\hst\\ images are aligned by matching two foreground stars in both images (\\S~\\ref{sec:astr}), which leads to the identification of a unique optical counterpart of X-1 (\\S~\\ref{sec:cp}). The photometry for the optical counterpart of X-1 and its nearby starforming region is performed in \\S~\\ref{sec:photo}.  A multiwavelength spectrum from optical to X-ray is presented in \\S~\\ref{sec:multi}. An emission line image obtained using \\hst\\ of the tooth nebula is shown in \\S~\\ref{sec:halpha}. All results are discussed in \\S~\\ref{sec:diss}.  We adopt a distance to IC~342 of $3.3 \\pm 0.3$~Mpc, which is estimated from observations of 20 Cepheids in the galaxy \\citep{sah02}. ",
        "conclusions": "\\label{sec:diss}  The stellar population within 300~pc ($20\\arcsec$) of X-1 appears to contain no O stars.  The color-magnitude diagram (Figure~\\ref{fig:cmd}), suggests that the minimum stellar age is $\\sim$10~Myr if the metallicity is close to $Z=0.019$.  Comparing to isochrones at a lower metallicity of $Z=0.004$, the minimum stellar age is about 20~Myr. The region appears to contain relatively young stars with ages of tens of Myr and also an older population with ages of several Gyr.  IC 342 X-1 appears at the center of the nebula (Figure~\\ref{fig:ha}), which strongly suggests an association. The relatively old age of the stellar population in the vicinity of the nebula makes it unlikely that the nebula is powered (even in part) by ionization from one or more O stars.  \\citet{abo07} suggested that the shock velocity in the nebula is 20 to 100 km/s, implying the nebula is younger than 2~Myr.  This age is an overestimate since the expansion would be slowing down.  This suggests that if the nebula was created in a supernova event, the progenitor star must has a lifetime close to 10~Myr.  The high luminosity of X-1 suggests it contains a black hole.  To form a black hole in an energetic explosion, the progenitor must be very massive \\citep[e.g., at least 40~$M_\\sun$ in][]{heg03} and thus has a lifetime significantly shorter than 10~Myr.  Therefore, it is unlikely that the nebula was powered by an unusually powerful explosion in which a black hole was formed.  A more plausible alternative is that the nebula is powered by an outflow from the accreting compact object in the ULX \\citep{pak02,rob03}.  The nebula appears to be predominantly shock powered with a shock velocity in the range 20--100~km/s and a total power of $3 \\times 10^{39}$~\\ergs, but with high excitation lines possibly powered by photoionization near the center \\citep{abo07}.  This power is 20\\% to 50\\% of X-ray luminosity, depending on the X-ray state \\citep{oka98}, and is reasonable in that X-ray binaries appear to have power in their outflows which is comparable to or may even exceed their X-ray emission \\citep{gal05}.  The emission line morphology in Figure~\\ref{fig:ha} indicates that the nebula is not a complete shell.  The morphology is somewhat similar to that of the W50/SS433 system \\citep{bou07} and may suggest a jet orientation along an axis slightly west of north.  Given the strong radio emission from W50/SS433, radio observations of IC 342 X-1 would be of great interest.  We find one optical source in the error circle of IC 342 X-1.  The source appears extended at long wavelengths and is either two stars or a star with an associated feature such as a jet.  We consider the relativistic jet interpretation unlikely because their spectra are quite steep, unlike the optical relativistic jets of black hole such as those of M87 \\citep{per01}.  The projected separation in IC 342 would be 2.6~pc, and the true separation may be larger.  This is too large to be an optical compact jet found in the hard state of stellar mass black holes, which usually has a spatial scale of AU \\citep{fen06}.  We find interpretation in terms of two superimposed stars more likely.  The brighter star (star A) lies well within the error circle for IC 342 X-1.  The dimmer star (star B) lies at the edge of the 90\\% confidence error circle of X-1 and could, instead, be the counterpart of X-1. Allowing either star as the counterpart, the X-ray to optical flux ratio indicates that the source is not a background blazar or quasar, but is consistent with an X-ray binary in IC 342.  If the companion star dominates the optical emission despite the high X-ray flux, then the spectral types of the stars can be determined.  The colors and absolute magnitude of star A are consistent with classification as an F8 to G0 Ib supergiant, which has a mass of about 10~$M_\\sun$.  There is some excess $I$ band emission, but this is often seen in supergiants and is a signature of a circumstellar disk or a stellar wind.  With this classification, the age of star A would be roughly 50~Myr.  The evolutionary status of star A, on the supergiant branch, would be compatible with the large mass transfer rates needed to power the inferred luminosity of the ULX.  The colors and absolute magnitude of star B are consistent with classification as a carbon star.  The progenitor mass of such a star is likely of intermediate mass, $1-5 M_{\\odot}$, and, thus, would require a rather old binary system.  The carbon star phase is short-lived, $\\lesssim 1$~Myr.  High luminosity carbon stars have high mass loss rates \\citep{jur91} which would be sufficient to power the ULX.  However, the optical emission could also be dominated by reprocessed X-rays on the companion star or on the accretion disk, or by the intrinsic emission of the disk.  In these cases, spectral classification of the companion star via its colors is not reliable.  We discuss these possibilities in following.  Assuming the X-ray emission is isotropic, the X-ray luminosity intercepted by the companion is \\begin{equation} L_{\\rm c} \\approx \\frac{\\pi R_2^2}{4\\pi a^2}L_{\\rm X} \\; , \\end{equation} where $R_2$ is the radius of the companion star, $a$ is the binary separation, and $L_{\\rm X}$ is the total X-ray luminosity.  The high luminosity of ULXs indicate that they accrete via Roche-lobe overflow.  Substituting equation~(2) in \\citet{egg83} in which $R_2/a$ is determined by the companion to compact object mass ratio $q$ to an accuracy better than 1\\% for any $q$, we have \\begin{equation} L_{\\rm c} = \\frac{L_{\\rm X}}{4} \\left[ \\frac{0.49q^{2/3}}{0.6q^{2/3} + \\ln(1+q^{1/3})} \\right]^2 \\; . \\end{equation} A plot of $L_{\\rm c}$ versus $q$ at different X-ray luminosities is presented in Figure~\\ref{fig:lc}. We can see that for a source like IC~342 X-1, which has $L_{\\rm X} \\approx 10^{40}$~\\ergs, the companion intercepted X-ray luminosity is above $10^3$~$L_\\sun$ when $q>0.001$ and above $10^4$~$L_\\sun$ when $q>0.01$.  These luminosities are as high as of massive main sequence stars, bright giants or supergiants, implying that the optical counterpart of a bright ULX will appear as bright as a massive star or giant/supergiant no matter what the companion really is.  The observed optical spectra suggest star A has a luminosity of about $10^4$~$L_\\sun$, and star B has a luminosity of about $10^3$ to $10^4$~$L_\\sun$. Thus, the optical emission may well be dominated by an X-ray ``heated'' companion star.  However, how much of the companion surface is shielded by the accretion disk is unknown.  For ULXs with $L_{\\rm X}>10^{40}$~\\ergs, if even a small portion ($>10\\%$) of the X-ray flux is not absorbed enroute to the companion surface, then the companion's temperature and luminosity will be notably raised.  \\begin{figure}[h] \\centering \\includegraphics[width=\\columnwidth]{f8.eps} \\caption{ Companion star intercepted luminosity ($L_{\\rm c}$ in unit of solar luminosity) as a function of the binary mass ratio ($q$, the companion to compact object mass ratio), assuming isotropic emission and Roche-lobe overflowing. Different lines are scaled with different X-ray luminosity as labeled. \\label{fig:lc}} \\end{figure}  Disk irradiation may be an important component of the optical emission if the X-ray luminosity is truly very high.  Gris\\'e et al.\\ (2007, in preparation) found significant variability from the optical counterpart of NGC~1313 X-2 on time scales from days to hours and even minutes, suggesting the optical emission is dominated by X-ray reprocessing.  Both star A and star B show no detectable variability in the four F625W images at time scales of a few hours. The X-ray to optical flux ratio taking either star A or B as the counterpart is consistent with those found in neutron-star and stellar mass black hole X-ray binaries (the low mass X-ray binaries; LMXBs), where the emission is thought to be dominated by reprocessed X-rays on the disk \\citep{van95,liu01}.  However, the color of either star is significantly redder than the colors, $(B-V)_0 \\sim 0$, typical of LMXBs.  Assuming a standard disk geometry, the disk temperature from X-ray irradiation varies as $T(R)^4 \\propto L_{\\rm X}/R^2$ at different radius, and the spectral energy distribution has a form of $S_\\nu \\propto \\nu^{-1}$ by integrating to infinite radius.  Therefore, for a finite disk, the X-ray reprocessing at the outermost radius ($R_{\\rm out}$) determines the peak of the energy distribution, which could be attributed to a simplified blackbody emission with an effective temperature $T_{\\rm eff}$.  We assume the disk size is scaled with the binary separation $a$, so the effective temperature of the heated disk varies as  $T_{\\rm eff}^4 \\propto T(R_{\\rm out})^4 \\propto L_{\\rm X}/a^2$.  If the system transfers mass via Roche-lobe overflow, then $a \\propto M^{1/3} P^{2/3}$, where $M$ is the compact object mass, $P$ is the orbital period and we have assumed that the companion mass is less than that of the compact object, also the companion density varies as $\\rho \\propto P^{-2}$.  Thus, we find $T_{\\rm eff} \\propto L_{\\rm X}^{1/4} M^{-1/6} \\rho^{1/6}$.  For most Galactic LMXBs, the companion is a main sequence star with a density in order of $1 \\rm \\, g \\, cm^{-3}$ and the $T_{\\rm eff}$ of the heated disk is around 25,000--30,000~K, or $(B-V)_0 \\sim 0$.  To produce a heated disk with a $T_{\\rm eff}$ around 5,500--6,000~K as star A at a high X-ray luminosity level, it requires either an IMBH or an well evolved companion that has a low density.  If the optical counterpart is star B, whose optical spectral index is much steeper than $-1$, it is unlikely that the optical emission is dominated by disk irradiation.  Parenthetically, we note that the relation found by \\citet{van94} between the visual luminosity $L_V$, X-ray luminosity, and orbital period of LMXBs has been used to make inferences about the orbital periods of ULXs.  This relation was based mainly on sources with relatively dense companions, and the correlation is inferred based on the assumption that $L_V \\propto T^2$, which is no longer correct when $T<10,000$~K.  Therefore, one should be cautious in applying this correlation in systems where the observed $T_{\\rm eff}$ is lower than 25,000--30,000~K, or $(B-V)_0$ is higher than 0.  The intrinsic disk emission in the optical band should present in a form of $S_\\nu \\propto \\nu^{1/3}$, which is totally different with the observed optical spectra.  Therefore, we suggest that the optical emission is not dominated by the viscous disk.  In sum, for a ULX system with $L_{\\rm X}>10^{40}$~\\ergs, it appears unlikely that the intrinsic emission of the companion star dominates the optical emission. The companion star is likely significantly heated by X-ray illumination.  Disk irradiation could also be important. The color of the optical counterpart is redder than typically found in systems, such as LMXBs, where a significant fraction of the optical light arises from an X-ray heated disk. The red color could be a consequence of a massive black hole or a low density companion."
    },
    "0711/0711.4307_arXiv.txt": {
        "abstract": "We consider the generation of large-scale magnetic fields in slow-roll inflation. The inflaton field is described in a supergravity framework where the conformal invariance of the electromagnetic field is generically and naturally broken. For each class of inflationary scenarios, we determine the functional dependence of the gauge coupling that is consistent with the observations on the magnetic field strength at various astrophysical scales and, at the same time, avoid a back-reaction problem. Then, we study whether the required coupling functions can naturally emerge in well-motivated, possibly string-inspired, models. We argue that this is non trivial and can be realized only for a restricted class of scenarios. This includes power-law inflation where the inflaton field is interpreted as a modulus. However, this scenario seems to be consistent only if the energy scale of inflation is low and the reheating stage prolonged. Another reasonable possibility appears to be small field models since no back-reaction problem is present in this case but, unfortunately, the corresponding scenario cannot be justified in a stringy framework. Finally, large field models do not lead to sensible model building. ",
        "introduction": "\\label{sec:introduction}  Magnetic fields are present on various scales in the Universe from planets, stars, galaxies, to clusters of galaxies, see {\\it e.g.}, Refs.~\\cite{Kronberg:1993vk, Grasso:2000wj, Widrow:2002ud} for reviews. In galaxies of all types, magnetic fields with the field strength $\\sim 10^{-6}$ G, ordered on $1-10$ kpc scale, have been detected \\cite{Widrow:2002ud,Sofue:1986wc}. There is also some evidence that they exist in galaxies at cosmological distances~\\cite{Kronberg:1992}.  Furthermore, in recent years, magnetic fields in clusters of galaxies have been observed by means of the Faraday rotation measurements (RMs) of polarized electromagnetic radiation passing through an ionized medium~\\cite{Kim:1990,Kim:1991,Clarke:2001}. Unfortunately, however, RMs give only the product of the field strength along the line of sight and the coherence scale, and we cannot find them independently.  In this situation, the strength and the scale are estimated on $10^{-7}-10^{-6}$ G and 10 kpc$-$1 Mpc, respectively.  It is interesting that magnetic fields in clusters of galaxies are as strong as galactic ones and that the coherence scale may be as large as $\\sim$ Mpc.  \\par  Since the conductivity of the Universe through most of its history is large, the magnetic field $B$ evolves conserving magnetic flux as $B \\propto a^{-2}$, where $a(t)$ is the scale factor. On the other hand, the average cosmic energy density $\\bar{\\rho}$ evolves as $\\bar{\\rho} \\propto a^{-3}$. Hence, the scaling of the magnetic field is given by $B \\propto {\\bar{\\rho}}^{2/3}$ and this relation should also be understood as giving the value of $B$ at different spatial locations characterized by different energy densities. The present ratio of interstellar medium density in galaxies ${\\rho}_\\mathrm{gal}$ to $\\bar{\\rho}$ and that of inter-cluster medium density in clusters of galaxies ${\\rho}_\\mathrm{cl}$ are ${\\rho}_\\mathrm{gal}/\\bar{\\rho} \\simeq 10^5-10^6$ and ${\\rho}_\\mathrm{cl}/\\bar{\\rho} \\simeq 10^2-10^3$, respectively.  Consequently, from these relations, we see that, using the law mentioned before, namely $\\bar{B}=B_{\\mathrm{gal, cl}} (\\rho _{\\mathrm{gal,cl}}/\\bar{\\rho})^{-2/3}$, the required strength of the cosmic magnetic field at the structure formation, adiabatically rescaled to the present time, is $10^{-10}-10^{-9} \\mathrm{G}$ in order to explain the observed fields mentioned before. On the other hand, in general, if the galactic dynamo mechanism is at play, then seed fields with a present strength of only $10^{-22}-10^{-16} \\mathrm{G}$ is required~\\cite{ParkerEN:1970,ParkerEN:1979,Zelodovitch:1983}.  \\par  The question addressed in this article is that of the origin of these magnetic fields and whether there exists simple, physically well-motivated mechanisms, which could provide the target values of $B$ just described above. Generation mechanisms of the seed magnetic fields fall into two broad categories. One is astrophysical processes and the other is cosmological physical processes in the early Universe. The former exploits the difference in mobility between electrons and ions. This difference can lead to electric currents and hence magnetic fields. If the scale of cluster magnetic fields is as large as $\\sim$ Mpc, however, it is unlikely that the origin of such a large-scale magnetic field is in astrophysical processes and we are tempted to assign the origin of cosmological magnetic field to the processes during inflation in the early Universe which stretches coherent scales exponentially~\\cite{Guth:1980zm,Sato:1980yn,Linde:2005ht,Olive:1989nu, Kolb:1990aa,Lyth:1998xn, Martin:2003bt,Martin:2004um,Martin:2007bw}.  \\par  During inflation quantum scalar (density) perturbations~\\cite{Mukhanov:1981xt,Hawking:1982cz, Starobinsky:1982ee, Guth:1982ec} and tensor (gravitational waves) perturbations~\\cite{Grishchuk:1974ny,Rubakov:1982df} are generated and stretched to cosmological scales by subsequent accelerated expansion. Usually, however, large scale fluctuations in electromagnetic fields are not generated, because they are conformally invariant~\\cite{Parker:1968}. Hence, in order to generate cosmological magnetic fields during inflation, one must break the conformal invariance of the theory in an appropriate manner.  A number of models have been proposed in this context so far~\\cite{Turner:1987bw,Ratra:1991av,Ratra:1991bn,Lemoine:1995vj, Giovannini:2001nh,Giovannini:2001xk,Giovannini:2002ki,Calzetta:1997ku, Kandus:1999st,Davis:2000zp,Dolgov:1993vg,Bamba:2003av,Bamba:2004cu}.  \\par  Here we propose a simple model making use of the arbitrariness of the gauge kinetic function in supergravity~\\cite{Bailin:1994}. Indeed, the general supergravity Lagrangian including vector super-multiplets contains the following term \\begin{equation} \\int {\\rm d}^2\\theta \\left[f_{ab}(\\Phi)W_a{}^{\\alpha }W_{\\alpha b} +\\mbox{h.c.}\\right]\\, , \\end{equation} where $W_a{}^{\\alpha }$ is the gauge field strength superfield with spinor index $\\alpha $ and gauge group index $a$. As already mentioned, the quantity $f_{ab}(\\Phi)$ is an arbitrary function of the chiral superfield $\\Phi$~\\cite{Cremmer:1982en}. This implies that the conformal invariance is naturally and generically broken as is most conveniently shown when the above term is expressed in terms of the superfields components \\begin{equation} \\frac14 \\Re f_{ab}(\\phi)\\left(F_{a}\\right)_{\\mu \\nu} \\left(F_{b}\\right)^{\\mu \\nu}\\, . \\end{equation} Since we are interested in the production of magnetic fields, we consider a $U(1)$ gauge field and, therefore, the gauge field strength do no longer carry a gauge index. As a consequence, there is also only one function $f(\\phi)$ left. Moreover, in order to consider the minimalist case with the least possible assumptions, we naturally identify the scalar part $\\phi $ of the chiral superfield $\\Phi$ to the inflaton field.  \\par  A related study was carried out in Refs.~\\cite{Ratra:1991av,Ratra:1991bn} but only in the particular case where $f(\\phi )\\propto {\\rm e}^{\\phi}$. Here, we study the most general case (at least, in the framework of single field inflation) and determine which form $f(\\phi )$ should take in order to obtain a power spectrum of the magnetic field in agreement with the currently available observational constraints on large scales. Then, for each category of single field inflationary scenarios, we explore whether there is a natural model in supergravity (possibly string-inspired) that could reproduce the required shape of the gauge coupling and, at the same time, leads to the correct form of the inflaton potential.  \\par  This paper is organized as follows. In Sec.~\\ref{sec:reverse}, we describe our model action and derive the equations of motion from it. Then, the gauge field is quantized in the Coulomb gauge and the vacuum energy density of the magnetic field at the end of inflation is determined. The evolution of the magnetic energy density through the subsequent stages of evolution, preheating, radiation and matter dominated eras is also computed. In a next step, for each single field slow-roll inflationary models, using a simple ansatz depending only on one free parameter, we calculate the required shape of the gauge coupling function $f(\\phi )$. In Sec.~\\ref{sec:couplingconstraint}, we use the currently available limits on the magnetic field to constraint the free parameter mentioned above, that is to say to constraint the shape of the gauge coupling function. We also investigate the consistency of the model and determine under which conditions this one does not suffer from a back-reaction problem. In Sec.~\\ref{sec:particlemodels}, we try to embed in particle physics the models that have been shown to be phenomenologically viable. Finally, in Sec.~\\ref{sec:conclusion}, we present our conclusions. Throughout this article, we use units in which $ c =\\hbar = M_G=k_\\mathrm{B} =1$. We also put the reduced Planck mass, $M_G\\equiv 1/\\sqrt{8\\pi G}$, to unity when it simplifies the notation. In this case, the full Planck scale takes $\\mpl \\equiv G^{-1/2}=\\sqrt{8\\pi}$. ",
        "conclusions": "\\label{sec:conclusion}  In this section, we recap our main results and discuss issues that should be studied in future works. In the present paper, we have studied the generation of large-scale magnetic fields in inflationary cosmology, breaking the conformal invariance of the electromagnetic field by introducing a coupling with the inflaton field as it is generic in a supergravity framework. We have determined the form of the coupling necessary in order to produce large-scale magnetic fields with the required strengths on the relevant astrophysical scales. We have shown that the scale-dependent magnetic energy density possess two branches. Among these two branches, one ($\\alpha \\sim -3$) seems to lead to sensible model building in the case of power-law inflation but only at the expense of having a low energy inflation scale and a long reheating stage. The other branch ($\\alpha \\sim 2$) is useful in the case of small field models and this case appears to be relevant since this does not require to fine-tune the reheating epoch. However, the model building condition $q\\gtrsim 57$ cannot be justified in a string inspired model and seems artificial. Finally, in the context of large field models, we have shown that the required coupling is quite difficult to justify.  \\par  Determining exactly all the consistent models in the framework envisaged here is a non-trivial issue because the parameter space is large (at least four-dimensional), in particular when a general reheating stage is considered. Here, we have just proven that consistent models exist. However, it would certainly be interesting to systematically explore the parameter space in order to have a more accurate idea of whether the consistent models are just peculiar or, on the contrary, quite generic. Another interesting avenue for the future is clearly the model building issue. Here, we have noticed that the branch $\\alpha \\sim 2$ can be used only in the context of small field models. It would be interesting to find other models where this can also been done. We hope that these issues will be addressed in the near future."
    },
    "0711/0711.0238_arXiv.txt": {
        "abstract": "We report on the first results of a new study aimed at understanding the diversity and evolutionary history of distant galactic bulges in the context of now well-established trends for pure spheroidal galaxies.  To this end, bulges have been isolated for a sample of 137 spiral galaxies within the redshift range 0.1\\,$<$\\,z$\\,<$\\,1.2 in the GOODS fields. Using proven photometric techniques we determine the characteristic parameters (size, surface brightness, profile shape) of both the disk and bulge components in our sample. In agreement with earlier work which utilized aperture colors, distant bulges show a broader range of optical colors than would be the case for passively-evolving populations.  To quantify the amount of recent star formation necessary to explain this result, we used the DEIMOS spectrograph to secure stellar velocity dispersions for a sizeable fraction of our sample. This has enabled us to compare the Fundamental Plane of our distant bulges with that for spheroidal galaxies in a similar redshift range.  Bulges of spiral galaxies with a bulge-to-total luminosity ratio ($B/T$)\\,$>$\\,0.2 show very similar patterns of evolution to those seen for pure spheroidals such that the stellar populations of all spheroids with $M$\\,$>$\\,$10^{11}$\\,$M_{\\odot}$ are homogeneously old, consistent with a single major burst of star formation at high redshift ($z_f$\\,$\\gtrsim$\\,2), while bulges with $M$\\,$<$\\,$10^{11}$\\,$M_{\\odot}$ must have had more recent stellar mass growth ($\\sim$\\,10\\% in mass since $z$\\,$\\sim$1).  Although further data spanning a wider range of redshift and mass is desirable, the striking similarity between the assembly histories of bulges and low mass spheroidals is difficult to reconcile with the picture whereby the majority of large bulges form primarily via secular processes within spiral galaxies. ",
        "introduction": "The history of galactic bulges remains a key issue in studies of the origin of the Hubble sequence. Originally thought to form at high redshift through dissipationless collapse (Eggen, Lynden-Bell \\& Sandage 1962), their continued growth, as predicted in hierarchical models (Baugh \\etal\\ 1998), is consistent with the diversity observed in their present-day stellar populations (Wyse, Gilmore \\& Franx 1997). Local data alone, however, cannot distinguish between quite different hypotheses for bulge formation, including secular processes triggered by interactions and the evolution of bars (Kormendy \\& Kennicutt 2004; Combes 2006).  The assembly history of bulges is also central to understanding the strong correlations observed between the nuclear black hole mass and bulge properties (Magorrian \\etal\\ 1998; Gebhardt \\etal\\ 2000). The different physical scales involved in these local scaling relations represent a major theoretical challenge (Miralda-Escud\\'e \\& Kollmeier 2005). As ambiguities remain even if fairly precise observations are available for local samples, an important route to understanding this puzzle lies with undertaking observations at intermediate redshift. This is especially challenging for measurement of black hole masses where the sphere of influence remains unresolved (c.f.\\@ Woo \\etal\\ 2006). However, as an alternative approach, it may be more practical to attempt to measure the growth history of the bulges in which they reside.  The arrival of deep multi-color imaging data from the {\\it Hubble Space Telescope} ({\\it HST}) provided the first glimpse of the photometric properties of bulges at intermediate redshift. In an early paper, Ellis, Abraham \\& Dickinson (2001, hereafter EAD), examined aperture colors of bulges in 68 suitably oriented isolated spirals with $I_{AB}$\\,$<$\\,24 in the northern and southern Hubble Deep Fields (HDF). The authors found a remarkable diversity in bulge colors over the redshift range 0.3\\,$<$\\,$z$\\,$<$\\,1 (using both spectroscopic and photometrically determined redshifts), with few as red as a passively-evolving track which matches the integrated colors of luminous spheroidal galaxies observed in the HDFs \\footnote{We will adopt the term `spheroidal galaxy' to denote pure elliptical systems, reserving `spheroidal component' where necessary to refer to bulges within galaxies also harboring a disk.}. EAD concluded that bulges have suffered recent periodic episodes of rejuvenation consistent with 15--30\\% growth in stellar mass since $z$\\,$\\simeq$\\,1. These conclusions were challenged by Koo \\etal\\ (2005, hereafter Koo05) who located 52 luminous ($I_{AB}$\\,$<$\\,24) bulges in the shallower but wider field Groth Strip Survey for which a more elaborate photometric decomposition was undertaken. They found that 85\\% of their field sample had uniformly red colors at $z$\\,$\\simeq$\\,0.8 ($\\delta (U-B)$\\,$\\simeq$\\,$\\pm$\\,0.03), as red as present-day and distant cluster E/S0s.  Only a minority (8\\%) showed blue rest-frame colors, most of which occurred in interacting or merging systems.  In the interim, much has been learned about the luminosity dependence of evolution in the field spheroidal galaxy population. Treu \\etal\\ (2002), van~Dokkum \\& Ellis (2003), Treu \\etal\\ (2005a\\&b, hereafter T05), van~der~Wel \\etal\\ (2005), and di Serego Alighieri \\etal\\ (2005) have undertaken comprehensive Fundamental Plane (FP) analyses of several hundred field galaxies to $z$\\,$\\simeq$\\,1.  Whereas the most luminous spheroidal galaxies studied support the long-held view of early collapse and subsequent passive evolution (\\eg\\ Bower, Lucey, \\& Ellis 1992), a surprising amount of recent star formation is necessary to explain the scatter and FP offsets for lower-luminosity galaxies. T05 find that as much as 20--40\\% of the present dynamical mass in systems with $M<10^{11}$\\,$M_{\\odot}$ formed since $z$\\,$\\simeq$\\,1.2.  For these systems, spectroscopic signatures of recent star formation are visible and several show resolved blue cores, consistent with recently-accreted gas-rich dwarfs.  The present paper aims to clarify the relationship between bulges and spheroidals in the light of the above work. The first part is motivated by the arrival of superior resolution multi-color ACS imaging data in the Great Observatories Origins Deep Survey (GOODS) fields; this represents a significant improvement over the WFPC2 data used by EAD and Koo05. The GOODS ACS dataset offers the potential of securing improved photometric parameters.  In particular, EAD chose not to undertake bulge-to-disk (B/D) decomposition, arguing by example with local data (de~Jong 1996) that aperture colors were adequate.  In the present paper, we revisit this discussion using a larger sample with equivalently-deep but higher spatial resolution, broader wavelength coverage, spectroscopic redshifts, and employing photometric decompositions of bulge and disk parameters.  The second component is concerned with securing dynamical estimates of the bulge stellar masses from resolved spectroscopy.  At intermediate redshifts, resolved spectroscopy of high quality is now possible thanks to instruments such as the Deep Imaging and Multi-Object Spectrograph (DEIMOS; Faber \\etal\\ 2003) on the Keck II telescope. Following recent progress in interpreting the mass assembly history of field spheroidals (\\eg\\ T05), mass-to-light ratios ($M/L$) determined from FP analyses are much superior to interpretations based solely on optical colors. Our goal is to secure the mass assembly history of a representative sample of intermediate redshift bulges and to compare trends found with the integrated properties of field spheroidals.  A plan of the paper follows. In \\S\\ref{sec:sample}, we discuss the selection criteria we adopt for both the photometric and spectroscopic sample of bulges in the GOODS fields. In \\S\\ref{sec:photom} we describe the photometric techniques used for bulge/disk decomposition, taking into account the effects of the ACS point spread function, and derive rest-frame properties (k-corrections).  In \\S\\ref{sec:colors} we compare our photometric results with earlier measures of the diversity and star formation history of intermediate redshift bulges. In \\S\\ref{sec:spec} we describe the spectroscopic measurements undertaken with DEIMOS and their reduction to the central stellar velocity dispersions necessary to construct the FP.  Key issues of disk contamination, systematic rotation, and relative aperture size are discussed.  In \\S\\ref{sec:spec_results} we compare the FP of bulges with the trends now well-established for more massive E/S0 galaxies.  We discuss the implications of both our photometric and spectroscopic results in the context of various formation hypotheses in \\S\\ref{sec:discuss}. We summarize our overall findings in \\S\\ref{sec:summary}.  Throughout this paper, for all distance dependent quantities we adopt a flat cosmological model with $\\Omega_M$\\,=\\,0.3, $\\Omega_{\\Lambda}$\\,=\\,0.7, and \\hub\\,=\\,65~\\hunit. We note that both diagnostics considered in this study -- colors and evolution of mass-to-light ratios -- are independent of the Hubble constant. All magnitudes are in the AB system (Oke 1974) unless otherwise noted. ",
        "conclusions": "\\label{sec:discuss}  Having discussed our photometric and spectroscopic results separately in Sections 4 \\& 5, we are now in a position to combine the inferences from the two diagnostics, compare them with previous studies in the literature, and discuss the physical interpretation.  Colors and mass-to-light ratios provide a consistent picture. The stellar populations of the most massive bulges are homogeneously old (formed at $z$\\,$>$\\,2--3), while an increasingly larger fraction of younger stars (formed below $z$\\,$\\sim$\\,1) is required as one moves down the mass function.  This ``downsizing'' trend is consistent within the errors with that observed for the pure spheroidals and for the spheroidal component of lenticular galaxies.  As discussed earlier, this finding allows us to reconcile the results of previous studies of intermediate redshift bulges based on WFPC2 data (EAD and Koo05).  The results of our study are also consistent with recent studies of the ``fossil record'' from local samples. Recently -- based on the analysis of absorption features -- Thomas \\& Davies (2006) found that, at a given velocity dispersion, local bulges are indistinguishable from elliptical galaxies in terms of their correlations with age and alpha-element enhancement ratios (as a proxy for star formation timescales).  Their sample is similarly limited to large $B/T$ spirals.  What does this all imply about bulge formation? It is important to emphasize that different scenarios could lead to the same star formation history. For example, for the most massive spheroids analyzed in T05, one cannot distinguish based on stellar populations alone whether the mass assembly happened at high redshift, perhaps concurrent with the major episode of star formation, or whether they were assembled at a later time via dry mergers (\\eg\\ van~Dokkum 2005; Bell \\etal\\ 2006). Independent data, such as those on the evolution of the stellar and dynamical mass functions (\\eg\\ Bundy, Treu \\& Ellis 2007), are needed to break the degeneracy and conclude that the {\\it assembly} happened at high redshift as well. For the less massive spheroids, a larger amount of recent star formation is required, but understanding whether this happened in situ or via the accretion of gas reach satellites of mass of order a few $10^9$\\,$M_{\\odot}$ (\\eg\\ T05), again requires external information. The study by Bundy Treu \\& Ellis (2007), points toward minor mergers or internal/secular process as being important alongside major mergers. As far as bulges are concerned then, the similarity in star formation need not necessarily reflect a similarity in mechanisms/timing of assembly. At the high mass end, the homogeneously old stellar populations of bulges seem hard to reconcile with recent formation from existing disk stars or gas.  Therefore it appears that we can safely conclude that massive bulges have been {\\it in place} well before the redshift range probed by this study, all but ruling out significant secular evolution (see also Thomas \\& Davies 2006). Hence the similarity with spheroids for the most massive bulges appears to be both in the star formation history and in the mass assembly history.  At lower masses, secular process may become more important and it is difficult to disentangle recent in situ star formation, mass transfer from the disk to the bulge, and merging with satellites. External evidence suggests that at this redshift and mass scale, bulges are growing significantly via mergers (Woo \\etal\\ 2006; Treu \\etal\\ 2007), although this may not be the only evolutionary mechanism. If bulges at masses below $10^{11}$\\,$M_{\\odot}$ grow by mergers, than the rightmost plot in Figure~\\ref{fig:DML} must be interpreted with caution: galaxies of a given mass at one redshift are not the progenitors of objects with the same mass at a lower redshift. Perhaps a way to identify the spheroid progenitors is to study bulges at fixed central black hole mass, assuming that growth by accretion is mostly negligible. This is difficult to do in practice and is limited at intermediate redshifts to samples of spheroids hosting active nuclei. However -- albeit with large uncertainties -- such an exercise (Treu \\etal\\ 2007) suggests that spheroids of a few $10^{10}$\\,$M_{\\odot}$ may grow in stellar mass by $0.20\\pm0.14$\\,dex in the past four billion years, while keeping the stellar mass-to-light ratio approximately constant as a result of adding young stars to an aging resident population.  The uncertainties are currently too large and the samples too small to perform a quantitative test of this scenario, but it is clear that the evolution of spheroids in this mass range is not purely passive (see also Hopkins \\etal\\ 2006) and that complementary information from stellar populations studies, spheroid demographics, and scaling relations between host galaxy and central black hole are needed to disentangle star formation, assembly history and black hole growth and feedback. The line of studies presented in this work is crucial to accurately pinpoint the star formation history and could in future provide additional independent information if extended to include tracers of chemical enrichment.    We have presented a comprehensive study of field galactic bulges, and of a comparison sample of lenticular and elliptical galaxies, in the redshift range $z$\\,=\\,0.1--1.2. The galaxies are selected from the GOODS survey based on visual classification and luminosity ($z_{AB}$\\,$<$\\,22.5). We measure accurate colors and structural parameters for 193 galaxies, fitting two components (bulge and disk) for 137 galaxies and a single \\sersic\\ profile for the remaining 56, classified as pure spheroids. Stellar velocity dispersions are measured from deep spectroscopic observations using DEIMOS on the Keck-II Telescope. By combining new observations with the measurements derived in T05, we compile a spectroscopic subsample of 147 galaxies suitable for Fundamental Plane analysis, \\ie\\ reliable stellar velocity dispersion and bulge to total luminosity ratio greater than 20\\%, to minimize disk contamination. The spectroscopic subsample includes 56 one component galaxies (pure spheroids) and 91 two component (bulge + disk) galaxies.  We use the two complementary diagnostics -- colors and mass-to-light ratios as determined from the evolution of the Fundamental Plane -- to derive the star formation history of bulges and compare it to that of the spheroidal component of lenticular galaxies and of elliptical galaxies. Although the uncertainties are larger for bulges than for pure spheroids (\\S~\\ref{ssec:cavecanem}) -- and more data both at intermediate redshift and in the local universe are needed to confirm the trends, the two diagnostics give consistent results that can be summarized as follows:  \\begin{itemize}  \\item The stellar populations of the more massive bulges ($M$\\,$>$\\,$10^{11}$\\,$M_{\\odot}$) are homogeneously old, consistent with a single major burst of star formation at redshift $\\sim$\\,2 or higher, and only minor episodes of star formation below $z$\\,$\\sim$\\,1 ($\\lesssim$\\,5\\% in mass).  \\item The colors and mass-to-light ratios of smaller bulges ($M$\\,$<$\\,$10^{11}$\\,$M_{\\odot}$) span a wider range, consistent with an increasing fraction of younger stars going to smaller masses ($\\sim$\\,10\\% in mass).  \\item The assembly history of bulges is consistent within the error with that of the spheroidal component of lenticular and elliptical galaxies of comparable mass.  \\end{itemize}  Our detection of a mass dependent star formation history allows us to reconcile the findings of two earlier studies of intermediate redshift bulges based on WFPC2 data (EAD and Koo05). The more massive bulges are as old and red as massive spheroids, but smaller bulges have quite diverse star formation histories, with significant star formation below $z$\\,$\\sim$\\,1.  The similarity between the mass assembly history of massive bulges and that of spheroids appears quite naturally explained in a scenario where the former are assembled at high redshift, rather than recently via secular process within spiral galaxies. At the lower mass end of our sample ($M$\\,$\\sim$\\,$10^{10}$\\,$M_{\\odot}$), the evolution of bulges below $z$\\,$\\sim$\\,1 becomes much more diverse, requiring perhaps a combination of merging and secular processes as observed for spheroids of similar mass (Bundy, Treu, \\& Ellis 2007).  This ``downsizing'' (\\eg\\ Cowie \\etal\\ 1996; Bundy \\etal\\ 2005; De~Lucia \\etal\\ 2006; Renzini 2007) picture for bulges extends previous findings based on the integrated properties of galaxies. In addition, this picture is qualitatively consistent with recent results on the co-evolution of black holes and spheroids (Treu, Malkan, \\& Blandford 2004; Walter \\etal\\ 2004; Woo \\etal\\ 2007; Peng 2007; Treu \\etal\\ 2007), which suggests that bulges in the mass range considered here may have completed their growth after that of the central black hole."
    },
    "0711/0711.0712_arXiv.txt": {
        "abstract": "We investigate the link between the warp function and the potential in DBI cosmologies in connection with the possibility they represent power-law solutions. A prescription is given to take advantage of the known result that given a warp factor there is always a choice of potential resulting in a constant ratio between pressure and energy density. The method is illustrated with examples with interesting models for either the warp factor or the potential. We complete this investigation by giving a recipe to exploit symmetries in order to generate new solutions from existing ones; this method can be applied, for instance, to the power-law cosmologies obtained using our prescription. ",
        "introduction": "Consistent flavours of string theory allowing for open strings include a class of extended objects called D-branes to  which the end points of the open strings are attached. Those extended physical objects are typically non-perturbative features of string theory, and have therefore played a substantial role in the development of the theory (see \\cite{Quevedo:2002fc} and references therein)   The idea that the inflaton might be an excitation of a D-brane in the form of an opening string has captivated quite a few researchers (see \\cite{Dvali:1998pa,Leigh:1989jq,Kachru:2003sx,Cline:2006hu,Kallosh:2007ig,HenryTye:2006uv,Silverstein:2003hf,Alishahiha:2004eh,Chen:2004hua,Chen:2005ad, Kecskemeti:2006cg,Lidsey:2006ia,Baumann:2006cd,Spalinski:2007qy,Huang:2007hh,Bean:2007eh,Gmeiner:2007uw,Spalinski:2007dv,Spalinski:2007kt} and references therein for regular papers and \\cite{Peiris:2007gz} for a review). The inflaton could be a mode describing the position of a D3-brane  wandering (radially) in a ten-dimensional space--time with a warped metric. The interest of this interpretation of inflation is not just that it would brush it with an aristochratic patina, it is a more pragmatic one, as in these stringy scenarios inflation happens to occur without such strict requirements as in the standard weekly coupled slow roll inflation model. Specifically, inflation can proceed with much steeper potentials.     The motion of the brane seems to admit an effective good description in terms of a Dirac-Born-Infeld action coupled to gravity \\cite{Alishahiha:2004eh}, and it results in a scalar field theory with non- canonical kinetic terms. The usual assumption is that the metric on the brane is a flat FRW one, and it is customary to take advantage of the usual perfect fluid formulation of the corresponding Einstein equations. The modifications with respect to canonical scalar field models appear in the form of factors in front of the square of the gradient of this scalar field, which in turn depend on both the scalar field and its gradient also. This factor is related to the speed of the wandering brane and investigations using inflationary parameters had been carried out to estimate how this non-canonicalities in the kinetic terms of the model affect the cosmological observables  \\cite{Silverstein:2003hf,Alishahiha:2004eh,Chen:2004hua,Chen:2005ad,Kecskemeti:2006cg,Lidsey:2006ia,Baumann:2006cd,Spalinski:2007kt,Huang:2007hh,Bean:2007eh,Gmeiner:2007uw}.   In this context of FRW DBI cosmologies there has been a considerable interest in possible realizations of power-law inflation, given the preeminence such evolutions have gathered over the years due to  their role as asymptotic equilibrium states \\cite{wain}.  An exceedingly well-known fact regarding power-law solutions is that, in terms of the customary perfect fluid interpretation (which applies also to DBI FRW cosmologies) power-law cosmologies correspond to a constant ratio between pressure and energy density. In these scenarios there are two degrees of freedom available which could be chosen so as to result is power-law cosmological models \\cite{Spalinski:2007dv} that for a given warp factor there is always a choice of potential resulting in a constant ratio between pressure and energy density:  one of these free functions is the warp factor of the metric, which is denoted as $f(\\phi)$, whereas the other is the inflaton potential $V(\\phi)$. In this paper we provide an algorithm to exploit that correspondence in order to generate DBI solutions. The novelty and  relevance of our method is that, as far as we are concerned, no exact (genuinely) DBI solutions exist in the literature, not even power-law ones. Of course, if one's aim is to enforce power-law solutions corresponding to a particular choice of either the warp factor or the potential, then some degree of compromise will be needed and approximations will perhaps be required \\cite{Silverstein:2003hf,Alishahiha:2004eh,Spalinski:2007kt,Spalinski:2007qy,Spalinski:2007dv}, this was indeed the route taken in precursor works. We advocate, however, that progress in the construction of power-law DBI cosmologies can be made along the path of exact formulations as well if the problem is posed in a different fashion and one rather seeks for a particular asymptotic form of the warp factor and potential rather than demanding it corresponds to all regimes.  Specifically, we present neat expressions leading to a constructive recipe for DBI power-law cosmologies which allows us for instance to show that the almost ubiquitous exponential potential can indeed accommodate that sort of solutions, provided one chooses an adequate warp factor. After analyzing this case we move on and conside others which under appropriate choices of parameters and in asymptotic regimes lead to for instance the AdS throat warp factor \\cite{Randall:1999vf}, or to  a generalization of the inverse power-law potential which was earlier studied in connection with tachyon cosmologies \\cite{Garousi:2004ph}.   In the course of the discussion we briefly comment on the conditions for the warp function $f$ to be positive (the contrary case is not consistent from the string theory point of view, but it is admissible in the field theory spirit). Finally, we also discuss about the possibility of generating further solutions from existing ones under the use of symmetries, and how phantom DBI cosmologies could fit into this picture. ",
        "conclusions": "The so called DBI cosmologies are scalar field models depending with one additional functional degree of freedom as compared to conventional scalar field configurations. This extra function has a geometrical meaning when DBI actions are interpreted as describing the motion of a brane in a warped space time, the warp factor being indeed that additional input required to specify the model. In this setup inflation is interpreted as the consequence of the motion of the brane in a background with extra dimensions, the compactification of which give rise to the scalar field's potential.  These scenarios have been profusely studied, and particular attention has been paid to the possibility they admit power-law solutions, as they seem to be favoured to play the role of equilibrium asymptotic states. In this paper we have shown, that despite earlier claims pointing in the direction of finding such solutions in an exact way, completing such task is indeed possible. The method relies on providing a given parametrization of the scalar field in terms of time, and then one can reconstruct the warp factor and the potential upon the sole requirement that the equation of state parameter be constant (in the fluid interpretation of the model). This powerful result is illustrated by resorting to some examples with interesting asymtptotic limits coinciding with some models studied in the literature: the AdS throat, cut-off throats,  and the inverse square potential.  Finally, in the last section, we show how to obtain new solutions  from existing ones using symmetries, specifically we give transformations rules for the scalar field, warp factor and potential. This method has interesting applications as is offers the possibility of realizing assisted inflation in a DBI context, but also permits playing with the idea of phantom DBI cosmologies (to be generated from non phantom ones)."
    },
    "0711/0711.2521_arXiv.txt": {
        "abstract": "We develop a new approach to study the nonlinear evolution in the large-scale structure of the Universe both in real space and in redshift space, extending the standard perturbation theory of gravitational instability. Infinite series of terms in standard Eulerian perturbation theory are resummed as a result of our starting from a Lagrangian description of perturbations. Delicate nonlinear effects on scales of the baryon acoustic oscillations are more accurately described by our method than the standard one. Our approach differs from other resummation techniques recently proposed, such as the renormalized perturbation theory, etc., in that we use simple techniques and thus resulting equations are undemanding to evaluate, and in that our approach is capable of quantifying the nonlinear effects in redshift space. The power spectrum and correlation function of our approach are in good agreement with numerical simulations in literature on scales of baryon acoustic oscillations. Especially, nonlinear effects on the baryon acoustic peak of the correlation function are accurately described both in real space and in redshift space. Our approach provides a unique opportunity to analytically investigate the nonlinear effects on baryon acoustic scales in observable redshift space, which is requisite in constraining the nature of dark energy, the curvature of the Universe, etc., by redshift surveys. ",
        "introduction": "Introduction }  Recent progress in cosmology is largely stimulated by precise measurements of the Universe. The redshift survey of galaxies is one of the most important methods in distinguishing cosmological models. The baryon acoustic oscillations (BAOs) in the early universe \\cite{BAOtheory}, which were first detected in the cosmic microwave background (CMB) anisotropy power spectrum \\cite{CMBobs}, were also detected in the large-scale structure probed by redshift surveys \\cite{BAOLSS}. The expansion history and curvature of the Universe can be geometrically measured by galaxy clustering in redshift space \\cite{AP79, GeomTest}, and the BAO signature in the large-scale structure yields a standard ruler \\cite{BAOSR} for this purpose. Therefore, accurately measuring the BAO scales in a statistical quantity such as the power spectrum or correlation function of galaxies, one can constrain the structure of the Universe, such as the nature of dark energy \\cite{BAODE}.  To utilize the BAOs as a standard ruler in the Universe, a theoretically precise description of the BAOs is of the utmost importance. The physics of the BAOs at sufficiently large redshift is well described by linear theory. Only the amplitude of the spatial pattern of clustering is increased, and the characteristic scale of the BAOs is unaltered in this regime. However, the nonlinear dynamics play an important role in a redshift range where we can measure the large-scale structure. Even though the BAO scales $\\simeq 100\\,h^{-1}{\\rm Mpc}$ are large, the nonlinearity deforms the BAO pattern in the power spectrum and correlation function \\cite{SE05, Spr05, BAONL, BAONLRD, ang07, ESW07}. Nonlinear deformations of the power spectrum are described by the standard perturbation theory (SPT) \\cite{jus81, vish83, Fry84, Goroff86, MSS, JB94, BCGS02} of higher orders as long as the density fluctuations on BAO scales are still in the linear regime at sufficiently large redshift \\cite{MWP99, JK06}. The last condition is not necessarily satisfied by realistic surveys. Moreover, when the cold dark matter (CDM) is present, the higher-order SPT does not give any sensible prediction for the correlation function because unphysical behavior in small wavelength limit prevents Fourier transform from converging.  Recently, several other approaches beyond the framework of SPT have been proposed \\cite{Sco01, NewPT, CS06a, CS06b, Val07, CS07, RGPT, TH07}. Among others, Scoccimarro's reformulation of fluid equations using the propagator, the vertex, and a source \\cite{Sco01} provides a way to use standard tools of field theory, yielding the renormalized perturbation theory (RPT) \\cite{CS06a, CS06b, CS07}, the large-$N$ expansions \\cite{Val07}, the renormalization group approach \\cite{RGPT}, and the closure theory \\cite{TH07} in the context of gravitational instability theory. In these newly developed approaches, the standard perturbative expansion is reorganized and partially resummed in various ways. It is still difficult to exactly calculate the power spectrum by these approaches, and different levels of approximations and ansatz are employed. The predictions of the power spectrum calculated by the 2-loop RPT and renormalization group method are shown to agree with measurements of the $N$-body simulations on BAO scales even at low redshift where the 1-loop SPT breaks down \\cite{CS07, RGPT}. It is still nontrivial why these approaches work well despite their approximate treatments and use of ansatz.  So far the predictions of the above nonlinear theories, such as SPT, RPT, and their variants, have been confined in real space. Since the large-scale structure is actually measured in redshift space, it is mandatory to make predictions in redshift space for applying the nonlinear theories to realworld observations. The redshift-space distortion effects in the linear regime are analytically given by first-order SPT (i.e., linear theory) \\cite{Kaiser87, Ham92, GeomTest}. However, investigations of nonlinear redshift-space distortions by higher-order SPT are quite limited in literature \\cite{HMV98,SCF99}, apart from numerical simulations and nonlinear modelings \\cite{MWP99, BAONLRD, ang07, Sco04, SE05, ESW07}.  In this paper, we consider resummations of the perturbative expansion by a different approach from RPT and its variants. Instead of relying on Scoccimarro's reformulation, we start from a Lagrangian picture of density perturbations. Why do we have to add yet another approach even though there are already enough? There are a number of reasons. First, our approach is simpler than other resummation methods. While it is straightforward to calculate the power spectrum in our approach, other resummation techniques are so complicated that approximate treatments and/or ansatz are necessary in deriving the final power spectrum. Second, as a result of that, our analytic expression of the resulting power spectrum is no more complicated than the SPT expression, and numerical evaluations are fast performed. Third, and most importantly, our approach is capable of evaluating the nonlinear power spectrum {\\em in redshift space}, which is beyond the scope of other resummation methods in literature. Fourth, unlike in SPT, the power spectrum in our approach can be Fourier transformed to predict the nonlinear effects in the correlation function.  An important ingredient of our approach is the perturbation theory in Lagrangian space, which is known as the Lagrangian perturbation theory (LPT) \\cite{LPT}. The first-order LPT corresponds to the classic Zel'dovich approximation \\cite{Zel70}. Since the observable power spectrum is given in Eulerian space, predicting the power spectrum by LPT is not a trivial task. We define a method of calculating the power spectrum from the LPT, partially expanding the Lagrangian variables in Eulerian space. The resulting expression contains an infinite series of SPT terms; that is, the SPT terms are automatically resummed in our formulation. The underlying motivation of this formulation is shared with RPT. Readers are recommended to go through an excellent introduction in Ref.~\\cite{CS06a} using the Zel'dovich approximation, which equally applies to our formulation since our formalism and RPT become equivalent in a case of Zel'dovich approximation.  The rest of this paper is organized as follows. In Sec.~\\ref{sec:Resum}, our basic formalism is described. A brief review of LPT is also included. In Sec.~\\ref{sec:RealSpace}, analytic results of the power spectrum and correlation function in real space are presented and compared with simulation results in the literature. The corresponding results in redshift space are derived in Sec.~\\ref{sec:RedshiftSpace}. The conclusions are given in Sec.~\\ref{sec:concl}. ",
        "conclusions": "}  In this paper, we have formulated a new resummation method of the cosmological perturbation theory, starting from the Lagrangian description of perturbations. The infinite series of perturbations in SPT is naturally contained in the exponential prefactor in our formalism, and expanding the exponential factor reproduces the SPT. We have presented 1-loop results of our formalism in this paper. Our method is particularly useful for studying nonlinear effects on BAO scales and always provides a better description of the $N$-body results than SPT on large scales.  Nonlinear corrections to the correlation function can be calculated by our formalism. This is a great advantage over SPT, since higher-order SPT cannot predict the correlation function. Our 1-loop results are comparable to the 1-loop results of the RPT \\cite{CS07}. In the RPT, it is shown that taking into account the 2-loop effects enlarges the applicable range toward small scales in the power spectrum. Since our formalism and RPT have common features in the resulting equations, we expect that 2-loop extensions of our calculation similarly enlarge the applicable range of scales in the power spectrum. The description of the BAO peak in the correlation function seems already sufficient by our 1-loop calculation.  It is an important and original outcome that our formalism gives an analytical description of nonlinear effects in redshift space. The large-scale structure is measured by redshift surveys of astronomical objects, such as galaxies, quasars, Lyman-alpha absorbers, 21 cm emission, etc. However, nonlinear corrections to the power spectrum and correlation function in redshift space by perturbation theory have not been sufficiently studied so far. Our formalism and results presented in this paper are unprecedented in this respect and are crucial steps toward applying the BAO measurements in redshift surveys to constraining the nature of dark energy, the curvature of the Universe, etc.  Considering the biasing effects in galaxy clustering is beyond the scope of this paper. A simple prescription of the biasing is to replace $f \\rightarrow \\beta = f/b$ in our equations, where $b$ is the linear bias parameter. However, nonlinearity and nonlocality of biasing might not be negligible even on BAO scales \\cite{BAONL, ang07}. Incorporating the nonlinear and nonlocal biasing into our formalism will be addressed in future work."
    },
    "0711/0711.2698_arXiv.txt": {
        "abstract": "The collapse of baryons into the center of a host dark matter halo is accompanied by radiation that may be detectable as compact ($\\lsim 10$ kpc) UV-continuum and Ly$\\alpha$ emission with Ly$\\alpha$ luminosities as high as $\\sim 10^{42}-10^{43}$ erg s$^{-1}$ in halos of mass $M=10^{11}-10^{12}M_{\\odot}$. We show that the {\\it observed} equivalent width (EW) of the Ly$\\alpha$ line emitted by these cooling clouds is EW$\\lsim 400$ \\AA\\hs(restframe). These luminosities and EWs are comparable to those detected in narrowband surveys for redshifted Ly$\\alpha$ emission. The rest-frame ultraviolet of Ly$\\alpha$ emitting cooling clouds radiation may be dominated by two-photon transitions from 2s$\\rightarrow$1s. The resulting spectrum can distinguish cooling clouds from a broad class of young star forming galaxies. ",
        "introduction": "\\label{sec:intro}  Gas collapse into the center of a host dark matter halo requires the release of gravitational binding energy in the form of cooling radiation. The efficiency of gas cooling therefore plays an important role in the formation of stars and galaxies. For example, it has been shown that trapping of cooling radiation can affect the fragmentation of a collapsing gas cloud \\citep{SS06}. Understanding cooling radiation is therefore essential in our understanding of galaxy formation. A large fraction of cooling radiation may emerge as Ly$\\alpha$ emission \\citep{HS00,Fardal01}, possibly accompanied by emission from helium \\citep{Yang06}. Galaxies may therefore be visible as extended Ly$\\alpha$ sources before star formation occurs. This realization coincided with the discovery of two large Ly$\\alpha$ 'blobs' (LABs) by \\citet{Steidel00}\\footnote{Extended Ly$\\alpha$ emission is quite common around radio galaxies, and is thought to be powered by some sort of jet-IGM interaction \\citep[e.g.][]{Chambers90,V07}. However, the LABs discovered by \\citet{Steidel00} have no associated radio-galaxies.}. Naturally, LABs have been associated with cooling radiation.  Since the discovery of LABs by \\citet{Steidel00} several tens of new LABs have been discovered \\citep{Matsuda04,Dey05,Saito06,Ouchi08b}. Various groups have investigated what powers Ly$\\alpha$ emission in LABs \\citep[e.g.][]{Matsuda06,Saito07}. Recent studies suggest that the majority is likely powered by (obscured) stars and/or AGN \\citep[e.g.][]{Basu04,Matsuda07}, or by a superwind \\citep[e.g.][]{Taniguchi00,Mori04}. However, in some LABs it is possible to rule out all these mechanisms, as each is expected to produce characteristic metal lines at flux levels that are not detected. This has led to the conclusion that cooling radiation may have been detected \\citep{Nilsson06,Saito07,Smith07}.  Cooling radiation may also manifest itself as compact Ly$\\alpha$ emission (rather than spatially extended) with luminosities comparable to those observed in high-redshift Ly$\\alpha$ emitting galaxies. \\citet{HS00} pointed out that $\\sim 30\\%$ of their total computed cooling luminosity may originate from the central $\\lsim 10\\%$ in radius. Furthermore, \\citet{Fardal01} and \\citet{Yang06} found that spatially extended cooling gas clouds in 3-D simulations contain dense compact clumps ($< 10$ kpc, $n_{\\rm H,max} \\sim 50$ cm$^{-3}$) which dominate the total Ly$\\alpha$ cooling emission. Since it is plausible that star formation takes place inside these densest clumps, a major challenge lies in determining which process provides the biggest contribution to the Ly$\\alpha$ luminosity emerging from clumps: cooling or young stars \\citep[e.g.][for a more detailed discussion]{Keres05}. \\begin{figure*} \\vbox{\\centerline{\\epsfig{file=fig1.eps,angle=270,width=11.0truecm}}} \\caption[]{The collapse of a $M_{\\rm tot}=10^{11.5}M_{\\odot}$ halo in a 1-D hydrodynamical simulation and its accompanying Ly$\\alpha$ cooling radiation are shown. The {\\it top panel} shows the radial evolution of mass shells initially enclosing log$M_{\\rm tot}/M_{\\odot}=10.5,11$ and $11.5$ (see labels). The lower horizontal axis shows the age of the Universe in Gyr, while the upper shows redshift. The {\\it lower three panels} show the total Ly$\\alpha$ luminosity {\\it outside} a sphere of radius $R$ as a function of $R$ (see text for a justification of this peculiar choice) at three times. The total Ly$\\alpha$ luminosity derived from the 1-D code is likely underestimated (see text). For comparison, we also show the luminosity in recombination ({\\it dotted line}) and bremsstrahlung ({\\it dashed line}) cooling radiation. These processes are clearly subdominant. In the early stages of the collapse ({\\it lower left panel}) shells at all radii contribute to the total Ly$\\alpha$ luminosity. However, during later stages in the collapse the Ly$\\alpha$ luminosity is dominated by emission from the central few kpc. This Figure illustrates a property of cooling radiation that is not unique to this spherical collapse model: namely that {\\it the total Ly$\\alpha$ luminosity of a cooling cloud may be dominated by emission from the central few kpc.}} \\label{fig:sim} \\end{figure*}  \\citet{D06} quantified the properties of cooling radiation by performing Monte-Carlo Ly$\\alpha$ radiative transfer calculations in collapsing optically thick gas clouds, and found that Ly$\\alpha$ cooling radiation should be systematically blueshifted. In practice it may prove difficult to actually observe this blueshift, as this requires the detection of other lines, which are often absent. \\citet{D06} also showed that Ly$\\alpha$ cooling radiation may reveal itself through a frequency dependent surface brightness profile: the surface brightness profile of the reddest Ly$\\alpha$ photons is expected to be flatter than that of the bluest Ly$\\alpha$ photons. However, the strength of this signal is model dependent and may be too subtle to be detectable. Therefore, it is desirable to have more diagnostics that characterize cooling radiation.  The main goal of this paper is to provide new diagnostics to identify (compact) cooling clouds in data generated by existing and upcoming telescopes\\footnote{ The term 'Ly$\\alpha$ cooling radiation' that is used throughout this paper refers to both real (i.e. collisionally excited) cooling radiation, as well as Ly$\\alpha$ radiation that is emitted as a by-product of recombination cooling (even though Ly$\\alpha$ radiation that is produced in radiative cascades following recombination is technically not cooling radiation, since it does not transport any heat from the gas). In both cases, the Ly$\\alpha$ radiation is emitted in a 'proto-galactic' gas cloud that is still collapsing to from stars. Furthermore, for neutral fractions as low as $x_{\\rm HI}=10^{-4}$, the 'real' Ly$\\alpha$ cooling rate due to collisional excitation, $\\Lambda_{{\\rm Ly}\\alpha}$ is equal to the Ly$\\alpha$ emission rate due to recombination (for $T_{\\rm gas}=3\\times 10^4$ K, see Appendix~\\ref{app:reccool}). Hence, even recombining \\ion{H}{2} regions can emit a significant amount of Ly$\\alpha$ cooling radiation.}. We compute the continuum emission of cooling clouds redward of Ly$\\alpha$, and show that this is dominated by two-photon transitions from 2s$\\rightarrow$1s. The spectrum of this continuum has a peculiar shape which may distinguish cooling clouds from young star forming galaxies.  The outline of this {\\it paper} is as follows. In \\S~\\ref{sec:sim} we illustrate that cooling radiation may be detectable as compact Ly$\\alpha$ emission. This is followed by a calculation of the spectrum of cooling radiation at wavelengths redward of Ly$\\alpha$ in \\S~\\ref{sec:spectrum}. Finally, in \\S~\\ref{sec:conc} we discuss the results of this paper, and present our main conclusions. ",
        "conclusions": "\\label{sec:conc}  The collapse and condensation of baryons into the center of their host dark matter halos is accompanied by emission of radiation that may be detectable as compact UV-continuum and Ly$\\alpha$ emission with Ly$\\alpha$ luminosities as high as $\\sim 10^{42}-10^{43}$ erg s$^{-1}$ in halos of mass $M=10^{11}-10^{12}M_{\\odot}$. These luminosities are comparable to those detected in narrowband surveys for redshifted Ly$\\alpha$ emission.  Cooling clouds emit continuum radiation redward of Ly$\\alpha$ that is dominated by two-photon transitions from 2s$\\rightarrow$1s. The {\\it emitted} Ly$\\alpha$ equivalent width (EW) relative to this continuum is EW$\\sim 1000-1300$ \\AA\\hs(restframe), which is comparable to the EW emitted by galaxies forming either metal-poor or very massive stars. The IGM is likely to transmit only $\\lsim 30\\%$ of this Ly$\\alpha$ flux and the {\\it observed} rest frame equivalent width is therefore EW$\\lsim 400$ \\AA. This is comparable to the observed equivalent widths among high-redshift Ly$\\alpha$ emitting galaxies. It is therefore possible that existing surveys have detected compact Ly$\\alpha$ emission that was emitted by cooling clouds.  Because of the dominance of the two-photon continuum, cooling clouds emit a spectrum that peaks at $\\lambda\\sim 1600$\\AA\\hs(rest-frame), and sharply drops down to $\\lambda=1216$\\AA. This peculiar spectral shape distinguishes cooling clouds from observed young star forming galaxies, whose continuum is dominated by radiation from stars/quasars \\citep[e.g.][]{Stanway05}. However, this is not true for metal-free galaxies that are forming stars according to top-heavy initial mass function. In these primordial galaxies, nebular emission may dominate the continuum redward of Ly$\\alpha$ during the first few Myr of their lifetime (Schaerer 2002). On the other hand, these galaxies also emit \\ion{He}{3} recombination lines such as the He 1640 line, which is not the case for cooling radiation when the gas is not shock heated to $T\\gsim 10^4$ K (see \\S~\\ref{sec:tracer}). Furthermore, deeper Ly$\\alpha$ observations of compact cooling clouds should find weaker emission on larger scales. However, the same applies to Ly$\\alpha$ emitting star forming galaxies, which are likely surrounded by faint Ly$\\alpha$ halos due to resonant scattering in the surrounding IGM \\citep[e.g.][]{LR99,DL08}. This scattered Ly$\\alpha$ emission is not accompanied by continuum radiation redward of the Ly$\\alpha$ line, which could distinguish it from spatially extended Ly$\\alpha$ cooling radiation (polarimetry may provide additional insights, see Dijkstra \\& Loeb 2008).  Ongoing and future surveys for redshifted Ly$\\alpha$ emission aim to detect Ly$\\alpha$ emitters out to redshifts as high as $z\\sim 6-12$ \\citep[e.g.][and references therein]{StarkLF,Nilsson07}. These observations will provide new insights into the epoch of reionization and galaxy formation at high redshift. Primordial galaxies -like cooling clouds- may emit very large EW Ly$\\alpha$ emission lines \\citep[e.g.][]{S03}. In order to measure the EW of the Ly$\\alpha$ line and to determine whether primordial galaxies are detected, the continuum redward of Ly$\\alpha$ must be detected. This is challenging; if the observed flux in the Ly$\\alpha$ line is $f_{\\alpha,17}$ (in units of $10^{-17}$ erg s$^{-1}$ cm$^{-2}$) and its equivalent width is EW (in \\AA), then the AB-magnitude of the continuum is $m_{\\rm AB}=30.4-2.5\\log f_{\\alpha,17}+2.5\\log ({\\rm EW}/200 {\\rm \\AA})$. For comparison, the Hubble Ultra Deep Field reached an $8-\\sigma$ limiting magnitude in the z' filter of $z'_{\\rm AB}=28.5$ \\citep{Bunker04}.  Detecting both the Ly$\\alpha$ line and the continuum is not only desirable for identifying primordial galaxies. An additional advantage of combining Ly$\\alpha$ and continuum measurements of high-redshift galaxies is that it can greatly improve the constraints on the epoch of reionization \\citep[e.g.][]{K06,McQuinn07} over constraints obtained using Ly$\\alpha$ alone. The work presented in this paper suggests that if primordial galaxies can be detected and identified, then the same applies to cooling clouds. Moreover, primordial galaxies may only be luminous in Ly$\\alpha$ for $\\lsim 10$ Myr \\citep{MR02,popIII}. Cooling clouds could be luminous in Ly$\\alpha$ for longer than 10 Myr, which would make them more easily detectable than primordial galaxies. The discovery of two-photon dominated Ly$\\alpha$ emitters would offer us an exciting glimpse of either the process of gas cooling in newly forming galaxies, or of star formation in the first primordial galaxies.  {\\bf Acknowledgments} I thank Marco Spaans for his comments on an earlier draft of this paper, and Zolt\\'{a}n Haiman for useful discussions. This work was supported by Harvard University funds."
    },
    "0711/0711.4182_arXiv.txt": {
        "abstract": "Suzaku TOO observation of the anomalous X-ray pulsar CXOU~J164710.2-455216 was performed on 2006 September 23--24 for a net exposure of 38.8 ks. During the observation, the XIS was operated in 1/8 window option to achieve a time resolution of 1 second. Pulsations are clearly detected in the XIS light curves with a barycenter corrected pulse period of 10.61063(2)~s. The XIS pulse profile is found to be highly non-sinusoidal. It shows 3 peaks of different amplitudes with RMS fractional amplitude of $\\sim$11\\% in 0.2--6.0~keV energy band. Though the source was observed with the Hard X-ray Detectors (HXD) of Suzaku, the data is highly contaminated by the nearby bright X-ray source GX~340+0 which was in the HXD field of view. The 1--10~keV XIS spectra are well fitted by two different models consisting of a power-law and a blackbody component and two blackbody components respectively. Although both the models are statistically acceptable, difference in the pulse profiles at soft (0.2--6.0~keV) and hard (6--12~keV) X-rays favors the model consisting of two blackbody components. The temperatures of two blackbody components are found to be 0.61$\\pm$0.01~keV and 1.22$\\pm$0.06~keV and the value of the absorption column density is 1.73$\\pm$0.03 $\\times$ 10$^{22}$ atoms cm$^{-2}$. The observed source flux in 1--10~keV energy range is calculated to be 2.6 $\\times$ 10$^{-11}$ ergs cm$^{-2}$ s$^{-1}$ with significant contribution from the soft blackbody component ($kT$ = 0.61~keV). Pulse phase resolved spectroscopy of XIS data shows that the flux of the soft blackbody component consists of three narrow peaks, whereas the flux of the other component shows a single peak over the pulse period of the AXP\\@. The blackbody radii changes between 2.2--2.7~km and 0.28--0.38~km (assuming the source distance to be 5~kpc) over pulse phases for the soft and hard components, respectively. The details of the results obtained from the timing and spectral analysis is presented. ",
        "introduction": "The Anomalous X-ray pulsars (AXPs) are a small group of X-ray pulsars which show common properties such as (i) known to present within 1$^\\circ$ of the Galactic plane, (ii) pulse period in a very narrow range of 5--12~s, unlike the radio pulsars and accreting X-ray pulsars, (iii) large and more or less steady spin-down or braking, in contrast to most accretion powered pulsars which show spin ups and downs, (iv) similar spectra with soft component characterized by a blackbody model with temperature below 1 keV, with additional harder component in some cases, and (v) relatively high X-ray luminosity ($\\sim$10$^{34}$--10$^{36}$~erg~s$^{-1}$) which cannot be obtained from the loss of rotational energy of a neutron star alone. Some of the AXPs also emit at optical and/or infrared wavelengths (Hulleman et al.\\ 2000; Durant \\& van Kerkwijk 2006b). The AXPs are considered to be very young (10$^{3}$--10$^{5}$~yr), some of which are associated with supernova remnants (SNRs). The AXPs are considered to be the neutron stars with the strongest known magnetic field (10$^{14}$--10$^{15}$~G), though a direct measurement to confirm the presence of such high magnetic field strengths is still needed. The source of energy for the radiative emission in AXPs is described by the magnetar model, in which the decay of an ultra-strong magnetic field powers the high-luminosity bursts and also a substantial fraction of the persistence X-ray emission (Thompson \\& Duncan 1996). The competing model for the mechanism of powering the X-ray emission in AXPs is that the AXPs are neutron stars surrounded by fossil disks that were acquired during supernova collapse or during a common-envelope interaction (Corbet et al.\\ 1995; van Paradijs et al.\\ 1995; Chatterjee \\& Hernqyist 2000). The properties of AXPs are summarized by Mereghetti et al.\\ (2002), Kaspi \\& Gavriil (2004), Woods \\& Thompson (2004).  Recently discovered AXP CXOU~J164710.2-455216 is located in the young, massive Galactic star cluster Westerlund 1 (Muno et al.\\ 2006). Coherent pulsations with a period of 10.61 s were detected from the 2005 May 22 and June 18 $Chandra$ observations of the cluster. The best-fit periods obtained from above two observations were 10.6112(4)~s and 10.6107(1)~s, which put a limit on the period derivative  of $\\dot{P} < 2 \\times 10^{-10}$~s~s$^{-1}$. The 0.5--8.0~keV $Chandra$ ACIS spectrum of the AXP was equally well described by a blackbody, power-law, or bremsstrahlung continuum model modified by the interstellar absorption (Muno et al.\\ 2006). Analysis with a more complex model of the magnetar atmosphere suggested that the emission likely arises in one or more hot spots covering a small fraction of the surface (Skinner, Perna, and Zhekov, 2006). Search for an infrared counterpart with the Son of ISAAC instrument on the European Southern Observatory (ESO) New Technology Telescope (NTT) yielded negative result within the \\timeform{0.3''} uncertainty in the location of the AXP (Muno et al.\\ 2006).  Following the detection of an intense ($\\sim$10$^{39}$~erg~s$^{-1}$) and short (20~ms) burst from CXOU~J164710.2-455216 with the Burst Alert Telescope (BAT) on Swift on 2006 September 21 (Krimm et al.\\ 2006), the AXP was observed with various X-ray observatories. Chandra Target of Opportunity (TOO) observation of the AXP on 2006 September 27 showed the presence of 10.61069~s pulsations in the ACIS-S event data (Gavriil et al.\\ 2006). Using the Swift, XMM-Newton, and Chandra data, Israel et al.\\ (2007) obtained a phase-coherent solution for the source pulsations after the burst. The solution required an exponential component decaying with a time scale of 1.4~d, which was interpreted to indicate a recovery stage following a glitch with $\\Delta P/P \\sim -10^{-4}$. They also detected a spin-down of $\\dot{P} \\sim 9 \\times 10^{-13}$~s~s$^{-1}$, which implies a magnetic field strength of $10^{14}$~G\\@. Suzaku performed a TOO observation of the AXP CXOU~J164710.2-455216 on 2006 September 23--24. The results obtained from the analysis of the Suzaku observation are presented in this paper. ",
        "conclusions": "The observed changes in the AXP CXOU~J164710.2--455216 before and after the intense and short burst in 2006 September is remarkable. The XMM-Newton observation on 2006 September 22 and the Suzaku observation on 2006 September 23 summaries various changes in the AXP, when compared with the pre-burst data of XMM-Newton on September 16 (Muno et al.\\ 2007). The pulse profile of the AXP was a single peaked before the burst, which changed to a three peaked profile on September 22. These peaks were seen in 0.5--7.0~keV range. Although only the pulse profile below 7~keV was presented in Muno et al.\\ (2007), current Suzaku data show that the pulse profile is single peaked at higher energy band, in 6--12 keV\\@. Because the spectral parameters are comparable between the XMM-Newton data and the Suzaku data, we consider that the nature of the pulse profile, especially the energy dependence, is basically same between these two sets of observations. Though the pulse profiles of some of the AXPs show changes before and after the burst (Kaspi et al.\\ 2003), the change from a single peak profile to a three peak profile is unique to CXOU~J164710.2--455216.  Before the intense and short burst, the spectrum was described by a single blackbody component and the pulse profile was single peaked (Muno et al.\\ 2007). The hardening of the source spectrum (after the burst) was seen from the XMM-Newton observations which was explained by adding a power-law component with an index of $\\sim$2.0 to the blackbody component to describe the spectrum. The pulse profile also changed to a three peaked profile as seen from XMM-Newton and Suzaku observations. However, energy resolved pulse profiles of Suzaku observation shows that the hard X-ray (6--12~keV) pulse profile is similar to that of the pre-burst profile obtained from XMM-Newton whereas the soft X-ray pulse profile is found to be three peaked. This energy dependent pulse profile of Suzaku observation ruled out the model with blackbody and power-law as spectral components and favored a model with two blackbody components. The phase resolved spectroscopy of Suzaku observation also showed that the change in the soft blackbody flux over pulse phase shows two narrow and intense peaks and a third minor peak exactly at the same phases of the peaks in the pulse profile. The second blackbody flux, however, shows only a single peak which is similar to the 6--12~keV pulse profile. These findings suggest that the single blackbody component which described both the spectrum and the pulse profile before the outburst is also required to describe the source properties after the outburst with an increased temperature and hence at hard X-ray band. The three peaked profile, seen after the outburst, is interpreted due to the emergence of a new spectral component which dominates at soft X-rays. This new spectral component may be originated from the giant flares in the AXPs/SGRs during which a significant amount of the magnetic energy is released.  The multi-peaked pulse profiles in AXPs/SGRs are attributed to the presence of multi-pole structures of the external magnetic fields. The rearrangement of these multi-pole magnetic fields changes the shape of the observed pulse profiles in these objects. Muno et al.\\ (2007) suggested that the observed changes in pulse profiles and spectral properties of the AXP before and after the burst are due to a change in the distribution of currents in the magnetosphere. This change in magnetosphere is triggered by a plastic motions in the crust of the neutron star. The observed three-peaked profile in the AXP CXOU~J164710.2--455216 suggests that three magnetic field foot points developed in response to the change in magnetoshepre. Above the foot points, there might be hot plasma columns, like an accretion column of the accretion powered pulsars. The temperatures distribution is hotter temperature at the lower altitude and lower temperature at the higher altitude. The electron-positron plasma within the column cannot move in the direction perpendicular to the magnetic filed. But those can move in the direction parallel to the magnetic field. Therefore, the X-ray photons pass through along with the magnetic field can escape freely. However, the X-ray photon passing in transversal direction to the magnetic field are restricted by the electron positron scattering. In that case, black body radiation from such foot point is beamed into the magnetic filed direction. This can explain the narrow three-peaked profile of the blackbody flux in CXOU~J164710.2--455216."
    },
    "0711/0711.2187.txt": {
        "abstract": "A Quintom universe with an equation-of-state crossing the cosmological constant boundary can provide a bouncing solution dubbed the Quintom Bounce and thus resolve the Big Bang singularity. In this paper, we investigate the cosmological perturbations of the Quintom Bounce both analytically and numerically. We find that the fluctuations in the dominant mode in the post-bounce expanding phase couple to the growing mode of the perturbations in the pre-bounce contracting phase. ",
        "introduction": "For years, it has been suggested that bouncing universe scenarios could provide a solution to the problem of the initial singularity in the evolution of the cosmos, a problem which afflicts both Standard Big Bang cosmology and inflationary universe models \\cite{Guth81,Sato,Steinhardt82,Linde82} \\footnote{An initial singularity occurs in inflationary models if scalar fields are used to generate  the accelerated expansion of space \\cite{Borde}. If inflation arises as a consequence of modifications of the gravitational part of the action, as in Starobinsky's original proposal \\cite{Starob}, it is possible that the evolution is singularity-free.}. In the literature there have been a lot of efforts towards constructing bouncing universes, for instance the Pre-Big-Bang scenario \\cite{PBB} or the cyclic/Ekpyrotic scenario \\cite{Ekp}. Also, in loop quantum cosmology the evolution might be singularity-free \\cite{Bojowald:2001xe}. In Refs. \\cite{Brustein:1997cv,Cartier,Tsujikawa} and \\cite{Biswas:2005qr,Setare:2004jx} models were considered in which the gravitational action is modified by adding high order terms. Since, in general, in these models the equations for cosmological perturbations are very involved, it is difficult to make definitive conclusions about the evolution of fluctuations, and the knowledge of the evolution is crucial in order to compare the predictions of the model with observations.  In some previous work \\cite{Cai:2007qw}, a subset of the present authors considered a bouncing universe model (``Quintom Bounce\") obtained within the standard 4-dimensional Friedmann-Robertson-Walker (FRW) framework, but making use of Quintom matter \\cite{Feng:2004ad}. In this model, it can be shown in \\cite{Cai:2007qw} that the null energy condition (NEC) is violated for a short period around the bounce point. Moreover, after the bounce the equation-of-state (EoS) of the matter content $w$ in our model is able to transit from $w<-1$ to $w>-1$ and then connect with normal expanding history.    A Quintom model \\cite{Feng:2004ad} was initially proposed to obtain a model of dark energy with an EoS parameter $w$ which satisfies $w>-1$ in the past and $w<-1$ at present. This model is mildly favored by the current observational data fitting \\cite{Zhao:2006qg}. The Quintom model is a dynamical scenario of dark energy with a salient feature that its EoS can smoothly cross over the cosmological constant barrier $w = -1$. It is not easy to construct a consistent Quintom model. A no-go theorem proven in Ref. \\cite{Xia:2007km} (also see Refs. \\cite{Feng:2004ad,Vikman:2004dc,Hu:2004kh,Caldwell:2005ai,Zhao:2005vj,Abramo:2005be,Kunz:2006wc}) forbids a traditional scalar field model with Lagrangian of the general form ${\\cal L} = {\\cal L}(\\phi,\\partial_{\\mu}\\phi\\partial^{\\mu}\\phi)$ from having its EoS cross the cosmological constant boundary. Therefore, it is necessary to add extra degrees of freedom with un-conventional features to the conventional single field theory if we expect to realize a viable Quintom model in the framework of Einstein's gravity theory. The simplest Quintom model involves two scalars with one being Quintessence-like and another Phantom-like \\cite{Feng:2004ad,Guo:2004fq}. This model has been studied in detail in \\cite{Quintom_tf}. In recent years there have been a lot of theoretical studies of Quintom models. For example, motivated from string theory, the authors of Ref. \\cite{Cai:2007gs} realized a Quintom scenario by considering the non-perturbative effects of a DBI model. Moreover, there are models which involve a single scalar with higher derivative terms in the action \\cite{Li:2005fm}, models with vector field \\cite{ArmendarizPicon:2004pm}, making use of an extended theory of gravity \\cite{Cai:2005ie}, of ideas from non-local string field theory theory \\cite{Aref'eva:2005fu}, and others (see e.g. \\cite{Quintom_1,Onemli:2004mb,Quintom_others}). Because of its brand new properties, Quintom models give many unexpected predictions. Most important for the subject of the present paper is that a universe dominated by Quintom matter can provide a bouncing cosmology which allows us to avoid the problem of the initial singularity. This scenario, developed in \\cite{Cai:2007qw}, is called  the Quintom Bounce. In this scenario, the EoS crosses the cosmological constant boundary twice around the bounce point.   In order to be able to compare the predictions of a Quintom bounce with observations, it is crucial to investigate the evolution of cosmological perturbations. In singular bouncing models the evolution of fluctuations is not under control classically. The fluctuations diverge at the bounce point, firstly invalidating the applicability of the theory of linear cosmological perturbations, and secondly making it impossible to reliably connect the fluctuations in the contracting phase with those in the expanding phase. In a Quintom bounce model the evolution of fluctuations is well-behaved. Our first main result is that cosmological perturbations can evolve smoothly from the contracting to the expanding phase. We show that the commonly used variable $\\Phi$, the metric perturbation in longitudinal gauge, passes smoothly both through the bounce point and through the points where the cosmological constant boundary of the EoS is crossed. Interestingly, we find that our result for the transfer of fluctuations through the bounce possesses features found both in some previous analyses of fluctuations in non-singular cosmologies, and features of the evolution in singular bouncing models. For small comoving wave numbers one finds an evolution similar to what occurs in singular bouncing cosmologies (the growing mode of $\\Phi$ in the contracting phase couples only to the decaying mode in the expanding phase), whereas the transfer of the perturbations is very different in the region of large comoving wave numbers (the growing mode in the contracting phase couples to the dominant mode in the phase of expansion). The physical reason is easy to see: small $k$ modes enter the sub-Hubble region very close to the bounce point and have no time to complete a single oscillation, whereas large $k$ modes undergo many oscillations which are in the bounce phase at sub-Hubble scales. Thus, whereas the large $k$ modes feel the resolution of the singularity, the small $k$ modes evolve as if there still were an abrupt transition between the contracting to the expanding phase. We will follow the fluctuations both numerically and analytically. In the analytical analysis, we apply matching conditions away from the bounce point, in regions where both the background and the perturbations can be matched consistently via the Deruelle-Mukhanov \\cite{Deruelle} (initially derived in \\cite{Hwang}) conditions \\footnote{For a discussion of the dangers in applying the matching conditions at the transition point between contraction and expansion see \\cite{Durrer}.}.   This paper is organized as follows. In Section 2, we review the equations of motion of double-field Quintom and derive the perturbation equations for this model. In Section 3, we study the behavior of a specific Quintom Bounce model systematically. In Subsection 3.1, we determine the complete evolution of the background using numerical integration. In Subsection 3.2, we investigate the perturbations, first analytically by solving the equations in separate periods by means of approximations valid during those periods, and matching the solutions at the transition points, and finally numerically. The numerical results support our analytical calculations consistently for all times. Section 4 contains discussion and conclusions. ",
        "conclusions": "Bouncing non-singular cosmologies, with an initial contracting phase which lasts until to a non-vanishing minimal radius is reached and then smoothly transits into an expanding phase, provide a possible solution to the singularity problem of Standard Big Bang cosmology, a problem which is not cured by scalar-field-driven inflationary models.   In this paper, we have considered the evolution of cosmological fluctuations in a particular bouncing cosmology, namely the Quintom Bounce. In this model, space-time continues to be described by Einstein's General Relativity. The violations of the null energy condition required to get a bounce are obtained by making use of Quintom matter, which allows a transition of the EoS $w$ through the cosmological constant boundary. We have studied the evolution of the gravitational potential both analytically and numerically in a particular model of a Quintom Bounce in which matter consists of two scalar fields, one of them a Quintessence-like field, the other a Phantom-like field. We have found that the evolution of the gravitational potential $\\Phi$ is completely non-singular. Whereas in models with a singular bounce such as the Pre-Big-Bang \\cite{PBB} and the Ekpyrotic \\cite{Ekp} scenarios, the perturbations must be matched across a singular space-like surface, which leads to serious conceptual problems \\cite{Durrer}, in our case no such problems arise \\footnote{Recently, a ghost condensation mechanism to smooth out the singularity in the Ekpyrotic scenario has been introduced \\cite{Senatore,Khoury}, and it has been shown in this context, using the techniques developed in \\cite{Riotto,Fabio3,Turok} that making use of iso-curvature modes, the dominant spectrum in the contracting phase can be passed on to the dominant spectrum in the expanding phase.}. In an analytical approximation scheme, we can apply matching conditions at non-singular surfaces which are then well justified, and we can also study the evolution numerically without any approximations.  Our second main result is that for short wavelength modes, modes which re-enter the Hubble radius in the slow contracting phase, the dominant fluctuation mode is determined by the growing mode in the contracting phase. This result is similar to what has recently been found in other non-singular bouncing cosmologies \\cite{Hassan,Stephon,Abramo}. The first of these is a bounce in the context of mirage cosmology, the second in the context of a specific higher derivative gravity model, and the third in the context of K-essence. Earlier work on fluctuations in models with a kinetic term with opposite sign generating a nonsingular bounce is in \\cite{Peter}. The authors of \\cite{Peter} observed that the way the post-bounce spectrum depends on the pre-bounce spectrum depends sensitively on the details of the bounce. In this light, it is not surprising that other bounces show that the dominant mode in the pre-bounce phase does not couple to the dominant mode in the post-bounce phase \\cite{Cartier,Tsujikawa,Bozza,Wands,Finelli2,Veneziano}. For long wavelength modes, however, we find that the dominant fluctuation mode in the contracting phase at leading order in $k$ only sources the decaying mode in the expanding phase, a result similar to what is obtained in singular bouncing models \\cite{BGGMV,Lyth,Fabio,Hwang2,Creminelli}.  Thus, the fluctuations in the Quintom Bounce model evolve in way which combines aspects found in some other recently studied non-singular bouncing cosmologies \\cite{Hassan,Stephon} (the behavior of short wavelength modes) with aspects found in singular bouncing cosmologies (the behavior of the long wavelength modes). Which category of modes determines scales accessible to current cosmological observations depends on the detailed parameters of the model. Our model may help us to understand the physics of singular bounce and non-singular bounce more clearly. As a side remark, let us add that in the Quintom Bounce model there is no trans-Planckian problem \\cite{RHBrev2,Martin} for fluctuations since, unless the periods of slow-climb contraction and slow-roll expansion last many Hubble expansion times, the wavelengths of all the observable perturbations will be larger than the Planck length at the bounce point and thus never enter the zone of trans-Planckian ignorance.   Let us close with some more general comments related to the Quintom Bounce. Although this theory is still in the stage of development, there have been works done in the context of a bouncing scenario related to predictions for observations. For example, in \\cite{Piao:2003zm} (see also \\cite{Piao:2003hh,Piao:2005ag}) it was pointed out that with a bounce followed by slow-roll inflation it is possible to give a reasonable explanation for the suppression of the low multi-poles of the CMB anisotropies. In \\cite{Cai:2007qw} and \\cite{Aref'eva:2007uk} the possibility of obtaining a bounce in a Quintom model with a single matter field with higher-derivative terms was explored. See also \\cite{Zhang:2007bi,Wei:2007rp} for other works."
    },
    "0711/0711.0903_arXiv.txt": {
        "abstract": "We investigated the characteristics of the shallow decay phase in the early X-ray afterglows of GRBs observed by $Swift$ X-Ray Telescope (XRT) during the period of January 2005 to December 2006. We found that the intrinsic  break time at the shallow-to-normal decay transition in the X-ray light curve $T^0_{\\rm brk}$ is moderately well correlated with the isotropic X-ray luminosity in the end of the shallow decay phase ($L_{X, \\rm end}$) as $T^0_{\\rm brk} = (9.39\\pm0.64) \\times 10^3 {\\rm s} ({L_{X,\\rm end}}/{10^{47} {\\rm ergs\\cdot s^{-1}}})^{-0.71\\pm0.03}$, while $T^0_{\\rm brk}$ is weakly  correlated with the isotropic gamma-ray energy of the prompt emission $E_{\\gamma, \\rm iso}$. Using  $T^0_{\\rm brk}-L_{X, \\rm end}$ relation we have determined the pseudo redshifts of 33 GRBs. We compared the pseudo redshifts of 11 GRBs with measured redshifts and found the rms error to be 0.17 in $\\log z$. From this pseudo redshift, we estimate that $\\sim 15\\%$ of the $Swift$ GRBs have $z > 5$. The advantages of this distance indicator is that (1) it requires only  X-ray afterglow data while other methods such as Amati and Yonetoku correlations require the peak energy ($E_p$) of the prompt emission, (2) the redshift is uniquely determined without redshift degeneracies unlike the Amati correlation, and (3) the redshift is estimated in advance of deep follow-ups so that possible high redshift GRBs might be selected for detailed observations. ",
        "introduction": "The optical afterglow light curves of most GRBs show the smooth power-law decays with time ($\\propto t^{-\\alpha}$ with a typical index value $\\alpha \\sim 1$). This is consistent with the prediction of the simplest model of a spherical blast wave propagating into a uniform medium, where the spectrum consists of several power-law segments in which $F_{\\nu}\\propto\\nu^{-\\beta}t^{-\\alpha}$ (Sari, Piran \\& Narayan 1998).  However, after the advent of the $HETE$-2 satellite, prompt localization of GRBs sent to ground-based telescopes made it possible to observe the early afterglows which revealed that some GRBs show deviations from a smooth power-law light curve. For example, GRB 021004 showed a highly variable light curve, with several bumps and wiggles from a simple power-law (e.g., Fox et al. 2002; Mirabal et al. 2002). GRB 030329 also showed such a short timescale variability (Uemura et al. 2003; Sato et al. 2003). Furthermore, the early slow decay was observed in GRB 021004 and was interpreted by Fox et al. (2003) as arising from delayed shocks and continuous energy ejection from the central engine (Rees \\& Meszaros 1998; Kumar \\& Piran 2000; Sari \\& Meszaros 2000).  The early X-ray afterglow was found to be even more complex by the $Swift$ observations. They generally consist of three distinct power-law segments (Nousek et al. 2006; Zhang et al. 2006): (1) an initial steep decay with $\\alpha_1 \\sim 3$, (2) a shallow decay with $\\alpha_2 \\sim 0.5$ and finally (3) a normal decay $\\alpha_3 \\sim 1$, where $\\alpha_1$, $\\alpha_2$ and $\\alpha_3$ are power-low indices of temporal variability ($\\alpha$ of $t^{-\\alpha}$). The initial steep decay is most likely the transition from the prompt emission to the afterglow (Nousek et al. 2006; Zhang et al. 2006). Such a steep decay is due to the ``curvature effect'' of the high-latitude emission expected for the emission ceasing abruptly (e.g., Kumar \\& Panaitescu 2000; Dermer 2004; Dyks et al. 2005; Panaitescu et al. 2006; Yamazaki et al. 2006). In the normal decay phase, the data for many bursts are consistent with an ISM medium rather than a wind medium. However, for the shallow decay phase, its physical background has not been understood and it is the most mysterious feature in the early X-ray afterglows.  As for shallow decay phase, Willingale et al. (2007) have studied the end-time of the shallow decay of the X-ray light curves of long GRBs, $T_{\\rm a}$. They suggested a possibility that $T_{\\rm a}$ depends on the total energy of the outflow. Nava et al. (2007) have investigated the correlation between the end-time in the GRB frame $T_{\\rm a}^{0} \\equiv T_{\\rm a}/(1 + z)$ and the isotropic gamma-ray energy of the prompt emission $E_{\\gamma, \\rm iso}$. They found that for the bursts in their sample, $T_{\\rm a}^0$ is weakly correlated with $E_{\\gamma,\\rm iso}$. However, this correlation disappears when considering all bursts of known redshift and $T_{\\rm a}$. Liang et al. (2007) have also showed that there is no significant correlation between the break time between shallow and normal decay segments and the $E_{\\gamma, \\rm iso}$.  In this paper we investigated the characteristics of the shallow decay phase in the early X-ray afterglows of GRBs observed by $Swift$ X-Ray Telescope (XRT) during the period of January 2005 to December 2006. We found that the intrinsic break time at the shallow-to-normal decay transition in the X-ray light curve $T^0_{\\rm brk}$ is correlated with the isotropic X-ray luminosity of the shallow decay phase, $L_{\\rm X}$. We tried to apply this relation to determine the redshifts of 33 GRBs and found that the distribution of the pseudo redshifts is similar to that of spectroscopically measured redshifts with more high pseudo redshift GRBs. Furthermore, we independently examine if there are any correlations among parameters of the prompt emission and the shallow decay phase for GRBs. We found that $T^0_{\\rm brk}$ is weakly correlated with the isotropic gamma-ray energy of the prompt emission $E_{\\gamma, \\rm iso}$.  Throughout this paper, we adopt the cosmological parameters $H_0 = 71$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm M} = 0.27$ and $\\Omega_{\\Lambda} = 0.73$. ",
        "conclusions": "Since $T^0_{\\rm brk} - E_{\\gamma, \\rm iso}$ correlation is weak we here argue mainly the physical implications of $T^0_{\\rm brk} - L_{X, \\rm end}$ correlation as described by $T^0_{\\rm brk}\\propto L_{X, \\rm end}^{-0.7}$. Here we note it is reasonable that $T^0_{\\rm brk} - L_{X, \\rm end}$ correlation is better than $T^0_{\\rm brk} - L_{X, \\rm ini}$ correlation since in the shallow decay phase more energy is emitted in the end phase than the initial phase. Interestingly this correlation is similar to the burning time $t_H$ and the luminosity $L$ relation of the hydrogen main sequence star ($t_H\\propto L^{-0.7}$).  From the theory of stellar evolution  $L$ and $t_H$ mainly depend on the mass of the star $M$ as \\begin{eqnarray} L\\propto M^3,\\\\ t_H\\propto \\frac{M}{L}\\propto M^{-2}. \\end{eqnarray} Eliminating $M$ from above two equations we have $t_H\\propto L^{-0.67}$.  \\subsection{Energy Injection Model}  The energy injection model (Nousek et al. 2006; Zhang et al. 2006; Granot \\& Kumar 2006) is thought to be that energy is injected continuously into the external shock so that the flux decay becomes slower than the usual $\\propto t^{-1}$. The injection may be caused by (1) a long-lived central engine or (2) a short-lived central engine ejecting shells with some range of Lorentz factors.  First we consider (1), the long-lived central engine model. This scenario requires the central engine to remain active until the end of the shallow decay phase ($T^0_{\\rm brk}$), which is in many cases $1,000-10,000$ s. Since the X-ray afterglow is a good indicator of the kinetic energy, the central engine injects energy with the kinetic luminosity proportional to the X-ray luminosity $L_{\\rm kin} \\propto L_{\\rm X} \\propto t^{-\\alpha_1}$. The kinetic energy ($E_{\\rm kin}$) of the afterglow increases as a function of time $\\propto t^{1-\\alpha_1}$, and the observed correlation suggests that the total kinetic energy is anti-correlated with the lifetime of the central engine, \\begin{eqnarray} E_{\\rm kin,end} \\propto L_{\\rm X,end} T_{\\rm brk}^{0} \\propto \\left(T_{\\rm brk}^{0}\\right)^{-\\frac{1}{0.71}+1} \\sim \\left(T_{\\rm brk}^{0}\\right)^{-0.41}. \\end{eqnarray} where $E_{\\rm kin, end}$ is the kinetic energy of afterglow in the end of the shallow decay phase.   Next we consider (2) a short-lived central engine with some range of Lorentz factors of ejected shells. After the internal shocks, shells are rearranged so that outer shells are faster and inner shells are slower. This configuration may also occur if the central engine ejects faster shell earlier. Outer shells are slowed down by making the external shock. Once the Lorentz factor of the leading shocked shell drops below that of a following slower shell, the slower shell catches up with the shocked shell, injecting energy into the forward shock. Since the Lorentz factor of the afterglow is proportional to $\\Gamma \\propto E^{1/8} T^{-3/8} n^{-1/8}$, the observed correlation suggests that the Lorentz factor of the slowest shell, which has almost all energy, is nearly proportional to the total kinetic energy, \\begin{eqnarray} \\Gamma_{\\rm slow} \\propto E_{\\rm kin,end}^{1.0} n^{-1/8}, \\label{eq:ga} \\end{eqnarray} where $n$ is the ambient density.  \\subsection{Inhomogeneous Jet Model}  In the inhomogeneous jet model, it is assumed that we observe more energetic components in the GRB jet at later times as the external shock decelerates and the visible region increases. The shallow decay phase is produced by the superposition of the afterglow emission from the off-axis components. This phase ends when the whole jet is observed in the ring-shaped jet model (Eichler \\& Granot 2006) or when the mini-jets merge and the inhomogeneities are averaged out in the multiple mini-jets model (Toma et al. 2006).  In the ring-shaped jet model, the X-ray luminosity at the end of the shallow phase is given by $L_{\\rm X, end} \\propto E_{\\rm kin, iso} (T^0_{\\rm brk})^{-1}$, and the opening angle of the whole jet is determined by $\\theta_j \\propto E_{\\rm kin, iso}^{-1/8} (T^0_{\\rm brk})^{3/8} n^{1/8}$, where $E_{\\rm kin, iso}$ is the isotropic kinetic energy of the whole jet. Thus the observed correlation suggests $\\theta_j \\propto E_{\\rm kin, iso}^{-1.0} n^{1/8}$. Since the collimation-corrected kinetic energy is calculated as $E_{\\rm kin} \\propto E_{\\rm kin,iso} \\theta_j^2$, we obtain $E_{\\rm kin} \\propto \\theta_j  n^{1/8}$.   \\subsection{Time-dependent Microphysics Model}  This model considers that the microphysical parameters, such as the energy fraction that is shared to electrons $\\epsilon_{\\rm e}$ and magnetic field $\\epsilon_{\\rm B}$, depend on time (Ioka et al. 2006; Fan \\& Piran 2006; Panaitescu et al. 2006b). We usually assume that the micro-physical parameters do not vary and in fact, constant $\\epsilon_{\\rm e}$ and $\\epsilon_{\\rm B}$ are consistent with the observation of late time afterglows (Yost et al. 2003). However, the behavior of these parameters in the early time afterglow is not yet known.  In the model, the microphysical parameters vary in the early afterglow. After reaching the equipartition value, the microphysical parameters remain constant as observed in the late time afterglow. The X-ray luminosity $L_{\\rm X}$ is given by the bolometric kinetic luminosity $L$ as $L_{\\rm X}\\sim\\epsilon_{\\rm e} L$. Since $L\\propto E_{\\rm kin}t^{-1}$, the shallow X-ray light curve $L_{\\rm X} \\propto t^{-1/2}$ suggests that $\\epsilon_{\\rm e}$ evolves as $\\epsilon_{\\rm e} \\propto t^{1/2}$ (Ioka et al. 2005). The observed correlation suggests that the saturation occurs at the Lorentz factor $\\Gamma_{\\rm sat} \\propto E_{\\rm kin,end}^{1.0} n^{-1/8}$ as in equation (\\ref{eq:ga}).   \\subsection{Pulsar Model}  Troja et al. (2007) showed that the abrupt drop of the X-ray light curve observed in GRB 070110 cannot be explained by an external shock as the origin of the shallow decay phase and implies the long-lived central engine. Furthermore, they suggest that the shallow decay phase might be powered by a spinning down central engine, possibly a millisecond pulsar. Motivated by this suggestion let us consider that the slowly changing usual shallow decay is also due to the pulsar activity. If the dipole magnetic field is constant, the luminosity of the pulsar decreases as $\\propto t^{-2}$, so that we need the increase of the dipole magnetic field to interpret the shallow decay phase. From the total energy conservation, we have \\begin{eqnarray} \\frac{d}{dt}\\left(\\frac{1}{2}I\\Omega^2+\\frac{B^2R^3}{6}\\right)=-L(t), \\end{eqnarray} where we assume that the rotational energy goes into the magnetic energy by some unknown mechanism satisfying the observed luminosity, \\begin{eqnarray} L(t)=\\frac{B^2\\Omega^4R^6}{6c^3} =L_0 \\left(\\frac{t}{t_0}\\right)^{-\\alpha_1}, \\label{eq:dipole} \\end{eqnarray} Integrating the above equations from the beginning of the shallow decay phase $t_0$ to $T^0_{\\rm brk}$, we have \\begin{equation} \\frac{1}{2}I(\\Omega^2-\\Omega_0^2)+\\frac{(B^2-B^2_0)R^3}{6} =-L_{\\rm X,end} T^0_{\\rm brk} \\frac{1}{1-\\alpha_1} \\left[1-\\left( \\frac{t_0}{T^0_{\\rm brk}}\\right)^{1-\\alpha_1}  \\right] \\end{equation} In this model $T^0_{\\rm brk}$ is essentially determined by the total energy conservation as \\begin{equation} \\frac{1}{2}I\\Omega_0^2\\sim L_{X, \\rm end}T^0_{\\rm brk}/(1-\\alpha_1)=E_s. \\end{equation} Therefore the energy of the shallow decay phase $E_s$ is essentially the total energy of the initial pulsar. Figure \\ref{fig:Ex-Tbrk} suggests that the initial rotational period is 1ms$\\sim$ 3ms, which is an appropriate value in this model. Using equation (\\ref{eq:dipole}) and $T^0_{\\rm brk}$-$L_{X, \\rm end}$ relation we have \\begin{equation} B_0^2\\Omega_0^4\\propto \\left(\\frac{T^0_{\\rm brk}}{t_0}\\right)^{\\alpha_1}(T^0_{\\rm brk})^{-\\frac{1}{0.71}}. \\end{equation} Since $\\Omega_0$ is known from $E_s$, the above equation tells us that the initial strength of the magnetic field determines $T^0_{\\rm brk}$."
    },
    "0711/0711.2906_arXiv.txt": {
        "abstract": "A number of theories, spanning a wide range of mass scales, predict dark matter candidates that have lifetimes much longer than the age of the universe, yet may produce a significant flux of gamma rays in their decays today.  We constrain such late decaying dark matter scenarios model-independently by utilizing gamma-ray line emission limits from the Galactic Center region obtained with the SPI spectrometer on INTEGRAL, and the determination of the isotropic diffuse photon background by SPI, COMPTEL and EGRET observations.  We show that no more than $\\sim$5\\% of the unexplained MeV background can be produced by late dark matter decays either in the Galactic halo or cosmological sources. ",
        "introduction": "Dark matter continues to live up to its name~\\cite{Zwicky:1933gu}, despite accumulated evidence of its existence from observations of large-scale structure formation, galaxy cluster mass-to-light ratios, and galactic rotation curves.  An attractive approach towards revealing dark matter's particle identity is to search for its signature in radiation backgrounds, either from the Milky Way or in the isotropic diffuse photon background (iDPB), which can contain both cosmological and Galactic halo contributions.  Dark matter might be comprised of particles that can decay with finite lifetimes much longer than the age of the universe. In such scenarios, the resultant fluxes of decay products depend on the amount of dark matter present alone, as opposed to self-annihilation, which, being dependent on particle density squared, is very sensitive to assumptions concerning details of dark matter clustering.  A wide variety of decaying dark matter models have been examined in regards to their observable implications~\\cite{decays}. Among late decaying dark matter models, sterile neutrinos with multi-keV masses have been extensively studied as dark matter candidates~\\cite{sterile}, with strong constraints placed on their decays, e. g. Refs.~\\cite{sterile,Watson:2006qb,Yuksel:2007xh} and references therein.  The decay of moduli dark matter~\\cite{moduli} with masses of several hundred keV may contribute to the sub-MeV iDPB. The dark matter model of Ref.~\\cite{Cembranos:2006gt}, inspired from minimal universal extra dimensions or supersymmetry~\\cite{muedsusy}, with a mass scale of hundreds of GeV, is advocated as the source of the iDPB in the MeV range~\\cite{Cembranos:2007fj}, which has yet to be accounted for with conventional sources (e.g., supernovae~\\cite{snia} or active galactic nuclei~\\cite{agn}) or more exotic mechanisms~\\cite{exotic}.  Similarly, decaying gravitino dark matter in R-parity breaking vacua~\\cite{rparity}, with multi-GeV masses, has been suggested as an explanation of iDPB spectral features in the GeV range~\\cite{Ibarra:2007wg,Bertone:2007aw}.     \\begin{figure}[b] \\includegraphics[width=3.25in,clip=true]{constraints} \\caption{Model-independent constraints on the product of mass and lifetime, $ m_\\chi \\tau$, versus the energy carried away by the monochromatic photon emission, $\\varepsilon$, for a generic late-decaying dark matter model: $\\chi \\rightarrow \\chi' + \\gamma$.  Regions excluded by either the gamma-ray line emission limits from the Galactic Center region or overproduction of the isotropic diffuse photon background are shown, together with preferred ranges of parameters from three well-studied models.} \\label{fig:constraints} \\end{figure}  Rather than focusing on a particular model, we first consider a generic decaying dark matter scenario in which the decay of the parent particle is dominated by a monochromatic photon emission. We assume that the lifetime of the parent particle, $\\tau$, is much longer than the age of the universe ($\\tau_0\\simeq4.5 \\times 10^{17}$~s), thus its cosmological abundance has not changed significantly since the time of dark matter decoupling. The decay under consideration is $\\chi \\rightarrow \\chi' + \\gamma$, where $\\gamma$ is a  monochromatic photon emitted with energy $\\varepsilon$. In general, $\\tau$ and $\\varepsilon$ will depend on the masses of the parent and the daughter particles ($m_{\\chi}$, $m_{\\chi'}$) and their splitting, $\\Delta m = m_{\\chi}-m_{\\chi'}$.   The flux of photons from dark matter decays is inversely proportional to both the particle lifetime (fixing the decay rate per particle as specified by a particular theoretical model) and the mass of an individual particle (yielding the total number of particles in a fixed amount of dark matter). Thus, gamma-ray observations allow us to place constraints only on the degenerate product $m_\\chi \\tau$ versus $\\varepsilon$, as we display in Fig.~\\ref{fig:constraints}. As we will discuss in detail, below the jagged line between 0.02--8~MeV, the gamma-ray line signal from the Galactic Center (GC) region due to dark matter decays violates the corresponding limit obtained with the SPI spectrometer on INTEGRAL satellite~\\cite{Teegarden}.  Additionally, the iDPB, as determined from SPI~\\cite{Churazov:2006bk}, COMPTEL~\\cite{Weidenspointner} and EGRET~\\cite{Sreekumar:1997un,Strong:2004ry} data, is overproduced (assuming it is fully accounted by late dark matter decays in a given energy band) in the triangular region, even disregarding any contributions from known astrophysical sources.  We also show three representative scenarios, inspired by the theories of sterile neutrinos, R-parity breaking vacua, and mUED.  Since $m_{\\chi}$, $m_{\\chi'}$ and $\\Delta m$ are not necessarily predetermined, they may be adjusted to yield the displayed curves relating $ m_\\chi \\tau$ and $\\varepsilon$.   \\begin{figure}[b] \\includegraphics[width=3.25in,clip=true]{dpb} \\caption{\\textbf{Top:} Limits on the diffuse gamma-ray line emission from the Galactic Center region (an angular region within a $13^{\\circ}$ radius) as adopted from Ref.~\\cite{Teegarden}. \\textbf{Bottom:} Representative measurements of the diffuse photon background from SPI~\\cite{Churazov:2006bk}, COMPTEL~\\cite{Weidenspointner} and EGRET~\\cite{Strong:2004ry} in the energy range around 0.01~MeV--100~GeV. The thick solid line, summarizing the overall trend of the data, is to be compared to predictions of decaying dark matter scenarios. } \\label{fig:dpb} \\end{figure} ",
        "conclusions": "Predictions of photon fluxes from dark matter decays are considerably more robust than those from annihilations, due to a lesser dependence upon theoretical uncertainties in the distribution and clustering of dark matter.  We have shown that the gamma-ray line emission limits from the Galactic Center region, along the isotropic diffuse photon background, allow for stringent constraints to be placed on late decaying dark matter scenarios that produce monoenergetic photons. We emphasize that the Galactic and cosmic constraints are not independent of each other, with the GC region providing stronger limits in its range of applicability due to new spectroscopic data. Rather than attempting to explain various gamma-ray phenomena with a specific model, we report model-independent bounds on the decaying dark matter parameter space (as defined by $ m_\\chi \\tau$ versus $\\varepsilon$).  Our general constraints are applicable to a number of models, and can be used as a guide for future model building.  Upcoming gamma-ray telescopes with improved energy resolution, such as GLAST~\\cite{Gehrels:1999ri} or ACT~\\cite{Boggs:2006mh}, can improve upon these bounds, particularly by making use of the unique spectral shape and directionality of decays from the Galactic halo.  One interesting application of our study is to assess the recent suggestion that cosmological late dark matter decays can explain the isotropic diffuse photon background in the 1-5~MeV range, whose origin remains a mystery~\\cite{Cembranos:2007fj}.  We plot $m_\\chi \\tau$ versus $\\varepsilon$ in Fig.~\\ref{fig:parameters} for ${m_\\chi}/{b} \\simeq 300$~GeV  and ${m_\\chi}/{b} \\simeq 50$~GeV to represent the aforementioned mUED and SUSY models of Ref.~\\cite{Cembranos:2007fj}, respectively.  We also show the range of parameters, $m_\\chi \\tau$ versus $\\varepsilon$, that can lead to a substantial ($>10$\\%) contribution to the iDPB (shaded diagonal band) or overproduce them (triangular region) through the sum of the local decays (Galactic halo) or decays from truly cosmological sources (all distant dark matter halos). The region excluded by the gamma-ray line emission limits from the GC region is below the jagged line. As seen in the figure, even the combined emission from the Galactic halo and cosmological sources due to either the mUED or SUSY inspired decaying dark matter models of Ref.~\\cite{Cembranos:2007fj} cannot make a significant contribution to the iDPB while abiding by the gamma-ray line emission limits from the GC region.  The mUED model can contribute to the iDPB only for $\\Delta m \\lesssim 1.5$~MeV with a contribution of $\\lesssim $ 5\\%, while the SUSY model is even more severly constrained.  Even relaxing our assumptions on the distribution of dark matter in the halo does not increase these fractions dramatically, thus, dark matter cannot decay in the late universe at a high enough rate to make a prominent contribution to the iDPB in the MeV range.   We thank Shmuel Nussinov, Gary Steigman, Louie Strigari, Ak{\\i}n Wingerter and especially John Beacom for fruitful discussions. HY is supported by National Science Foundation under CAREER Grant PHY-0547102 to JB; MDK is supported by the Department of Energy grant DE-FG02-91ER40690; both by CCAPP at OSU."
    },
    "0711/0711.0418_arXiv.txt": {
        "abstract": "We report on the first global three-dimensional (3D) MHD simulations of disk accretion onto a rotating magnetized star through the Rayleigh-Taylor instability. The star has a dipole field misaligned relative to the rotation axis by a small angle $\\Theta$. Simulations show that, depending on the accretion rate, a star may be in the stable or unstable regime of accretion. In the unstable regime, matter penetrates deep into the magnetosphere through several elongated ``tongues\" which deposit matter at random places on the surface of the star, leading to stochastic light-curves. In the stable regime, matter accretes in ordered funnel streams and the light-curves are almost periodic. A star may switch between these two regimes depending on the accretion rate and may thus show alternate episodes of ordered pulsations and stochastic light-curves. In the intermediate regime, both stochastic and ordered pulsations are observed. For $\\Theta > 30^\\circ$, the instability is suppressed and stable accretion through funnel streams dominates. ",
        "introduction": "\\begin{figure*}[t] \\epsscale{1.6} \\plotone{f1.eps} \\caption{(a) An example of stable accretion. The surface is a constant density surface, and the lines are sample magnetic field lines. The magnetic axis is misaligned relative to the rotational axis by $\\Theta=15^\\circ$. (b) An example of unstable accretion for $\\Theta=5^\\circ$. The lower panels show light curves from the hot spots. Time is measured in orbital periods at $r=1$. Only the inner part of the simulation region is shown.} \\label{Figure 1} \\end{figure*}  Magnetospheric accretion occurs in different stars, for example in classical T Tauri stars (CTTSs), which are the progenitors of Solar-type stars (e.g., Hartmann 1998; Bouvier et al. 2007), in magnetized white dwarfs in some cataclysmic variables (e.g., Warner 1995; Warner \\& Woudt 2002), and in accreting millisecond pulsars, which are weakly magnetized neutron stars (e.g., van der Klis 2000). The matter close to the star may either flow around the magnetosphere, forming ordered funnel streams going towards the magnetic poles of the star (e.g., Ghosh \\& Lamb 1978; Camenzind 1990; K{\\\"o}nigl 1991), or may accrete directly through the magnetosphere due to the Rayleigh-Taylor (RT) instability (e.g., Arons \\& Lea 1976; Elsner \\& Lamb 1977; Scharlemann 1978). Earlier 2D and 3D numerical simulations have shown accretion through funnel streams (Koldoba et al. 2002; Romanova et al. 2002, 2003, 2004). Recent 3D MHD simulations performed for a {\\it wider} range of parameters have shown that in many cases  the disk-magnetosphere boundary is RT unstable.   Previous theoretical models and non-global simulations gave useful but restricted analyses of this problem. Arons and Lea (1976) investigated magnetospheric accretion  through the RT instability in the non-rotating case with a spherical accretion geometry. Spruit and Taam (1990) investigated the stability of an infinitely thin rigidly rotating disk. The RT instability in magnetized disks was studied by Kaisig et al. (1992); Spruit et al. (1995); Lubow \\& Spruit (1995). Li and Narayan (2004) investigated disk-magnetosphere interaction in the case of an infinitely thick disk with a vertical magnetic field. 2D simulations have been performed by Wang and Nepveu (1983) and Wang \\& Robertson (1985), while Rast{\\\"a}tter and Schindler (1999) performed both 2D and 3D simulations, but in a patch (see also Stone and Gardiner 2007). Such simulations, while shedding light on many important features of the instability, do not take into account many factors which are present in global simulations, such as the possibility of matter flowing through funnel streams to the magnetic poles of the star.  In this paper we show results from global 3D MHD simulations, where the simulation region includes the disk and the whole magnetosphere of the star. This paper summarizes these new results, while in the following paper (Kulkarni \\& Romanova 2008, hereafter - KR08) we present results for a larger range of parameters. ",
        "conclusions": ""
    },
    "0711/0711.2071_arXiv.txt": {
        "abstract": "We find that all classes of galaxies, ranging from disks to spheroids and from dwarf spheroidals to brightest cluster galaxies, lie on a two dimensional surface within the space defined by the logarithms of the half-light radius, $r_e$, mean surface brightness within $r_e$, $I_e$, and internal velocity, $V^2 \\equiv ({1\\over 2}v_c^2 + \\sigma^2)$, where $v_c$ is the rotational velocity and $\\sigma$ is the velocity dispersion. If these quantities are expressed in terms of kpc, $L_\\odot$ pc$^{-2}$, and km s$^{-1}$, then this surface is described by the equation $ \\log r_e -  \\log V^2 + \\log I_e + \\log \\Upsilon_e + 0.8 = 0$, where we provide a fitting function for $\\Upsilon_e$, the mass-to-light ratio within $r_e$ in units of $M_\\odot/L_\\odot$, that depends only on $V$ and $I_e$. The scatter about this surface for our heterogeneous sample of 1925 galaxies is small ($<$ 0.1 dex), and both the scatter within one of the galaxy subsamples  (1319 disks) and the analysis of subsamples with independently derived mass-to-light ratios suggest that the intrinsic scatter could be as low as $\\sim 0.05$ dex, or 10\\%, prior to any correction for observational errors.  This small scatter has three possible implications for how gross galactic structure is affected by internal factors, such as stellar orbital structure, nuclear activity, or mass loss history, and by external factors, such as environment or accretion history. These factors either 1) play no role beyond generating some of the observed scatter, 2) move galaxies along the surface, or 3) balance each other to maintain this surface as the locus of galactic structure equilibria. We cast the behavior of $\\Upsilon_e$ in terms of the fraction of baryons converted to stars, $\\eta$, and the concentration of those stars within the dark matter halo, $\\xi \\equiv R_{200}/r_e$, where $R_{200}$ is the standard estimate of the virial radius. We derive expressions for $\\eta$ and $\\xi$, use an independent measurement of $\\eta$ to evaluate leading constant terms, obtain $\\eta = 1.9 \\times 10^{-5} (L/L^*)\\Upsilon_* V^{-3}$ and $\\xi = 1.4 V r_e^{-1}$, and relate these to each other via $\\log \\eta + \\log \\xi = -\\log \\Upsilon_e + \\log \\Upsilon_* + const$. Finally, we present the distributions of $\\eta$ and $\\xi$ for the full range of galaxies and conclude that the high $\\Upsilon_e$'s of dSphs are due to low $\\eta$ rather than any differences in $\\xi$, that $\\eta$ is similar for spheroids and disks of a given $V$, and that $\\eta$ decreases with increasing $V$ for systems with $V > 30$ km sec$^{-1}$. For systems with internal velocities comparable to that of the Milky Way ($149 < V <  163$ km s$^{-1}$) , $\\eta = 0.14 \\pm 0.05$, and $\\xi$ is, on average, $\\sim$ 5 times greater for spheroids than for disks. ",
        "introduction": "\\label{sec:intro}  Are galaxies fundamentally a simple family of collapsed objects, whose gross structure is describable by a few basic parameters, or are they highly complex systems whose structural properties are determined by a myriad of internal and external factors?  If the former, there must be an analogous construct to the stellar Hertzprung-Russell diagram that testifies to deep, systemic structural patterns among galaxies and serves as a guide to a simple, if not entirely complete, analytic description of galactic structure. The study of stellar structure offers a beautiful example of the power of reductionism in astrophysics. By focusing on the HR diagram, investigators solved the problem of stellar structure without needing to address other unsolved problems, such as the origin of the initial mass function. The observation that the position of main sequence stars on the HR diagram is insensitive to their location in the Galaxy indicates that their structure does not depend sensitively on parameters that vary from one place to another. We now know that mass is the primary determinant of where a star lies on the main sequence. Other physical characteristics, such as age, metallicity, and rotation, affect stellar colors and magnitudes (and therefore should be included in a complete model of stellar structure), but they are relatively minor factors along the main sequence and can be neglected in the interest of isolating the basic physics.  Among galaxies there are hints of analogous ``sequences\". These are referred to as galaxy scaling laws and include the Faber-Jackson \\citep[FJ;][]{fj} and Tully-Fisher relations \\citep[TF;][]{tf}, the Fundamental Plane \\citep[FP;][]{dd87, d87}, and the Fundamental Manifold \\citep[FM;][]{zgz,zgzb}. Although it is not yet evident that any of these is as fundamental for galaxies as the main sequence is for stars, they do imply that a limited number of parameters  characterize the gross properties of at least certain subsets of galaxies.  There are two arguments against using existing galaxy scaling laws as guides to a fuller description of galactic structure. First, existing scaling laws work only over a limited range of galaxy types and luminosities. While this is not an insurmountable obstacle --- not all stars lie on the main sequence --- it suggests that the current scaling laws are incomplete and that they will not lead to a description of all galaxies. Second, for historical reasons related to the empirical nature of the scaling laws, their current formulation is not optimal with respect to possible theoretical constructs. For example, one determination of the Fundamental Plane has $\\log r_e = 1.24 \\log \\sigma - 0.82 \\log I_e + \\gamma$ \\citep{jorg}, where $r_e$ is the half-light radius, $\\sigma$ is the velocity dispersion, $I_e$ is the surface brightness within $r_e$, and $\\gamma$ is a constant. It is unlikely that a simple theory would reproduce the 1.24 and 0.82 coefficients.  In this paper, we attempt to address both of these shortcomings. Although we are not the first to hope to identify a unifying description of galaxies \\cite[see $\\kappa$-space;][]{burstein}, we achieve three of our key goals: 1) to find an empirical relationship for {\\sl all} galaxies that has comparable scatter to those relations identified previously for limited subsets of galaxies (TF, FJ, FP, and FM), 2) to isolate the critical additional knowledge beyond the virial theorem that is needed to derive this relationship, and 3) to begin constructing the bridge between a purely empirical relationship that utilizes observables and a theoretical one that is based on physical parameters.  In summary, we begin from basic dynamical principles and examine the dimensionality of the family of galaxies, ranging from dSph's to brightest cluster galaxies and from disks to spheroids. We address our basic question --- how uniform are the gross structural properties of galaxies? --- by determining that a single scaling relation exists that spans all luminosities and galaxy types and by quantifying its scatter. We employ an extended version of the fundamental plane formalism that reproduces the structural properties of all galaxies at a level comparable to that achieved with either the TF relation for disks or the FP for spheroids. We establish that connecting this new scaling relation to the virial theorem requires knowledge only of the mass-to-light ratio within $r_e$, $\\Upsilon_e$, and that $\\Upsilon_e$ can be accurately modeled as a function of the observed structural parameters themselves. We proceed to describe $\\Upsilon_e$ as a combination of the fraction of baryons converted into stars, $\\eta$, and the degree to which those stars are packed within the dark matter halo, $\\xi$. Using our empirical findings, we then  calculate these two physical parameters for all of our galaxies and compare $\\eta$ with independent measurements. The aim of this work is to define simple expressions for basic physical parameters of galaxies  that may illustrate which physical processes drive the observed patterns of galactic structure, with the expectation that this will focus subsequent, more detailed theoretical work. ",
        "conclusions": "\\label{sec:discussion} \\subsection{Proceeding from the Virial Theorem} \\label{sec:beyond}  In this section, we revisit the standard derivation of the FP to provide a framework and physical intuition for our observational results. We begin with the tensor virial theorem, use simplifying assumptions to rewrite the virial theorem in terms of observed quantities where possible, discuss the resulting equation and its implications for the nature of galactic structure, and finally suggest a way to proceed even though some terms in the resulting equation cannot be expressed in terms of observed quantities. We apply this suggestion and explore it in quantitative detail in \\S \\ref{sec:simplicity} and \\S \\ref{sec:validity}.  We begin with the tensor virial theorem, which is \\begin{equation}\\label{eq:1} {1\\over 2}{d^2I_{jk}\\over dt^2} = 2T_{jk} + \\Pi_{jk} + W_{jk}, \\end{equation} where $\\bf{T}$ and ${\\bf \\Pi}$ are the contributions to the kinetic energy tensor from the ordered and random motions, ${\\bf W}$ is the potential energy tensor, and ${\\bf I}$ is the moment of inertia tensor.  In steady state, the left hand side of Eq. \\ref{eq:1} is zero. We evaluate the trace of this equation and express the ordered component of the kinetic energy as ${1 \\over 2}A_0 M v_c^2$ and the random as $A_1 M\\sigma^2$,  where $v_c$ is the circular velocity in disk galaxies, $\\sigma$ is the line-of-sight velocity dispersion for spheroidals, $M$ is the mass of the system, and the $A$'s represent the correction factors obtained by fully evaluating the appropriate integrals. Similarly, the potential energy is expressed as $-B_0GM^2/R$, where $R$ is a characteristic radius that we define to be the half-light radius, $r_e$, and $B_0$ is a correction factor obtained by fully evaluating the appropriate integral. Hence, without loss of generality, \\begin{equation}\\label{eq:2} A_0v_c^2 + A_1\\sigma^2 = B_0{GM \\over r_e}. \\end{equation} All of the possible real-world complications are encapsulated in the yet unspecified $A$'s and in $B_0$, and, in principle, these could be extremely complicated functions of the formation history and environment of  galaxies. The only assumption that we have made so far is that the virial theorem holds over these radii, which is reasonable for galaxies because $r_e \\ll r_{vir}$.  To numerically evaluate Eq. \\ref{eq:2}, we now introduce two sets of simplifying assumptions that we will eventually test by determining whether we reproduce the observations. First, is the  {\\sl kinematic} simplification. We reduce the number of $A$ parameters, by requiring $A_0v_c^2 + A_1\\sigma^2 \\equiv AV^2$, where $V \\equiv ({1 \\over 2} v_c^2 + \\sigma^2)$. This simplification is accurate if we are dealing with isothermal spheres and isotropic velocity dispersions. In such systems, at large radii, $v_c$ for a purely rotationally-supported population equals $\\sqrt{2}\\sigma$ for a purely pressure-supported population. Furthermore, in such a pressure-supported system, the internal velocity dispersion is equal to the line-of-sight velocity dispersion. To evaluate $V$ for disk galaxies, we will use the measured $v_c$ and set $\\sigma=0$, while for spheroidal galaxies we will use the measured $\\sigma$ and set $v_c = 0$.  Although the kinematic simplification relies on highly specific assumptions, both disks and spheroids satisfy the relevant conditions well, and this conversion has been used previously in various contexts \\citep{burstein,kassin}. Optical disks are characterized by flat rotation curves, which imply that the mass profile is that of an isothermal sphere over these radii, and that the velocity tracers, H II regions or neutral hydrogen, are on circular orbits \\citep[see][]{faber}. Spheroids also lie in mass distributions that are consistent with being isothermal spheres \\citep{gavazzi}, and their stellar velocity dispersions are nearly isotropic if the system is a slow rotator \\citep{cappellari2007}. Due to the nature of our spheroid samples, we expect that there are strong selection biases against fast rotators among the more luminous systems (i.e., they would often be removed from Fundamental Plane studies), and the lowest luminosity systems show little rotation \\citep{walker}. Therefore, our sample is likely to satisfy the assumptions involved in the kinematic simplification, but we discuss possible signatures of failure in \\S \\ref{sec:revisiting}.  \\begin{figure} \\plotone{figure1.eps} \\caption{The distribution of galaxies (444 spheroids, 1481 disks) in the ($r_e,I_e, V$) space. This projection was chosen by eye to demonstrate that galaxies lie on a fairly well defined surface. The data are discussed in \\S \\ref{sec:data}. The axes in the lower right show the orientation of the 3-space. The poorly populated tail consists of Local Group dwarfs and would be, we expect, well populated in a volume limited sample. Units for the three axes are kpc, $L_\\odot$/pc$^2$, and km sec$^{-1}$, where $V={v_c \\over \\sqrt{2}}$ for disks and $V=\\sigma$ for spheroids.} \\label{fig:all} \\end{figure}  Continuing in our attempt to convert Eq. \\ref{eq:2} into an equation that we can numerically evaluate, our second simplification involves the replacement of the difficult-to-measure $M$ with $M_e$, the mass enclosed at $r_e$. We refer to this as the {\\sl mass} simplification. We rewrite $M_e$ as $\\Upsilon_e L_e$ so that it is expressed as a function of observable quantities: $L_e$ and $\\Upsilon_e$, the luminosity ($\\equiv \\pi r_e^2I_e$) and the mass-to-light ratio within $r_e$, respectively. We then replace $B_0$ with $B$ to account for the unknown difference between $M$ and $M_e$.   We now apply the kinematic and mass simplifications to Eq. \\ref{eq:2}, rewriting it as \\begin{equation}\\label{eq:3} AV^2 = B{G\\pi\\Upsilon_e r_e I_e}. \\end{equation} Finally, we take the logarithm of both sides and rearrange terms to obtain \\begin{equation}\\label{eq:4} \\log r_e -  \\log V^2 + \\log I_e + \\log \\Upsilon_e - \\log A +  \\log B = const. \\end{equation}  This equation leads to the rather dispiriting conclusion that galaxies populate at least a six dimensional parameter space --- more if yet unspecified parameters, such as $\\Upsilon_e$, are actually functions of additional parameters, like age, metallicity, formation history, bulge-to-disk ratio, or environment. Surprisingly, as shown in Figure \\ref{fig:all}, galaxies populate a limited region of the $(r_e,I_e,V)$ space, indicating a much lower dimensionality.    One way in which the dimensionality of the galaxy family might be reduced from that suggested by Eq. \\ref{eq:4} is if $\\Upsilon_e$, $A$, and $B$ are functions {\\sl only} of $r_e, I_e$, and $V$. A simple variant of this scaling is referred to as ``homology\", in which the functional forms are assumed to be power laws. Because of the logarithms, the end effect of rewriting Eq. \\ref{eq:4} in such a variant is a change in the coefficients of the $\\log V^2$ and $\\log I_e$ terms. Therefore, the assumption of homology results in a prediction that galaxies lie on a plane in the $(r_e,I_e,V)$ space. The values of the coefficients describe the tilt of that plane. The success of the Fundamental Plane description for giant ellipticals \\citep{dd87,d87,jorg, bernardi} demonstrates that, over the limited mass range of these galaxies, the homology assumption holds surprisingly well. This success was extended in the $\\kappa-$space formalism of \\cite{burstein}, where different classes of objects were found to lie on different planes. However, the failure of a single plane to describe the distribution of all spheroidal galaxies demonstrates that over a more extended mass range,  which includes the most and least massive spheroids, homology does not hold \\citep{zgz}.  Here we take a different approach in that we (1) assume that {\\sl all} galaxies (faint, luminous, disk, and spheroid) fall on a single manifold in the $(r_e,I_e,V)$ space and (2) examine the behavior of $\\Upsilon_e$ that would make that possible. This approach is motivated by our earlier finding that the behavior of $\\Upsilon_e$ is both simple and qualitatively reasonable for all spheroids \\citep{zgz,zgzb} --- leading to what we termed the Fundamental Manifold (FM) of spheroids. So emboldened, we now assert that for all galaxies deviations from homology are dominated by the behavior of  $\\Upsilon_e$, ignoring variations in $A$ and $B$ among galaxies\\footnote{An alternate treatment that eventually leads to the same conclusion is to group together $\\Upsilon_e$, $A$, and $B$ into a generic unknown, $\\Delta$, fit for $\\Delta$, and then use the argument in \\S \\ref{sec:validity} to demonstrate that $\\Delta \\propto \\Upsilon_e$.}. Our approach here represents a philosophical departure from ours and others' earlier work, which usually focused on establishing or quantifying the tight empirically-derived scaling relationships (e.g., the FP or FM), because we posit the existence of a fundamental manifold of all galaxies and then examine the implications.   \\begin{figure*} \\plotone{figure2.eps} \\caption{Projections of $\\log \\Upsilon_e$ and $\\log \\Upsilon_e^f$. We plot the projections of $\\log \\Upsilon_e - const.$, determined from Eq. \\ref{eq:5},  vs. $\\log V$ and $\\log I_e$ in the upper panels. In the lower panels, we plot the values of $\\log \\Upsilon_e^f$ for every galaxy in our sample using the fit given in Table 1. The lower panels illustrate how even with no intrinsic scatter in $\\log \\Upsilon_e^f$ the projections show significant apparent scatter. We conclude that the bulk of the apparent scatter in the upper panels is due to the effects of projecting the complicated surface onto these axes rather than observational errors or intrinsic scatter. } \\label{fig:projections} \\end{figure*}    \\begin{figure*} \\plotone{figure3.eps} \\caption{Testing an equation of galactic structure, Eq. \\ref{eq:5}. We replace $\\log \\Upsilon_e - const.$ with $\\Upsilon_e^f$ to evaluate Eq. \\ref{eq:5} for the entire sample (left panel) and for spheroids and disks separately (right panels).   By construction, the mean relation should lie along the 1:1 line. The low scatter and the lack of systematic deviations for galaxy subsamples testifies to the universal nature of this simple relation. } \\label{fig:fm} \\end{figure*}   \\subsection{The Simplicity of Galaxies} \\label{sec:simplicity}  The treatment described in \\S \\ref{sec:beyond} and culminating in Eq. \\ref{eq:4} is incomplete. The simple theoretical approach fails because it does not predict the low dimensionality of the data seen in Figure \\ref{fig:all}. On the other hand, the purely empirical treatment of fitting a manifold to the data in the $(r_e,I_e,V)$ space fails because it does not connect the actual functional form to a physical framework. Much like the case with the FP coefficients, directly fitting the data will subsume the behavior of $\\Upsilon_e$, $A$, and $B$ in Eq. \\ref{eq:4} into the coefficients of the various structural terms (see axes in Figure \\ref{fig:all}). Instead, we merge the two approaches by retaining the values of the coefficients derived from the virial theorem treatment as given in Eq. \\ref{eq:4}, assert that the most distinct break from homology occurs in $\\Upsilon_e$, set $A$ and $B$ to be constants, and then solve for $\\Upsilon_e$, \\begin{equation}\\label{eq:5} \\log \\Upsilon_e = \\log V^2 - \\log I_e - \\log r_e  + const., \\end{equation} This approach may seem like only mathematical sleight-of-hand, but we will quantitatively test our association of $\\Upsilon_e$ with the dominant departures from homology in \\S \\ref{sec:validity}.  To proceed, we evaluate $\\log \\Upsilon_e -\\ const.$ using Eq. \\ref{eq:5} and plot the results in Figure \\ref{fig:projections}. We then fit for the function, $\\log \\Upsilon_e^f$, that describes these data and also plot the calculated values using this fit in Figure \\ref{fig:projections}. Because we are fitting to a distribution of points in a 3-space, the fitting function will depend on two variables, and the natural choices are those that are distance independent, $V$ and $I_e$. We want to minimize the fitting order, while still capturing the behavior of the distribution. As demonstrated by \\cite{zgz}, there is at least a second-order dependence on $\\log \\sigma$ and some dependence on $\\log I_e$, and so we fit to second-order in both $\\log V$ and $\\log I_e$ and include cross-terms. For this fit, we use only a randomly selected one-sixth of the \\cite{springob} sample to avoid having that sample dominate the fit. We present the coefficients of our fit in Table 1, but, because of the heterogeneous nature of the data and our avoidance of any type of Malmquist-like corrections \\citep{willick}, these numbers are far from definitive.  The distinction between this work, with its complex characterization of $\\Upsilon_e$, and either FP or $\\kappa$-space, with their assumption of homology, becomes evident when examining Figure \\ref{fig:projections}. The projections of the data and the fitting function in Figure \\ref{fig:projections} illustrate how even in projection the functional form that describes $\\Upsilon_e$ deviates from power laws. The upper panels contain the inferred values of $\\log \\Upsilon_e - const.$ from Eq. \\ref{eq:5} vs. either $\\log V$ or $\\log I_e$. The lower panels, which show the fitted values, $\\log \\Upsilon_e^f$, and therefore have no intrinsic scatter, illustrate how the bulk of the observed scatter in the upper panels comes simply from the projection of a complicated surface onto these axes. In other words, the reason why galaxies of the same $V$ have a range of $\\Upsilon_e$'s is not primarily because there is intrinsic scatter --- say, due to age or metallicity --- but rather because galaxies have a range of $I_e$. For a given $V$ and $I_e$, the scatter in $\\Upsilon_e$ is much smaller than that observed in Figure \\ref{fig:projections}. To be specific, the scatter for the entire sample about the fit is 0.094 dex (24\\% rms in $\\Upsilon_e$). In contrast, the observed scatter in $\\log \\Upsilon_e$ in the upper left panel of Figure \\ref{fig:projections} for $1.9 < V < 2$ is 0.22 dex (66\\% rms in $\\Upsilon_e$).   \\begin{deluxetable}{rrrrrr} \\tabletypesize{\\scriptsize} \\tablecaption{$\\Upsilon_e^f$ Fit Coefficients} \\tablewidth{0pt} \\tablehead{ \\colhead{Constant} & \\colhead{$\\log V$} & \\colhead{$\\log I_e$} & \\colhead{$\\log^2 V$ } & \\colhead{$\\log^2 I_e$} & \\colhead{$\\log V \\log I_e$} \\\\ } \\startdata 2.12&$-$0.01&$-$1.05&0.07&0.13&0.14\\\\ \\enddata \\label{tab:ml} \\end{deluxetable}  We are now ready to evaluate the degree to which Eq. \\ref{eq:5} describes our set of galaxies. Replacing $\\log \\Upsilon_e - const$ with $\\log \\Upsilon_e^f$, we evaluate Eq. 5 and plot a rearrangement of the terms  in Figure \\ref{fig:fm}.  By construction, Eq. \\ref{eq:5} is satisfied on average when $\\log \\Upsilon_e - const.$ is replaced by $\\log \\Upsilon_e^f$, which is evident in Figure \\ref{fig:fm}. The actual test of our approach comes from examining the scatter about the mean and whether distinct galaxy populations fall off the mean trend. If galactic structure depends strongly on parameters not included in this simple description, then the scatter will be large. In other words, two galaxies that are identical in the quantities $V$, $r_e$, and $I_e$ could, in principle, have very different values of $\\Upsilon_e$ due to a dependence of $\\Upsilon_e$ on accretion history, age, varying degrees of mass loss, or many other possible physical effects. These differences in $\\Upsilon_e$ are not accounted for in $\\Upsilon_e^f$, thereby potentially leading to a large scatter about the mean. However, the scatter is only 0.094 for the entire sample. For reference, the scatter in this new relation for {\\sl all} galaxies is comparable to the scatter observed in either FP or TF studies for the relevant subset of galaxies. The scatter can be reduced slightly, to 0.084, if we correct each galaxy sample separately for zero point shifts. These inferred zero point shifts, obtained by calculating the mean offset relative to the 1:1 line, are all comparable to plausible photometric errors, and generally correspond to a few hundredths of a magnitude.  The success of placing all galaxies onto a single surface in the $(r_e,I_e, V)$ space demonstrates that, to within the scatter ($< 25\\%$), the family of galaxies is a two parameter sequence, i.e., measuring two of these structural parameters specifies the third. This implies that  potentially important factors in galaxy development, such as environment, nuclear activity, star formation history, and accretion history, do not play a significant role in determining galactic structure {\\sl unless} they either move galaxies along the surface in $(r_e,I_e,V)$ space or act in concert to preserve the manifold as the locus of equilibrium points of galactic structure. To reiterate, the important aspect of Figure 3 is not its linear nature, which is a result of our assertion that a fundamental manifold exists for all galaxies, but rather the low scatter and lack of any systematic departures for specific classes of galaxies, both of which imply that the assumptions that we have made so far are appropriate to this level of precision. In other words, random or systematic variations of $A$ and $B$ across either the mass range of galaxies or galaxy types, variations from isothermality, and any other factors that are not considered contribute scatter that is at most the observed scatter, which is comparable to that measured in TF or FP studies. We have achieved the first of the goals described in \\S \\ref{sec:intro} and now proceed to examine whether our attribution of the departures from homology to $\\Upsilon_e$ is correct.  \\begin{figure} \\plotone{figure4.eps} \\caption{A comparison of mass-to-light ratios derived from independent means, $\\Upsilon_e^{IND}$ and our estimates of $\\Upsilon_e$ using $\\Upsilon_e^f$.  For normal ellipticals (circles), we compare the mass-to-light ratio derived using the Schwarschild method of dynamical modeling for a set of nearby spheroidal galaxies \\citep{cappellari} to our estimates of $\\Upsilon_e$ using $\\Upsilon_e^f$. Open circles represent galaxies that \\cite{cappellari} note are fast rotators and filled circles represent those that are not. The line is the 1:1 correspondence. The filled circles, which are the most appropriate comparison sample, show only 0.06 dex scatter (15\\% in $\\Upsilon_e$). For the Milky Way dSph galaxies (triangles), we compare the mass-to-light ratio derived from fitting NFW profiles to kinematic data \\citep{walker} with our estimates of $\\Upsilon_e$ using $\\Upsilon_e^f$. The unknown constant relating $\\Upsilon_e^f$ to $\\Upsilon_e$ is set to ensure agreement in the mean values of $\\Upsilon_e^{IND}$ and $\\Upsilon_e^f$ and that value ($-$0.8) is then adopted for Eq. \\ref{eq:6}. } \\label{fig:ml2} \\end{figure}  \\subsection{The Physical Validity of  $\\Upsilon_e^f$} \\label{sec:validity}  The mathematical trick of placing all of the galaxy formation physics beyond the virial theorem into $\\Upsilon_e$ potentially masks the importance of $A$ and $B$. To check whether $\\Upsilon_e^f$ truly reflects $\\Upsilon_e$, or whether it is in actuality a  composite of various terms,  we compare $\\Upsilon_e^f$ to independent determinations of $\\Upsilon_e$. We do this for both normal ellipticals and dSphs.  First, we compare to values of $\\Upsilon_e$ derived from a full Schwarschild analysis of the 2D line-of-sight velocity distributions, $\\Upsilon_e^{Sch}$, of normal ellipticals \\citep{cappellari, cappellari2007}. Because of the unknown constant in the definition of $\\Upsilon_e^f$ relative to $\\Upsilon_e$, we have the freedom to normalize $\\Upsilon_e^f$ to best match the \\cite{cappellari} data, which we do below (Figure \\ref{fig:ml2}). Figure \\ref{fig:ml2} illustrates the excellent correspondence between $\\Upsilon_e^f$ and $\\Upsilon_e^{Sch}$. The agreement is particularly good (0.06 dex rms, 15\\% in $\\Upsilon_e$) for the galaxies that are most appropriate for our construction, namely those where the velocity dispersion dominates over systemic rotation and anisotropy measures are small ($|\\beta, \\gamma, \\delta| < 0.15$ as measured by \\cite{cappellari2007}). The scatter for those with large anisotropies is significantly greater (0.17 dex rms, 48\\% in $\\Upsilon_e$), suggesting that a full knowledge of $A$ and $B$ would decrease the scatter in Figure \\ref{fig:fm} among those galaxies that do not fully satisfy the basic assumptions of our approach.  Second, we examine whether this correspondence holds across the range of galaxy masses. In Figure \\ref{fig:ml2} we include values of $\\log \\Upsilon_e$ for Galactic dSphs estimated using NFW model fits to the extensive kinematic data of \\cite{walker}\\footnote{The calculated $\\Upsilon_e$'s use masses enclosed within $r_e$ as calculated from the published fits, courtesy of Matthew Walker.}. Unlike the Schwarschild analysis of normal ellipticals, which has the freedom to include anisotropic velocity distributions, the \\cite{walker} analysis does not. The dSph data suggest a slightly different constant offset between $\\log \\Upsilon_e^f$ and $\\log \\Upsilon_e$ ($-$0.78 rather than $-$0.82), but this change is modest and has a nearly undetectable effect on the scaling relation when ignored (see Figure \\ref{fig:mlfm}). We adopt the average offset, $-$0.8, as the normalization constant and suggest an uncertainty in this number of $ \\sim$ a few hundredths. Using this correspondence, we replace the constant in Eq. \\ref{eq:5} to obtain \\begin{equation} \\label{eq:6} \\log r_e -  \\log V^2 + \\log I_e + \\log \\Upsilon_e + 0.8 = 0, \\end{equation} where one can either evaluate $\\Upsilon_e$ in some independent manner or express it in terms of $V$ and $I_e$ using the fit given in Table 1 and replacing $\\log \\Upsilon_e + 0.8$ with $\\log \\Upsilon_e^f$.  The applicability of the same normalization constant for both normal ellipticals and dSphs supports the contention that structural variations, as reflected by changes in $A$ and $B$, are modest over most of the mass scale covered in Figure \\ref{fig:fm}. Using these independently derived measures of $\\Upsilon_e$, we now return to Eq. \\ref{eq:6}, use the literature values for $\\Upsilon_e$ rather than our fitting function, and plot the results in Figure \\ref{fig:mlfm}. The difference between the left and right panels is the exclusion of the ellipticals that show evidence for rotation or anisotropic velocity dispersions. The data in both panels follow the 1:1 correspondence well, although the correspondence in the right panel is striking. The scatter in that panel is 0.04 dex, or less than 10\\% in the parameter values themselves. We conclude that to a level of precision between 10 and 25\\% (the scatter measured using these independently measured values of $\\Upsilon_e$ and the scatter measured using our fitting function, $\\Upsilon_e^f$, respectively), $\\Upsilon_e$ encompasses all of the additional physics necessary to proceed from the virial theorem to a description of galactic structure. Thus, we achieve the second goal listed in \\S \\ref{sec:intro}.  \\begin{figure} \\plotone{figure5.eps} \\caption{The scaling relationship, Eq. \\ref{eq:6}, using $\\Upsilon_e^{IND}$ in place of $\\Upsilon_e$. Data and symbols are as in Figure \\ref{fig:ml2}. The left panel includes all of the spheroids with $\\Upsilon_e^{IND}$, and the scatter  is 0.09 dex about the 1:1 line with a mean offset of 0.005. In the right panel, we have removed the ellipticals with either significant rotation or anisotropy \\citep{cappellari,cappellari2007}, and the scatter drops to 0.04 dex about the 1:1 line with a mean offset of $0.004$.} \\label{fig:mlfm} \\end{figure}   \\subsection{Evolving onto the Manifold} \\label{sec:evolving}  For various reasons, the small scatter seen in Figures \\ref{fig:fm} and \\ref{fig:mlfm} is remarkable. Even if $r_e$ and $V$ are the same in two similar galaxies, one might expect variations in the stellar mass-to-light ratios, $\\Upsilon_{*}$, of more than 50\\%, which would introduce scatter via variations in $I_e$. We suspect that at least part of the reason for the low observed scatter lies in the selection of galaxies in TF and FP studies, which generally favor  evolved, dynamically-relaxed galaxies, which are unlikely to exhibit dramatic variations in $\\Upsilon_{*}$.  Therefore, we now return to the last of our disk galaxy samples \\citep{geha2006}, which contains disks with extremely high gas mass fractions, and might therefore be expected to harbor systems with dramatically different values of $\\Upsilon_{*}$ and $\\Upsilon_e$ than those included in our analysis so far.  When the majority of the baryons in a galaxy are in the gas rather than in stars, one might expect that our treatment as described above will fail because the connection between optical luminosity and mass via $\\Upsilon_e$ becomes tenuous. In fact, the galaxies in the \\cite{geha2006} sample do fall off the surface, as shown in the upper panel of Figure \\ref{fig:geha}. However, as demonstrated with regards to the Tully-Fisher relation \\citep{mcgaugh05, geha2006},  gas-rich and gas-poor galaxies have consistent scaling relationships if one considers total baryonic content instead of just that in the stellar component. Reviewing Eqs. \\ref{eq:2} and \\ref{eq:3}, it is evident that the derivation of later equations, such as Eq. \\ref{eq:6}, depended on a proxy for the enclosed mass. In light of  previous studies, while optical luminosity may be an adequate proxy when most of the baryons are in stars, it is clearly inadequate when the majority of baryons are in the gas (or, more precisely, our fitting formula for $\\Upsilon_e$ fails when different fractions of the baryons are luminous in otherwise similar galaxies). The reason optical luminosity works well for most disks is that the majority of their cold gas has been turned into stars \\citep{roberts, bothun, kannappan, read}. Given that we do not have measured colors and gas masses for most of our sample, we cannot reformulate everything in terms of baryonic mass, but we can ask where these gas-rich galaxies would lie with respect to the 1:1 relationship in Figure \\ref{fig:geha} if they turned their gas into stars. Can they evolve onto the manifold?  \\begin{figure} \\plotone{figure6.eps} \\caption{Results for the \\cite{geha2006} sample of disk galaxies with high gas mass fractions. Upper panel shows the galaxies  (points) in comparison to the fitted relationship (line). Lower panel shows to where the galaxies might evolve as their gas is turned into stars.} \\label{fig:geha} \\end{figure}  To complete this exercise, we make two questionable assumptions. First, we assume that $r_e$ does not change during the conversion of gas to stars. Because the gaseous and stellar radial distributions are likely to be different, this assumption must fail at some level.  Second, we assume that the stellar population formed from the gas eventually has the same stellar mass-to-light ratio (i.e., for stars only) as the late-type galaxies in our sample \\citep[$\\Upsilon_*= 0.97$ in the $I$-band for a typical Sbc galaxy with $B-V$ = 0.57;][]{fukugita,bell}. For these two assumptions, we then find the fraction of gas turned into stars that produces the best match to our scaling relation (Eq. \\ref{eq:6}). We exclude one galaxy with a rotation velocity of 5 km sec$^{-1}$, which appears to be unphysically small, and one galaxy with $b/a > 0.9$, for which it is difficult to deproject the velocity width.  The result of this exercise, that these high gas mass fraction galaxies can evolve onto the relationship defined by the larger sample, is shown in the lower panel of Figure \\ref{fig:geha}. The resulting scatter is 0.16 dex, which is slightly lower than that found for the dSphs, but larger than for all of the galaxies combined. The difficulty with our scenario is that the agreement requires turning about  20\\% more gas than is available throughout the galaxy into stars. Perhaps this failure reflects the need for infalling gas, but it may also mean that as the galaxy evolves there is a corresponding change in $r_e$. For example, in a model where the gas is funneled efficiently into the center so that $r_e$ is only one-third of its current value, and 75\\% of the current gas is turned into stars, the scatter about the relationship is only 0.14 dex. We conclude that there are plausible scenarios in which these gas-rich galaxies lie on the observed relationship and that one might identify galaxies that are still strongly evolving or those with a significant reservoir of cold gas as outliers from the manifold. Nevertheless, if these galaxies either convert their gas to stars or if we properly account for their entire baryonic content within $r_e$, then we expect that they will satisfy the same scaling relation as all other galaxies.  \\subsection{Revisiting our Simplifications} \\label{sec:revisiting}  Before we proceed to discuss further implications of these results, we step back to explore how a failure to satisfy the assumptions invoked in our simplifications of the virial theorem would manifest itself in our evaluation of Eq. \\ref{eq:6}. The potential ``failures\" fall into three classes.  First, we might have introduced errors that are constant  across the galaxy population. An example of such an error would be if we always underestimated the potential energy in our evaluation of the virial theorem by a fixed factor. Such an error would manifest itself as a zero point shift of the data relative to the expectation. Because we do not calculate the specific constant in Eq. \\ref{eq:6} from any physical argument, this type of error will be transparently absorbed into the constant term when we determine it using independent measurements, as done in \\S \\ref{sec:validity}. For almost all of our discussion, this type of error is difficult to detect but irrelevant.  Second, we might have introduced errors that vary systematically across the galaxy population. An example of such an error would be if we underestimated the potential energy by a certain factor for low mass systems but overestimated it by a similar factor for high mass systems. Such errors, to the degree that they correlate with at least one of the structural parameters, will lead to changes in the coefficients in Eq. \\ref{eq:6} or that describe $\\Upsilon_e^f$ (Table \\ref{tab:ml}), but would not introduce scatter. This effect is analogous to introducing a ``tilt\" in FP analyses. Identifying this type of error is critical if one aims to understand the specific nature of the fitted relationships such as that describing $\\Upsilon_e^f$ or to compare with simulations. We implicitly tested for such effects across galaxy types in \\S \\ref{sec:simplicity} and across galaxy mass in \\S \\ref{sec:validity}. The coefficients derived from the virial theorem and a mass-to-light ratio that scales directly with independently-derived mass-to-light ratios successfully produce a tight scaling relation (Eq. \\ref{eq:6} and Fig. \\ref{fig:fm}). This result demonstrates that there is no effective ``tilt\" either with morphological type or across the full range of galaxy luminosities.  Third, we might have introduced errors that are variable and not systematic across the galaxy population. An example of such an error would be if we have ignored a key determining factor of galactic structure, {\\sl e.g.} the number of nearby neighbors. In naive models, close passages affect the luminosity of the system but do not affect the size or internal kinematics --- leading to potential outliers in Eq. \\ref{eq:6}. Such effects, to the degree that they do not correlate with the remaining structural parameters, will introduce scatter at each point in the $(r_e,I_e,V)$ space. The low scatter in both the entire sample (Fig. \\ref{fig:fm}) and the subsample with independently-measured mass-to-light ratios (Fig. \\ref{fig:mlfm}) demonstrates that any such errors introduce little noise.   \\subsection{Connecting $\\Upsilon_e$ to Physical Parameters} \\label{sec:connecting}  What we have discussed so far are the end-products, or observables $(r_e,I_e,V,\\Upsilon_e)$, rather than inputs, or true physical quantities, that determine the structure of galaxies. Two natural candidates for the parameters that drive galactic structure are the mass and angular momentum of a galaxy. For disks, analytic treatments of galaxy formation using these two variables have been relatively successful \\citep{fall, dalcanton}. Although these models require a few key assumptions that may not be entirely accurate, such as the conservation of angular momentum during collapse, their success, in combination with the results presented here, suggest that simple dynamical models may be able to reproduce the observable properties of galaxies. However, proceeding from such difficult to measure quantities as mass and angular momentum to the observed structure of galaxies in a single step is likely to prove difficult.  We focus instead on what we have learned so far from our analysis. Given that the virial theorem plus $\\Upsilon_e$ are all one needs to generate a gross description of galactic structure (\\S \\ref{sec:beyond} --- \\ref{sec:validity}), all of the interesting astrophysics of galaxy formation --- at least as related to determining the current, gross structure of a galaxy --- is encapsulated in $\\Upsilon_e$.  What determines differences in $\\Upsilon_e$ among galaxies?  First, galaxies might convert a different fraction, $\\eta$, of their baryons to stars. Assuming a universal baryon-to-dark matter ratio for halos, this fraction can be measured using the total mass-to-light ratio, $\\Upsilon_{200}$, evaluated at $R_{200}$, the radius within which the mass density is 200 times the critical density and the system is roughly virialized. Systematic  variations in $\\eta$ have already been noted in studies of halo occupation statistics \\citep{vdB,yang}, lensing mass measurements \\citep{hoekstra, mandelbaum}, and direct baryon measurements \\citep{lin,baryons}. All of these studies show that $\\Upsilon_{200}$, and therefore $\\eta$, depend on halo mass.  Second, the stars might be concentrated to varying degrees relative to the dark matter due to differences in the assembly history. We quantify the stellar concentration, $\\xi$, using $\\xi \\equiv R_{200}/r_e$. As the stars become more concentrated in the halo, $\\xi$ increases and $\\Upsilon_e$ decreases.  \\begin{figure} \\plotone{figure7.eps} \\caption{Fraction of baryons that are converted to stars, $\\eta$, as a function of $L/L^*$ for spheroids (upper) and disks (lower). Our data are the small dots. The values from \\cite{mandelbaum} are shown as large filled circles with error bars and represent average values for bins of $L/L^*$. This plot includes the cluster spheroids of \\cite{zgz} as open triangles, so that we extend the range of $L/L^*$ to the higher luminosities (see text) probed by \\cite{mandelbaum}.} \\label{fig:mandelbaum} \\end{figure}  We now return to Eq. \\ref{eq:2} with the aim of extracting from it expressions for the mass fraction of baryons that are converted to stars, $\\eta$, and the degree to which the stars are concentrated relative to the dark matter, $\\xi$. We rewrite Eq. \\ref{eq:2} as appropriate at the virial radius, $R_{200}$, \\begin{equation} \\label{eq:7} A_{200} V_{200}^2 = {B_{200} GM_{200}\\over R_{200}}, \\end{equation} where $M_{200} \\equiv {4\\over 3}\\pi R_{200}^3\\rho_{200}$, $\\rho_{200} \\equiv 200\\rho_{crit}$, and $R_{200} \\equiv \\xi r_e$, where $\\rho_{crit}$ is the universal critical mass density at the present epoch. To make further progress, we set $V = V_{200}$. This is patently incorrect both because the dark matter potential itself is unlikely to be isothermal out to $R_{200}$ \\citep{nfw} and any central concentration of the baryons will affect $V$. However, any constant fractional differences --- for example, $V = 1.2 V_{200} $ for all galaxies --- will be absorbed later into our normalization. What do concern us, but are ignored here, are differences in this velocity scaling that depend on the properties of the galaxy \\citep{courteau}. This problem might be correctable in an iterative scheme (i.e., assume a non-isothermal potential, estimate $\\xi$, evaluate the difference between $V$ and $V_{200}$, correct $V_{200}$, and iterate until convergence) or in a more sophisticated model of galaxy formation \\citep{somerville}, but both of these remedies require some model assumptions and lead us from the analytic descriptions that we aim to explore. Work is needed to determine the magnitude of the error introduced by our simple treatment.    Continuing, we define $\\eta$ through the equation  \\begin{equation} \\label{eq:8} L \\equiv {\\eta f_B M_{200}\\over \\Upsilon_{*}}, \\end{equation} where $\\Upsilon_{*}$ is the mass-to-light ratio of the stellar population, $L$ is the total luminosity, and $f_B$ is the baryon mass fraction. Algebra enables us to derive equations for $\\xi$ and $\\eta$:   \\begin{equation} \\label{eq:9} \\xi = \\Big({{3A_{200} \\over 800\\pi \\rho_{crit} B_{200} G}}\\Big)^{1 \\over 2} {\\rm \\ }{V\\over r_e} \\end{equation}  and  \\begin{equation} \\label{eq:10} \\eta = \\Big({{800 \\pi B_{200}^3 G^3   \\rho_{crit}  \\over 3 f_b^2 A_{200}^3}}\\Big)^{1 \\over 2} {\\rm \\ } {L \\Upsilon_* \\over V^3}. \\end{equation}  We express the combination of these two quantities as \\begin{equation} \\label{eq:11} \\log \\xi + \\log \\eta =  \\log {KV\\over r_e} + \\log {JL \\Upsilon_* \\over V^3}, \\end{equation} where all constants, as well as the structural factors $A_{200}$ and $B_{200}$, in Eqs. \\ref{eq:9} and \\ref{eq:10} are contained in  $K$ and $J$. $A_{200}$ and $B_{200}$ are not necessarily constant, although their analogs at $r_e$ are well-behaved (\\S \\ref{sec:simplicity}). We assume that $A_{200}$ and $B_{200}$ are similarly well-behaved at $R_{200}$. Rewriting, we get \\begin{equation} \\label{eq:12} \\log \\xi + \\log \\eta =  -2\\log V +\\log r_e + \\log I_e + \\log \\Upsilon_* + const. \\end{equation} We know from \\S \\ref{sec:simplicity} that we can replace the three leading terms on the right hand side with $-\\log \\Upsilon_e$ to within a constant, so \\begin{equation} \\label{eq:13} \\log \\xi + \\log \\eta =  -\\log \\Upsilon_e + \\log \\Upsilon_* + const. \\end{equation} Because  $\\Upsilon_e = f(V,I_e)$, as defined in Table 1, $\\log \\eta + \\log \\xi$ is also a function $V$ and $I_e$.  To provide a more direct example of the possible use of these equations, we use the results from \\cite{mandelbaum} to evaluate the constant terms (including the assumed constant terms $A_{200}$ and $B_{200}$) in Eqs. \\ref{eq:9} and \\ref{eq:10}. \\cite{mandelbaum} use results from weak lensing to evaluate the fraction of baryons that are turned into stars as a function of galaxy luminosity. They provide empirical values for $\\eta$ as a function of galaxy luminosity and morphology. We use their results for $\\eta$ for $L^*$ spirals, and assume $\\Upsilon_{*} = 1.7$ in the $I$-band and a universal baryon mass fraction of 0.175 \\citep{spergel}, to set the values of the leading coefficients in our Eqs. \\ref{eq:9} and \\ref{eq:10}: $\\xi =  1.4 V r_e^{-1}$ and $\\eta = 1.9 \\times 10^{-5} (L/L^*)\\Upsilon_* V^{-3}$. The adopted value of $\\Upsilon_*$ for an $L^*$ spiral (assuming that it is an Sab type with $B-V$= 0.8; \\citep{fukugita}) is calculated using Table 1 of \\cite{bell}. We then plot our calculated values of $\\eta$ for all of our galaxies (Figure \\ref{fig:mandelbaum}), assuming that $\\Upsilon_*$ is 1.7 and 2.5 for disks and spheroids, respectively.  For the purpose of this comparison, we augment our galaxy sample with measurements of cluster spheroids (CSph), the brightest cluster galaxy plus the intracluster stars of groups and clusters \\citep{zgz}. Our previous work shows that these systems lie on the FM and, in this context, they allow us to extend the range of $L/L^*$ over which we can compare to the \\cite{mandelbaum} results. We decided against including the CSphs throughout our current study because our focus is on galaxies, but CSphs do indeed fall on the 1:1 line in Figure \\ref{fig:fm}.   Our values for $\\eta$ are in agreement over the range of luminosities and galaxy types presented by \\cite{mandelbaum}, except possibly for the highest luminosity spheroids: brightest cluster galaxies and CSphs. Our data suggest a qualitatively similar drop in $\\eta$ to that found by \\cite{mandelbaum}, although at somewhat larger values of $L/L^*$. There are many technical reasons (such as the use of different bandpasses, subtleties in the definition of total magnitudes, and complications introduced by intracluster light for these most massive spheroids) that preclude any conclusion about whether there is a true discrepancy. In general, we agree quite well both qualitatively and quantitatively with their results. This agreement, in turn, suggests that $A_{200}$ and $B_{200}$ do not vary strongly as a function of mass or morphological type.  We proceed now to calculate $\\eta$ and $\\xi$ as a function of $V$ for all of our galaxies and show the results in Figure \\ref{fig:structure}. Here, to within the flaws in our simple derivation and the heterogeneous sample, is a full description of how baryons turn into stars and distribute themselves in all galaxies ranging from dSphs to BCGs. There are several striking results. First, the relatively large values of $\\Upsilon_e$ for dwarf spheroidals are primarily driven by $\\eta$ rather than by $\\xi$, which is surprisingly constant across the full range of $V$. Second, at a given $V$, ellipticals and spirals have similar values of $\\eta$. Therefore, the difference in their values of $\\Upsilon_e$ are due primarily to differences in $\\xi$. Third, there is a steady decline in $\\eta$ for $\\log V > 1.5$ ($V >$ 32 km s$^{-1}$) even for the spheroids, among which large variations in $\\Upsilon_*$ are less likely. Systems with comparable $V$ to the Milky Way, $149 < V < 163$ km s$^{-1}$ or alternatively $210 < v_c < 230$ km s$^{-1}$, have $\\eta =  0.14 \\pm 0.04$, with the spheroids being on average 5 times more concentrated than the disks. All of these results await resolution of two key open questions in the evaluation of $\\xi$ and $\\eta$: 1) how to treat the difference between $V$ and $V_{200}$ and 2) whether the structural terms are as well-behaved at $R_{200}$ as they are at $r_e$. Nevertheless, this analysis illustrates how we might construct a bridge between the empirical relations based on observables and more theoretical ones based on fundamental physical parameters. Thus, we achieve the third goal listed in \\S \\ref{sec:intro}.  \\begin{figure*} \\plotone{figure8.eps} \\caption{Mass fraction of baryons that are converted to stars, $\\eta$, (top) and stellar concentration, $\\xi$, (bottom) as a a function of $V$ for our entire sample (spheroids represented with filled circles, disks with open circles).  Various results are in evidence, including the dramatic decrease in $\\eta$ for the dSphs, the similarity in $\\eta$ among spirals and spheroids in the regime where they overlap in $V$, the systemic decline in $\\eta$ with increasing $V$ for $\\log V > 1.5$, and the generally greater $\\xi$ for spheroids relative to disks in that same velocity range.} \\label{fig:structure} \\end{figure*}     We have shown that all classes of galaxies, ranging in mass from dwarf spheroidals to brightest cluster galaxies, and in type from spheroids to disks, fall on a two dimensional surface in the observable space  $(r_e,I_e,V)$, where $ V^2 \\equiv {1\\over 2} v_c^2 + \\sigma^2$, over three orders of magnitude in $r_e$ with $<$ 25\\% scatter. The scatter about that surface is comparable to that observed in Fundamental Plane and Tully-Fisher studies in which the range of galaxy types and luminosities is much more limited. The TF and FP relationships are subsets of the manifold presented here. This finding alone demonstrates that the structure of all galaxies can be described with a highly limited set of parameters. The observational ones, $(r_e,I_e,V)$, may not be optimal, even though they do a remarkably good job. The small scatter about the mean relation implies that a host of potential physical phenomena such as environmental effects, star formation history, nuclear activity, accretion history, and feedback are either (1) relatively unimportant in determining the structure of galaxies, (2) move galaxies along this well-defined relationship, or (3) balance each other so as to define the mean relation as the locus of galactic structure equilibria.  We developed a simple analytic treatment in which we asserted the existence of a fundamental manifold of galaxies. By requiring the simple virial theorem derivation to result in a two dimensional manifold in observed space, we specify the behavior of the mass-to-light ratio within $r_e$, $\\Upsilon_e$. We then tested this assertion by comparing our inferred values of $\\Upsilon_e$ to those derived independently from much more sophisticated modeling for both normal ellipticals and dSphs. The agreement is quantitatively excellent, with less than 15\\% scatter in mass-to-light ratios for those galaxies that satisfy our dynamical criteria. This result demonstrates that the principal additional ingredient necessary in proceeding from the virial theorem to a description of galactic structure is knowledge of the mass-to-light ratio within $r_e$. Additional factors, which could have been important, such as internal kinematic anisotropy or variations in the radial profile of the gravitational potential from one galaxy to another, must play a role at less than the 25\\% level. The observed manifold is described by \\begin{equation} \\label{eq:14} \\log r_e -  \\log V^2 + \\log I_e + \\log \\Upsilon_e + 0.8 = 0, \\end{equation} where we also provide a fitting function for $\\log \\Upsilon_e$  in terms of $V$ and $I_e$. The equations presented here are numerically appropriate for $r_e$, $V$, $I_e$, and $\\Upsilon$ in units of kpc, km s$^{-1}$, $L_\\odot/$pc$^{-2}$, and solar units, respectively, and based primarily on $I$-band observations.  We then discuss what the inferred behavior of $\\Upsilon_e$ may mean for the physical characteristics of the galaxies. In particular, we speculate that the two principal determinants of $\\Upsilon_e$ are the mass fraction of baryons that are turned into stars, $\\eta$, and the degree to which the stars are spatially concentrated relative to the dark matter, $\\xi \\equiv R_{200}/r_e$. We derive equations for both quantities in terms of unknown structural parameters and the observables. We relate the two quantities using the expression that we derived for $\\Upsilon_e$. Finally, we use independent measures of $\\eta$ \\citep{mandelbaum} to solve for the unknown structural terms for one set of galaxies and then compare the behavior of $\\eta$ across other luminosities and galaxy types as determined both from our analysis and that independent weak lensing study. This comparison leads to simple expressions for $\\eta$ and $\\xi$, \\begin{equation} \\eta = 1.9 \\times 10^{-5} {(L/L^*) \\Upsilon_* \\over V^3} \\label{eq:15} \\end{equation} and \\begin{equation} \\label{eq:16} \\xi = 1.4 {V \\over r_e}. \\end{equation} We are then able to extend the measurements of $\\eta$ and $\\xi$ to the full range of galaxies. As rough guides, we find that,  for most galaxies, $0.04< \\eta < 0.6$ and $10 < \\xi < 200$, although these can be evaluated on a galaxy-by-galaxy basis. Systems with comparable $V$ to the Milky Way, $149 < V < 163$ km s$^{-1}$ or alternatively $210 < v_c < 230$ km s$^{-1}$, have $\\eta =  0.14 \\pm 0.04$, with the spheroids being on average 5 times more concentrated than the disks. Overall, we reach a set of general conclusions. First, the relatively large values of $\\Upsilon_e$ for dwarf spheroidals are primarily driven by $\\eta$ rather than by $\\xi$, which is surprisingly constant across the full range of $V$. Second, at a given $V$, ellipticals and spirals have similar values of $\\eta$ ($<$ 10\\% difference for spheroids and disks with $149 < V < 163$ km s$^{-1}$). Therefore, the difference in their values of $\\Upsilon_e$ is due primarily  to differences in $\\xi$. Third, there is a steady decline in $\\eta$ for $\\log V > 1.5$ ($V >$ 32 km s$^{-1}$).  The data used here fall short of the ideal sample from which to properly derive the quantitative values that mathematically describe the manifold, primarily due to the heterogeneous nature of the amalgamated sample. Nevertheless, the sample does demonstrate that the range of galaxy structure is dominated by only two parameters. The lower scatter obtained either for a single sample (0.06 dex for the \\cite{springob} disk sample) or for independently derived $\\Upsilon_e$'s (0.04 dex when using both the \\cite{cappellari} data for ellipticals and the \\cite{walker} data for dSphs) suggest that a homogeneous sample might show that the myriad of possible influences on galactic structure (environment, accretion history, AGN activity, star formation history, and others) contribute at most a $\\sim$ 10\\% scatter to the scaling relationship presented in Eq. \\ref{eq:14}.  The existence of a highly constrained surface on which galaxies lie does not eliminate the need for additional physics. In particular, as we have hinted throughout, many physical effects might move galaxies along the surface or perhaps counter-balance to move galaxies back to the equilibrium surface described by the manifold. As such, future galaxy models may be more constrained by the distribution of galaxies on the surface rather than perpendicular to it. Our heterogeneous sample is ill-suited to say much about the distribution of sources on the surface. Much work still remains.  We close by returning to the analogy of stellar structure. It is evident that we are still far from a physical theory of galactic structure, but that we have progressed in several key aspects. First, we have now demonstrated that the entire family of galaxies can be described by sets of two parameters (e.g., $V$ and $I_e$ or $\\eta$ and $\\xi$). This finding motivates the search for relatively simple expressions of galactic structure that are connected to a small set of physical parameters, such as mass and angular momentum. Second, we have identified the principal characteristic that remains to be explained, namely $\\Upsilon_e$. The virial theorem plus an understanding of $\\Upsilon_e$ are all that are necessary to predict the size, internal kinematics, or luminosity of a galaxy, when given the other two. This in turn places the focus on understanding what determines the fraction of baryons that are turned into stars and how those stars are packed within the dark halo. If those quantities can then be connected to more fundamental parameters, such as mass and angular momentum, then one could proceed from the physical parameters directly to the observables. At that point, we will have indeed produced equations of galactic structure."
    },
    "0711/0711.1390_arXiv.txt": {
        "abstract": "We examine the connections between the current global star formation activity, luminosity, dynamical mass and morphology of galaxies in the Local Volume, using H$\\alpha$ data from the 11 Mpc H$\\alpha$ and Ultraviolet Galaxy Survey (11HUGS). Taking the equivalent width (EW) of the H$\\alpha$ emission line as a tracer of the specific star formation rate, we analyze the distribution of galaxies in the $M_B$-EW and rotational velocity ($V_{max}$)-EW planes.  Star-forming galaxies show two characteristic transitions in these planes.  A narrowing of the galaxy locus occurs at $M_B\\sim-15$ and $V_{max}\\sim$ 50 km s$^{-1}$, where the scatter in the logarithmic EWs drops by a factor of two as the luminosities/masses increase, and galaxy morphologies shift from predominately irregular to late-type spiral. Another transition occurs at $M_B\\sim-19$ and $V_{max}\\sim$120 km s$^{-1}$, above which the sequence turns off toward lower EWs and becomes mostly populated by intermediate and early-type bulge-prominent spirals.  Between these two transitions, the mean logarithmic EW appears to remain constant at 30\\AA.  We comment on how these features reflect established empirical relationships, and provide clues for identifying the large-scale physical processes that both drive and regulate star formation, with emphasis on the low-mass galaxies that dominate our approximately volume-limited sample. ",
        "introduction": "It is well established that many of the global properties of present-day galaxies scale coherently with one another.  Of particular interest is the star formation activity; understanding and deciphering trends with this property provides insight into the astrophysics that control the pace and extent of a galaxy's evolution.  Perhaps the most well-known trend is that the current star formation rate (SFR) varies monotonically along the Hubble sequence (Roberts 1963; Kennicutt 1998).  The scale begins at the ellipticals, which generally have red colors, little gas and essentially no star formation, and proceeds through the lenticulars, spirals and irregulars, which progressively have bluer colors, larger gas fractions and greater relative levels of star formation.  The quantification of these patterns began with the development of SFR measurement diagnostics (e.g., Tinsley 1968; Kennicutt 1983a), and is now reaching full maturity with the application of such techniques to large, statistically complete samples of local galaxies (e.g., Brinchmann et al. 2004, Salim et al. 2007), as well as of galaxies at higher redshift (e.g. Noekse et al. 2007a).  Progress has been rapid in quantifying relations for galaxies in the intermediate and high-end of the luminosity function, with the significant result that the galaxy distribution is bimodal and may divide into distinct populations of old, red spheroids and blue star forming disks at a stellar mass of $\\sim3\\times10^{10}$M$_{\\odot}$ (Kauffmann et al. 2003).  In contrast, progress on the dwarf galaxy population has been slower, primarily due to the lack of complete and statistically robust samples in the low-mass regime.  To address this need, we have focused on our nearest neighbors and are taking an appoximately volume-limited (and hence dwarf-dominated) inventory of star formation in the Local Volume.  Our 11 Mpc H$\\alpha$ and UV Galaxy Survey (11HUGS) has obtained narrowband H$\\alpha$ emission-line imaging, which traces massive O-star formation, and is also collecting GALEX UV imaging, which traces the photospheric emission of O-type as well as longer-lived B-stars.  11HUGS also serves as the basis for the Spitzer Local Volume Legacy program, an IRAC (mid-IR) and MIPS (far-IR) imaging survey of 258 galaxies, which will capture the energy which is absorbed and re-radiated by dust. The 11HUGS parent sample contains $\\sim$400 galaxies and includes all known spirals, irregulars, as well as star-forming early-type galaxies with $d < $ 11 Mpc, $|b| >$ 20 \\degr and $B <$ 15 mag.  This Letter takes a first look at results from the recently completed H$\\alpha$ imaging component of the survey (Lee 2006, Kennicutt et al. 2008), and explores trends in the magnitude vs. H$\\alpha$ equivalent width (EW) plane, a close cousin of the more commonly studied color-magnitude diagram (e.g. Baldry et al. 2004 and references therein).  \\begin{figure*}[t] \\epsscale{1.05} \\plottwo{f1a.eps}{f1b.eps} \\caption{The Local Volume star-forming galaxy sequence in the (a) $M_B$-EW and the (b) rotational velocity-EW planes.  Galaxies in the core sample of 11HUGS (i.e. those with T$\\geq$0, D$<$11 Mpc and $|b|>20$\\degr) are shown.  Gray dashed lines are drawn at $M_B = -19$ and $V_{max}$=120 \\kms, and $M_B =-15$ and $V_{max}$=50 \\kms to indicate the two transition regions discussed in the text.} \\end{figure*} ",
        "conclusions": ""
    },
    "0711/0711.3395_arXiv.txt": {
        "abstract": "{}{To reconcile the measurements of $^3$He/H in Galactic HII regions with high values of $^3$He in a couple of planetary nebulae, we propose that thermohaline mixing is inhibited by a fossil magnetic field in red giant stars that are descendants of Ap stars. } {We examine the effect of a magnetic field on the salt-finger instability, using a local analysis.} {We obtain a threshold for the magnetic field of $10^4-10^5$ Gauss, above which it inhibits thermohaline mixing in red giant stars located at or above the bump. Fields of that order are expected in the descendants of the Ap stars, taking into account the contraction of their core.} {We conclude that in a large fraction of the descendants of Ap stars thermohaline mixing does not occur. As a consequence these objects must produce $^3$He as predicted by the standard theory of stellar evolution and as observed in the planetary nebulae NGC~3242 and J320. The relative number of such stars with respect to non-magnetic objects that undergo thermohaline mixing is consistent with the statistical constraint coming from observations of the carbon isotopic ratio in red giant stars. It also satisfies the Galactic requirements for the evolution of the $^3$He abundance.} ",
        "introduction": "Numerous observations provide compelling evidence for a non-canonical mixing process which modifies the surface abundances of Li, C and N of low-mass red giants\\footnote{i.e., stars with initial masses below $\\sim$ 2 - 2.5~M$_{\\odot}$ which evolve along the Red Giant Branch (RGB) to high luminosities until helium is ignited in their core under degenerate conditions.} when they reach the bump in the luminosity function (see references in Charbonnel \\& Zahn 2007, hereafter CZ07). The quest for a mechanism which could lead to that extra-mixing between the hydrogen-burning shell and the convective envelope of red giant branch (RGB) stars has initially been unsuccessful. For example rotational mixing, i.e. rotation-induced meridional circulation and shear turbulence (Zahn 1992; Maeder \\& Zahn 1998), is insufficient to explain the observed abundance pattern, as was shown by Charbonnel \\& Palacios (2004) and Palacios et al.~(2006). But recently Eggleton et al.~(2006) suggested a possible cause of such mixing, namely the molecular weight inversion created by the $^3$He($^3$He,2p)$^4$He reaction in the upper part of the hydrogen-burning shell. Based on 3D hydrodynamic simulations of a low-mass star at the RGB tip (Dearborn et al. 2006),  they found that such a $\\mu$-profile is unstable and produces very efficient mixing. They claimed that this mixing was due to the well-known Rayleigh-Taylor instability, which occurs in incompressible fluids when there is a density inversion. In stellar interiors, which are stratified due to their compressibility, a similar dynamical instability occurs when the Ledoux criterion for convection is satisfied (Ledoux 1947), but it acts to render the temperature gradient adiabatic rather than to suppress the density inversion. Presumably it is this instability that Eggleton and colleagues observed with their 3D code, as attested by the high velocities they quote, because they started their simulation with a unrealistic 1D model which did not include any mixing, in which the $\\mu$-inversion had already attained its maximum value.  In reality, the first instability to occur in a star, as the inverse $\\mu$-gradient gradually builds up, is  the thermohaline instability, as was pointed out by CZ07. This double-diffusive instability is observed in salted water in the form of elongated fingers, when the temperature is stably stratified, but salt is not, with fresh water at the bottom and salted at the top, the overall stratification being dynamically stable (Stern 1960).  It was Ulrich (1972) who first noticed that the $^3$He($^3$He,2p)$^4$He reaction would cause a $\\mu$-inversion, and he was the first to derive a prescription for the turbulent diffusivity produced by the thermohaline instability in stellar radiation zones. This prescription is based on a linear analysis, and it is certainly very crude, but it has the merit to exist. When it is applied to the $\\mu$-inversion layer in RGB stars, with the shape factor recommended by Ulrich and sustained by laboratory experiments (Krishnamurti 2003), it yields a surface composition that is compatible with the observations of the carbon isotopic ratio as well as of the abundances of lithium, carbon and nitrogen in RGB stars (Fig.~3 of CZ07).  Simultaneously thermohaline mixing leads to the destruction of most of the $^3$He produced during the star's lifetime (Fig.~4 of CZ07). It seems thus that we have finally identified the mechanism responsible for abundance anomalies in RGB stars, which simultaneously accounts for the measurements of $^3$He in Galactic HII regions (Balser et al.~1994, 1999a; Bania et al.~1997, 2002)\\footnote{According to classical theory of stellar evolution which predicts a large production of $^3$He by low-mass stars, Galactic HII regions should be highly enriched in $^3$He since they formed out of matter that has undergone 12 billion years of chemical evolution. However, their $^3$He content is similar to that of the Sun, and thus no evidence for any enrichment of $^3$He during the last 4.5 Gyr was found. In addition, the ``best\" determination in a Galactic HII region (namely S209) by Bania et al.~(2002) has yielded a $^3$He abundance almost identical to the WMAP value, indicating that the stellar contribution to $^3$He is very small. }.  A problem remains however, namely that all RGB stars brighter than the bump should undergo thermohaline mixing since they develop the same $\\mu$-inversion. As underlined by Balser et al.~(2007), this was obviously not the case at least in two stars, i.e., in the RGB progenitors of the Planetary Nebulae NGC~3242 and J320. Slightly more massive than the Sun, these PNe are indeed caught in the act of returning to the interstellar matter fresh elements synthesized in their womb among which $^3$He with the amount predicted by classical stellar models (Balser et al.~1997, 1999b, 2006; Galli et al.~1997). On the other hand the carbon isotopic ratio data indicate that a few RGB stars escape that extra-mixing - about 4$\\%$ according to Charbonnel \\& do Nascimento (1998). This is consistent with the predictions of the chemical evolution models, which suggest that fewer than 10$\\%$ of all planetary nebulae enrich the ISM with $^3$He at the level of classical stellar models (Galli et al.~1997; Tosi 1998, 2000; Palla et al.~2000, 2002; Romano et al.~2003), thus supporting the self-consistency of the extra-mixing mechanism.  It is tempting to confront these numbers with the fraction of the Ap stars (suspected to all host magnetic fields) relative to all A stars, which is $5 - 10\\%$ as estimated by Wolff (1968). Since A-type stars are the progenitors of the more massive RGB stars, we are led to conjecture that in those giants that show no sign of extra-mixing, the thermohaline instability is inhibited by a deeply buried magnetic field.  Here we check that possibility by examining the effect of a magnetic field on the salt-finger instability, using a simple local analysis. ",
        "conclusions": "The observational uncertainties in effective temperature, luminosity and metal content of Ap stars hamper a precise determination of their mass (see Bagnulo et al.~2006). According to a survey covering the Ap/Bp stars located within 100 pc from the Sun (Power et al.~2006), magnetic stars range at solar metallicity between 3.6 and 1.5~M$_{\\odot}$  (with one exception at 6~M$_{\\odot}$). The lower limit coincides with the appearance of developed surface convection zones, which depends on the initial stellar metallicity. If this correlation held in the past independently of the bulk stellar metallicity, metal-poor magnetic stars could have extended to lower initial mass: standard models predict indeed that the convection zone of a 0.9 M$_{\\odot}$ main sequence star with [Fe/H]=-2 has the same thickness as that of a 1.5 M$_{\\odot}$ main sequence star of solar metallicity (Talon \\& Charbonnel 2004). As a consequence, one can conjecture that below solar metallicity, Ap star counterparts formed with  initial masses much below 1.5~M$_{\\odot}$. This expected dependancy of the mass threshold for the magnetic inhibition of thermohaline mixing with metallicity is of direct relevance to the Galactic evolution of $^3$He and will have to be taken into account in future chemical evolution models which consider metal-dependant yields.  Galli et al.~(1997) have determined a mass of 1.2~$\\pm 0.2$~M$_{\\odot}$ for the progenitor of the planetary nebulae NGC~3242. Although such a determination is quite a feat and is undermined by several observational and theoretical difficulties, we find very encouraging that the initial mass derived for this object is so close to the lower mass limit for Ap stars at solar metallicity. We have to note indeed that the metallicity of NGC~3242 is half-solar (Barker 1985), so that the lower mass for Ap stars may already be shifted towards a lower value.  It remains to be checked whether the magnetic stars are in sufficient number on the main sequence to account for the $\\approx 4\\%$ of red giants which are spared from non-canonical mixing as estimated from the carbon isotopic ratio data (Charbonnel \\& Do Nascimento 1998). The classical value of the incidence of Ap/Bp stars, which is $5-10\\%$ based on a magnitude-limited sample (Wolff 1968), has been challenged by the aforementioned survey performed by Power et al.~(2006): in the solar neighborhood they find only 1.7\\% of magnetic stars. This value is close to that (i.e., $\\sim 3.5 \\%$) derived by North (1993) for Ap stars with masses between 1.7 and 2.5~M$_{\\odot}$ that are members of Galactic open clusters, and by Paunzen et al.~(2006; $\\sim 2.2 \\%$) for chemically peculiar stars in young LMC open clusters. We note that the bulk incidence of Ap stars is very sensitive to the chosen value of the lower mass limit for these objects in a given sample (see Fig.~4 of Power et al.~2006 or Table VII of North~1993). In conclusion, we consider that the present statistics support our suggestion that thermohaline mixing is inhibited by a fossil magnetic field in those red giant stars that are the descendants of main sequence Ap stars. Within this framework we reconcile the long-standing problem of $^3$He overproduction on the Galactic timescale with the high values of $^3$He observed in a couple of planetary nebulae. We encourage the observers to look for magnetic fields in ``classical\" RGB and in NGC~3242 and J320 which we propose to call ``thermohaline deviant stars\"."
    },
    "0711/0711.3730_arXiv.txt": {
        "abstract": "% We calculate the formation of dust clouds in atmospheres of giant gas-planets.  The chemical structure and the evolution of the grain size distribution in the dust cloud layer is discussed based on a consistent treatment of seed formation, growth/evaporation and gravitational settling. Future developments are shortly addressed. ",
        "introduction": "Clouds are a common feature in gas giant and brown dwarf atmospheres, and observation start to spectroscopically infer their existence.  The description of the actual formation of clouds in substellar atmospheres has been a challenge in particular if consistently calculated with radiative transfer and large-scale hydrodynamics. Various branches pressed ahead in the recent years to advanced the field of planet atmosphere modelling: i) detailed radiative transfer calculations with simplified cloud chemistry ({\\it Fortney} and {\\it Barman} this volume), ii) large-scale multi-dimensional hydrodynamical simulations but without dust chemistry ({\\it Showman} and {\\it Menou}, this volume), iii) photochemical gas-phase models for a given atmosphere stratification ({\\it Tinetti}, this volume), and iv) detailed non-equilibrium dust-cloud modelling for a prescribed atmosphere structure. Our efforts are centred on point iv) which allows us to study details of the cloud formation process like e.g. the cloud's chemical composition and the grain size distribution function.  The basic idea is that a phase-transition can only take place if the gas is supersaturated. Supersaturation requires a considerable super-cooling below the thermal stability temperature. Once seed particles have formed (nucleation), chemical surface reactions quickly form a mantle. The grain mantle consists of a variety of compounds in an oxygen-rich environment because compounds can become thermally stable in a very narrow temperature interval.  The newly formed dust particles drift into deeper layers due to the high gravity of the objects and grow further on their way into warmer regions. If the temperature becomes too high, the compounds are not any more thermally stable and, hence, they evaporate. Would these processes occur in a truly static environment, no clouds would be present (Woitke \\& Helling 2004)\\nocite{Woitke2004}. But substellar objects posses a considerable convective over-shooting (Ludwig et al. 2006) which brings up uncondensed material and which keeps the circle of dust formation running.  Cloud formation due to the condensation of solids/liquids reduces the remaining gas phase element abundances and their gravitational settling introduces an additional dynamic process, both changing the atmospheres appearance considerably.  In the following, we will demonstrate the results of our kinetic dust-cloud model for gas-giant atmospheres and shortly address some future developments.   \\begin{figure}[h] \\centering \\epsfig{file=VolFrac13003.0.eps, scale=0.9}\\\\[0.2cm] \\epsfig{file=GrainSizseDist13003.0_solar_Exp_small.eps, scale=0.9} \\vspace*{-0.1cm} \\caption{Atmospheric dust cloud in a giant gas-planet with  T$_{\\rm eff}=1300\\,$K, $\\log$\\,g$=3.0$. {\\bf Top:} Material composition (in volume fractions) of the dust cloud. {\\bf Bottom:} Grain size distribution functions for each atmospheric layer in the dust cloud. } \\label{meanSize} \\vspace*{-0.3cm} \\end{figure} ",
        "conclusions": ""
    },
    "0711/0711.1735_arXiv.txt": {
        "abstract": "{More and more evidence lead to consider classical Be stars as stars rotating close to the critical velocity. If so, then the question which naturally arises is the origin of this high surface velocity.} {We determine which are the mechanisms accelerating the surface of single stars during the Main Sequence evolution. We study their dependence on the metallicity and derive the frequency of stars with different surface velocities in clusters of various ages and metallicities.} {We have computed 112 stellar models of four different initial masses between 3 and 60 $M_{\\sun}$, at four different metallicities between 0 and 0.020, and with seven different values of the ratio $\\Omega/\\Omega_\\mathrm{crit}$ between 0.1 and 0.99. For all the models, computations were performed until either the end of the Main Sequence evolution or the reaching of the critical limit.} {The evolution of surface velocities during the Main Sequence lifetime results from an interplay between meridional circulation (bringing angular momentum to the surface) and mass loss by stellar winds (removing it). The dependence on metallicity of these two mechanisms plays a key role in determining for each metallicity, a limiting range of initial masses (spectral types) for stars able to reach or at least approach the critical limit. Present models predict a higher frequency of fast rotating stars in clusters with ages between 10 and 25 Myr. This is the range of ages where most of Be stars are observed. To reproduce the observed frequencies of Be stars, it is necessary to assume first that the Be star phenomenon occurs already for stars with $\\upsilon/\\upsilon_\\mathrm{crit}\\ge 0.7$ and second, that the fraction of fast rotators on the Zero Age Main Sequence is higher at lower metallicities. Depending on the stage at which the star becomes a Be star, the star at this stage may present more or less important enrichments in nitrogen at the surface.} {} ",
        "introduction": "The aim of the present work is to study how the surface velocity varies as a function of time during the Main Sequence (MS) evolution for stars of various initial masses, metallicities and rotation. Although predictions of the evolution of the surface velocity can already be found in literature \\citep[see][]{Hel00,MMV, MMXI}, the published results lack homogeneity in their physical ingredients and cover too narrow ranges of initial velocities to provide valuable answers to the following questions: \\begin{itemize} \\item When is very fast rotation encountered in the course of the evolution of massive single stars on the MS phase? \\item What is the percentage of stars rotating faster than a given value of $\\upsilon_\\mathrm{eq}/\\upsilon_\\mathrm{crit}$ in stellar clusters of various ages and metallicities? \\item What are the expected chemical enrichments  at the surface of these very fast rotating stars? \\item What are the initial conditions required for single stars to reach the critical velocity limit (defined as the rotational velocity such that the centrifugal acceleration exactly balances the gravity at the equator)? \\item Do these conditions depend on the mass and metallicity and if yes how? \\end{itemize}  To know which are the stars which can encounter the critical limit during the MS evolution is interesting in at least two astrophysical contexts: first, the study of Be stars, second, the very metal-poor stars representative of the stellar generations in the early Universe. Let us briefly describe the link with these two research areas.  \\subsection{Be stars and the critical limit}  It is now widely accepted that the ``Be phenomenon'' is linked to fast rotation. Indeed, classical Be stars are B-type stars that exhibit line emission over the photospheric spectrum. This emission probably arises from an outwarding equatorial disk \\citep{PR03} whose formation is (at least in part) due to the fast rotation of the star. If the star is rotating with $\\upsilon_\\mathrm{eq}/\\upsilon_\\mathrm{crit} \\sim 0.95$, the additional velocity necessary to launch material into orbit is of the same order as the speed of sound in the outer layers \\citep{Town04}. Therefore many instabilities, as {\\it e.g.} non-radial pulsations, begin to be effective for orbital ejection \\citep{Owocki}. Very interestingly, recent interferometric observations brought strong support to the idea that these stars are very fast rotators: the oblateness of the Be star Achernar \\citep{domi} suggests indeed that the star is rotating near the critical limit; \\citet{Mei07} obtained for the first time direct evidence that the disk around the Be star $\\alpha$ Arae is in Keplerian rotation. The wind geometry is compatible with a thin disk wind and a polar enhanced wind. These authors also found that $\\alpha$ Arae is rotating very close to its critical rotation. Even the mystery behind the $\\beta$ Ceph case seems to be elucidated: $\\beta$ Ceph is a known Be star but also a known slow rotator, thus challenging the fast rotation scenario. However $\\beta$ Ceph is a multiple system. \\citet{Sch06} have shown that $\\beta$ Ceph, the slow rotator, produces most of the light but is not responsible for the emission lines. The emission lines originate from a much fainter companion, the Be star, which they consider as probably fast rotating.  If the above observations do support the view that Be stars are at the critical limit, there are still some doubts that all Be stars are indeed at this limit. For instance \\citet{Cra05}, on the basis of the observed distribution of the $\\upsilon\\sin i$ for Be stars, conclude that while late type Be stars could be at the critical limit, early Be type stars might rotate at only 40\\% of their critical velocity. Let us note that the interpretation of the observations is stained with some theoretical uncertainties. A good example is the way the deduced velocities are affected by the fact that for a fast rotating star most of the radiative flux originates from the polar regions \\citep{vZ24}. The low velocity regions, near the pole, will thus have a stronger impact on the spectrum than the dimmer faster regions near the equator. This implies that the rotational velocity of very fast rotators have been probably systematically underestimated \\citep{coll04,Town04}. \\citet{Cra05} has accounted for this effect in his work, but the problem is sufficiently complex that the results are uncertain.  Anyway, at the critical limit or not, Be stars represent wonderful laboratories for studying the effects of extreme rotation.  \\subsection{Very metal poor stars and the critical limit}  The link between metal-free or very metal-poor stars and the question of the evolution of the surface velocity certainly deserves great attention, since it might deeply affect the evolution of these stars. Let us recall that at low metallicity, the radiation driven stellar winds are much weaker than at high metallicity, thus it is commonly accepted that mass loss for very metal-poor stars is very inefficient. This might not be true however when rotation is taken into account in the models for at least two reasons \\citep[see][to get a more detailed discussion of these points]{NICIX}: first and somewhat paradoxically, since the stars are losing very little amount of mass by radiatively driven stellar winds, they have more chance to lose large amount of mass by rotational ejection during the MS evolution. Indeed, because of the low line driven winds, they will not be able to lose their angular momentum. They are thus more prone to reach the critical limit during their evolution \\citep{MMVII,MEM06}. At the critical limit, as in the Be case, stellar matter is launched into an equatorial disk which will eventually dissipate by radiative effects. In that case the material will be lost by the star. If correct, such an evolution would thus promote mass loss through a mechanical wind triggered by rotation. This ``mechanical wind'', as stressed above, appears to be more important in metal-poor regions where the radiation driven stellar winds are expected to be very low. Second, rotational mixing enhances the surface CNO content which will lead to some increase of the outer layers opacities. This may revive the line driven winds and lead to strong mass loss.  Such mass loss, either triggered by rotational ejection or through the enhancement of the CNO content of the outer layers, have (if realized) important consequences on the evolution of the first stellar generations as well as on their nucleosynthesis \\citep{EkTar,Hi07,De06}. Also such mass loss may be important in the process leading to the formation of a collapsar which is the most favoured model for the progenitors of the long Gamma Ray Bursts \\citep{W93,YL05,WH06,MMGRB07}.  In the present work, we use rotating stellar models to address the questions of: \\begin{itemize} \\item the possible origin of Be stars from a purely theoretical point of view: can models reproduce the way the number of Be stars varies with the metallicity or with the age? Is their number sufficient to account for the observed number of Be stars (supposing that Be stars are star rotating with values of $\\upsilon_\\mathrm{eq}/\\upsilon_\\mathrm{crit}$ superior to a given value)? \\item the behaviour of rotating stars at very low or zero metallicity: how does the lack of metals influence the evolution of the equatorial velocity? \\end{itemize} Of course the answers given to these questions will depend strongly on the stellar models used. However we think that the exercise deserves to be undertaken for the following reason: the stellar models we use were computed with physical ingredients that made predictions consistent with many observational features of massive stars \\citep[see][and the references therein on the series of papers by the Geneva group on rotating models]{MMXI} and thus it is interesting to explore in more detail what these models say about fast rotating stars. For very fast rotation, the present models probably need still to be improved by the incorporation of additional physics (e.g. pulsation), and we hope that the results obtained here will guide us towards improvements in the future.  The paper is organized as follows: in Sect.~\\ref{inpphys} the physical ingredients of the models are presented. We also recall some basic relations from the Roche model useful for discussing some properties of our models on the ZAMS. In Sect.~\\ref{zams}, we present various characteristics of our models on the ZAMS as a function of the initial mass, velocity and metallicity. The evolution of the models on the MS is presented in Sect.~\\ref{evolms}. The question of the initial conditions required to reach a given value of $\\upsilon_\\mathrm{eq}/\\upsilon_\\mathrm{crit}$ during the MS phase is the subject of Sect.~\\ref{bul}. The frequency of stars with velocities exceeding a given value of $\\upsilon_\\mathrm{eq}/\\upsilon_\\mathrm{crit}$ is presented in Sect~\\ref{freqbe}. The sensitivity of the predicted frequencies on various hypotheses made in this work are discussed in Sect.~\\ref{discuss}. Possible links with observational features of Be stars will be presented in Sect~\\ref{becomp}. Conclusions and future perspectives are the subject of Sect.~\\ref{conclu}.  Let us mention that electronic tables are available for all our ZAMS and evolutionary models at \\url{http://obswww.unige.ch/Recherche/evol/Critical-limit-and-Be-stars} ",
        "conclusions": "} From the above computations we can deduce that \\begin{enumerate} \\item As long as the mass loss rates are not too important, meridional currents, by advecting angular momentum from the inner regions to the outer layers, bring the $\\upsilon/\\upsilon_\\mathrm{crit}$ ratio close to 1 during the MS phase. \\item Since the velocity of the meridional currents in the outer layers scales with the inverse of the density, the process becomes more efficient for stars of higher initial mass and/or higher initial metallicity. \\item At high metallicity however, mass loss becomes more and more important and can prevent the stars from reaching the break--up limit. \\item In stellar clusters, one expects that the fraction of stars having $\\upsilon/\\upsilon_\\mathrm{crit} \\ga 0.8$ becomes maximum for ages between 20-32 Myr at $Z=0.020$. This range of ages shifts to older ages at higher metallicities and to younger ages at lower metallicities. Such a shift results from the dependence on the metallicity and the mass of the intensity of the mass loss rates and of the velocities of the meridional currents. If we relax the condition and look at stars with $\\upsilon/\\upsilon_\\mathrm{crit} \\ga 0.7$, the maximum number will be found in clusters of ages around 10 Myr and below, and this at all non-zero metallicities. \\item Be stars might be the natural outcome of stars with an initial rotational velocity in the upper tail of the initial velocity distribution. To explain their frequency from the present rotating stellar models, minimum values of $\\upsilon/\\upsilon_\\mathrm{crit} \\ga 0.7$ have to be reached during the MS for the Be star phenomenon to occur. \\item Depending on when this limit is reached, one expects more or less high surface enrichments. If the limit is reached very early during the MS phase, no enrichment is expected, while if this limit is reached at the end of the MS phase, high N/C and N/O ratios are expected. For Be stars originating from small initial mass stars (typically from 3 $M_{\\sun}$ at standard metallicity) one expects no or small surface enrichments. \\item To reproduce the higher fractions of Be stars at low metallicity, we must assume that in metal poor regions, a higher number of stars are born with high values of $\\Omega/\\Omega_\\mathrm{crit}$. \\end{enumerate} In a future work, it will be interesting to explore the effects of a magnetic field. As already mentioned above, this will probably change significantly the results presented here, allowing the stars to approach more easily the critical limit. It will be interesting to see if, in that case, it would still be necessary to invoke a higher proportion of fast rotators at low $Z$."
    },
    "0711/0711.4110_arXiv.txt": {
        "abstract": "The dissipation of turbulent gas motions is one of the likely mechanisms that has been proposed to heat the intracluster medium (ICM) in the cores of clusters and groups of galaxies. We consider the impact of gas motions on the width of the most prominent X-ray emission lines. For heavy elements (like iron) the expected linewidth is much larger than the width due to pure thermal broadening and the contribution due to turbulent gas motions should be easily detected with the new generation of X-ray micro-calorimeters, such as the Spektr-RG calorimeter (SXC).  For instance in the Perseus cluster the turbulent velocity required to balance radiative cooling (as derived by Rebusco et al. 2006), would imply a width of the 6.7 keV Fe line of 10-20 eV, while the pure thermal broadening is $\\sim$4 eV. The radial dependence of the linewidth is sensitive to i) the radial dependence of the velocity amplitude and ii) the \"directionality\" of the stochastic motions (e.g. isotropic turbulence or predominantly radial gas motions). If the width of several lines, characteristic for different gas temperatures, can be measured, then it should be possible to probe both the \"directionality\" and the amplitude of the gas motions. Moreover a measurement of the width would put a lower limit on the amount of the  kinetic energy available for dissipation, giving a constraint on the ICM models. ",
        "introduction": "Clusters of galaxies are the largest gravitationally bound and nearly virialized systems in the Universe. High resolution X-ray surveys have revealed that the ICM (with temperatures in the range $\\sim 2-10$ keV) is not fully relaxed.  The substructures in the surface brightness and temperature are observed both on large scales and in the cluster cores indicating that the gas is not at rest. At the same time strong shocks are rarely observed (see Markevitch \\& Vikhlinin, 2007 for a review) suggesting that the cluster gas motions are predominantly subsonic. The shape of the emission lines as a source of information on the ICM velocity distribution has been discussed in detail in Sunyaev, Norman \\& Bryan (2003) and Inogamov \\& Sunyaev (2003).  These papers mostly consider the gas motions induced by cluster mergers. In the present paper we focus on the linewidth as a diagnostic of gas motions in the cores of the galaxy clusters.  In the cores of most clusters and many groups (e.g. Stewart et al 1984, Nulsen et al 1984, Edge \\& Stewart 1991, Mulchaey et al 1993, Fabian 1994), the radiative cooling time of the gas is short compared to the age of the cluster and an external source of energy is needed to avoid catastrophic cooling (e.g. David et al. 2001, Matsushita et al. 2002, Peterson et al. 2003, Kaastra et al. 2004). The dissipation of turbulent gas motions is one of the highly plausible sources of the required energy (e.g. Pedlar et al. 1990, Churazov et al. 2002, Fabian et al. 2003a, Chandran 2005). These motions can have distinctive effects on the properties of the X-ray lines (e.g. Br\\\"{u}ggen, Hoeft \\& Ruszkowski 2005). There are several indirect ways of probing the gas velocities, i) using the resonant scattering of the brightest emission lines (e.g. Gilfanov, Sunyaev \\& Churazov, 1987, Churazov et al., 2004); ii) considering the spreading of metals ejected by a central galaxy through the ICM (Rebusco et al., 2005,2006, Graham et al., 2006, Roediger et al., 2007); iii) using H$\\alpha$ emitting filaments as tracers of the gas motions (e.g. Fabian et al., 2003b, Hatch et al., 2006). In particular in Rebusco et al. (2005,2006) the characteristic spatial and velocity scales of the turbulent eddies were estimated assuming that the dissipation of turbulent motions can generate enough heat to prevent the gas from extreme cooling and that the same motions provide moderate mixing of metals through the ICM. Using these constraints we calculated the expected width of the brightest X-ray lines for the brightest clusters (e.g., A426 - Perseus).   A micro-calorimeter with energy resolution of less than $10$ eV and effective area of $\\sim~200~$ cm$^2$ has been built for the Suzaku mission (Kelley et al., 2007).  The effective areas of the upcoming calorimeters may reach thousands and even tens of thousands of square cm in the next 10-15 years and will have a spatial resolution between a few arcseconds and a few arcminutes (see e.g. \\verb#http://www.astro.isas.ac.jp/future/NeXT#, \\verb#http://www.rssd.esa.int/index.php?project=XEUS#,  \\verb#http://constellation.gsfc.nasa.gov/# ).  The Spektr-RG Calorimeter (SXC) (McCammon et al., 2007; Mitsuda et al., 2007), proposed for the Spektr-RG Mission, has a peak effective area of 300 cm$^2$ at $\\sim 1.5$ keV, decreasing to about $30~$ cm$^2$ at $6.7$ keV and an energy resolution (goal) of $\\sim 4$ eV.  As we show in Section $3$, an instrument with these characteristics is well suited to study the line broadening associated with turbulent gas motions in the cores of the brightest galaxy clusters.  The structure of the paper is the following.  In section \\ref{sec:turb} we describe the line broadening under different assumptions for the character of the gas motions; in section 3 we consider the specific case of the Perseus cluster as an example and we investigate the capabilities of SXC in this context. Our findings are discussed and summarized in sections 4 and 5.  Throughout the paper when we mention turbulence, we mean all types of macroscopic gas motions. ",
        "conclusions": "\\label{sec:conclusions}  The dissipation of turbulent gas motions driven by AGN activity is a plausible source of heat for the cooling gas in cluster cores. In this case, we show that the expected width of the iron 6.7 keV line is well above the thermal broadening for most plausible values of the turbulent velocity. In our fiducial model of the Perseus cluster, the 6.7 keV linewidth is larger than 10 eV. Hence the proposed SXC micro-calorimeter is well suited for firmly establishing the level of microturbulence in the brightest galaxy clusters.\\\\ As discussed in Sections 2.1 and 2.2 the linewidth is sensitive to both the radial dependence of the velocity amplitude and the \"directionality\" of the stochastic motions. This may complicate an unambiguous determination of the properties of the gas motions. However this degeneracy can be removed if the width of several lines, characteristic for different gas temperatures, can be measured. Assume for example that in addition to the 6.7 keV iron line, the width of the iron L-shell line is detected. This line is bright at lower temperatures, that are present only in the central region (cool core). Therefore the combined measurement of the width of two or more lines should be sufficient to discriminate among the simplest models with different \"directionalities\"."
    },
    "0711/0711.0443_arXiv.txt": {
        "abstract": "The deconfinement phase transition will lead to the release of latent heat during spins down of neutron stars if the transition is the first-order one. We have investigated the thermal evolution of neutron stars undergoing such deconfinement phase transition. The results show that neutron stars may be heated to higher temperature.This feature could be particularly interesting for high temperature of low-magnetic field millisecond pulsar at late stage. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.0396_arXiv.txt": {
        "abstract": "\\begin{center} {\\bf ABSTRACT} \\end{center} We study the possible phenomenology of a three-family Pati-Salam model constructed from intersecting D6-branes in Type IIA string theory on the $\\mathbf{T^6/(\\Z_2\\times \\Z_2)}$ orientifold with some desirable semi-realistic features. In the model, tree-level gauge coupling unification is achieved automatically at the string scale, and the gauge symmetry may be broken to the Standard Model (SM) close to the string scale. The small number of extra chiral exotic states in the model may be decoupled via the Higgs mechanism and strong dynamics. We calculate the possible supersymmetry breaking soft terms and the corresponding low-energy supersymmetric particle spectra which may potentially be tested at the Large Hadron Collider (LHC). We find that for the viable regions of the parameter space the lightest CP-even Higgs boson mass usually satisfies $m_H \\leq 120$~GeV, and the observed dark matter density may be generated. Finally, we find that it is possible to obtain correct SM quark masses and mixings, and the tau lepton mass at the unification scale.  Additionally, neutrino masses and mixings may be generated via the seesaw mechanism. Mechanisms to stabilize the open and closed-string moduli, which are necessary for the model to be truly viable and to make definite predictions are discussed. ",
        "introduction": "Although string theory has long teased us with her power to encompass all known physical phenomena in a complete mathematical structure, an actual worked out example is  still lacking. Indeed, the major problem of string phenomenology is to construct {\\it at least one} realistic model with all moduli stabilized, which completely describes known particle physics as well as potentially being predictive of unknown phenomena. With the dawn of the Large Hadron Collider (LHC) era, new discoveries will hopefully be upon us. In particular, supersymmetry is expected to be found as well as the Higgs states required to break the electroweak symmetry. Therefore, it is highly desirable to have complete, concrete models derived from string theory which are able to make predictions for the superpartner spectra, as well as describing currently known particle physics.  In the old days of string phenomenology, model builders were primarily focused on weakly coupled heterotic string theory. However, with the advent of the second string revolution, D-branes~\\cite{JPEW} have created new interest in Type I and II compactifications.  In particular, Type IIA orientifolds with intersecting D6-branes, where the chiral fermions arise at the intersections of D6-branes in the internal space~\\cite{bdl}, with T-dual Type IIB description in terms of magnetized D-branes~\\cite{bachas}, have shown great promise during the last few years.  Indeed, intersecting D-brane configurations provide promising setups which may accommodate semi-realistic features of low-energy physics. Given this, it is an interesting question to see how far one can get from a particular string compactification to reproducing the finer details of the Standard Model as a low-energy effective field theory.  In order to construct globally consistent vacua with intersecting D-branes, conditions must be imposed which strongly constrain the models.  In particular, all Ramond-Ramond (RR) tadpoles must be cancelled and K-theory \\cite{Witten9810188} conditions for cancelling the nontrivial $Z_2$ anomally also must be imposed. Despite the clear benefits of supersymmetry, there have been many three-family standard-like models and Grand Unified Theories (GUT) constructed on Type IIA orientifolds~\\cite{Blumenhagen:2000wh, Angelantonj:2000hi, Blumenhagen:2005mu} which are not supersymmetric.  Although these models are globally consistent, they are generally plagued by the gauge hierarchy problem and vacuum instability which arises from uncancelled Neveu-Schwarz-Neveu-Schwarz (NSNS) tadpoles.  Later, semi-realistic supersymmetric Standard-like, Pati-Salam, unflipped $SU(5)$ as well as flipped $SU(5)$ models in Type IIA theory on $\\mathbf{T^6/(\\Z_2\\times \\Z_2)}$~\\cite{CSU1, CSU2, Cvetic:2002pj, CP, CLL, Cvetic:2004nk, Chen:2005ab, Chen:2005mj} and $\\mathbf{T^6/(\\Z_2\\times \\Z_2')}$~\\cite{Dudas:2005jx, Blumenhagen:2005tn, Chen:2006sd} orientifolds were eventually constructed, and some of their phenomenological consequences studied~\\cite{CLS1, CLW}. Other supersymmetric constructions in Type IIA theory on different orientifold backgrounds have also been discussed~\\cite{ListSUSYOthers}.  Nonperturbative D-instanton effects have also been receiving much attention of late, and may play an important role~\\cite{Blumenhagen:2006xt}~\\cite{Ibanez:2006da}~\\cite{Cvetic:2007ku}~\\cite{Ibanez:2007rs}.  In addition to satisfying the above consistency conditions, all open and closed-string moduli must be stabilized in order to obtain an actual vacuum.  Unstabilized moduli are manifest in the low-energy theory as massless scalar fields, which are clearly in conflict with observations. Given a concrete string model, the low-energy observables such as particle couplings and resulting masses are functions of the open and closed string moduli. In a fully realistic model, these moduli must therefore be stabilized and given sufficiently large masses to meet the astrophysical/cosmological and collider physics constraints on additional scalar fields.  Although satisfying the conditions for $\\mathcal{N}=1$ supersymmetry in Type IIA (IIB) fixes the complex structure (K\\a\"ahler) moduli in these models, the K\\a\"ahler (complex structure) and open-string moduli generally remain unfixed. To stabilize some of these moduli, supergravity three-form fluxes \\cite{Kachru:Blumen} and geometric fluxes \\cite{Grimm:Zwirner} were introduced and  flux models on Type II orientifolds have been constructed~\\cite{Cascales:2003zp, MS, CL, Cvetic:2005bn, Kumar:2005hf, Chen:2005cf, Blumenhagen:2006ci, Camara:2005dc, Chen:2006gd, Chen:2006ip}.  Models where the D-branes wrap rigid cycles, thus freezing the open-string moduli have also been studied~\\cite{Dudas:2005jx, Blumenhagen:2005tn, Chen:2006sd}.  Despite substantial progress, there have been other roadblocks in constructing phenomenologically realistic intersecting D-brane models, besides the usual problem of moduli stabilization.  Unlike heterotic models, the gauge couplings are not automatically unified.  Additionally, there has been a rank one problem in the Standard Model (SM) fermion Yukawa matrices, preventing the generation of mass for the first two generations of quarks and leptons. For the case of toroidal orientifold compactifications, this can be traced to the fact that not all of the Standard Model fermions are localized at intersections on the same torus. However, one example of an intersecting D6-brane model on Type IIA $\\mathbf{T^6/(\\Z_2\\times \\Z_2)}$ orientifold has recently been discovered where these problems may be solved~\\cite{CLL,Chen:2006gd}.  Thus, this particular model may be a step forward to obtaining realistic phenomenology from string theory. Indeed, as we recently discussed ~\\cite{Chen:2007px}, it is possible within the moduli space of this model to obtain correct quark mass matrices and mixings, the tau lepton mass, and to generate naturally small neutrino masses via the seesaw mechanism. Furthermore, it is possible to generically study the soft supersymmetry breaking terms, from which can be calculated the supersymmetric partner spectra, the Higgs masses, and the resulting neutralino relic density.   This paper is organized as follows. First, we will briefly review the intersecting D6-brane model on Type IIA $\\mathbf{T^6/(\\Z_2\\times \\Z_2)}$ orientifold which we are studying and discuss its basic features.  We then discuss the low-energy effective action, and show that the tree-level gauge couplings are unified near the string scale. We also find that the hidden sector gauge groups will become confining at a high energy scale, thus decoupling chiral exotics present in the model. Next, we study the possible low-energy superpartner spectra which may arise. We also calculate the Yukawa couplings for quarks and leptons in this model, and show that we may obtain the correct quark masses and mixings and the tau lepton mass for specific choices of the open and closed string-moduli VEVs.  We should emphasize that for the present work, we will not focus on the moduli stabilization problem, as our goal is only to explore the possible phenemological characteristics of the model. However, we do comment on this issue and discuss how it may potentially be solved for this model. We also should note that models with an equivalent observable sector have been constructed in Type IIA and Type IIB theory as Ads and Minkowski flux vacua~\\cite{Chen:2006gd, Chen:2007af}, so that the issue of closed-string moduli stablization has already been addressed to some extent. ",
        "conclusions": "We have analyzed in detail a three-family intersecting D6-brane model where gauge coupling unification is achieved at the string scale and where the gauge symmetry can be broken to the Standard Model. In the model, it is possible to calculate the supersymmetry breaking soft terms and obtain the low energy supersymmetric particle spectra within the reach of the LHC. Finally, it is possible to obtain the SM quark masses and CKM mixings and the lepton masses, and the neutrino masses and mixings may be generated via the seesaw mechanism.  Clearly, this model cannot be regarded as being fully realistic until the moduli stablization issue has been fully addressed. There are many free parameters which have been fixed in order to obtain the desired values for the Yukawa mass matrices and the value of the gauge couplings at the unification scale, although it should be kept in mind that these parameters are tightly constrained and it is not possible to tune them to just any value. In the case of the Yukawa matrices, the free parameters are the K\\a\"ahler moduli, the brane positions on each torus (open-string moduli) and the specific linear combination of states with which have been indentified the two pairs of Higgs eigenstates. Although we have chosen specific values for the moduli fields to obtain agreement with experiments, it may be possible to uniquely predict these values by introducing the most general fluxes.  It might also be possible to fix the open string moduli if we require the D-branes wrap rigid cycles. However, it seems likely that the Higgs eigenstates would still need to be fine-tuned.  On the other hand, this does appear to be the first such string-derived model where it is possible to give mass to each family of quarks and leptons.  Even if we cannot at present {\\it uniquely} predict these values, it must still be regarded as highly significant in that it is {\\it possible} to come very close to getting correct mass matrices and mixings at the unification scale. This suggests that the model may be a candidate for a phenomenological description of elementary particle physics in much the same way as the MSSM.  It is also very appealing that the tree-level gauge couplings are unified at the string scale, although it is still an open question if the running of the gauge couplings can be maintained all the way down to the electroweak scale.  The reason for this is that there are chiral exotic states present in the spectrum which are bifundamentals under the observable and hidden sector gauge groups. We should note that most of these chiral exotic states can be decoupled at the string scale and the rest may may be decoupled at an intermediate scale.  Even if this were not the case, we have found that the hidden sector gauge interactions will become confining at around $10^{7}$~GeV and $10^{13}$~GeV respectively, and so states charged under these groups will not be present in the low-energy spectrum.  However, there are exotic states in the spectrum which transform as representations of both the hidden and observable sectors which may effect the RGE running of the gauge couplings, although we do not expect them to effect the running that much.  If we are optimistic, then it is possible that these would amount to threshold corrections which might push the unification scale up to the string scale.  If the model does turn out to be a realistic effective description of the observed elementary particle physics, then it should be possible to predict the low-energy superpartner mass spectra from the model.  Besides the D-brane wrapping numbers and closed-string moduli, the superpartner spectra depend strongly on the exact way in which supersymmetry is broken. In principle, it should be possible to completely specify the exact mechanism, whether through gaugino condensation in the hidden sector, flux-induced soft terms, or via instanton induced couplings. In the present analysis, we have studied supersymmetry breaking generically via a generic parameterized F-term and have shown that it is possible to constrain the phenomenologically allowed parameter space by imposing experimental limits on the neutralino relic density and mass limits coming from LEP. We have found that the viable parameter space is quite large.  Once the experimentally determined superpartner mass spectrum begins to take shape it may be possible to find a choice of F-terms which will correspond to the observed spectrum.  It would be very interesting to explore the collider signatures which this model may produce at LHC for the regions of the parameter space which satisfy all phenomenological constraints.  In summary, the model we have studied may produce a realistic phenomenology, once the issue of moduli stabilization has been fully addressed.  The model represents the first known intersecting D-brane model for which mass may be given to each generation of quarks and leptons.  Furthermore, the supersymmetry breaking soft terms may be studied in the model and may yield realistic superpartner spectra.  Certainly, the model and the current theoretical tools are not presently developed to the point where specific predictions for the known fermion masses and superpartner spectrum may be made.  However, it is clearly possible for the model to describe the known physics of the Standard Model, as well as potentially describing new physics."
    },
    "0711/0711.2393_arXiv.txt": {
        "abstract": "Treating protons as de Broglie waves shows that up to a few MeV energies protons experience total external reflection using the index of refraction concept for the target earlier applied to electrons. Angular scattering distributions can be explained by random surface scattering as known for X-rays. Applied to the {\\it{Chandra}} and {\\it{XMM-Newton}} X-ray telescopes the calculated reflection efficiencies can explain the observed degradation of the X-ray CCDs for both missions. Some discussion about the possibility of realizing imaging sub-MeV and MeV proton optics is presented. ",
        "introduction": "\\label{sec:intro}  % The X-ray astronomy observatory {\\it{Chandra}} was launched into orbit in July 1999. Very early in the mission a performance degradation, predominantly with the front illuminated CCD detector ACIS-FI was noticed. Apparently some damage was done to the CCD which resulted in degradation of spectral performance due to an increase of charge transfer inefficiency (CTI)$\\sp 1$. Investigation of the problem revealed that the most likely source for the damaging were protons with energies between 100 keV and a few MeV$\\sp 2$. Analysis also showed that the only path for protons getting to the focal plane is through the mirror system of the X-ray telescope as the structure has more than sufficient shielding against such protons. The X-ray astronomy observatory {\\it{XMM-Newton}} at that time was prepairing for the flight readiness review (FRR) in October 1999, and the team was badly concerned because the observatory carrys similar CCDs. Meanwhile measures have been taken to prevent the protons reaching the CCDs, like moving the device out of the focal plane ({\\it{Chandra}}) or cover it with a solid absorber implemented on the filter wheel in front ({\\it{XMM-Newton}} EPIC), where possible.  The scattering of protons off the mirror reflecting surfaces came totally unexpected, and the question is what physical process is at work. Rutherford scattering was considered because of the very pronounced forward scattering, and scattering by an angle of something like 0.5$\\sp{\\circ}$ to 1$\\sp{\\circ}$ is just needed. However, for a proton of 100 keV energy to get scattered by 0.5$\\sp{\\circ}$ it needs an electrical field of about 900 V. This is equivalent to the binding energies of the electrons in the N-shells of a Au atom, whereby Au is the coating of the {\\it{XMM-Newton}} mirrors. Most of the interatomic space is found outside of the N-shell region. Therefore it is much more likely that the scattering angle is realized by multiple scattering processes in the coating. Computer based Monte-Carlo routines like TRIM and GEANT4 have been available but had to be substantially modified to simulate multi-scattering of a charged particle in matter under grazing incidence$\\sp{2, 3}$. The dedicated measurements of proton scattering off an {\\it{XMM-Newton}} representative flat mirror sample undertaken by Rasmussen et al.$\\sp{4}$ are quite well reproduced by the GEANT4 simulations$\\sp{4}$ within a factor of 2 to 3 as far as the efficiency is concerned. I note here that the telescopes employ two reflections or scatterings, which might increase the difference to some larger level. Details of the scattering distributions are not reproduced like a turn down of the efficiency towards low incidence angles. Measurements show a maximum at some scattering angle for fixed incidence angle$\\sp{4}$. However, at the very low angles the measurements might have had some problems, so that more detailed studies are needed to confirm the trend.  In this paper I would like to present a totally different approach to the proton scattering phenomenon using the dualism between particle and wave, treating the proton as a wave. Much of the material has been presented at several {\\it{XMM-Newton}} emergency meetings in October 1999 prior to FRR. ",
        "conclusions": "Describing the proton as de Broglie wave in conjunction with an index of refraction less than one, predicts total external reflection of protons. Using the values for the inner potential observed in electron measurements allows, together with the Fresnel equations, computing of reflectance values. Total reflection at incidence angles of less than one degree occurs for protons with energies less than a few MeV. Ordinary targets, even fairly well polished mirrors, are likely to have surface mircroroughness values which lead to scattering distributions one degree or more wide. Targets with much lower microroughness for spatial wavelengths in the sub-micron region are needed, which might be realized with crystals or other materials, to use this effect for building proton optics or a proton scatterometer, with which structures in the nanometer range might be revealed. The total external reflection occurs not only for protons but for any positively charged particle, including positrons, muons and ions. In principle, there is no limit to the particle energy for total external reflection inherent to the theory presented but some experimental practical limits are coming up with the total external reflection angle as small as one arcmin, which would correspond to an energy of about 175 MeV."
    },
    "0711/0711.4676_arXiv.txt": {
        "abstract": "{} {The appearance of folding ion rays in cometary comae is still not very well understood, so our aim is to gain more insight into the role of the local solar wind in the formation of these structures.} {Comet C/2004 Q2 (Machholz) was intensively monitored during its closest approach to Earth (January 2005) with the CCD camera Merope mounted on the Flemish 1.2m Mercator telescope, in three different bands (Geneva U and B and Cousins I). Spectacular ion rays, thin ionic structures rapidly folding tailward, were recorded in the U band during one night, January 12th.} {Data from the SOHO satellite that was extrapolated corotationally to the position of the comet showed that the ion rays were formed during a sudden change in the in-situ solar wind state. We were able to succesfully correlate a high-speed solar wind stream with the appearance of folding ion rays.} {To our knowledge, this is the first clear observational evidence that folding ion rays in cometary comae are produced by a sudden change in the local solar wind state.} ",
        "introduction": "\\begin{figure*} \\includegraphics[width=\\textwidth]{all_together_merged2.pdf} \\caption{Azimuthal renormalised frames of 6 of 8 observations in the U band of comet C/2004 Q2 (Machholz) on 2005 Jan. 12, together with three observations of the night before (upper left), and one on the night after (lower right). The trails on the images are due to background stars and indicate the comet's motion. All images taken on the 12th show clear ion rays, while on the 11th, only the typical sharp ion tail is visible; no ion rays appear. In the right lowerhand corner, it is clear that the ion rays gradually disappear. All images have the same scale and color coding, the direction of the Sun is upwards, and the field of view radius is approximately 4.3\\,$\\times$10$^4$\\,km.}\\label{fig:pdegroote:foldings_comp} \\end{figure*}  A comet can be considered to consist of a small nucleus ($\\sim$10\\,km) surrounded by a dusty gas cloud, dubbed the coma ($\\sim$10$^5$\\,km). On even larger scales, the coma is embedded in the halo, extending up to $\\sim$10$^7$\\,km. Two other typical components of comets in the vicinity of the Sun are the dust tail, consisting of relatively large neutral grains, and the ion tail, made of ionized gas. Basically, the two tails are seperated because the outward solar radiation pressure is higher than the inward gravitational pull. \\citet{bierman51} showed that radiation pressure alone could not be responsible for the outward force on the ions. In this way, the morphology of comets was one of the first proofs for the existence of the solar wind.  Thus, variations in the solar wind result in a variable ion tail. Conversely, through certain properties of the cometary plasma structures, it may be possible to deduce valuable information on the solar wind. For example, a recent survey of Chandra comet observations has demonstrated that X-rays provide a quantitative diagnostics of the interacting solar wind \\citep{bodewits07}. Other properties that have been exploited are tail disconnection events (DEs); see for example \\citet{nieder78}, \\citet{snow04}, and the recent tail rip-off from comet 2/P Encke \\citep{vourlidas07} and statistical studies of the aberration angle of the main tail axis \\citep{brandt73}. This sharply contrasts with the study of cometary folding rays, as they appear closer to the nucleus and thus require a higher spatial resolution. Current theories suggest that folding rays are tailward extensions of ion pile-up regions (receding envelopes) on the sunward side of the nucleus, but it is still not well known what exactly causes this phenomenon. Possible explanations include external origins: some studies invoke the crossing of the comet with IMF (interplanetary magnetic field) sector boundaries, others denote the IMF as just a source mechanism and put more emphasis on the cometary ionosphere. Still other studies claim they have a completely internal origin and are dependent on the nucleus' ion production rate \\citep[see][for a review]{ip04}. This letter supports the idea that the solar wind is largely responsible for the appearance for cometary folding rays, as we found a clear indication that a fast cold wind (originating from coronal holes) triggers tail rays.  We observed comet C/2004 Q2 (Machholz) while it encountered a change in the solar wind regime. Such observations are rare, as there are a lot of variables that have to be set: firstly, we must have accurate solar wind data and detailed images of the comet with an adequate spatial and temporal resolution. Secondly, the observed comet has to be relatively close to the Earth, preferably in or near the ecliptic, to be able to extrapolate solar wind information to the position of the comet (as the solar wind is highly variable in both time and place). Furthermore, a solar event has to occur during the observing run and the effects have to be detected. It is quite a challenge to prepare for such an observation run, as many of these variables are not predictable on a long term and thus constant monitoring of solar wind data is necessary. In this light, the passage of comet C/2004 Q2 (Machholz) in January 2005 was an exceptional event since all of these conditions were set. % ",
        "conclusions": "Comet C/2004 Q2 (Machholz) reached a visual brightness of $\\sim$3.5\\,mag. at its closest approach to Earth on 2005 Jan. 5, making it one of the most spectacular astronomical events of that year. % The comet was visible during a significant part of the night, making it ideally suited to performing a dedicated ground-based monitoring. The first objective of our intensive monitoring % was to derive a rotation period \\citep{reyniers08}. In addition to this, coma structures were also studied in three different bands. During one night, 12 January, clear ion rays were detected in the U band images after enhancement. Due to the comet's position relatively close to the Lagrangian point L1 and close to the ecliptic, we were able to extrapolate the solar wind velocity data from SOHO to the comet's position. It turned out that the development of the ion rays coincide perfectly  with a sudden change in the local solar wind from the ``slow'' ($\\sim$\\,400\\,km\\,s$^{-1}$) to the ``fast'' ($\\sim$\\,700\\,km\\,s$^{-1}$) regime. To our knowledge, this is the first time that the relatively rare event of folding ion rays in the coma of a comet could be correlated with a change of the in-situ solar wind state."
    },
    "0711/0711.3783_arXiv.txt": {
        "abstract": "{Precise determinations of the chemical composition in early B-type stars constitute fundamental observational constraints on stellar and galactochemical evolution. Carbon, in particular, is one of the most abundant metals in the Universe but analyses in early-type stars are known to show inconclusive results. Large discrepancies between analyses of different lines in \\ion{C}{ii}, a failure to establish the \\ion{C}{ii/iii} ionization balance and the derivation of systematically lower abundances than from other indicators like \\ion{H}{ii} regions and young FG-type stars pose long-standing problems.} {We discuss improvements to the non-LTE modelling of the visual line spectrum and to the spectral analysis of early B-type stars, as well as their consequences for stellar parameter and abundance derivations. The most relevant sources of systematic uncertainies and their effects on the analysis are investigated. Consequences for the present-day carbon abundance in the solar vicinity are discussed.} {We present a comprehensive and robust \\ion{C}{ii/iii/iv} model for non-LTE line-formation calculations based on carefully selected atomic data. The model is calibrated with high-S/N spectra of six apparently slow-rotating early B-type dwarfs and giants, which cover a wide parameter range and are randomly distributed in the solar neighbourhood. A self-consistent quantitative spectrum analysis is performed using an extensive iteration scheme to determine stellar atmospheric parameters and to select the appropriate atomic data used for the derivation of chemical abundances.} {We establish the carbon ionization balance for all sample stars based on a unique set of input atomic data. Consistency is achieved for all modelled carbon lines of the sample stars. Highly accurate atmospheric parameters and a homogeneous carbon abundance of $\\log$\\,(C/H)\\,$+$\\,12\\,$=$\\,8.32\\,$\\pm$\\,0.04 with reduced systematic errors are derived. Present evolution models for massive stars indicate that this value may require only a small adjustment because of the effects of rotational mixing, by $<$\\,$+$0.05\\,dex per sample star. This results in a present-day stellar carbon abundance in the solar neighbourhood, which is in good agreement with recent determinations of the solar value and with the gas-phase abundance of the Orion \\ion{H}{ii} region. Our finding of a homogeneous present-day carbon abundance also conforms with predictions of chemical-evolution models for the Galaxy. Moreover, the present approach allows us to constrain the effects of systematic errors on fundamental parameters and abundances. This suggests that most of the difficulties found in previous work may be related to large systematic effects in the atmospheric parameter determination and/or inaccuracies in the atomic data. } {} ",
        "introduction": " ",
        "conclusions": "The motivation of this work was the solution of a long-standing problem in stellar astrophysics: the reliable determination of carbon abundances from early-type stars. For this purpose we constructed a sophisticated \\ion{C}{ii/iii/iv} model atom for non-LTE line-formation calculations based on input atomic data that were carefully selected by an empirical calibration process. This was performed through an extensive iteration scheme that not only allowed us to constrain the input atomic data but also simultaneously the atmospheric parameters of our calibration stars, the basis for all further studies. The calibration sample consists of six bright and apparently slow-rotating early B-type stars in the solar neighbourhood, with high-S/N, high-resolution spectra and a broad wavelength coverage.  The self-consistent analysis provides atmospheric parameters with unprecedented accuracy and with reduced systematic error: for effective temperature the uncertainties are as low as $\\sim$1\\% and for surface gravity $\\sim$10\\%. The atmospheric parameters derived from the carbon ionization equilibria are highly consistent with our previous spectroscopic analyses: results from a non-LTE study of H and He in the visual and in the near-IR, including the establishment of the \\ion{He}{i/ii} ionization equilibrium in the hotter stars (Paper~1), and the reproduction of the spectral energy distributions from the UV to the near-IR (Paper~2).  All C lines considered, from up to three ionization stages, indicate similar abundances. The linelist includes 40 transitions suitable for analysis over a wide wavelength range. In particular, the strongest features, of highest importance for extragalactic applications, are consistently modelled. The statistical 1$\\sigma$-uncertainties of the C abundance in each star are of the order $\\sim$0.05-0.10\\,dex. We estimate the systematic 1$\\sigma$-uncertainties to be $\\sim$0.10-0.15\\,dex. The average carbon abundance from \\ion{C}{ii/iii/iv} in the sample stars is highly uniform, $\\varepsilon$(C)\\,$=$\\,8.32$\\pm$0.04.  These results of unprecedented precision provide important constraints for stellar evolution and the chemical evolution of the Galaxy, despite the small sample size: {\\sc i}) They suggest that carbon depletion due to rotational mixing in the course of stellar evolution is small for stars without excessive initial angular momentum (by $<$0.05\\,dex) up to the giant stage, in agreement with theoretical predictions. {\\sc ii}) In consequence they suggest that the present-day carbon abundance in the solar neighbourhood is higher and much more homogeneous than indicated by previous work on early-type stars. This is consistent with the uniform abundance found in studies of the ISM and with predictions from models of the chemical evolution of the Galaxy. Moreover, our re-evaluated stellar value is in agreement with the gas-phase abundance derived for the Orion \\ion{H}{ii} region and the recently revised solar abundance.  This work shows that one should not underestimate the importance of a careful choice of input atomic data for non-LTE analyses. Not only accurate data for radiative transitions are required but also for collisional transitions. The carbon abundance analysis also turned out to react highly sensitively to the choice of atmospheric parameters, in particular to effective temperatures. If not accounted for properly, both factors may result in systematic errors in the interpretation of observed spectra, which cannot be reduced by statistics, i.e. an increase of the sample size of stars to be analysed. Detailed comparisons and calibrations of models with observed spectra are of highest importance, as these constitute the only empirical constraints for quantitative analyses. Only then can stellar and Galactochemical evolution models be verified in a meaningful way, and the studies be extended to stars in other galaxies."
    },
    "0711/0711.1862_arXiv.txt": {
        "abstract": "We have performed an extensive Voigt profile analysis of the neutral hydrogen (\\ion{H}{i}) and metal absorption present in a sample of 18 high resolution, high signal-to-noise QSO spectra observed with VLT/UVES. We use this analysis to separate the metal contribution from the \\ion{H}{i} absorption and present an improved measurement of the flux probability distribution function (PDF) due to  \\ion{H}{i} absorption alone at $<\\!z\\!>\\,=$ 2.07, 2.52, and 2.94.   The flux PDF is sensitive to the continuum  fit in the normalised flux range $0.8 < F < 1.0$ and to metal absorption at $0.2 < F < 0.8$.   Our new measurements of the flux PDF due to \\ion{H}{i} absorption alone are systematically lower at $0.2 < F < 0.8$ by up to 30\\% compared to the widely used measurement of \\cite{mc00}, based on a significantly smaller sample of Keck/HIRES data. This discrepancy is  probably due to a combination of our improved removal of the metal absorption and cosmic variance, since variations in the flux PDF between different lines-of-sight are large. The \\ion{H}{i} effective optical depth $\\tau_{\\ion{H}{i}}^{\\rm eff}$ at $1.7 < z < 4$ is best fit with a single power law, $\\tau_{\\ion{H}{i}}^{\\mathrm{eff}} = (0.0023 \\pm 0.0007) (1+z)^{3.65 \\pm 0.21}$, in good agreement with  previous measurements from comparable data.  As also found previously, the effect of noise on  the flux distribution is not significant in high resolution,  high signal-to-noise data. ",
        "introduction": "The Ly$\\alpha$ forest refers to the numerous, narrow absorption features bluewards of the Ly$\\alpha$ emission line in the spectra of high-redshift QSOs.  Our understanding of the origin of the Ly$\\alpha$ forest has made major progress with the advent of high resolution spectroscopy \\citep{cowie95,hu95,lu96,kirk97,rau97,cri00,kim01,jan06} and the availability of hydrodynamic cosmological simulations \\citep{cen94,miralda96,zhang97,weinberg98, cro98,theuns98,bryan99,dave99,cro02,jena05}.  Comparison of the predictions from hydrodynamic  cosmological simulations with observations have established a picture where the Ly$\\alpha$ forest is due to the neutral hydrogen (\\ion{H}{i}) component of the warm ($\\sim 10^4$K) photoionised intergalactic medium (IGM) which traces moderate-amplitude density  fluctuations of the  (dark) matter distribution in a simple manner. In this picture the spectra of high-$z$ QSOs provide a one-dimensional probe of the matter  density along the line-of-sight to high redshift QSOs.  The flux distribution in QSO absorption spectra   thus encodes  information on the underlying matter  distribution and evolution and the  Ly$\\alpha$ forest has been recognised as a powerful cosmological tool, complementing other cosmological  probes.  So far it is mainly the power spectrum of the flux distribution which has been used to quantitatively probe the properties of the large scale distribution of matter and its dependence on cosmological parameters. The flux power spectrum is sensitive to the matter power spectrum over a wide redshift range ($2<z<4$) and on scales of 1--$50\\hm$, significantly smaller than those probed by CMB fluctuations, galaxy surveys and gravitational lensing (Croft et al. 2002; Seljak, McDonald \\& Makarov 2003; Viel, Haehnelt \\& Springer 2004b; McDonald et al. 2006; Viel, Haehnelt \\& Lewis 2006; Lesgourges et al. 2007).  The probability distribution of the flux is also sensitive to the spatial distribution of dark matter and cosmological parameters, in particular, the amplitude of matter fluctuations \\citep{rau97,weinberg98,theuns00,mek01}.  However, the effect of the thermal state of the IGM (Theuns et al. 2000) and  other uncertainties make it more difficult to extract this information from the flux PDF. First attempts at a joint analysis of the flux PDF and the flux power spectrum have been made \\citep{des05, lid06, des07}. The analysis of Desjaques et al.  gives a somewhat smaller amplitude for the matter power spectrum compared to studies based on the flux power spectrum alone. Note, however, that these results were based on dark-matter-only numerical simulations in which the gas distribution was modelled in an approximate way rather than on full hydrodynamical simulations.  The measured flux PDF is sensitive not only to the \\ion{H}{i} absorption, but also, unfortunately, to continuum level uncertainties, contaminating heavy element absorption, and noise.  The effect of continuum fitting uncertainties on the \\ion{H}{i} effective optical depth has been discussed in detail \\citep{pre93, kim02, ber03, tyt04, kirk05}, but their impact on the flux PDF has not been considered as extensively (see Rauch et al. 1997 and McDonald et al. 2000 for notable exceptions).  Very little, if anything, has been done to quantify the effect of metal absorption and noise on the flux PDF.  To remedy these omissions, we perform a detailed Voigt profile analysis of the metal and \\ion{H}{i} absorption in 18 high resolution, high signal-to-noise (S/N) VLT/UVES QSO absorption spectra to assess the impact of continuum fitting uncertainties, metal absorption, and noise properties on the \\ion{H}{i} flux probability distribution at $1.7 < z < 3.2$ and the \\ion{H}{i} effective optical depth at $1.7<z<4$.  Our study is based on a factor of two larger sample and probes towards lower $z$ than similar work based on 8 Keck/HIRES spectra by \\citet{mc00}. It also differs significantly in the way we perform the continuum fitting and, more importantly, in the level of the characterisation of metal absorption.  Most previous estimates of the effective optical depth and the flux PDF due to \\ion{H}{i} absorption have either removed the metal lines {\\it statistically} \\citep{tyt04, kirk05}, excluded the metal-contaminated regions from the study \\citep{mc00, lid06} or included absorption due to metal lines \\citep{son04}.  We have carefully identified and {\\it fitted} metal and \\ion{H}{i} absorption lines with Voigt profiles for our sample of 18 spectra. The line fitting is essentially complete in the Ly$\\alpha$ forest region and redwards of the Ly$\\alpha$ emission lines, but less complete in the higher-order Lyman forest where the corresponding spectral regions were not always available.  During the fitting process, we have adjusted the initially estimated continuum level by a localised continuum fit in order to get an overall satisfactory fit to multiple-transition metal lines and higher-order Lyman series lines. This approach enables us to remove the metal contribution in the forest {\\it directly} and to recover one long continuous metal-free forest spectrum, and also to obtain an improved estimate of the continuum level.  The paper is organised as follows.  In Section \\ref{sec2}, we describe the data used for our analysis.  In Section \\ref{sec3}, we present the Voigt profile fitting analysis used to remove the metal contribution in the forest.  We investigate systematics of uncertainties affecting the flux PDF in Section \\ref{sec4}.  The observed flux PDF function and effective optical depth are presented in Section \\ref{sec5}.  Our conclusions are given in Section \\ref{sec6}. ",
        "conclusions": "\\label{sec6}   We have obtained improved measurements of the flux probability distribution at $1.7 < z < 3.2$ and effective optical depth due to the intergalactic \\ion{H}{i} absorption at $1.7 < z < 4$, based on a sample of 18 high resolution, high signal-to-noise VLT/UVES QSO absorption spectra for which we have performed an extensive Voigt profile analysis of the \\ion{H}{i} and metal absorption.  The main results are as follows:  \\begin{enumerate}  \\item{The normalised flux probability distribution (PDF) is affected mainly by metals at the level of 10--20\\% at flux levels of $0.2 < F < 0.8$ depending on redshift, and by continuum fitting uncertainties at the level of 5--20\\% at flux levels of $0.8 < F < 1$, depending on the signal-to-noise of the spectra. The effects of pixel noise and zero-level offset are very small, only noticeable at $F\\sim 0$ and $F\\sim 1$. The metal contribution to the absorption varies from a few percent to up to 30 percent between different lines of sight. A careful individual removal of the metal absorption is therefore essential for an accurate determination of the shape of the PDF.}  \\item{Our new measurements of the flux PDF due to \\ion{H}{i} alone are systematically lower at $0.2 < F < 0.7$ than the widely-used PDFs measured by \\cite{mc00} based on a sample of Keck/HIRES spectra half the size of our sample. The difference has the same sign as would be expected if the removal of the metal absorption by \\cite{mc00} was less complete than ours, but it is actually somewhat larger and does not have the correct redshift dependence if this were the sole cause. Given the rather small sample size cosmic variance is likely to be responsible for part of the discrepancy.}  \\item{The effect of our improved removal of metal absorption on the measured effective optical depth due to \\ion{H}{i} absorption, compared to previous measurements where the metal absorption has been taken into account in simpler ways, is small. Our new measurements of $\\tau_{\\ion{H}{i}}^{\\rm eff}$ are in good agreement with other measurements from high resolution, high signal-to-noise spectra.  In the redshift range   $1.7 < z < 4$  the redshift  evolution of the of our measurements of the effective optical depth and other measurements from high resolution spectra is  best fit with a single power law $\\tau_{\\ion{H}{i}}^{\\mathrm{eff}} = (0.0023 \\pm 0.0007) (1+z)^{3.65 \\pm 0.21}$.} \\end{enumerate}"
    },
    "0711/0711.1323_arXiv.txt": {
        "abstract": "We focus on the initial state spatial anisotropies, originating at the thermalization stage, for central collisions in relativistic heavy-ion collisions. We propose that a plot of the root mean square values of the flow coefficients $\\sqrt{\\overline {v_n^2}} ~~\\equiv v_n^{rms}$, calculated in a laboratory fixed coordinate system, for a large range of $n$ from 1 to about 30, can give non-trivial information about the initial stages of the system and its evolution. We also argue that for all wavelengths $\\lambda$ of the anisotropy (at the surface of the plasma region) much larger than the acoustic horizon size $H_s^{fr}$ at the freezeout stage, the resulting values of $v_n^{rms}$ should be suppressed by a factor of order $2H_s^{fr}/\\lambda$. For non-central collisions, these arguments naturally imply a certain amount of suppression of the elliptic flow. Further, by assuming that initial flow velocities are negligible at thermalization stage, we discuss the possibility that the resulting flow could show imprints of coherent oscillations in the plot of $v_n^{rms}$ for sub-horizon modes. For gold-gold collision at 200 GeV/A center of mass energy, these features are expected to occur for n $\\ge$ 5, with $n < 4$ modes showing suppression due to being superhorizon. This has strong similarities with the physics of the anisotropies of the cosmic microwave background radiation (CMBR) resulting from inflationary density fluctuations in the universe (despite important differences such as the absence of gravity effects for the heavy-ion case). It seems possible that the statistical fluctuations due to finite multiplicity may not be able to mask such features in the flow data, or, at least a non-trivial overall shape of the plot of $v_n^{rms}$ may be inferred. In that case, the successes of analysis of CMBR anisotropy power spectrum to get cosmological parameters can be applied for relativistic heavy-ion collisions to learn about various relevant parameters at the early stages of the evolving system. ",
        "introduction": "In the experimental search of the deconfined phase of QCD, namely the quark-gluon plasma (QGP), one of the most important results has been the observation of the elliptic flow \\cite{flow0,flow1}. This has given strong evidence of very early thermalization and of the collective behavior of the partonic matter produced in RHICE \\cite{flowex}. (We use RHICE to denote general class of relativistic heavy-ion collision experiments, to distinguish from the Relativistic Heavy-Ion Collider, RHIC, at Brookhaven). Much work has been done to extract physical information about the system from the behavior of the elliptic flow, e.g., equation of state, thermalization time, freezeout time etc.  Elliptic flow results from the spatial anisotropy of the thermalized region at the initial stage in a given event with nonzero impact parameter. Anisotropic pressure gradients then lead to anisotropic fluid velocity which results in anisotropic momentum distribution of particle momenta. Elliptic flow measures the second Fourier coefficient of the angular distribution of the particle momenta in the transverse plane. It has also been noticed \\cite{hijing,cntrl,v2cntrl} that even in central collisions, due to initial state fluctuations, one can get non-zero anisotropies in particle distribution (and hence in final particle momenta) in a given event, though these will be typically much smaller in comparison to the non-central collisions. These will average out to zero when large number of central events are considered. Fluctuations in the elliptic flow resulting from these initial state fluctuations, as well as the resulting modifications in the eccentricity of the initial region have been investigated in the literature \\cite{eccflct,v2hyd}. In this paper, we present a different approach to analyze the flow anisotropies for central events, resulting from these initial state fluctuations.  It is sometimes mentioned in popular terms that attempts to learn about early stages of phases of matter in RHICE from the observations of hadrons is in some sense similar to the attempts to understand the early stages of the universe from the observations of the cosmic microwave background radiation (CMBR). The surface of last scattering for CMBR is then similar to the freezeout surface in RHICE. We will argue below that this correspondence is in fact much deeper. There are strong similarities in the nature of density fluctuations in the two cases (with the obvious difference of the absence of gravity effects for RHICE). Following the successes of the analysis of the CMBR anisotropy power spectrum, we argue below that, for central events in RHICE, a plot of the root mean square values of the flow coefficients $\\sqrt{\\overline {v_n^2}} ~~\\equiv v_n^{rms}$, calculated in a laboratory fixed coordinate system, for a large range of $n$ from 1 to about 30, can give non-trivial information about the initial stages of the system and its evolution. In addition, we recall that one of the most important aspects of the density fluctuations in the universe is the coherence effect which eventually results in the remarkable acoustic peaks in the power spectrum of CMBR anisotropies. The source of this lies in the inflationary origin of the density fluctuations leading to production of super-horizon density fluctuations. Some of these eventually re-enter the horizon around the decoupling stage and leave these imprints on CMBR anisotropies \\cite{cmbr}. We will argue below that quite similarly, super-horizon fluctuations are present in RHICE as well. However, here they will result in non-zero flow coefficients as spatial anisotropies lead to anisotropic pressure gradients. The anisotropies in the momentum distributions of the particles, especially at large orders of the Fourier coefficients, will capture information about the nature  of initial spatial anisotropies, their evolution, and freezeout.  The paper is organized as follows. In section II we discuss the nature of the anisotropies present initially at the stage of equilibration. Section III discusses the basic physics of our model exploring the correspondence with CMBR physics. Section IV incorporates the presence of a finite acoustic horizon in RHICE and discusses the expected suppression of the values of $v_n^{rms}$. Section V discusses expected features in the plot of $v_n^{rms}$ based on the physics of our model. Section VI presents numerical results and conclusions are given in section VII. ",
        "conclusions": "Important aspects of our model are that we argue that important information about initial anisotropies of the system and their evolution in relativistic heavy-ion collisions can be obtained by plotting the root mean square values of the Fourier coefficients $v_n^{rms}$ of the anisotropies in the fluctuations $\\delta p/p$ of the particle momenta, calculated in a fixed laboratory frame, starting from $n = 1$ upto large values of $n \\simeq 30$. Note that $n = 30$ almost corresponds to wavelength of fluctuation $\\lambda$ at the surface of the region, at $\\tau_{fr}$, being of order 1 fm. (At $\\tau_{eq}$ it will correspond to $\\lambda \\simeq$ 0.7 fm). Fluctuations with wavelengths smaller than 1 fm cannot be treated within hydrodynamical framework, so we restrict attention within this range of $n$. (It may be useful to plot $v_n^{rms}$ for a much larger range of $n$. One will expect that beyond a critical value of $n$ the nature of the curve should change in some qualitative manner indicating breakdown of underlying hydrodynamical description for smaller modes. The wavelength corresponding to that critical value of $n$ will determine the smallest scale below which hydrodynamical description is not valid.) Important thing to note is that by taking very large number of (central) events, the root mean square values of $v_n$s can show any possible systematic variation, for example initial rapid decrease of $v_n^{rms}$ with $n$ which eventually flattens out for large $n$. One needs to develop a proper understanding of this behavior, (which presumably suggests a qualitative change in the nature of the fluctuations for small wavelengths)  and of various factors affecting this. Further, there is the possibility of a peak near $n \\sim 5$ (for $Au-Au$ collision at 200 GeV) and of subsequent peaks for larger $n$. If any of these non-trivial features are detected in the particle momentum spectra then it can open up a new way of accessing the information about initial stages of the matter produced. Just as for CMBR where the location of the first peak refers to the decoupling stage, here also the location of the first peak will give information about the freezeout stage, including the all important equation of state which could  distinguish a QGP phase from a hadronic phase. For CMBR, successive peaks yield important information about baryon content etc. which couple to the radiation and hence contribute to the acoustic oscillations. In the same way for RHICE, if successive peaks are present, then they may give information about the detailed properties of the matter present at that stage, and any dissipative factors etc. One important factor which can affect the shapes of these curves, especially the peaks, is the nature and presence of the quark-hadron transition. Clearly the duration  of any mixed phase directly affects the freezeout time and hence the location of the first peak. More importantly, any softening of the equation of state near the transition may affect locations of any successive peaks and their relative heights. Effects of mass ordering may also be important in this respect, especially for the behavior of these plots at large $n$ corresponding to small $\\lambda$ probing very small scale fluctuations of the order of hadronic scale.  If one does see even the first peak for RHICE then one very important issue relevant for CMBR can be studied with controlled experiments. It is the issue of horizon entering. For example, by changing the nuclear size and/or collision energy, one can arrange the situation when first peak occurs at different values of $n$. As we have emphasized above, even the overall shape of the curve (as given by the solid plot in Fig.3) will contain valuable information about the presence and evolution of the initial state anisotropies of the system resulting from ultra-relativistic heavy-ion collisions. It will be interesting if one can find a way to analyze the effects of such fluctuations along the longitudinal direction, i.e. with rapidity, by taking fixed azimuthal angle window. Due to rapid longitudinal expansion, presumably one will need to confine to a narrow window of rapidity.  As we discussed earlier, longitudinal fluctuations are not expected to be superhorizon. Hence one may expect qualitatively different behavior in the plots of $v_n^{rms}$ for this case compared to the transverse case (this is apart from the other important differences due to very different dynamics of longitudinal and transverse expansions). However, the situation may be different in the context of color glass condensate models \\cite{cgc}. In such models, initially a color field is established between receding nuclei, which subsequently decays in particles. If certain fluctuations are imprinted on such a color field in the longitudinal direction, then even if those are causal for the color field, for the resulting parton system, they may constitute fluctuations of super-horizon wavelengths (as the decay of the color field will involve different time scale than the time scale of establishing the initial color field). This issue needs to be investigated in detail as it may provide a distinctive signal for color glass condensate initial conditions. All the considerations discussed in this paper can be straightforwardly applied to non-central events, as well as for collision of nuclei which are deformed. The only difference is that for these cases certain Fourier coefficients (with small values of $n$) will be most dominant, as discussed above. As for the case of elliptic flow, here also it is important to estimate and measure higher moments of the flow coefficients to have a detailed understanding of the nature of the original fluctuations. Recall that these play crucial role for the CMBR analysis where non-Gaussian effects are very tightly constrained in inflationary models of density fluctuations. Similarly, in RHICE also, the physical origin of original fluctuations will constrain and determine various moments of the flow coefficients. Other issues for CMBR physics such as the effects of the thickness of the surface of last scattering on CMBR anisotropy power spectrum will have obvious implications for RHICE where also the width of the freezeout surface will affect the plots of values of $v_n^{rms}$."
    },
    "0711/0711.0864_arXiv.txt": {
        "abstract": "% We review the current state of the art in double degenerate merger simulations to better understand the role this phenomenon plays in type Ia progenitors.  Because the fate of a merged system may well depend on the exact evolution of the matter temperature (as well as mixing of the merged system), precision simulations are required to determine the true fate of these systems.  Unfortunately, if we compare the results of current simulations, we find many-order of magnitude differences in quantities like mass-transfer rates in the merger process.  We discuss these differences and outline an approach using verification and validation that should allow us to achieve a level of precision sufficient to determine the true fate (thermonuclear vs. collapse) of double degenerate mergers. Understanding the fate of lower-mass systems (e.g. R Coronae Borealis stars) may be key in our final testing phase. ",
        "introduction": "Supernovae (SNe) and their Gamma-Ray Bursts cousins mark the most powerful explosions in the universe.  The quest for an understanding of the driving mechanism behind these explosions has been nearly as turbulent and heated as the events themselves.  It is now nearly universally agreed that these explosions are produced in one of two engines: the collapse of a massive stellar core (the gravitational potential energy released in the collapse being the energy source of the supernova explosion), the thermonuclear explosion of a white dwarf (where the fusion of carbon and oxygen into heavier elements releases the energy to power the explosion).  Nearly as fractious has been the discussion of the progenitors of these mechanisms.  Understanding the progenitors of supernovae requires an understanding of stellar evolution, something we are far from doing.  For ``core-collapse'' supernovae, SN 1987A provided our first direct glimpse of a supernova progenitor because the progenitor was observed prior to the explosion (Kirshner et al. 1987).  It also profices an example of how wrong stellar evolution can be (past models insisted that only red supergiant stars would collapse to form supernovae).  But for thermonuclear explosions (type Ia Supernovae), we have not been so lucky.  All current type Ia SNe progenitors require mass transfer from a binary companion.  To understand these progenitors, we must not only understand stellar evolution, but also binary interactions, introducing an entirely new set of uncertainties into our understanding of the type Ia progenitor.  One of the persistent progenitor scenarios of type Ia supernovae is the merger of two white dwarfs (double degenerate scenario), producing a single white dwarf with mass above the Chandrasekhar limit that will contract and explode.  The primary drawback of this scenario is that current theory argues that such a merger will not produce a type Ia supernovae.  But its advantages, mostly that theory predicts that the rate of such mergers is consistent with the supernova rate (for many other proposed scenarios, this is not the case), has kept this proposed progenitor alive.  To understand why this scenario is not believed to work, we must understand what happens when a white dwarf accretes matter.  Nomoto and Kondo (1991) summarized this understanding in a single plot (Fig. 4 of that paper).  They determined the fate of a Carbon/Oxygen (CO) white dwarf as a function of its birth mass and the rate at which we accrete material on this white dwarf.  For the double degenerate scenario, the region of concern is at the topmost accretion rates.  When the accretion rate is above a few times 10$^{-6} {\\rm M_\\odot s^{-1}}$, the carbon in the white dwarf ignites at its edge and burns inward, transforming the CO white dwarf into an Oxygen/Magnesium/Neon white dwarf.  When such a white dwarf approaches the Chandrasekhar limit and contracts, neutrino emission cools the white dwarf sufficiently to prevent a thermonuclear runaway until the matter has contracted too much for the explosion to escape its own potential well.  This material will continue to collapse until nuclear forces and neutron degeneracy halt the collapse at the formation of a neutron star.  The fate of such an object is interesting (it is termed Accretion Induced Collapse ``AIC'' and has been used to explain a number of neutron star populations), but definitely not a type Ia supernova (Fryer et al. 1999).  As we shall discuss below, all current results showing the merger of 2 white dwarfs suggest that this process is rapid and that the ultimate accretion rate onto the white dwarf will be very close to the Eddington rate.  Hence, theory currently predicts that the double degenerate scenario produces AICs and not type Ia supernovae.  There are a number of caveats to this result.  The Nomoto \\& Kondo result assumed constant accretion rates and did not account for the fact that the white dwarf could have a very complex rotation profile (Saio \\& Nomoto 2004; Yoon et al. 2007).  This has led to small windows of opportunity for the double degenerate scenario to still produce type Ia supernovae.  But to determine what the true fate of these objects is, we must simulate the merger.  And we must be able to believe the results of our merger in detail.  This brings us to an important concept in the scientific method.  In astronomy, we have two types of theoretical investigations: predictive and, for symmetry sake, `` post-dictive''.  Predictive science is what we strive for - to be able from first principles to determine how something should behave.  Post-dictive is what happens with a lot of science.  We know the answer we must get (e.g. explaining the solar abundance pattern) and we fit in some free parameters in our model to make sure to get this answer.  Of course, we can not predict errors in the observations, and when the solar abundance pattern changes, we must then simply accept this change.  Post-dictive science has many virtues: first, if we can match data with a reasonable set of parameters, we show that our basic model may also be reasonable; second, the parameters required for such a model to fit the data might teach us something about the underlying physics requirements. Unfortunately, scientists often forget that they have made parameter assumptions and start to believe they have predicted the answer.  In such scenarios, post-dictive models can do more harm than good (scientists may argue they have solved a problem prematurely, preventing continued work on a subject).  Stellar evolution is rife with examples of such misuse of post-dictive models.  But truly predictive models are very difficult to do in astrophysics. Generally the problems we are interested in astrophysics are too complex to be solved with a simple analytic derivation.  Once we resort to numerical models, we must be wary about numerical artifacts in our simulations.  The national laboratories have focused their testing procedure on a process of Verification and Validation (V\\&V). Verification is the process by which scientists test to make sure their numerical models are solving correctly the physics in their code. Validation is the process by which scientists confirm that the physics in their code is the correct physics for the problem they are solving. Most of our time is spent on Verification, which can take many forms: comparison to analytic solutions, convergence studies, code comparison, and even comparison to laboratory experiments.  Validation is almost entirely focused on comparison to some observation or experiment.  In this paper, we review the current status of simulations of WD mergers.  Recent simulations suggest that understanding the details will require extremely accurate ($\\sim$10\\% in temperature) simulation results.  Such accuracies will require a focused V\\&V effort and I will outline a basic approach for this problem.  Validation requires an observational constraint similar to the problem we are solving and we will discuss the potential of hydrogen deficient stars (R Coronae Borealis stars) as a validation test for type Ia progenitors. ",
        "conclusions": "So what can we say at this point in time?  First, no calculation has yet accurately calculated the mass transfer rate in double degenerate mergers.  We know that it is probably not as fast as that shown in most SPH calculations of type Ia progenitors, but it is probably faster than those results using unshocked matter.  If we only had to distinguish between a mass transfer rate above or below the Nomoto \\& Kondo (1991) line, these simulations would be sufficient: the mass transfer rate is almost certainly above this line, possible even above the Eddington accretion rate (that means it will form a thick disk). But if we want to get accurate temperature evolution profiles, we will have to do more accurate calculations.  \\begin{figure} \\plotone{fryer05.ps} \\caption{Star masses and accretion rates as a function of time.} \\end{figure}  Figure 5 shows the accretion rates for our current simulations. Although we have much more testing to do before such quantities like accretion rate can be trusted, we can still study a number of trends that may have implications in the observations to help constrain the simulations.  For example, the accretion is periodic caused by slight errors in the initial orbit.  How does this alter the nuclear burning?  It will affect the temperature as well as the mixing.  We may be able to use observations of R Coronae Borealis stars to place limits on the level of eccentricity in nature.  Such high-precision calculations can not be done without systematic verification and validation.  Here we have shown some of the tests that can be used in V\\&V: comparing (and understanding) the differences in code results in concert with analytic comparisons.  This requires a step-by-step process, adding physics one piece at a time to understand its effect.  Convergence studies will also prove useful in this problem.  This all fits under the ``verification'' of a code.  It would be nice to also conduct some validation tests - comparison to a very similar problem where the data is more plentiful.  It could well be that stars like R Coronae Borealis stars provide such a validation experiment."
    },
    "0711/0711.2791_arXiv.txt": {
        "abstract": "In this letter we investigate the deep Newtonian regime of the MOND paradigm from a purely phenomenological point of view by exploiting the least-square estimated corrections to the secular rates of the perihelia of the inner and of some of the outer planets of the Solar System by E.V. Pitjeva with the EPM2004 ephemerides. By using $\\mu(x)\\approx 1-k_0(1/x)^n$ for the interpolating MONDian function, and by assuming that $k_0$, considered body-independent so to avoid violations of the equivalence principle, experiences no spatial variations throughout the Solar System we tightly constrain $n$ with the ratios of the perihelion precessions for different pairs of planets. We find that the range $1\\leq n\\leq 2$ is neatly excluded at much more than $3-\\sigma$ level. Such a test would greatly benefit from the use of extra-precessions of perihelia independently estimated by other groups as well. ",
        "introduction": "The MOND scheme (MOdified Newtonian Dynamics) was put forth by Milgrom in Ref.~\\cite{Mil83} in order to phenomenologically explain two basic facts concerning spiral galaxies without resorting to the concept of hidden dark matter: the asymptotic flatness of the rotation curves of spiral galaxies \\ct{Rub01} and the Tully-Fisher law which is a well-defined relationship between the rotation velocity in spiral galaxies and their luminosity \\ct{Tul77}.  Viewed as a modification of gravity\\footnote{It can also be considered as a modification of the inertia of a particle under the action of a generic force $\\bds F$.}, MOND predicts that the gravitational acceleration $\\bds A_g$ felt by a particle in the field of a distribution of mass is \\eqi\\bds A_g=\\rp{\\bds A_{\\rm N}}{\\mu\\left(\\rp{A_g}{A_0}\\right)},\\eqf where $\\bds A_{\\rm N}$ is the Newtonian acceleration, $A_0$ is an acceleration scale which different, independent ensembles of observations set to \\ct{San02} $A_0=1.2\\times 10^{-10}$ m s$^{-2}$, and $\\mu(x)$ is an interpolating function which approximates 1 for $x\\gg 1$, i.e. for accelerations larger than $A_0$; for $x\\ll 1$ $\\mu(x)=x$, so that in such a strongly MONDian regime $A_g\\approx \\sqrt{A_{\\rm N}A_0}$.  Here we wish to investigate the deep Newtonian regime ($x\\gg 1$) in view of recent advances in planetary orbit determination occurred for the inner planets of the Solar System. For a quite general class of interpolating functions, $\\mu(x)$ can be cast into the form \\ct{Mil83} \\eqi\\mu(x)\\approx 1-k_0\\left(\\rp{1}{x}\\right)^n,\\lb{milgrom}\\eqf which yields a modified gravitational acceleration \\ct{Tal88} \\eqi \\bds A_g\\approx \\bds A_{\\rm N}\\left[1+k_0\\left(\\rp{A_0}{A_{\\rm N}}\\right)^n\\right].\\lb{accmilgrom}\\eqf Note that the most commonly used expressions for $\\mu(x)$, i.e. \\ct{Mil83} \\eqi\\mu(x)=\\rp{x}{\\sqrt{1+x^2}},\\eqf and \\ct{Fam05} \\eqi\\mu(x)=\\rp{x}{1+x},\\eqf reduce to \\rfr{milgrom} for $k_0=1/2, n=2$ and $k_0=1, n=1$, respectively, in the appropriate limit. For a recent review of many aspects of MOND as various attempts to theoretically justify it see Ref.~\\cite{Bek06} and Ref.~\\cite{Bru07}. ",
        "conclusions": "A major outcome of our analysis is that $n=2$ is neatly ruled out independently of $k_0$ and $A_0$, contrary to what obtained in Ref.~\\cite{Ser06} in which only the perihelion of Mars was used by keeping fixed $k_0\\approx 1$ (and using the commonly accepted value \\ct{San02} $A_0=1.2\\times 10^{-10}$ m s$^{-2}$). The results of the analysis presented here seem to suggest that a reconsideration of the matching between the deep Newtonian and MONDian regimes should be looked for. However, caution is in order because the present analysis is based upon the extra-precessions of perihelia estimated by only one team; it would be of great importance to have at disposal  corrections to the perihelia rates determined independently by other groups as well."
    },
    "0711/0711.2044_arXiv.txt": {
        "abstract": "Primary GCR interact with the Earth's atmosphere originating atmospheric sho\\-wers, thus giving rise to fluxes of secondary particles in the atmosphere. Electroma\\-gnetic and hadronic interactions interplay in the production of these particles, whose detection is performed by means of complementary techniques in different energy ranges and at different depths in the atmosphere, down to the Earth's surface.  Monte Carlo codes are essential calculation tools which can describe the complexity of the physics of these phenomena, thus allowing the analysis of experimental data. However, these codes are affected by important uncertainties, concerning, in particular, hadronic physics at high energy. In this paper we shall report some results concerning inclusive particle fluxes and atmospheric shower properties as obtained using the FLUKA transport and interaction code. Some emphasis will also be given to the validation of the physics models of FLUKA involved in these calculations. \\\\ \\\\ {\\keywords GCR EAS, inclusive fluxes of secondary particles, Monte Carlo $\\mathrm {\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,}$ models and codes } \\\\ {\\pacs 98.70.Sa, 13.85-t, 13.75-t} ",
        "introduction": "GCR interactions in the Earth's atmosphere, induced by particles with energies high enough (E $\\gsim 10^{14} - 10^{15}$ eV), may originate Extended Air Showers (EAS). Both the EM and the $\\mu$ component of EAS can be detected and used to infer properties concerning GCR primary spectrum energy and composition. At present, a special emphasis is put in the investigation of Very High and Ultra High Energy Cosmic Rays (VHECR and UHECR), through dedicated experiments devoted to the detection of EAS originated by these particles, due to the still open questions concerning the shape of the spectrum ($2^{nd}$ knee, ankle), the composition ($\\gamma$, $p$, heavy ions), and the origin of cosmic rays (galactic / extragalactic), their transport, acceleration and the GZK cut-off.  Information on GCR primary flux and composition are inferred from the experimental observables by using Monte Carlo (MC) simulations. In particular, while the fluo\\-re\\-scen\\-ce technique, used by many experiments, allows with some uncertainties to estimate the EM energy deposited in the atmosphere, the total primary energies can be obtained only after estimating the missing energy carried by other shower components ($\\mu$, nucleons...), and this requires indeed MC simulations. Even the search for particular estimators, such as the S variable, allows~to obtain primary energy information in a way nearly independent of composition~\\cite{rebel}, but needs MC for calibrations.   In particular, in the AUGER experiment~\\cite{auger} information on primary energies are inferred by hybrid measurements and the fluorescence technique as far as the EM sector is concerned (with an uncertainty amounting to $\\sim$ 20\\%), whereas the number of $\\mu$s at the Earth's surface, $N_\\mu$, is still related to total primary energies by means of MC simulations. On the other hand, in the KASCADE and KASCADE-Grande experiments~\\cite{kascade} information on primary energies are inferred from measures of $\\mu$ and $e$ with the aid of MC shower simulations. Both in AUGER and in KASCADE / KASCADE-Grande, as well as in many other EAS experiments, MCs are heavily needed to analyse data in terms of GCR mass composition. At present there are however still important uncertainties, since different MC models produce important differences on the interpretation of the same data. The critical ingredients in these MC codes are the physical models for hadronic and nuclear interaction, and their implementation.  In this work we summarize some of the results which can be obtained by MC simulation of EAS, giving examples taken from the use of the FLUKA code. ",
        "conclusions": ""
    },
    "0711/0711.0107_arXiv.txt": {
        "abstract": "We present a coincidence search method for astronomical events using gravitational wave detectors in conjunction with other astronomical observations. We illustrate our method for the specific case of the LIGO gravitational wave detector and the IceCube neutrino detector.    LIGO trigger-events and IceCube events which  occur  within a  given time window are selected as time-coincident events.  Then the spatial overlap of the reconstructed event directions is evaluated using an unbinned maximum likelihood method. Our method was tested  with Monte Carlo simulations  based on realistic LIGO and IceCube event distributions. We estimated a typical false alarm rate for the analysis to be 1 event per 435 years.  This is significantly smaller than the false alarm rates of the individual detectors. ",
        "introduction": "In this paper, we present an analysis method to look for astrophysical sources which produce both gravitational wave and high energy neutrino bursts using data from LIGO and IceCube.  One example of a possible source is a gamma-ray burst (GRB).  Thanks to the {\\it Swift} satellite\\,\\cite{2004ApJ...611.1005G}, there is accumulating observational evidence suggesting the association of long GRBs with the death of massive stars and supernova-like events (e.g. SN2006aj and GRB060218\\,\\cite{2006Natur.442.1008C}, see also\\,\\cite{2005tmgm.meet..860L}).  The collapsar model\\,\\cite{1999A&AS..138..499W} is widely accepted for explaining long GRBs and stellar collapse.  During the gravitational collapse of rapidly rotating stars, gravitational waves are emitted (see \\cite{2006RPPh...69..971K} for a review).  First, fireballs heated by neutrinos from the accretion disk are thought to produce the prompt gamma-ray emissions \\,\\cite{2005RvMP...76.1143P}. Subsequently in the prompt and afterglow phases, high energy neutrinos ($\\sim 10^5-10^{10}$\\,GeV) are expected to be produced by accelerated protons in relativistic shocks (see \\cite{1997PhRvL..78.2292W, 1998PhRvL..80.3690V} for reviews).  High energy neutrinos could be emitted also from short-duration GRBs, which are thought to be the outcome of neutron star mergers\\,\\cite{2007NJPh....9...17L}.  There is currently limited knowledge both observationally and theoretically about the details of the astrophysical process connecting  the gravitational collapse/merger of compact objects and  black hole formation  with the formation of fireballs. Coincident observations of gravitational waves and neutrinos from those events could therefore make an important contribution to the understanding of such phenomena.  Apart from GRBs, there may be other (unknown) classes of sources which produce bright bursts both in gravitational waves and neutrinos. Since our proposed method is not specific to any source type, our search will be able to set an upper limit for the population of any sources that produce nearly simultaneous bursts of gravitational waves and high energy neutrinos within the detection range of LIGO and IceCube.  We may also discover a previously unknown astrophysical phenomenon,  if correlated events are found at a high confidence level.  There are several interferometric gravitational wave (GW)\\,\\cite{2001astro.ph.10349H} detectors around the world, such as LIGO\\,\\cite{2006CQGra..23S..51S}, TAMA\\,\\cite{2004CQG..21..403T}, GEO\\,\\cite{2006CQGra..23S..71L} and VIRGO\\,\\cite{2006CQGra..23S..63A}, currently in operation.  These detectors monitor the relative displacement of mirrors (test masses) in response to distortions induced by gravitational waves. There are also several high-energy neutrino detectors operating, including AMANDA\\,\\cite{ahrens_search_2004}, IceCube\\,\\cite{2004APh....20..507A}, and ANTARES\\,\\cite{1999astro.ph..7432A}, which look for the Cherenkov light of charged particles emitted by neutrino interactions in water or ice. We illustrate the coincidence search method for the specific case of LIGO and IceCube.  LIGO is a network of interferometric gravitational wave detectors consisting of three interferometers\\footnote{From now on, we treat the network of the three LIGO interferometers as one detector and use the word ``detector'' to refer to them as a whole. To refer to individual LIGO interferometers, we always use the word ``interferometer'' to avoid confusion.} in the USA\\,\\cite{1992Sci...256..325A}. Two interferometers (4\\,km and 2\\,km long ones) are co-located in Hanford, WA and another 4\\,km interferometer is located in Livingston, LA. They have now achieved the design sensitivity\\,\\cite{2007AIPC..928...11F}.  Since the interaction of gravitational waves with matter is extremely weak, expected signals even from very strong gravitational wave sources are very small.  In order to declare a detection, we have to find a small signal in an overwhelming noise background  with high confidence. Generally, the output from the detector contains glitches which are not associated with gravitational waves but rather caused by various local disturbances such as laser noises, seismic excitations, etc. In order to search for GW bursts, which are gravitational waves of short duration, it is therefore important to distinguish gravitational wave signals from noise glitches without prior knowledge of signal waveforms.  One way to pick  out gravitational wave signals of unknown waveform from the noise background is to find coincidences between independent detectors. We can reject a large fraction of background events by comparing the arrival time and other properties (frequency, duration, etc) of the signals detected by independent gravitational wave detectors\\,\\cite{2006CQGra..23S..29A, 2004PhRvD..69j2001A, 2005PhRvD..72l2004A}.  Additionally, event lists from other astronomical observations, such as Gamma-Ray Bursts (GRB), optical supernovae, neutrinos, etc., can be used to find events that may be associated with GW bursts with an increased likelihood\\,\\cite{2006AIPC..836..605M, 2007PRD..76S..062003T, Marka:2005nz, Mohanty:2004tk, 2004APh....21..201A}.  Moreover, strict coincidence requirements allow us to investigate candidate events at lower Signal-to-Noise Ratios (SNRs) while maintaining a low False Alarm Rate (FAR).  Here, we propose a method for coincidence analysis of gravitational wave data with other detectors and illustrate it for the case of the LIGO gravitational wave detector and the IceCube neutrino detector.  IceCube is a cubic kilometer-scale neutrino detector under construction at the geographic South Pole. Its primary mission is the search for high energy extraterrestrial neutrinos. When completed, IceCube will consist of an array of 4800 digital optical modules, attached to 80 strings submerged within the Antarctic ice.  Currently the detector is taking data with more than 90\\,\\% livetime, except during a few months each year for construction and commissioning of new strings.  IceCube is optimized to look \"down,\" using the Earth as a screen to block all particles except neutrinos; thus its field of view is the northern hemisphere.  Neutrino arrival directions are resolved with a median error between 1$^\\circ$ and 2$^\\circ$\\,\\cite{2007ICRC...IPC}. The threshold neutrino energy for the IceCube detector is ~ 100\\,GeV.  The full energy range of observed events depends primarily on the competition between the unknown, but presumably falling, source flux, versus the rising neutrino cross-section. A flux with an $E^{-2}$ differential energy spectrum, for example, results in an energy distribution of neutrino events that peaks in the range of $10^4$ to $10^5$\\,GeV.   In our search method, the data streams from the LIGO interferometers are processed by a trigger generation pipeline, which generates a list of gravitational wave triggers for each interferometer. Then we compare the trigger lists from LIGO interferometers to generate a coincident LIGO event list, which contains the arrival time and the source direction of each  event. The LIGO event list is compared with an event list from the IceCube detector which also contains the timing and source direction information of the events. From the event lists we choose pairs of LIGO-IceCube events which lie within a certain time interval as time-coincident events. Then the spatial overlap between the LIGO and IceCube events is statistically evaluated to obtain the significance of the coincident event.  Because of the very different nature and geographical location of the two detectors, it is extremely unlikely that  the coincident triggers are due to the same source of noise. Therefore, the remaining possibility for time coincident trigger generation in both detectors, other than real astronomical events, is accidental coincidence.  Furthermore, the chance for two time-coincident noise triggers to generate overlapping reconstructed directions on the sky is also  small.  By the combination of  timing and directional coincidence discrimination, we can expect that most background events will be rejected and the FAR will be significantly reduced. ",
        "conclusions": "For each LIGO event, the expected number of IceCube events found within a time window ($\\pm T_\\mathrm{W}$) is $2T_\\mathrm{W}\\cdot R_\\mathrm{\\nu}$, where $R_\\mathrm{\\nu}$ is the event rate of IceCube. Therefore, using the event rate $R_\\mathrm{GW}$ of LIGO, the overall rate for LIGO-IceCube  time-coincident events can be calculated by $2T_\\mathrm{W}\\cdot R_\\mathrm{GW}\\cdot R_\\mathrm{\\nu}$. Using a p-value threshold $p_0$ for spatial coincidence, the FAR of this analysis method can be expressed by the following formula, \\begin{equation} \\mathrm{FAR}= 2T_\\mathrm{W}\\cdot R_\\mathrm{GW} \\cdot R_\\mathrm{\\nu}\\cdot p_0. \\end{equation} More specifically in the case of the Monte Carlo simulation explained in the previous section, the FAR is given by the following formula, \\begin{equation} \\mathrm{FAR}=\\frac{1}{435}\\left(\\frac{p_0}{1\\,\\%}\\right) \\left(\\frac{T_\\mathrm{W}}{1\\,\\mathrm{sec}}\\right)\\,\\mathrm{[events / year]}. \\end{equation}  The obtained FAR is 1 false alarm in 435 years for one-second coincidence time window and spatial coincidence p-value threshold of 1\\,\\%. If we allow a higher FAR, for example 1 event per 100 years used by SNEWS (SuperNova Early Warning System)\\,\\cite{Antonioli04}, we can relax the LIGO or IceCube event selection thresholds to search for weaker signals in the background noise.  In the case of long GRBs, high energy neutrinos from relativistic shocks are expected to be emitted  between a few hours (for the internal shocks\\,\\cite{1997PhRvL..78.2292W, GUPTA07, 2006PhRvL..97e1101M}) to a few days (for the external shocks\\,\\cite{2000..Wax..ApJ.Neut}) after gravitational wave emission caused by core bounce.  In order to look for this type of event, we have to use a large time window of order of days.  In this case, the FAR may be unacceptably large because most LIGO events will be able to find at least one companion IceCube event (and vice versa) within a day. However, if the discrimination power of spatial coincidence can be improved, this would offset the larger time overlap. Such improvement would result from, for example, the continued enlargement of the IceCube detector, or the addition of another gravitational wave detector operating in conjunction with LIGO. On the other hand, the time coincidence is effective to search for GW and neutrino bursts with small time delay.  Our method can also be applied to coincidence analyses with other neutrino detectors such as Super-Kamiokande\\,\\cite{Fukuda03}, Lake Baikal\\,\\cite{2006astro.ph..9743W}, Baksan\\,\\cite{1998PPN....29..254A} etc, without  significant modification. Combinations with the low energy-threshold detectors would enable us to search for supernova events. Moreover, our method can be used with any astronomical detectors which provide timing and source location information of burst events. Coincidence search with a large number of detectors will increase the confidence of detections.  We shall extend our method to include the VIRGO gravitational wave detector. The use of three geographically separated interferometers will enable us to constrain possible source locations of a gravitational wave event to two points on the sky\\,\\cite{1989AsAph..217..381B}. Additionally, time coincidence discrimination between VIRGO and LIGO interferometers will further reduce the background event rate of the gravitational wave detectors network.  Both of these changes in the time and spatial coincidence rates will work together to provide a much lower FAR and/or better sensitivity.  \\ack  The authors are grateful for the support of the United States National Science Foundation under cooperative agreement PHY-04-57528 and Columbia University in the City of New York. We are grateful to the LIGO collaboration for their support. We are indebted to Jamie Rollins for his useful comments on the manuscript. The authors gratefully acknowledge the support of the United States National Science Foundation for the construction and operation of the LIGO Laboratory. This work was also supported in part by the Office of Polar Programs of the National Science Foundation and a Grant-in-Aid for Scientific Research from the Ministry of Education, Science and Culture of Japan through No.S 19104006. This paper has been assigned LIGO Document Number LIGO-P070115-01-Z."
    },
    "0711/0711.4038_arXiv.txt": {
        "abstract": "We present {\\it Spitzer} Infrared Spectrograph spectra of 28 Class I protostars in the Taurus star-forming region. The 5 to 36 $\\mu$m spectra reveal excess emission from the inner regions of the envelope and accretion disk surrounding these predecessors of low-mass stars, as well as absorption features due to silicates and ices. Together with shorter- and longer-wavelength data from the literature, we construct spectral energy distributions and fit envelope models to 22 protostars of our sample, most of which are well-constrained due to the availability of the IRS spectra. We infer that the envelopes of the Class I objects in our sample cover a wide range in parameter space, particularly in density and centrifugal radius, implying different initial conditions for the collapse of protostellar cores. ",
        "introduction": "In a now widely accepted evolutionary sequence based on the shape of the spectral energy distribution (SED) in the infrared, there are four stages in the pre-main-sequence life of a low-mass star: Class 0 objects are deeply embedded protostars with large submillimeter to bolometric luminosity ratios, which are very faint at near-infrared wavelengths; Class I objects are surrounded by infalling envelopes, but are less embedded; Class II objects have cleared their envelopes and are surrounded by accretion disks; and Class III objects have dispersed almost all of their circumstellar material \\citep{lada84, lada87,andre93}.  Therefore, Class I protostars are thought to be in one of the earliest evolutionary stages. After the gravitational collapse of a molecular cloud core starts, the infalling material forms an envelope around the central object. Since cloud cores are usually rotating slowly, the material falls preferentially onto a disk, which serves as a mass reservoir, most of which is eventually accreted by the protostar. Class I objects are thus surrounded by accretion disks, in addition to envelopes extending to several 1000 AU. The radius of the disk is roughly equal to the centrifugal radius ($R_c$), which defines the centrifugal barrier encountered by matter falling in on the equatorial plane, if the specific angular momentum is conserved; inside this radius the density distribution flattens, and the matter follows non-radial trajectories, falling onto a disk. Material tends to accumulate close to $R_c$, which is close to the protostar at the beginning of the collapse, but increases rapidly over time \\citep{terebey84,adams86}.  Besides accreting matter, young stars generate powerful outflows which are launched along magnetic field lines; mass accretion onto the star and mass loss in the form of outflows seem to be correlated \\citep[e.g.,][]{hartigan95}. In addition to collimated outflows, protostars might also emit a wide-angle wind, which is expected to open up cavities in the envelope \\citep{arce02}. The presence of both collimated outflows and wide-angle stellar winds explains the fact that some objects with large cavities have well-collimated outflows. A partially evacuated cavity also forms as a natural by-product of the collapse of a sheet \\citep{hartmann94, hartmann96}; since molecular cloud cores are rarely spherically symmetric, this scenario might be more realistic. In some cases, both stellar winds and the flattening of the envelope are responsible for cavities oriented along the rotational axis of a protostar.  The spectral index $n$ ($\\lambda F_{\\lambda} \\propto {\\lambda}^n$) from 2 to 25 $\\mu$m is used for the SED classification of protostars and pre-main-sequence stars \\citep{lada87,adams87,andre94}; Class I objects have $n > 0$, Class II objects have $ -2 < n < 0 $, and Class III objects have $ -3 < n < -2 $ ($n=-3$ is the spectral index of a photosphere in the Rayleigh-Jeans limit). Therefore, Class I objects are characterized by a strong infrared excess; the dust in the envelope absorbs radiation from the central protostar and disk and reemits it at infrared to mm wavelengths. The peak of the SED of this type of objects usually lies in the mid- to far-infrared, depending on the density of the envelope (and thus mass and infall rate), as well as the source luminosity \\citep{kenyon93a}. Flat spectrum sources are likely objects surrounded by envelopes and disks and therefore protostars, but with wide cavities in their envelopes and/or seen close to pole-on \\citep[e.g., DG Tau, T Tau;][]{furlan06a, calvet94}. Some of the flat-spectrum sources could be transition objects (Class I/II), consisting of a T Tauri star and its disk surrounded by remnant envelope material (e.g., IRAS 04154+2823). Their SEDs typically show a double peak, one centered in the near-IR and one in the far-IR.  In addition to a rising SED in the infrared due to this excess emission, Class I protostars display various absorption features in their near- and mid-infrared spectra due to silicates and ices in the envelope. However, many parts of the infrared spectrum are not accessible from ground-based telescopes, which in addition often lack the sensitivity to observe faint objects. The {\\it Spitzer Space Telescope} \\citep{werner04} opened a new window onto low-mass star formation; with unprecedented sensitivity, the precursors of low-mass stars can be studied in various evolutionary stages.  Here we present 28 spectra from 5 to 36 $\\mu$m of Class I objects in the Taurus star-forming region observed with the Infrared Spectrograph\\footnote{The IRS was a collaborative venture between Cornell University and Ball Aerospace Corporation, funded by NASA through the Jet Propulsion Laboratory and the Ames Research Center.} \\citep[IRS;][]{houck04} on board of {\\it Spitzer}. The 5--20 $\\mu$m spectra of 5 of the 28 Class I objects were already introduced in \\citet{watson04}; these objects are IRAS 04016+2610, 04108+2803B, 04181+2654 B, 04239+2436, and DG Tau B. Our IRS observations are part of a larger Infrared Spectrograph guaranteed-time observing program; we presented the Class II and III objects from our sample of 150 young stellar objects in Taurus in an earlier paper \\citep{furlan06a}.  Since the Taurus-Auriga star-forming region is young (1--2 Myr) and relatively nearby \\citep[140 pc;][]{kenyon94b, bertout99}, many of its objects have been fairly well-studied at wavelengths covering almost the entire electromagnetic spectrum. Here we use SEDs constructed with data from the literature and our IRS spectra to constrain envelope models and to determine which parameters describe the circumstellar environment of the Class I objects in our sample.  This paper is structured as follows: in \\S\\ \\ref{obs_data_reduction} we lay out our {\\it Spitzer} IRS observations and data reduction; in \\S\\ \\ref{sample} we present the IRS spectra and SEDs of the 28 Class I objects in our sample; in \\S\\ \\ref{env_models} we describe the envelope models used to fit 22 of our Class I objects and give detailed descriptions of the individual model fits; in \\S\\ \\ref{discussion} we discuss the different appearances of Class I objects and give a summary of our models; and in \\S\\ \\ref{conclusions} we present our conclusions. ",
        "conclusions": "\\label{conclusions}  After analyzing the IRS spectra and SEDs of 28 Class I objects in the Taurus star-forming region, and generating envelope models for 22 of these objects, we conclude the following:  ${\\bullet}$ Almost all protostars we observed display ice absorption features in their mid-IR spectra, which originate in the envelopes around these young stars. In particular, the CO$_2$ ice feature at 15.2 $\\mu$m is ubiquitous and strong among the Class I objects in our sample, signifying an origin in the cold outer envelope regions. The envelope is also generally responsible for a deep silicate absorption feature at 10 $\\mu$m. The three objects with silicate emission features in their spectra, IRAS 04248+2612, 04264+2433, and CoKu Tau/1 are likely Class I objects seen through low-density envelopes, which are in an advanced stage of dispersal and thus considered to be more evolved.  ${\\bullet}$ The SEDs of Class I objects peak in the mid- to far-IR, where the emission is dominated by the envelope. Objects with flat SEDs in the infrared spectral range could be in transition between the Class I and II stage. The inclination angle of protostellar systems strongly affects the appearance of their mid-IR SED; Class I objects seen pole-on could appear as Class II objects with additional long-wavelength excess emission, while edge-on Class I objects have SED shapes similar to Class 0 objects, but with less flux in the sub-mm and mm wavelength region.  ${\\bullet}$ The IRS spectra are instrumental in constraining envelope model parameters, in particular the reference density $\\rho_1$ (which is linked to the mass infall rate) and centrifugal radius $R_c$. The depth of the silicate feature and the shape of the spectrum between 12 and 30 $\\mu$m provide constraints which can be further tightened by model fits to near-IR, far-IR, sub-mm, and mm observations. The median values for $\\rho_1$ and $R_c$ for the Class I objects we modeled are $3 \\times 10^{-14}$ g cm$^{-3}$ and 60 AU, respectively, with a relatively large range for both parameters. The disk contribution dominates the emission in the 3-8 $\\mu$m range for most of the modeled Class I objects, confirming the importance of disk accretion luminosity at this early evolutionary stage. Most SEDs can be reproduced by TSC models, thus not requiring an initially flattened density distribution of the collapsing cloud core; only 5 of the 22 objects we modeled needed sheet-collapse models for an adequate fit of their SEDs.  The IRS spectra of Class I objects in Taurus display a variety in appearances and add valuable constraints to envelope models. Fits to the observed SEDs yield the parameters describing their physical structure, and thus aid in confirming and refining our understanding of protostellar cloud collapse and evolution."
    },
    "0711/0711.4454_arXiv.txt": {
        "abstract": "{ {\\it Aims.} Stellar parameters and abundances of Na, Mg, Al, K, Ca, Sr, Ba, and Eu are determined for four very metal-poor (VMP) stars ($-2.66 \\le$ [Fe/H] $\\le -2.15$). For two of them, HD\\,84937 and HD\\,122563, the fraction of the odd isotopes of Ba derived for the first time.  \\noindent {\\it Methods.} Determination of an effective temperature, surface gravity, and element abundances was based on non-local thermodynamic equilibrium (non-LTE) line formation and analysis of high-resolution (R $\\sim60000$ and 90000) high signal-to-noise (S/N $\\ge 200$) observed spectra. A model atom for \\ion{H}{i} is presented. An effective temperature was obtained from the Balmer H$_\\alpha$ and H$_\\beta$ line wing fits. The surface gravity was calculated from the {\\sc Hipparcos} parallax if available and the non-LTE ionization balance between \\ion{Ca}{i} and \\ion{Ca}{ii}. Based on the hyperfine structure affecting the \\ion{Ba}{ii} resonance line $\\lambda\\,4554$, the fractional abundance of the odd isotopes of Ba was derived from a requirement that Ba abundances from the resonance line and subordinate lines of \\ion{Ba}{ii} must be equal.   \\noindent {\\it Results.} For each star, non-LTE leads to a consistency of $\\Teff$ from two Balmer lines and to a higher temperature compared to the LTE case, by up to 60~K. Non-LTE effects are important in spectroscopic determination of surface gravity from the ionization balance between \\ion{Ca}{i} and \\ion{Ca}{ii}. For each star with a known trigonometric surface gravity, non-LTE abundances from the lines of two ionization stages, \\ion{Ca}{i} and \\ion{Ca}{ii}, agree within the error bars, while a difference in the LTE abundances consists of 0.23~dex to 0.40~dex for different stars. Departures from LTE are found to be significant for all investigated atoms, and they strongly depend on stellar parameters. For HD\\,84937, the Eu/Ba ratio is consistent with the relative solar system $r-$process abundances, and the fraction of the odd isotopes of Ba, $f_{odd}$, equals 0.43$\\pm$0.14. The latter can serve as an observational constraint on $r-$process models. The lower Eu/Ba ratio and $f_{odd}$ = 0.22$\\pm$0.15 found for HD\\,122563 suggest that the $s-$process or the unknown process has contributed significantly to the Ba abundance in this star. ",
        "introduction": "Very metal-poor stars (hereafter VMP stars, with [Fe/H]\\footnote{[A/B] = $\\log(N_A/N_B)_* - \\log(N_A/N_B)_\\odot$ where $N_X$ are number densities.} $< -2$) belonging to the halo population of the Galaxy were formed at the earliest epoch of star formation. These stars preserve in the atmosphere the chemical composition produced by the early generation of stars. Determination of element abundances in VMP stars is important for understanding nucleosynthesis processes in the early Galaxy. The more metal-poor stars are, on average, more distant, and they are more difficult to analyze than nearby stars with metal abundances close to solar. In particular, the methods of surface gravity determination are of limited usefulness. For distant field halo stars, trigonometric parallaxes are unknown. Progress in this direction is expected due to the upcoming ESA Gaia satellite mission (Perryman et al.~2001). The strongest lines of neutral magnesium (the \\ion{Mg}{i}b lines) and neutral calcium (the \\ion{Ca}{i}\\, triplet $4p - 5s$) show no wings, and, therefore, they are not sensitive to surface gravity. In fact, surface gravities of VMP stars can only be found from the ionization balance between neutral and ionized species of a selected atom. One mostly selects iron and Fe peak elements. In two recently found hyper metal-poor stars (hereafter HMP stars, with [Fe/H] $< -5$), only Ca is observed in two ionization stages (Christlieb et al. \\cite{chri02}, Frebel et al. \\cite{frebel05}) and, therefore, can be a potential tool for the derivation of $\\log g$.  The atmospheres of VMP stars are characterized by low electron number densities and low opacities, which are particularly low in the ultraviolet spectral range. Hence, local thermodynamical equilibrium (LTE) is not fulfilled in the atmospheric layers where the lines are formed. As was proven in the past decade, the surface gravity based on the ionization balance \\ion{Fe}{i}/\\ion{Fe}{ii}\\, (Thevenin \\& Idiart \\cite{fe99}, Korn et al. \\cite{Korn03}) and element abundances (for review, see Asplund \\cite{asplund05}) are significantly affected by departures from LTE. In this study, we select four metal-poor stars observed with the High Dispersion Spectrograph of the Subaru Telescope and perform detailed analysis of stellar parameters and element abundances based on the non-local thermodynamic equilibrium (non-LTE) line formation for 10 important chemical species, including \\ion{H}{i}, \\ion{Ca}{i}, \\ion{Ca}{ii}, etc. One of the stars, HD\\,122563, is a giant, and the remaining three, HD\\,84937, BD\\,$+3^\\circ$740, and BD\\,$-13^\\circ$3442, are at the hot end of the stars that evolve on time scales comparable to the Galaxy lifetime. Similar intrinsically bright objects are mainly found among VMP stars due to the selection effect. For example, one of two HMP stars, the Christlieb's star is a giant with $\\log g$ = 2.2 (Christlieb et al. \\cite{chri04}) and another one has an effective temperature $\\Teff$ = 6180~K (Aoki et al. \\cite{frebel06}). In the selected stars, it is possible to study the formation of some of the strongest spectral lines, such as \\ion{Ca}{ii}\\, $\\lambda8498$, \\ion{Mg}{i}\\, $\\lambda5172$ and $\\lambda5183$, which are among the few spectral lines observed in HMP stars. At [Fe/H] $< -2$, they are presumably of purely photospheric origin. Thus, the stars of our small sample can serve as reference stars in studies of VMP and HMP stars.  The paper is organized as follows. Observations of the selected stars are described in Sect.\\,\\ref{obs}. Non-LTE calculations are found in Sect. \\,\\ref{NLTE}. The solar element abundances necessary for further differential abundance analysis are derived in Sect.\\,\\ref{line_list} from the selected spectral lines. In Sect.\\,\\ref{st_param}, we determine effective temperatures from the wings of Balmer lines and check surface gravities from the ionization balance \\ion{Ca}{i}/\\ion{Ca}{ii}. Non-LTE abundances of 8 chemical elements are derived in Sect.\\,\\ref{Abundance}. For two stars, HD\\,84937 and HD\\,122563, both the \\ion{Ba}{ii}\\, resonance and subordinate lines are available, and we determine in Sect.\\,\\ref{isotope} the fraction of the odd isotopes of barium that gives an independent estimate of relative contribution of the $r-$process and main $s-$process to heavy element abundances in the star. Our recommendations and conclusions are given in Sect.\\,\\ref{conclusion}. ",
        "conclusions": "In this paper, we determine for the first time stellar parameters $\\Teff$ and $\\log g$ of four VMP stars ($-2.15 \\le$ [Fe/H] $\\le -2.66$) based on non-LTE line formation. Effective temperature is derived from the Balmer H$_\\alpha$ and H$_\\beta$ line wing fits. For each star, non-LTE leads to a consistency of $\\Teff$ from two lines and to a higher temperature compared to the LTE case, by up to 60~K.  We find a clear advantage of the non-LTE approach in spectroscopic determination of surface gravity compared to the LTE case. For each of three stars with {\\sc Hipparcos} parallax available, a surface gravity obtained from the non-LTE ionization balance between \\ion{Ca}{i} and \\ion{Ca}{ii} agrees within the error bars with a trigonometric one. The \\ion{Ca}{i}/\\ion{Ca}{ii} line intensity ratio involving \\ion{Ca}{ii} $\\lambda8498$ is particularly sensitive to a variation of surface gravity. However, the wings of the \\ion{Ca}{ii} resonance lines can also be used for spectroscopic determination of $\\log g$ of VMP stars. We conclude that non-LTE calculations based on ${\\rm 1D}$ model atmospheres provide reliable results for $\\Teff$ and $\\log g$ of VMP stars in the wide range of temperatures ($\\Teff$ = 4600~K to 6400~K) and gravities ($\\log g$ = 4.0 down to 1.5).  Due to strong influence of non-LTE effects in the atmospheres of VMP stars abundances of Na, Mg, Al, K, Ca, Sr, Ba, and Eu cannot be determined in a reliable way assuming LTE. Departures from LTE depend strongly on stellar parameters such that a precise abundance analysis is only possible on the base of non-LTE calculations for each individual star.  We derive the fraction of the odd isotopes of Ba in two stars, HD\\,84937 and HD\\,122563, and, thus, increase the number of halo stars with a measured $f_{odd}$ to four. Two stars among four have low $f_{odd}$: HD\\,122563 with $f_{odd}$ = 0.22$\\pm$0.15 (this study) and HD\\,140283 with $f_{odd}$ = 0.31 (Lambert \\& Allende Prieto \\cite{iso140283}). Both stars reveal the lower Eu/Ba ratios compared to the relative solar system $r-$process abundances. These results suggest that the $s-$process has significantly contributed to the Ba abundance in these stars, or that the unknown process that produced high Sr/Ba ratios in these stars yielded the low $f_{odd}$ value. Two remaining stars have the larger fraction of the odd isotopes of Ba: $f_{odd}$ = 0.43$\\pm$0.14 in HD\\,84937 (this study) and $f_{odd}$ = 0.42 in HD\\,103095 (Mashonkina \\& Zhao \\cite{ba_iso}). In both stars, the Eu/Ba ratios are consistent with the relative solar system $r-$process abundances. We conclude therefore, that $f_{odd}$ found in these stars can serve as observational constraint on $r-$process models."
    },
    "0711/0711.1792_arXiv.txt": {
        "abstract": "Due to the lack of low redshift long Gamma-Ray Bursts (GRBs), the circular problem has been a severe obstacle for using GRBs as cosmological candles. In this paper, we present a new method to deal with such a problem in Markov Chain Monte Carlo (MCMC) global fitting analysis. Our methodology is similar to that of self-calibrations in using clusters of galaxies as cosmological probes. Assuming that a certain type of correlations between different observables exists in a subsample of GRBs, for the parameters involved in the correlation relation, we treat them as free parameters and determine them simultaneously with cosmological parameters through MCMC analysis on GRB data together with other observational data, such as SNe Ia, cosmic microwave background radiation (CMB) and large-scale structure (LSS). Then the circular problem is naturally eliminated in this procedure. To demonstrate the feasibility of our method, we take the Ghirlanda relation ($E_{\\gamma}\\propto CE_{\\rm peak}^A$) as an example while keeping in mind the debate about its physical validity. Together with SNe Ia, WMAP and SDSS data, we include 27 GRBs with the reported Ghirlanda relation in our study, and perform MCMC global fitting. We consider the $\\Lambda$CDM model and dynamical dark energy models with equation of state (EoS) $w_{\\rm DE}=w_0+w_1(1-a)$ and the oscillating EoS $w_{\\rm DE} = w_{0} + w_{1}\\sin(w_2\\ln(a))$, respectively. We also include the curvature of the universe in our analysis. In each case, in addition to the constraints on the relevant cosmological parameters, we obtain the best fit values as well as the distributions of the correlation parameters $A$ and $C$. We find that the observational data sets other than GRBs can affect $A$ and $C$ considerably through their degeneracies with the cosmological parameters. With CMB+LSS+SNe+GRB data included in the analysis, the results on $A$ and $C$ for different cosmological models are in well agreement within $1\\sigma$ range. The best fit value of $A$ in all models being analyzed is $A\\sim 1.53$ with $\\sigma \\sim 0.08$. For $C$, we have the best value in the range of $0.94-0.98$ with $\\sigma\\sim 0.1$. It is also noted that the distributions of $A$ and $C$ are generally broader than the priors used in many studies in literature. Our method can be readily applied to other GRB relations, which might be better physically motivated. ",
        "introduction": "\\label{introduction}  Searching for the nature of dark energy has been one of the most challenging tasks in cosmological studies. Because of the existence of degeneracies between dark energy parameters and the other cosmological parameters in different observables, multi-probe analysis are essential in constraining tightly the properties of dark energy. In this regard, exploring new probes has its great importance. On the other hand, thorough investigations on different probes both observationally and theoretically are equally important so that we can understand their validity and limitations in cosmological applications.  GRBs are the most powerful events observed in the cosmos, and can potentially be used to probe the high-redshift universe. Recently, several empirical correlations between GRB observables were reported, which have triggered intensive studies on the possibility in using GRBs as cosmological known candles (Norris et al. (2000), Lloyd-Ronning $\\&$ Ramirez-Ruiz (2002), Ghirlanda et al. (2004a), Ghirlanda et al. (2004b), Dai et al. (2004), Xu et al. (2005), Firmani et al. (2005), Friedman $\\&$ Bloom (2005), Firmani et al. (2006), Schaefer (2007)). Constraints on cosmological parameters from GRBs alone and in conjunction with other geometrical probes, including SNe Ia, the shift parameter from CMB measurements (Wang $\\&$ Mukherjee (2006)) and the $A$ parameter for the signals of baryon acoustic oscillation (BAO) from galaxy redshift surveys (Eisenstein et al. (2005)), have been analyzed (Su et al. (2006), Wright (2007), Wang et al. (2007)). Li et al. (2008), for the first time, performed global fitting on the GRB data with the MCMC technique together with SNe Ia data (Riess et al. (2007)), and data from WMAP (Spergel et al. (2007)), SDSS (Tegmark et al. (2004)) and 2dFGRS (Cole et al. (2005)). On the other hand, the physics behind the empirical correlations are poorly understood. There are also observational indications that some of the reported correlations may have potential problems. Thus there is an ongoing debate for the validity of using GRBs as cosmological candles. The circular problem has been recognized as another obstacle in the cosmological applications of GRBs. Up to now, there have been about $100$ GRBs with measured redshifts, and few are at low redshifts with known distances. Thus it is lack of observational data to calibrate, in a cosmology-independent way, the correlation relations. The reported relations are often given assuming an input cosmology. Applying such relations to constrain cosmological parameters leads to the circular problem. Different methods have been put forward to avoid this problem (Firmani et al. (2005), Schaefer (2007)). All of them are discussed in the context of using geometrical constraints only.  In this paper, we present a method to deal with the circular problem in the MCMC global fitting. It is known that for constraining cosmological parameters, the most reliable way is to perform global fitting from observational data directly. Li et al. (2008) made a first effort to integrate GRBs in the MCMC chains. However, the GRB data they used are released by Schaefer (2007) where the distance moduli of GRB samples are not independent of the input cosmology model and are still subject to the circular problem. Due to this reason, here we aim at introducing a new method to get rid of the circular problem of GRBs in order to avoid biases arising from it, so that the advantage of the MCMC global fitting can be fully realized. We are aware of the current debate regarding the cosmological applicability of GRBs (Bloom et al. (2003), Friedman $\\&$ Bloom (2005)). On the other hand, with both observational and theoretical advances, reliable correlation relations from sub classes of GRBs may eventually emerge. With these considerations in mind, we have made our method as general as possible. It is not limited to any specific correlations. However, to demonstrate its feasibility, we have to work on a concrete example. We choose the Ghirlanda relation (Firmani et al. (2006)) in our present study. Our method can be readily applied to other correlations. The cosmological models to be analyzed include the $\\Lambda$CDM model and dynamical dark energy models with the EoS following $w_{\\rm DE}=w_0+w_1(1-a)$ and $w_{\\rm DE}= w_{0} + w_{1}\\sin(w_2\\ln(a))$, respectively. The dark energy perturbations are fully taken into account. ",
        "conclusions": "In this section, we present our main results. Firstly we consider the constraint on the non-flat $\\Lambda$CDM model using GRB data only. In order to study the dependence of $A$ and $C$ on the cosmology models, we combine the GRB data with other cosmological observational data, such as CMB, LSS and SN, to constrain different dark energy models, namely the $\\Lambda$CDM model and the dynamical dark energy models, Para I and Para II. Considering the degeneracy between the dark energy parameters and the curvature of universe (Zhao et al. (2007b), Clarkson et al. (2007)), we also consider the non-flat case in the dynamical dark energy model Para I. During these calculations, we focus on discussions of constraints on $A$ and $C$, and the effects of GRB data on the determination of cosmological parameters.  In Figure 1, we show the one-point likelihood function for $A$ and $C$, respectively. The black solid line in each panel represents the result for the non-flat $\\Lambda$CDM model using GRB data only. The colored lines are the results from combined analysis of CMB+LSS+SNe+GRB for different cosmological models. The green solid, red dashed, blue dash-dotted, and magenta dotted lines are for the $\\Lambda$CDM, flat model with dark energy parametrization I, non-flat model with dark energy Para I, and flat model with dark energy Para II, respectively.  \\begin{figure}[htbp] \\begin{center} \\includegraphics[scale=0.4]{Fig2a.eps} \\includegraphics[scale=0.4]{Fig2b.eps} \\includegraphics[scale=0.4]{Fig2c.eps} \\caption{2-d joint $68\\%$ and $95\\%$ confidence regions for the parameters $w_0$ and $w_1$ of flat model with dark energy Para I (first panel), non-flat model with dark energy Para I (second panel), and flat model with dark energy Para II (third panel), respectively. The black solid line is given by using WMAP3+SNIa+LSS+GRBs while the red dashed line comes from WMAP3+SNIa+LSS without GRBs. For both cases we considered the dark energy perturbation. \\label{fig2}} \\end{center} \\end{figure}  It is immediately seen that through constraining the cosmological parameters, the data sets other than GRB have some effects on the likelihood of GRB parameters, especially on $C$. Both concerning the non-flat $\\Lambda$CDM model, the black solid line in the right panel of Figure 1 is much broader than the green solid line. This is due to the degeneracy between the correlation parameters of GRB and the cosmological parameters. With GRB data only, cosmological parameters cannot be tightly constrained, and their large uncertainties in turn broaden the likelihood of $C$. Including other data sets greatly reduces the error ranges of cosmological parameters, resulting a much narrower likelihood distribution for $C$. It is also noted that the peak of the distribution of $C$ shifts its position as well. This comparison clearly demonstrates the importance of our methodology in dealing with $A$ and $C$ and the importance of multi-probe analysis.   On the other hand, given the observational data sets, the dependence of $A$ and $C$ on cosmological models is rather weak, as seen from the results shown by the colored lines in Figure 1. The mean value of $A$ in all four cosmological models is $A\\sim 1.53$ with $1\\sigma$ error about $0.08$. For the parameter $C$, its variation for different cosmological models is slightly larger than that of $A$, with the average value of $C=0.94, 0.96, 0.98$, and $0.94$, for the four models, respectively. They are in well agreement with each other within $1\\sigma$ range with $\\sigma \\sim 0.1$. Keeping in mind the hot debate regarding GRBs as cosmological candles due to the lack of thorough understanding of GRB physics and the quality of observational data, the consistency in $A$ and $C$ for different cosmological models seen from our global fitting analysis may hint that the Ghirlanda relation could be intrinsic to a subsample of GRBs.  Because we determine $A$ and $C$ simultaneously with the cosmological parameters, their likelihood distributions shown in Figure 1 have included the effects of uncertain cosmology constrained by the current observational data. Therefore our global fitting analysis give a consistent evaluation on the contribution of GRBs to the determination of cosmological parameters. Both the distributions in $A$ and $C$ are broader than those estimated with a fixed cosmological model. Thus inappropriate use of the narrow uncertainties in $A$ and $C$ resulting from a given cosmology can lead to an overestimate of the power of GRBs in cosmological studies.  The results on the cosmological parameters constrained from our analysis are listed in Table II. In Figure 2, we show, respectively, the constraints on the dark energy parameters $w_0-w_1$ for the flat model with dark energy Para I (first panel), non-flat model with dark energy Para I (second panel), and flat model with dark energy Para II (third panel), respectively. The black solid and red dashed lines represent the results with and without the $27$ GRBs included. For all the cases considered, the flat $\\Lambda$CDM model is well consistent with the observational data with or without GRBs. For the case of flat model with Para I for the dark energy EoS, quintom A type of models with $w_0<-1$ and $w_1>0$ are mildly favored by the data. Including the GRB data in the analysis, the $1\\sigma$ contour shifts more toward the quintom A region. Relaxing the strong prior on the flatness of the universe, the $\\Lambda$CDM gives a better fit to the data than the above case. Considering the oscillating dark energy model, it is seen that the $\\Lambda$CDM model with $w_0=-1$ and $w_1=0$ is again in excellent agreement with the observational data. Comparing the black solid lines with the red dashed lines in Figure 2, we see some shrinkage of the error contours when including GRB data in the analysis, which is largely attributed to the high redshift range of GRBs. This indicates the possible potential in using GRBs as high-redshift cosmological candles. The contribution from the current GRB data is however not greatly significant. Nevertheless, given the apparent advantage of GRBs as the tracers of the high-redshift universe, it is important to perform detailed analysis as we did here to investigate their usefulness in cosmological studies."
    },
    "0711/0711.4348_arXiv.txt": {
        "abstract": "The microquasar SS~433 is an interacting massive binary consisting of an evolved mass donor and a compact companion that ejects relativistic jets. The mass donor was previously identified through spectroscopic observations of absorption lines in the blue part of the spectrum that showed Doppler shifts associated with orbital motion and strength variations related to the orbital modulation of the star-to-disk flux ratio and to disk obscuration. However, subsequent observations revealed other absorption features that lacked these properties and that were probably formed in the disk gas outflow. We present here follow-up observations of SS~433 at orbital and precession phases identical to those from several previous studies with the goals of confirming the detection of the mass donor spectrum and providing more reliable masses for the two system components.  We show that the absorption features present as well as those previously observed almost certainly belong to the mass donor star, and we find revised masses of $12.3\\pm 3.3$ and $4.3\\pm 0.8 M_\\odot$ for the mass donor and compact object, respectively. ",
        "introduction": "The unique system SS~433 is an X-ray binary star that falls in the class of microquasars \\citep{fab04}. The system exhibits relativistic jets originating from a compact object surrounded by an accretion disk. The companion star in the system contributes only a small fraction of the total light and is thus difficult to detect.  Because the accretion disk and its wind are the dominant source of light in the system, reasonable measurements have been made of the radial velocity amplitude of the compact object \\citep*[e.g.][]{fab90,gie02a}.  Models of the precessing jets and eclipses have also provided a very accurate system inclination \\citep[78\\fdg8,][]{mar89}.  Therefore, identifying the mass donor and accurately determining its radial velocity amplitude will lead to direct kinematical masses for both components.  The first potential observations of the mass donor star were reported by \\citet*{gie02b} with additional observations published later by \\citet{che05}.  These studies detected faint absorption lines present in blue spectra of SS~433 that exhibited the Doppler shifts expected for the donor star.  Finding these lines is difficult, not only due to the small relative flux contribution of the A-star, but also to the active emission spectrum from the accretion disk, jets, and strong disk wind in the system. The most convincing observations to date of the mass donor spectrum were obtained by \\citet{hil04}.  Their spectra suggested that the donor star has a spectral classification of A3--7~I, and their radial velocity measurements led to component masses of $M_O=10.9\\pm3.1~M_\\odot$ and $M_X=2.9\\pm0.7~M_\\odot$ (where the subscripts represent the optical companion star {\\it O} and the X-ray emitting compact object {\\it X}).  This identification of the mass donor star was challenged by \\citet{cha04}, \\citet{bar06}, and \\citet*{cla07}.  These authors present spectra of SS~433 that show features found in an A-supergiant, but that do not follow a coherent orbital radial velocity curve. Their observations show a larger scatter, were performed over a greater variety of orbital and precessional phases, and have a much lower average velocity than those presented by \\citet{hil04}.  \\citet{bar06} and \\citet{cla07} suggest that these A-supergiant absorption lines arise in an accretion-driven outflow. Thus, their work demonstrates that not all the absorption features in the spectrum of SS~433 form in the photosphere of the mass donor star.  Some of the differences between these investigations are due to the timing of the observations and the inherent time variability of the spectrum. Given the relative faintness of the mass donor star compared to the super-Eddington disk in the system, \\citet{hil04} planned observations for orbital phases when the disk is partially eclipsed and for precessional phases when the donor star appears above the disk plane at orbital inferior conjunction.  This strategy maximizes the opportunity to reduce the glare of the disk and avoid obscuration by the extensive disk gas.  On the other hand, \\citet{bar06} made observations at a variety of other orbital and precessional phases and made plausibility arguments that the donor star spectrum might appear at some of these (but was not in fact detected).  Given the need to resolve the differing interpretations in the papers of \\citet{hil04} and \\citet{bar06}, we undertook follow-up observations to repeat the observations of \\citet{hil04}.  If the absorption features identified do originate in a disk outflow, then the highly variable nature of the outflow in SS~433 suggests that repeat observations will show significant scatter from prior observations, if in fact the absorption features are still present. However, if the absorption lines originate from the mass donor star, then the highly stable and periodic nature of the orbit should produce observations consistent with the data of \\citet{hil04}. ",
        "conclusions": "We have shown that the behavior of the A-supergiant absorption spectrum observed in SS~433 during eclipse and near precessional phase zero is repeatable. This suggests that the origin of the spectrum is a consistent, well-behaved source. We suggest that the source is the mass donor star in the system rather than the disk or a disk wind.  It is interesting that SS~433 would show such similar spectra from both the mass donor star and the disk/disk wind.  However, the appearance of such absorption might be anticipated if the disk outflow material is clumpy and has a temperature in the line forming region comparable to that of the supergiant.  The radial velocity curve resulting from observations of the mass donor star then allow calculations of the kinematical masses of the two components.  We find masses of $M_O=12.3\\pm3.3~M_\\odot$ and $M_X=4.3\\pm0.8~M_\\odot$ for the mass donor and compact object.  It is possible that the inner hemisphere of the mass donor is irradiated, which would change the mass of both components, potentially reducing the compact object mass to below the neutron star limit.  The behavior of the CCF amplitude suggests that a strong irradiation effect is not present in the data presented here.  If this is the case, then the compact object in SS~433 can be firmly identified as a black hole candidate."
    },
    "0711/0711.3809_arXiv.txt": {
        "abstract": "We place constraints on the dynamics of the Local Group (LG) by comparing the dipole of the Cosmic Microwave Background (CMB) with the peculiar velocity induced by the 2MRS galaxy sample. The analysis is limited by the lack of surveyed galaxies behind the Zone of Avoidance (ZoA). We therefore allow for a component of the LG velocity due to unknown mass concentrations behind the ZoA, as well as for an unknown transverse velocity of the Milky Way relative to the Andromeda galaxy. We infer extra motion along the direction of the Galactic center (where Galactic confusion and dust obscuration peaks) at the $95\\%$ significance level.  With a future survey of the ZoA it might be possible to constrain the transverse velocity of the Milky Way relative to Andromeda. ",
        "introduction": "The amplitude of the dipole of the cosmic microwave background (CMB) is two orders of magnitude larger than the characteristic amplitude of the higher-order multipoles of its anisotropies \\citep{Kogut}.  It is therefore widely believed that the CMB dipole originates from the Doppler effect of our peculiar velocity, which is induced by inhomogeneities in the local universe \\citep{Huchra,Conklin,Henry} rather than by a primordial origin \\citep{Gunn,Piran}. Indeed, 21cm surveys employing the Tully-Fisher relation for distance calibration have inferred that the peculiar velocity of the Local Group (LG) relative to distant galaxies converges within a distance of $\\sim 5,000~{\\rm km~s^{-1}}$ or $\\sim 70$Mpc \\citep{Giovanelli,DG} \\footnote{We note that the distance at which dipole convergence is achieved is still controversial. Surveys of galaxy clusters imply large convergence distances (Plionis et al 2000; Basilakos \\& Plionis 2006; Ebeling et al. 2002, 2005; Kocevski 2005) but are affected by the strong bias and Poisson fluctuations of clusters relative to the underlying matter distribution. In this paper we assume that convergence is reached within the maximum 2MRS distance of 280Mpc or $\\sim 20,000 \\kms$ (see \\S 4.2 for the justification of this assumption).}.  This important result confirms the notion that the peculiar velocity is induced within that distance, since otherwise the distant galaxies would also be moving relative to the CMB together with the LG.  Surveys of galaxies in the local universe have attempted over the past two decades to explain the amplitude and direction of the CMB dipole within a distance of $\\gsim 100$Mpc \\citep{LB,Strauss,conf1,conf2,conf3,conf4} The adopted method assumes that: {\\it (i)} the LG peculiar motion is induced by gravity; and {\\it (ii)} the amplitude of inhomogeneities in the distribution of the observed light from galaxies traces the underlying mass distribution on large spatial scales with a constant bias factor $b$. The latest results, based on the 2 Micron All-Sky Redshift Survey (2MRS, Erdo{\\u g}du et al. 2006), show convergence of the flux-weighted dipole in the galaxy survey out to $\\sim 150~{\\rm Mpc}$ but still indicate a discrepancy of $24^\\circ \\pm 8^\\circ$ with the direction of the CMB dipole.  The main limitation of the 2MRS sample results from the lack of sample galaxies behind the Zone of Avoidance (ZoA), a strip around the Milky Way disk where confusion and dust obscuration compromise the survey efficiency. For lack of better information, the 2MRS analysis is also based on the assumption that the Milky-Way galaxy is moving radially towards the Andromeda galaxy (M31) with no transverse motion \\citep{Vanden}. Our goal in this paper is to constrain the unknown peculiar velocity of the LG within the ZoA as well the unknown transverse speed of the Milky-Way relative to Andromeda, by requiring a match between the 2MRS and CMB dipoles.  The contribution of mass concentrations outside the survey volume of 2MRS can be ignored based on the success of Tully-Fisher surveys in converging to the CMB dipole within the same volume \\citep{Giovanelli,DG}.  The outline of this paper is as follows. We first summarize the existing data on the LG velocity from the CMB and Galactic measurements (\\S 2) as well as from the 2MRS analysis (\\S 3). We then compare the results from the CMB and 2MRS data sets and interpret our results in the context of transverse motion of the Milky Way relative to the LG (\\S 4.1), structure beyond the extent of 2MRS (\\S 4.2), and nearby structure within the ZoA (\\S 4.3).  We conclude that the last effect is the most likely explanation for the discrepancy between the CMB and 2MRS dipoles.  We also derive the likelihood function for the bias parameter $b$ of the 2MRS galaxies. Finally, we discuss the implications of our results in \\S 5. Throughout our analysis, we use Galactic Cartesian coordinates in which the $x$-axis is oriented towards the Galactic Center, i.e., towards longitude $l=0$ and latitude $b=0$, $y$ is in the direction $l=90^\\circ,b=0$, and $z$ is in the direction $b=90^\\circ$. ",
        "conclusions": "We have obtained new constraints on the dynamics of the Local Group by comparing the CMB and the 2MRS dipoles. The analysis is limited by the lack of surveyed galaxies behind the Zone of Avoidance (ZoA).  In order to match the CMB and 2MRS dipoles, we have inferred excess motion (that is not acounted for by 2MRS) along the direction of the Galactic center (\\S 4.2). This happens to be the most natural direction for hiding mass concentrations because it is associated with enhanced confusion and dust obscuration by the Galaxy.  The implications of our findings have a simple interpretation.  To acquire an excess peculiar velocity $\\Delta v$ towards the ZoA over the age of the Universe, $\\Delta t=1.4\\times 10^{10}~{\\rm yr}$, requires an average acceleration of order $g \\sim \\Delta v/\\Delta t$.  Assuming that this acceleration is induced gravitationally (i.e. $g\\sim GM/d^2$) by a hidden object of mass $M$ at a characteristic distance $d$, we get the simple scaling relation for the required mass: $M_{12} \\sim 1.7 v_7 d_0^2$, where $M_{12}=(M/10^{12}M_\\odot)$, $v_7=(\\Delta v/100~{\\rm km~s^{-1}})$ and $d_0=(d/1~{\\rm Mpc})$.  According to Fig. 2, the most likely value of the excess velocity towards the GC is $\\Delta v\\equiv v_{\\rm ZOA,x}\\sim 120~{\\rm km~s^{-1}}$, which requires a hidden galaxy comparable in mass to Andromeda ($M\\sim 2\\times 10^{12}M_\\odot$) at a distance of $\\sim 1$ Mpc, or a hidden galaxy cluster comparable in mass to the Coma cluster ($M\\sim 10^{15}M_\\odot$) at a distance of $\\sim 20$ Mpc. At these distances, the inner 10 kpc diameter of a hidden galaxy would occupy an angle of $\\sim 0.6^\\circ$, and the inner 1Mpc diameter of a hidden cluster would occupy $\\sim 2.5^\\circ$.  An extended supercluster might not be fully hidden behind the ZoA.  As argued in \\S 4.2, it is very unlikely that structure beyond the maximum distance of 2MRS of $\\sim 280$ Mpc ($\\sim 20,000 \\kms$) accounts for the discrepancy. The LCDM power spectrum would typically account for a velocity offset that is an order of magnitude smaller than the central value we infer.  The above possible objects are already constrained by existing data. In particular, a new galaxy cluster must have escaped detection by dedicated ZoA searches in the X-ray band (Ebeling et al 2002, 2005; Kocevski 2005), optical galaxy searches (Roman et al 1998; Wakamatsu et al. 2005; Hasegawa et al. 2000) and 21cm surveys (Meyer et al 2004; Henning et al.2005; Kraan-Korteweg et al. 2007). Similarly, a nearby galaxy must have escaped detection by 2MRS as well as existing 21cm surveys.  We are not in a position to evaluate how likely this is.  We note that the excess mass we predict behind the GC will have an associated infrared flux which is independent of the attractor's mass and is linearly proprtional to $\\Delta v$, since it scales as $\\propto M/d^2$ for a fixed mass-to-light ratio. Erdo{\\u g}du et al. (2006) infer a net luminosity density of $(7.67\\pm 1.02)\\times 10^8 h L_\\odot~{\\rm Mpc^{-3}}$ for the 2MRS galaxies. When compared to the average matter density of the Universe for $\\Omega_m=0.24$ and $h=0.7$, this results in a predicted excess radiation flux of $\\sim 2.4\\times 10^{9} v_7 L_\\odot{\\rm ~Mpc^{-2}}$ behind the GC.  Unfortunately, the current error budget is too large to provide a useful constraint on the transverse velocity of the Milky Way relative to the Andromeda galaxy (\\S 4.1).  A future survey of the ZoA might allow to determine this transverse velocity by requiring a match between the peculiar velocity inferred for the local group and the CMB dipole.  The bias parameter of the surveyed galaxies, $b$, could be determined by requiring that the inferred radial velocity of Andromeda relative to the Milky-Way will match its observed value. It would then be possible set a lower limit on the local group mass so that the two galaxies will be gravitationally bound.  Measuring the relative transverse velocity of the Milky-Way and Andromeda is of great interest, since it would affect the future trajectory of the two galaxies (Cox \\& Loeb 2007) and will determine whether the LG is likely to be gravitationally bound (Binney \\& Tremaine 1986). Current methods for inferring the transverse speed (Loeb et al. 2005; van der Marel \\& Guhathakurta 2007) are indirect and highly uncertain.  A future ZoA survey for infrared or 21cm galaxies (Henning et al. 2005) would provide a new elegant path for constraining this unknown velocity component, which is difficult to measure otherwise.   \\bigskip  {\\bf Acknowledgments} The authors thank John Huchra and Mark Reid for useful discussions. We also thank an anonymous referee for a number of very valuable comments and for pointing out the existing constraints on nearby objects behind the ZoA.  \\newcommand{\\noopsort}[1]{}"
    },
    "0711/0711.4104_arXiv.txt": {
        "abstract": "At the intersection of galactic dynamics, evolution and global structure, issues such as the relation between bars and spirals and the persistence of spiral patterns can be addressed through the characterization of the angular speeds of the patterns and their possible radial variation.  The Radial Tremaine-Weinberg (TWR) Method, a generalized version of the Tremaine-Weinberg method for observationally determining a single, constant pattern speed, allows the pattern speed to vary arbitrarily with radius.  Here, we perform tests of the TWR method with regularization on several simulated galaxy data sets.  The regularization is employed as a means of smoothing intrinsically noisy solutions, as well as for testing model solutions of different radial dependence (e.g. constant, linear or quadratic).  We test these facilities in studies of individual simulations, and demonstrate successful measurement of both bar and spiral pattern speeds in a single disk, secondary bar pattern speeds, and spiral winding (in the first application of a TW calculation to a spiral simulation).  We also explore the major sources of error in the calculation and find uncertainty in the major axis position angle most dominant.  In all cases, the method is able to extract pattern speed solutions where discernible patterns exist to within $20\\%$ of the known values, suggesting that the TWR method should be a valuable tool in the area of galactic dynamics.  For utility, we also discuss the caveats in, and compile a prescription for, applications to real galaxies. ",
        "introduction": "One of the prime unresolved issues in the dynamics and evolution of galaxy disks remains the origin and evolution of large-scale bar and spiral structure.  Though the persistence of grand-design spirals has been tied observationally to the presence of bars or companions \\citep{kn79}, virtually nothing is known about the actual lifetimes of spiral patterns.  Additionally, despite indications that the relation between bar and spiral pattern speeds (which \\citet{ss88} first argued may not be equal) may be important for understanding the role of bars in angular momentum transfer during secular disk evolution (e.g. \\citealt{ds98} and \\citealt{ds00}), there are as yet unanswered questions about the connections between multiple patterns in different radial zones.  While mode-coupling between patterns, which allows efficient outward angular momentum transfer in disks (\\citealt{syg}; \\citealt{mt97}) seems a most promising link, in 2D N-body simulations with a dissipative gas component \\citet{rs} find evidence for spiral structure in the absence of a bar, bar-spiral mode coupling, spiral-spiral mode coupling, and multiple pattern speeds without mode coupling.  Clearly, to address questions about the persistence of spiral patterns and the relation between bars and spirals requires not only determination of the pattern speed but how it varies with radius; only with accurate measurement of bar and spiral or inner and outer spirals pattern speeds in the same galaxy can we confirm whether spiral structure is steady or winding, whether bars and spiral pattern speeds are equal or are unrelated, whether mode-coupling exists, and the domain and number of patterns that can be sustained in a disk.  Because they are not directly accessible through observation, pattern speeds are often determined with indirect means such as the identification of predicted behavior at resonance radii (e.g. \\citealt{elm89}; \\citealt{elm96}) or kinematic and morphological comparisons of simulated and observed structure (e.g. \\citealt{raut2005}; \\citealt{gb}).  It is also clearly desirable to employ methods for estimating pattern speeds which do not rely on theoretical models or simulation.  Many other pattern speed determinations have therefore centered on the use of the model-independent method of Tremaine \\& Weinberg (1984; hereafter TW) which presents a rigorous derivation for the pattern speed $\\Omega_p$ based on the requirement of continuity using observationally accessible quantities.  The determination of $\\Omega_p$ involves surface density-weighted position and velocity line integrals parallel to the galaxy major axis under several essential assumptions.  Specifically, the method requires that the disk of the galaxy is flat (unwarped) and contains a single, well-defined rigidly-rotating pattern; that the surface density of a kinematic tracer of a disk component, which must obey continuity, becomes negligibly small at some radius and all azimuths within the map boundary (thereby critically yielding converged integrals); and that the relation between the emission from this component and its surface density is linear, or if not, suspected deviations from linearity can be modeled.  In that information from all sampled radii is associated with a single, constant pattern speed, the TW calculation poses a challenge for extracting multiple distinct or radially varying pattern speeds. Non-axisymmetric structures beyond a dominant pattern such as a bar in the disk of a galaxy will interfere with the measurement of $\\Omega_p$ of the bar; when a non-axisymmetric disk can be decomposed into two components with different pattern speeds, then the TW estimate is a luminosity- and asymmetry-weighted average of the two patterns \\citep{dgs02}.  The TW estimate for the secondary pattern speed in the inner disk of NGC 6946 observed in H$_\\alpha$, for example,  is estimated to be limited to an uncertainty of as much as 50\\% given the primary pattern's contribution to the TW integrals \\citep{fathi}.  Such issues notwithstanding, there have been several recent adaptations of the TW method to positive effect.  Applications of the TW method to SB0 galaxies using stellar light as a tracer extend the limits of integration in the TW calculation to just past the end of the bar in order to minimize contributions to the TW integrals from non-axisymmetric (and several magnitudes dimmer) features beyond the bar; in such cases, integrating past the structure of interest is found to be sufficient for achieving converged integrals.  The bar pattern speeds in NGC 7079 \\citep{dw04} and NGC 1023 \\citep{dca02}, for example, have both been successfully measured in this way.  When there exist more than one pattern in distinct radial zones, however, arguments about the convergence of the TW integrals are less straight-forward.  To measure the secondary bar pattern speed in NGC 2950, \\citet{cda03} and \\citet{mac06} explore decoupling the inner secondary and outer primary bar pattern speeds by associating each component with unique surface brightness contributions in the TW calculation.  The TW integrals are modified by the presence of the inner pattern based on an assumption about how and where the two patterns decouple (to first order).  This analysis has confirmed the existence of, if not measured, a unique secondary bar pattern speed possibly indicating counter-rotation with respect to the primary bar.  At their best--improving the accuracy of pattern speed estimates to about 20\\% \\citep{GD}--these kinds of adaptations of the TW method for measurement of single or multiple patterns still require assumptions about bar extent based on morphological or kinematic signatures. To separate the observed surface brightness in NGC 2950 into secondary and primary bars components, \\citet{mac06} must assume that the patterns indeed decouple at the inner-bar end, or that the outer pattern is axially symmetric at least within the inner's extent.  There, the transition between the two is inferred from the location of a plateau in the TW surface brightness-weighted position integral as the limits of integration are extended from zero. Perhaps more critically, as investigated in $\\S$ \\ref{sec:SIM1}, the direct association of bar length measured in this manner (or perhaps others) with pattern extent may introduce error into TW calculations.  Furthermore, adaptations of the TW method based on this type of identification are likely to be inapplicable for spiral pattern estimation.  In spiral galaxies, not only can identifying transitions between patterns be less clear, but non-axisymmetric motions are significantly smaller (e.g. \\citealt{rs87}) than in bars (at least those typically analyzed with the TW method).  Indeed, the TW method has yet to be tested on a simulated spiral.  A recent modification to the TW calculation (Merrifield, Rand \\& Meidt 2006; hereafter MRM) in which $\\Omega_p$ is allowed radial variation promises to be an invaluable resource for tests of long-lived density wave theories and for understanding the connection, if any, between bar and spiral pattern speeds. Like the TW method, the so-called radial TW (TWR) method uses measurements of observables to extract radially varying pattern speeds. As first applied using the BIMA Survey of Nearby Galaxies CO observations of the grand-design Sb galaxy NGC 1068 (MRM), the TWR method returned a spiral pattern speed solution that declines with radius, allowing a winding time for the pattern to be estimated (e.g. MRM).  As described in MRM, the nature of the discretized calculation presents numerical solutions that are highly susceptible to fluctuations as a result of compounded noise in the data.  In this paper ($\\S$ \\ref{sec:Regularization}), we develop the TWR method with regularization as a means of smoothing intrinsically noisy solutions as well as testing model solutions of different radial dependence (described in $\\S$ \\ref{sec:Smooth}).  These and other commendations notwithstanding, with regularization one may risk introducing an unrealistic prejudice to TWR solutions.  In $\\S$ \\ref{sec:Bias}, we address a means of identifying when this is likely to occur, and in $\\S$ \\ref{sec:Fourier} describe a scheme for minimizing such regularization-induced bias.  Using evidence that arises from these considerations, as well as other a priori information, theoretically and observationally motivated models for $\\Omega_p(r)$ can be developed which constrain the number and extent of patterns present in a disk.  Once solutions for these models are calculated, the goodness of each must be assessed in order to identify the best-fit solution.  In $\\S$ \\ref{sec:weight}, we outline the criteria with which the models are judged and describe our concept of error evaluation in the final solutions.  Beginning in $\\S$ \\ref{sec:Nbody}, we analyze three simulated galaxies with known pattern speeds (a barred spiral, a slowly winding spiral, and a double barred spiral, marking the first application of the TW method to simulated spiral patterns) in order to develop a general stratagem for application to real galaxies.  As applied to these simulations, we find that the TWR method is able to extract multiple pattern speeds with accuracies on the order of (and, as we will see in some cases, better than) the traditional TW method.  With the barred spiral simulation in $\\S$ \\ref{sec:SIM1} we show that TWR bar pattern speed measurement presents an improvement over traditional TW bar estimates, particularly when there is evidence of a significant contribution from the spiral pattern to the TW integrals.  We find that the regularized TWR method can recover information from both patterns effectively by identifying and treating the bar-to-spiral transition radius (which the TW values themselves may not indicate) as a free parameter in the calculation.  In $\\S$ \\ref{sec:iError}, we analyze the results of the method in detail, particularly with regard to morphological limitations.  We compare our TWR results with TW estimates in $\\S$ \\ref{sec:TWTWR} and examine the influence of systematic errors due to the assumed disk position angle and inclination (shown to be crucial for TW estimates in \\citet{debPA}) on both in $\\S$ \\ref{sec:Err}.  We also explore the reliability of a fixed parameterization for the bar-to-spiral transition radius. ($\\S$ \\ref{sec:Trans}).  In the last third of the paper, we investigate the prospects for extracting spiral pattern speed solutions that are winding in nature ($\\S$ \\ref{sec:SIM2}), marking the first application of the TW method to a spiral simulation, and in $\\S$ \\ref{sec:SIM3} we address the use of the TWR method for the purposes of parameterizing an independently rotating nuclear bar in a double barred simulation. There, the techniques we employ for decoupling and extracting measurements of the pattern speeds of both the primary and secondary bar components may present an interesting corollary to recent attempts with the TW method to measure secondary bar pattern speeds in the presence of a strong primary bar pattern.  We note here that the TWR method is a generalized version of the procedure used on NGC 2950 (see \\citealt{cda03, mac06}); separating the surface brightness into two components can be thought of as the coarsest version of the discretization that is the back-bone of numerical TWR solutions.  Based on our experience with these simulations, we conclude with comments on the applicability of the method to observations of real galaxies in $\\S$ \\ref{sec:caveats} where we also outline a general prescription for using the TWR calculation with regularization. ",
        "conclusions": "In this paper we have shown that regularizing the TWR calculation is an effective means of smoothing intrinsically noisy solutions for more precise measurement of $\\Omega_p(r)$.  Specifically, (barring a large, limiting resolution) regularization admits the use of a much smaller bin width than that required to achieve comparable smoothness in the unregularized calculation.  This affords improved assessment of radial variation as well as more accurate determination of the transitions between multiple pattern speeds (and thus of the values of the pattern speeds themselves, in principle).  Moreover, with the regularized TWR calculation, different theoretically and observationally motivated models for the radial dependence of $\\Omega_p(r)$ can be tested in fairly short time and with only the minor addition of information compared to the unregularized TWR and TW methods.  With a simple scheme for generating $n^{th}$ order polynomial solutions which can be incorporated into step models, we have shown that the TWR method can be used to parameterize the radial domains of multiple pattern speeds. Together with a priori information identifying zones in the disk which may be incompatible with measurement (either because they are characterized by low signal-to-noise or show no evidence for a pattern), we can further constrain the extent of patterns while optimally reducing regularization-induced bias in pattern speed solutions.  As applied to three simulated galaxies, we find that the TWR method developed in this manner performs with a high degree of accuracy (with less than 15\\% error) both in measurement and in extracting information about the true functional form for the pattern speed.  Tests on a simulation of a barred spiral galaxy indicate that not only can the constant pattern speed for a relatively weak spiral be reliably reproduced, but information about both the pattern speed and the radial extent of the bar pattern can also be extracted.  (Indeed, we find that the bar pattern speed estimate is strengthened by the proper use of information from beyond the bar end.)  And though the bar pattern speed estimates is highly susceptible to systematic errors--with PA errors introducing the largest uncertainty to TWR pattern speeds, as with TW estimates--we find that the identification of the transition between the two is relatively stable.  The TWR method can also be effectively employed to measure patterns that are winding in nature. In a simulation of a two-armed spiral, the best-fit TWR solutions from several slice orientations correctly reproduce the high-order radial variation of the pattern speed, despite modest indication that not all orientations supply the same authority (this, of course, would seem to depend on the morphology of this spiral, in particular).  Indeed, though the TWR method can, in principle, handle any (presumably random) alignment of patterns, in all of the simulations studied, slice orientations which provide the most uniform coverage of the patterns are preferred.  This is of particular importance for nuclear bar pattern speed measurement, as found in tests of the method on a double barred simulation.  Since the innermost bins which provide the foremost leverage on the nuclear bar are also the most susceptible to errors from throughout the disk, confident measurement requires all other patterns to be well constrained.  In principle, comparable accuracy should be achievable on real galaxies.  However, these tests do not constrain how well the TWR method can perform under severe observational limitations which may commonly arise.  Not only can determinations of the PA, inclination, and dynamical center be subject to considerable errors given low-quality data, but identifying constraints on the patterns present in the disk to be incorporated into models for $\\Omega_p(r)$ could prove challenging. Additionally, though regularization can reduce the impact of noise on solutions, large measurement errors for each slice could make discriminating between several possible models for $\\Omega_p(r)$ difficult.  And most critically, since the nature of the numerical calculation relies on a relatively small bin width to achieve its greatest accuracy, without high resolution, some observations may not afford practical solutions.  Nevertheless, if restricted to high resolution observations with adequate sensitivity, and given radially stable kinematic parameters, TWR solutions can be used to study the connection between multiple patterns and the nature of spiral winding.  So, too, can we expect progressively more satisfactory applications of the method; though the number of galaxies to which the method can be successfully applied is limited by the current generation of instruments, in the future, larger IFUs, ALMA, and eventually, SKA should yield much higher quality data with larger areal coverage and higher angular resolution.  This prospect in itself should warrant future studies with the TWR method.  \\\\    This research was supported in part by NASA through the American Astronomical Society's Small Research Grant Program.  S. E. M. acknowledges support from a NASA-funded New Mexico Space Grant Consortium Graduate Research Fellowship.  V. P. D. is supported by an RCUK Fellowship at the University of Central Lancashire.  J. S. acknowledges support from a Harlan J. Smith fellowship.  This material is based on work partially supported by the National Science Foundation under grant AST 03-06958 to R. J. R."
    },
    "0711/0711.1498_arXiv.txt": {
        "abstract": "{} % {During the first 40 s after their birth, proto-neutron stars are expected to be subject to at least two types of instability. The first one, the \\textit{convective} instability, is excited in the inner regions, where the entropy gradient produces a Rayleigh-type convection. The second one, the \\textit{neutron-finger} instability, is instead excited in the outer layers where the lepton gradients are large. Both instabilities involve convective motions and hence can trigger dynamo actions that may be responsible for the large magnetic fields in neutron stars and magnetars. However, because they have rather different mean turbulent velocities, they are also likely to give rise to different types of dynamo.} {We have solved the mean-field induction equation in a simplified one-dimensional model of both the convective and the neutron-finger instability zones. Although very idealized, the model includes the nonlinearities introduced by the feedback processes that tend to saturate the growth of the magnetic field ($\\alpha$-quenching) and suppress its turbulent diffusion ($\\eta$-quenching). The possibility of a dynamo action is studied within a dynamical model of turbulent diffusivity where the boundary of the unstable zone is allowed to move. A large number of numerical simulations have been performed in which the relevant parameters, such as the spin-period, the strength of the differential rotation, the intensity of the initial magnetic field, and the extent of the neutron finger instability zone, have been suitably varied.} {We show that the dynamo action can also be operative within a dynamical model of turbulent diffusivity and that the amplification of the magnetic field can still be very effective. Furthermore, we confirm the existence of a critical spin-period, below which the dynamo is always excited independently of the degree of differential rotation, and whose value is related to the size of the neutron-finger instability zone. We provide a relation for the intensity of the final field as a function of the spin of the star and of its differential rotation.} {Although they were obtained by using a toy model, we expect that our results are able to capture the qualitative and asymptotic behaviour of a mean-field dynamo action developing in the neutron-finger instability zone. Overall, we find that such a dynamo is very efficient in producing magnetic fields well above equipartition, and thus that it could represent a possible explanation for the large surface magnetic fields observed in neutron stars.} ",
        "introduction": "}  The present understanding of the processes that produce the large magnetic field strengths observed in neutron stars (NSs) is still far from being complete. Most of the information about their magnetic fields is in fact derived either from their X-ray spectra, or from their spin-down when these NSs are seen as pulsars. While the former reflects a measure of the local surface field $B_{\\rm surf}$, the latter provides information on the global dipolar magnetic field $B_{\\rm d}$, if the spin-down is assumed to be solely due to dipolar electromagnetic emission. These two measures are not always in agreement, showing that the measured magnetic fields may have very different length-scales and intensities. In particular, they seem to suggest the presence of more intense and small-scale surface magnetic fields, together with less strong globally dipolar ones. For the pulsar 1E 1207.4-5209, for instance, the dipolar magnetic field estimated from the spin-down rate is $B_{\\rm d} \\sim 2-4\\,\\times\\,10^{12}\\,\\rm G$ (Pavlov et al. \\cite{pavlov02}), while the surface field estimated from the absorption features in its spectrum is $B_{\\rm surf} \\sim 1.5\\,\\times\\,10^{14}\\,\\rm G$ (Sanwal et al. \\cite{sanwal02}). Similarly, observations of the pulsar RBS B1821-24 (Becker et al. \\cite{becker03}) indicate that $B_{\\rm d} \\sim 10^9\\,\\rm G$, while $B_{\\rm surf} \\sim 10^{11}\\,\\rm G$.  The existence of magnetic fields with different strengths and distributed on different length-scales can be explained in terms of a dynamo mechanism driven by the simultaneous presence of rotation and turbulent motions. During the first $\\sim 40$ s after their birth, proto-neutron stars (PNSs) are expected to develop hydrodynamical instabilities (Epstein \\cite{epstein79}; Livio et al. \\cite{livio80}; Burrows \\& Lattimer \\cite{burrows86}), which can excite a hydro-magnetic dynamo. Such instabilities could be essentially of two types. The first one, driven by the entropy gradient, is a Rayleigh-type {\\em convective instability} (CI) that operates in the inner regions of the star. The second one is a double diffusive instability, driven by both the entropy and leptonic gradients. This is usually referred to as the {\\em neutron-finger instability} (NFI); it operates in the outer regions of the PNS (Miralles et al. \\cite{miralles00}; \\cite{miralles02}) and is expected to evolve by creating finger-like downflows when the neutrinos are still confined (Mezzacappa et al.\\cite{mezzacappa98}). Although some authors have recently raised doubts about the existence of the NFI (Buras et al. \\cite{buras06}, Dessart et al. \\cite{dessart06}, Bruenn et al. \\cite{bruenn07}), no firm conclusion has yet been reached, leaving the debate open. Here, we do not attempt to enter this debate but rather, because of the interesting astrophysical implications that it may have, we will consider the NFI as taking place and having the dynamical properties as described by Miralles et al. (\\cite{miralles00}, \\cite{miralles02}).  We note that the co-existence of the two instability mechanisms produces both a local dynamo process (Thompson \\& Duncan \\cite{thompson93}; Xu \\& Busse \\cite {xu01}) and a mean-field one (Bonanno et al. \\cite{bonanno03}). As shown in Miralles et al. (\\cite{miralles00}, \\cite{miralles02}), the growth-times of the instabilities in the two regions differ by 2 or 3 orders of magnitudes, being $\\tau_{\\rm {\\mbox{\\tiny CI}}} \\sim 0.1\\,\\rm ms$ in the CI zone and $\\tau_{\\rm {\\mbox{\\tiny NFI}}} \\sim 30-100\\,\\rm ms$ in the NFI zone.  Since the typical spin period of a PNS is $P \\sim 100\\,\\rm ms$, the turbulent eddies created by the CI are not influenced by the rotation and therefore they can only excite a local dynamo. On the other hand, the Rossby number in the NFI zone, defined as $R_{\\rm o}\\simeq P/\\tau_{\\rm {\\mbox{\\tiny NFI}}}$, is about unity and the turbulent motions can therefore be influenced significantly by the rotation, favoring the excitation of a global mean-field dynamo. Because of the large difference in the growth-rates of the two instabilities, the two processes, i.e. the local dynamo and the global mean-field one, are essentially decoupled.  Here, we focus our attention on the turbulent mean-field dynamo action that may be excited in the NFI zone. More specifically, we exploit a simple one-dimensional (1D) toy model that aims at capturing, at least qualitatively, the features of the dynamo action.  The model, which includes the nonlinearities introduced by the feedback processes, which in turn tend to saturate the growth of the magnetic field, i.e. $\\alpha$-quenching (Bonanno et al. \\cite{bonanno05}, \\cite{bonanno06}; R\\\"udiger \\& Arlt \\cite{ruediger96}), and suppress its turbulent diffusion, i.e. $\\eta$-quenching (R\\\"udiger \\& Arlt \\cite{ruediger96}), is evolved numerically with a very large variety of initial conditions. These include varying the spin period of the PNS, the strength of the differential rotation between the core and the surface, the intensity of the primordial (seed) magnetic field, and the extent of the NFI zone.  Overall, we find that increasing the extent of the neutron-finger instability zone favours the dynamo excitation, and that the combined action of differential rotation and diffusion can produce an increase in the strength of the generated toroidal field by several orders of magnitude. We also confirm the existence of a critical spin-period, below which the dynamo is always excited independently of the differential rotation strength, and whose value is related only to the size of the neutron-finger instability zone. Since the numerical simulations show that the turbulent mean-field dynamo is very efficient in producing magnetic fields well above the equipartition value, we confirm the results of Bonanno et al. \\cite{bonanno06}; even in presence of both a moving boundary of the instability zones and of an $\\eta$-quenching, i.e. the proposed mechanism could represent an explanation for the large surface magnetic fields observed in several neutron stars.  The paper is organized as follows: in Sect. \\ref{MODEL} we formulate the model and derive the system of partial differential equations, which are then solved with the numerical approach described in Sect. \\ref{NMandT}. In Sect. \\ref{AandR} we describe the methodology and report numerical results. Finally in the conclusions we summarize our results. ",
        "conclusions": "}  We have presented a toy model to describe the amplification of the magnetic field inside a proto-neutron star (PNS) via a dynamo action. The model assumes that a neutron-finger instability (NFI) develops in the outer regions of a PNS during the early stages of its life as discussed by Miralles et al.~(\\cite{miralles00}, \\cite{miralles02}), and that the conditions for the generation of a mean-field dynamo process are met.  Although highly simplified by being only one-dimensional and by adopting the kinematic approximation, our model aims to capture the qualitative features of the dynamo action by including a moving boundary of the instability zone and the nonlinearities introduced by the feedback processes, which saturate the growth of the magnetic field (i.e. $\\alpha$-quenching) and suppress its turbulent diffusion (i.e. $\\eta$-quenching).  In essence, the amplification of the magnetic field is described in terms of a system of coupled partial differential equations of mixed hyperbolic-parabolic type, which are solved numerically for a very large variety of initial conditions. These include varying the spin period of the PNS, the strength of the differential rotation between the core and the surface, the intensity of the primordial (seed) magnetic field, and the extent of the NFI zone.  Overall, we have found that, independently of whether the size of the NFI zone varies in time or not, the amplification of the magnetic field undergoes a first exponential increase with growth-time that is the same for both the toroidal and the poloidal components of the magnetic field. The exact value of the growth-time depends on several parameters and it is roughly in the range $\\tau_{amp} \\sim \\left[ 0.5-2.5 \\right]\\, \\tau_D$, with $\\tau_D$ being the diffusion timescale. The exponential growth then stops through the back reaction of the $\\alpha$-quenching and the magnetic field reaches saturation. The final magnetic field produced at the end of the 40 s of evolution does not depend sensitively on the initial magnetic field, but it does depend on whether the $\\eta$-quenching is active or not, becoming 2-3 orders of magnitude larger in the first case.  Despite its crudeness, our model is also able to capture another important feature of the dynamo mechanism, namely the existence of a critical rotation period $P_c$, above which no dynamo action is possible and the magnetic field simply decays (Bonanno et al. \\cite{bonanno03}). For periods near the critical one, on the other hand, the dynamo is just able to sustain the magnetic field close to its initial value, thus avoiding its decay. However, as the spin rate (or the degree of differential rotation) is increased, the dynamo becomes more and more efficient, amplifying the magnetic field up to values several orders of magnitude larger than the equipartion magnetic field. These very high intensities ($10^{18}\\,G$ for the toroidal component and $10^{14}\\,G$ for the poloidal one) may seem unphysical at first sight; however, here we are considering magnetic fields still inside the neutron star and only the poloidal component is thought to emerge afterwards. Determining the critical period accurately is important to constrain the fraction of neutron stars that may undergo this magnetic-field amplification at birth, and we have found that $P_c$ is in the range $33-600 \\,\\rm ms$ for rigidly rotating PNSs, becoming larger as the degree of differential rotation is increased.  As a result, as long as we have no general constraints on the strength of the differential rotation, a lower limit of $P \\simeq 30$ ms can be taken as the generic threshold below which a mean-field dynamo may be active in a newly born PNS.  Another interesting result of this investigation is that, despite the large parameter space considered, the final value of the magnetic field seems to follow a surprisingly robust dependence with the spin period and the degree of differential rotation, both of which can be summarized in a phenomenological expression of the type \\begin{equation} \\label{scaling} {\\cal B}^{fin} \\propto (C_\\alpha)^\\delta |q|^\\gamma \\end{equation} which holds only for $|q|$ larger than a transition level of differential rotation $|q^*|$, whose exact value depends on the configuration and on the other parameters (overall it is in the range $10^{-2}-10^1$). The exponents $\\delta,\\ \\gamma$ are different for the toroidal and poloidal magnetic field components and depend only very weakly on all of the parameters varied in this analysis. In particular, for the toroidal magnetic-field component we have found $\\delta \\sim 1$ and $\\gamma\\sim 1/2$, while for the poloidal one $\\delta \\sim 0$ and $\\gamma\\sim -1/2$. The exact values of these exponents are likely to be modified by more realistic and multidimensional calculations, but we also expect that the scaling in expression (\\ref{scaling}) will persist in further refinements of this treatment.  The work presented here can be improved in a number of different ways. A first possibility is a more realistic description of the geometry of the problem, with a two or three-dimensional description of the PNS. A second and computationally less expensive alternative is that of improving the nonlinear feedback of the magnetic field on the dynamics of the matter. This can be done by using the same geometry adopted here, but coupling the mean-field induction equation (\\ref{induction}) with the solution of the MHD equations for the conservation of energy and momentum."
    },
    "0711/0711.1588_arXiv.txt": {
        "abstract": "We present an analysis of HI, H$\\alpha$,  and oxygen abundance data for NGC~3741. This galaxy has a very extended  gas disk ($\\sim$ 8.8 times the Holmberg radius), and a dark to luminous (i.e. stellar) mass ratio of $\\sim 149$, which makes it one of the ``darkest'' dwarf irregular galaxies known. However its ratio of baryons (i.e. gas + stellar) mass to dark mass is typical of that in galaxies. Our new high-resolution HI images of the galaxy show evidence for a large scale (purely gaseous) spiral arm and central bar. From our HI data, a rotation curve can be derived out to $\\sim$ 37$-$ 44 disk scale lengths in the J and B band respectively. This is just slightly short of the radius at which one would expect an NFW type rotation curve to start falling. The galaxy has an integrated star formation rate of $\\sim$ 0.0034 M$_\\odot$~yr$^{-1}$, while the average star formation rate within the optical disk is $\\sim$ 0.0049 M$_\\odot$~yr$^{-1}$~kpc$^{-2}$. Despite the gaseous spiral feature and the on-going star formation, we find that the global gas density in NGC~3741 is significantly lower than the Toomre instability criterion. This is consistent with the behaviour seen in other dwarf galaxies. We also find that the star formation rate is consistent with that expected from the observed correlations between HI mass and SFR and the global Kennicutt-Schmidt law respectively. We measure the oxygen abundance to be 12+ log(O/H)=7.66$\\pm$0.10, which is consistent with that expected from the metallicity-luminosity relation, despite its extreme gas mass ratio. We also examine the issue of chemical evolution of NGC~3741 in the context of closed-box model of chemical evolution. The effective oxygen yield of NGC~3741 is consistent with recent model estimates of closed-box yields, provided one assumes that the gas has been efficiently mixed all the way to edge of the HI disk (i.e. $> 8$ times the optical radius). This seems a priori unlikely. On the other hand, using a sample of galaxies with both interferometric HI maps and chemical abundance measurements, we find that the effective yield is anti-correlated with the total dynamical mass, as expected in leaky box models. ",
        "introduction": "\\label{sec:intro}   The processes that govern the conversion of gas to stars in galaxies are complicated and poorly understood. Nonetheless, there do exist empirical star formation recipes which are widely applied to models of galaxy evolution. Typically, these recipes allow one to compute a global star formation rate given the gas column density or, alternatively, some estimate of the dynamical timescale (e.g. the angular rotation frequency $\\omega$). While the existence of such empirical relationships is fairly well established in large galaxies (e.g. Kennicutt 1998), it is unclear whether these are applicable to the smallest dwarf galaxies (e.g. \\citealp{begum06}). However, star formation in such small galaxies is of particular interest in studies of galaxy formation, because, in hierarchical models of galaxy formation, small galaxies form first and then merge to form larger galaxies. Nearby, metal poor dwarf galaxies are, in this sense, analogs of first units of star formation in the universe.  Detailed studies (at resolutions that are not achievable for high redshift objects) of the star formation  in nearby dwarf galaxies could hence give some insight into the star formation processes occurring in primordial galaxies at high redshifts.  \\begin{figure*} \\psfig{file=montage_mom0.ps,width=7.0truein} \\caption{{\\bf{[A]}}The optical DSS image of NGC~3741 (greyscale) with 71$^{''}\\times71^{''}$ resolution integrated HI column density map (contours) overlayed. The contour levels are 0.1, 1.34, 2.59, 3.83, 5.08, 6.32, 7.57, 8.82, 10.05, 11.31, 13.79 and 15.04 $\\times 10^{20}$ cm$^{-2}$. {\\bf{[B]}} J band image of NGC~3741 (grey scale) overlayed on  HI column density map (contours) at 11$^{''}\\times9^{''}$ resolution. The contour levels are 2.8,10.73, 18.66, 16.57, 34.5, 42.4 $\\times 10^{20}$ cm$^{-2}$. {\\bf{[C]}}Greyscale HI image of NGC~3741 at 11$^{''}\\times9^{''}$ resolution . The range of grey scale  is from 2.5$\\times~10^{19}$ cm$^{-2}$ to 4.2 $\\times~10^{21}$ cm$^{-2}$. The image shows one sided spiral arm emanating from a central bar like region on the western side of the galaxy. } \\label{fig:overlay} \\end{figure*}  We have been conducting a systematic multi-wavelength survey (the Faint Irregular Galaxy GMRT Survey -- FIGGS (\\citealp{begum06}) of nearby dwarf galaxies. This survey revealed the truly unusual properties of NGC~3741, with an HI disk that extends to more than $\\sim$8 times its Holmberg radius (\\citealp{begum05}, henceforth Paper~I).  NGC~3741 has the most extended gas disk known. Further, with the observed M${\\rm{_{HI}/L_B=5.8\\pm1.4}}$, NGC~3741 is also one of the most gas-rich dwarf galaxies known. In Paper~I we presented a rotation curve for NGC~3741 that extended out to $\\sim 38$ optical B band scale lengths (Bremnes et al. 2000) and showed that within the last measured point of the rotation curve, the dark matter in NGC~3741 contributes to 92 \\% of its total dynamical mass, whereas stars contribute only 1\\% of the total mass. On the other hand, we also found that its total baryonic (i.e. HI + stellar mass) to dark matter ratio of ${\\rm{M_{bar}/M_T}}\\sim0.08$ falls within the range that is typical of galaxies ($\\sim0.07-1.0$; Paper~I). NGC~3741 thus seems to be a case of inefficient star formation, that is worthy of further study.  In this paper we present new HI and H$\\alpha$ data for NGC~3741. We also estimate the gas-phase oxygen abundance of this galaxy using the spectrophotometric  data from \\cite{mou06}. We use these data to study the the relationship between gas and star formation  and  abundance properties of this very gas-rich galaxy. Throughout the paper, we adopt a distance of 3.03 Mpc, as derived from the tip of the red giant branch (Karachentsev et al. 2003). At this distance 1$^{\\prime \\prime}$ corresponds to 14.7 pc. ",
        "conclusions": ""
    },
    "0711/0711.1407_arXiv.txt": {
        "abstract": "Here we present Australia Telescope Compact Array \\HI\\ maps of 16 \\HI\\ sources in six Group Evolution Multiwavelength Study (GEMS) groups that were previously observed with the Parkes telescope. The higher spatial resolution of the ATCA allows us to clearly identify the optical counterparts for the first time -- most being associated with low surface brightness late-type galaxies. New integrated \\HI\\ maps and velocity fields for each source are presented. We find several interacting systems; one of which contains three galaxies within a common \\HI\\ envelope. Extended \\HI\\ structures in the sample are more consistent with tidal effects than ram pressure stripping. We identify two \\HI\\ detections with previously uncatalogued optical galaxies, and add a total of six newly identified group members to the NGC 3923, 5044 and 7144 groups. ",
        "introduction": "Approximately half of all galaxies reside in galaxy groups (Eke et al.~2004).  Despite being so abundant in the universe, the properties of groups and their constituent galaxies have been less well studied compared to rare galaxy clusters.  Most detailed multiwavelength studies to date (e.g. Mulchaey \\& Zabludoff 1998) have been restricted to a small number of groups.  Galaxies in the dense environments of cluster cores exhibit suppressed star formation rates, neutral hydrogen (\\HI) deficiencies, bulge-dominated morphologies, and redder colors than their field counterparts (e.g. Wilman et al.~2005a,b). The environment of galaxy groups is also thought to have a strong impact on the evolution of constituent galaxies. It has been argued to be the most effective environment for galaxy interactions and mergers due to the low relative velocities in groups (e.g. Mulchaey \\& Zabludoff 1998; Mamon 2006). Given these low velocities, and low intragroup gas densities, it is expected that ram pressure stripping plays a much reduced role in groups compared to clusters, however examples do exist (e.g. Rasmussen, Ponman \\& Mulchaey 2006).   \\HI\\ mapping is an extremely useful technique in understanding galaxy groups, as most \\HI\\ envelopes around galaxies tend to extend further out than the stellar component, and as such, are the first component of the galaxies to react to external forces (Broeils \\& van Woerden 1994). Therefore, the \\HI\\ properties of group galaxies serve as a sensitive indicator of past and present interactions, as well as internal kinematics, mass, and orientation parameters of the galaxies themselves.  Several studies of \\HI\\ in galaxy groups have been made, showing that group galaxies often display \\HI\\ material that is extended and stripped away from their host galaxies (e.g. Gordon et al.~2003; Kantharina et al.~2005; Koribalski \\& Manthey 2005; Kilborn et al.~2006). In the case of compact groups, galaxies in such environments appear to be deficient in \\HI\\ (Verdes-Montenegro et al. 2001) but see Stevens et al. (2004) for a different view point.  To gain a better understanding of the processes operating within groups, the Group Evolution Multiwavelength Study (GEMS) was established. For details of the selection criteria and science goals, see Osmond \\& Ponman~(2004) and Forbes et al.~(2006). Briefly, this is a multiwavelength study of 60 nearby galaxy groups, using X-ray, optical, infrared and \\HI\\ data. All have deep ROSAT X-ray maps available. We have also obtained wide-field  \\HI\\ observations of 16 southern GEMS groups using the multibeam receiver on the Parkes Radio telescope (Kilborn et al.~2007). This Parkes survey, which is about twice as sensitive as the HIPASS survey (Barnes et al. 2001), has discovered several new galaxy group members (McKay et al.~2004; Kilborn et al.~2005) and interesting cases of extended \\HI\\ gas (Kilborn et al.~2007). In the Kilborn et al.~(2007) survey there are $\\sim$50 sources for which the \\HI\\ detection could not be confidently assigned to a single (optical) galaxy. This `confusion' is due to the relatively large beamsize of the Parkes telescope, i.e. 15.5${\\arcmin}$. Here we report the first of our higher resolution (beamsize $\\sim$ 2${\\arcmin}$) \\HI\\ maps with the Australian Telescope Compact Array (ATCA) of 16 confused sources in six GEMS groups. Our aim is to identify the correct optical counterpart and map the \\HI\\ at higher spatial resolution. ",
        "conclusions": "We have followed up, at higher spatial resolution with the ATCA, sixteen \\HI\\ sources in 6 southern GEMS groups that were previously detected by the Parkes telescope for which the optical counterpart was not clearly identified (due to the large beamsize of the Parkes telescope). The superior spatial resolution of the ATCA allows us to clearly identify an optical counterpart in each case. We find a variety of \\HI\\ structures and velocity fields.  These optical counterparts are mostly low surface brightness late-type galaxies, some are clearly tidally interacting galaxies.   Three Parkes sources are clearly associated multiple optical counterparts. The Parkes source GEMS\\_HCG90\\_11 is resolved into \\HI\\ associated with NGC 7204 and ESO 467-G002. There is no sign of an interaction between them, however we note that NGC 7204 itself is a strongly interacting close pair of galaxies. Two Parkes sources in the NGC 7144 group appear to be associated with interacting systems. Optical counterparts KTS 65 A\\&B and ESO 236-G039 are in an interacting system, which reveals a bridge of \\HI\\ within a common \\HI\\ envelope. We suggest it will merge into a single gas-rich galaxy on a $\\sim$100 Myr timescale. The second Parkes source in this group is associated with  ESO 236-G035 and ESO 236-G036. The \\HI\\ contours of both galaxies are roughly aligned with each other. Thus they may also form an interacting system with a relative velocity of only 46 \\kms\\ between them. Interestingly, a friends-of-friends analysis indicates that the NGC 7144 group has yet to form a coherent group structure and may be an example of a group that is still collapsing (Brough et al.~2006).  We find no obvious cases for ram pressure stripping in the \\HI\\ structures. Most galaxies are at projected radii in the outskirts of the groups, where any intragroup medium would be minimal. When irregular, extended \\HI\\ structures were seen they tended to be in interacting systems. Kilborn et al. (2007) measured the \\HI\\ deficiency of the Parkes sources in our 6 groups and 10 other GEMS groups, concluding that ram pressure stripping was not the dominant process in removing \\HI\\ gas.   For two of the \\HI\\ sources (GEMS\\_N3923\\_11 and GEMS\\_N5044\\_18) we have succeeded in identifying previously uncatalogued optical counterparts. In both cases, the optical galaxies appear as faint, low surface brightness objects on the DSS. The \\HI\\ velocity, and spatial location, of these galaxies indicates that they are new members of the NGC 3923 and NGC 5044 groups. We have also confirmed six new group members in the NGC 3923, 5044 and 7144 groups.    Deeper \\HI\\ observations will no doubt reveal additional group members and hence further populate the dwarf end of the group mass (and optical luminosity) function. The recent very deep \\HI\\ observations of Morganti et al. (2006) indicate that the \\HI\\ mass function has a very long tail to low \\HI\\ masses. This suggests that more dwarf galaxies in groups await discovery with future deep \\HI\\ observations."
    },
    "0711/0711.2632_arXiv.txt": {
        "abstract": "{I calculate the emission expected from a Poynting-flux-dominated gamma-ray burst (GRB) flow in which energy is dissipated gradually by magnetic reconnection. In this picture, the energy of the radiating particles is determined by heating and cooling balance (slow heating model). Detailed radiative transfer calculations show that, at Thomson optical depths of order of unity, the dominant radiative process is inverse Compton scattering. Synchrotron-self-absorbed emission and inverse Compton dominate in the Thomson thin parts of the flow. The electrons stay thermal throughout the dissipation region because of Coulomb collisions (Thomson thick part of the flow) and exchange of synchrotron photons (Thomson thin part). The resulting spectrum naturally explains the observed sub-MeV break of the GRB emission and the spectral slopes above and below the break. The model predicts that the  $\\gamma$-ray power-law tail has a high-energy cutoff typically in the $\\sim 0.1-1$ GeV energy range that should be observable with {\\it GLAST}. The model also predicts a prompt emission component in the  optical and UV associated with the GeV emission. Observations of the prompt emission of GRB 061121 that cover the energy range from the optical to $\\sim 1$ MeV are explained by the model. \\\\  {\\bf Key words:} Gamma rays: bursts -- radiation mechanisms: general -- methods: statistical} ",
        "introduction": "\\label{intro}   The GRB emission is the likely result of internal energy release in an ultrarelativistic flow. The dissipative and radiative mechanisms for the GRB largely remain uncertain. A popular model for the energy dissipation invokes internal shocks in an unsteady flow (Paczynski \\& Xu 1994; Rees \\& M\\'esz\\'aros 1994). An alternative proposal is magnetic dissipation in a strongly magnetized flow (Thompson 1994; Spruit et al. 2001)  Internal shocks are efficient in dissipating a large fraction of the kinetic energy of the flow provided that it is highly variable, i.e., composed of distinct ejection events with strong variation in their bulk Lorentz factor $\\gamma$ (e.g. Kobayashi et al. 1997). Energy is dissipated by the shocks at the location of the collision of the shells. Particles are assumed to be accelerated on a very short timescale at the shock front to ultrarelativistic speeds and non-thermal distributions. Subsequently, they radiate a fraction of the dissipated energy via synchrotron and inverse Compton processes. The relevant radiative mechanisms and the emitted spectra depend, to a large extent, on the shock microphysics and the corresponding Thomson optical depth of the flow at the radius of the collision (see, e.g., Pe'er et al. 2006)  On the other hand, the energy dissipation that powers the GRB emission may be gradual and distributed over a large part of the volume of the flow. The energy of the radiating particles is determined by the heating/cooling equilibrium (slow heating model; Ghisellini \\& Celotti 1999; Stern \\& Poutanen 2004). Such an energy balance is expected to lead to sub-relativistic or mildly relativistic temperatures in the flow. Magnetic dissipation in a strongly magnetized flow can provide a possible physical setup where gradual dissipation is realized. As shown in Drenkhahn (2002) and Drenkhahn \\& Spruit (2002; hereafter DS02), dissipation though reconnection takes place over several decades in radius; typically in both Thomson thick and thin conditions. Dissipation in the reconnection model is responsible for both the acceleration of the flow and the prompt emission.  \\subsection{Emission from magnetic dissipation}  In previous works (Giannios 2006; Giannios \\& Spruit 2007; hereafter G06 and GS07 respectively), we studied the radiative transfer close to the Thomson photosphere of the flow in the context of the magnetic reconnection model. The detailed Monte Carlo calculations have shown that, due to energy release, the flow develops a hot photosphere with comoving electron temperatures of tens of keV. In the photospheric region, Coulomb collisions are fast enough to thermalize the electron distribution. Upscattering of photons that are produced deeper in the flow by those hot electrons leads to a powerful photospheric emission; it accounts for $\\sim3-30$\\% of the luminosity of the flow. The resulting  $E\\cdot f(E)$ spectrum has a characteristic $\\sim 1$ MeV peak followed by a flat high-energy power-law emission.  Here, I extend the radiative transfer calculation to the Thomson thin region of the flow where, in the reconnection model, there can be substantial energy release and associated emission. Because of the strong magnetic fields synchrotron self absorption results in efficient energy exchange of the electrons, keeping their distribution thermal (Ghisellini et al. 1998). I show that heating and cooling balance in that region leads to electron temperatures of the order of the electron rest mass. Under these conditions, synchrotron-self-absorbed (SSA) emission is an important radiative mechanism in addition to inverse Compton. SSA dominates the observed emission in the soft X-rays and softer bands.  The efficient thermalization of the emitting particles throughout the flow reduces the dependence of the model on the, poorly understood, mechanisms of particle acceleration that operate in magnetic reconnection. The model is defined by just 3 main parameters (luminosity, baryon loading and a reconnection-rate parameter; see Sect. 3). In contrast to the internal shock model, no quantities have to added to parameterize the particle distributions and the amplification of magnetic fields; the field strength is an integral part of the reconnection model. Because the model is so well defined, it makes direct and stable predictions for the emitted spectrum. In the following, I show how this spectrum can be computed.  The structure of the paper is the following. In the next section I summarize and contrast the main aspects of internal shock and magnetic reconnection models for the prompt GRB emission. The dynamics of the flow in the reconnection model and the radiative transfer in the flow are the topic of Sects. 3 and 4 respectively. The resulting spectra and their direct comparison with multi-frequency observations (from optical to $\\gamma$-rays; see Page et al. 2007) of GRB 061121 are presented in Sect. 5. I discuss the results and conclude in Sect. 6. ",
        "conclusions": "\\label{conclude}  The GRB emission may be result of internal collisions in a variable flow (Rees \\& M\\'esz\\'aros 1994). In the internal shock model energy is dissipated at the location of the shell collision. Particles are accelerated at the shock front to ultrarelativistic speeds and non-thermal distributions. Magnetic fields are assumed to be amplified because of plasma instabilities. The emitted spectrum depends on the distribution of the emitting particles and the strength of the shock amplified magnetic field, both of which are not understood from first principles. The emitted spectrum also depends on the properties of the flow (such as optical depth) at the radius of the collision (e.g. Pe'er et al. 2006).  As an alternative to internal shocks,  magnetic dissipation in a strongly magnetized flow can power the GRB (Thompson 1994). Magnetic dissipation may lead to gradual release of energy over a wide range of radii (e.g. Drenkhahn 2002; DS02). It typically proceeds in both Thomson thick and thin regions of the flow. The released energy can be distributed to a large fraction of the particles of the flow leading to the slow heating scenario for the GRB emission (Ghisellini \\& Celotti 1999). The emitting particles (i.e. electrons) are heated up to mildly relativistic speeds. Because of the strong magnetic fields exchange of synchrotron photons provides an efficient mechanism for the thermalization of the electron distribution (Ghisellini et al. 1998). Since the emitting particles are thermal the resulting emission does not depend sensitively on poorly understood physics of particle acceleration in magnetic reconnection. The model is defined by just the luminosity of the flow, its baryon loading and a reconnection-rate parameter and makes direct and stable predictions for the electromagnetic spectrum.  In previous works (G06 and GS07), we calculated the photospheric emission expected from the reconnection model. The radiative transfer study was made with Monte Carlo simulations. Those calculations have shown that the  flow is characterized by powerful photospheric emission  with most of the energy appearing in the hard X-rays and $\\sim1$ MeV $\\gamma-$rays. This emission is the result of photons, produced deep into the flow, that are inverse Compton scattered by sub-relativistic electrons at Thomson optical depths of order of unity.  In the reconnection model, for low enough baryon-loading of the flow, there are substantial amounts of energy dissipated in its outer, Thomson thin, parts. Here I have extended the G06 calculations to include the Thomson thin emission. Energy release at large radii leads to mildly relativistic electrons that cool down though emitting synchrotron radiation and inverse Compton scattering soft photons. SSA emission dominates the observed radiation in the soft X-ray and softer bands.  This soft (e.g. optical) emission is very weak in models where dissipation is limited below or around the Thomson photosphere (see, e.g., M\\'esz\\'aros \\& Rees 2000; Pe'er et al. 2006; Ioka et al. 2007). Inverse Compton in the Thomson thin region leads to a flat high-energy spectrum that extends up to GeV energies.  The resulting spectra from the radiative transfer calculations naturally explain the observed sub-MeV break of the GRB emission and the spectral slopes above and below the break (Band et al. 1993). Furthermore, the model makes rather robust predictions for the emission in other energy bands. The flat $\\gamma$-ray spectrum is expected to show a cutoff in the $\\sim 0.1-1$ GeV energy range that should be observable with {\\it GLAST}. The optical and ultra violet emission can be powerful and the optical spectrum hard with photon number index $0\\simless \\Gamma\\simless 1$. Bright prompt optical emission is predicted to be accompanied by powerful $\\sim$ GeV emission and rather soft spectra below the sub-MeV break. Comparison with multi-frequency observations of the the prompt emission from GRB 061121 that span from the optical to the $\\sim$ MeV range (Page et al. 2007) supports the model."
    },
    "0711/0711.0771_arXiv.txt": {
        "abstract": "{We demonstrate the power of the local correlation tracking technique on stellar data for the first time. We recover the spot migration pattern of the long-period RS\\,CVn-type binary $\\sigma$\\,Gem from a set of six Doppler images from 3.6 consecutive rotation cycles. The resulting surface flow map suggests a weak anti-solar differential rotation with $\\alpha\\approx-0.0022\\pm0.0016$, and a coherent poleward spot migration with an average velocity of $220\\pm10$\\msec. This result agrees with our recent findings from another study and could also be confirmed theoretically.} ",
        "introduction": " ",
        "conclusions": "\\label{disc}  We have demonstrated that LCT technique can be used to map flow fields on stellar surfaces. Applying LCT for time-series Doppler images of $\\sigma$~Gem we reconstruct large-scale surface flows in the order of several hundreds of \\msec\\ and observe complex network of convergent flows around spots. Zonal flow components suggest rigid-body rotation or anti-solar DR with a very weak surface shear of $\\alpha=-0.0022\\pm0.0016$, while meridional flow components show a coherent poleward spot migration at an average velocity of $220\\pm10$\\msec. For a more detailed description of the method as well as for numerical tests on artificial data, we refer to a forthcoming paper (\\v{S}vanda et al. \\cite{lctlq}).  Our findings agree well with the results from a recent application of the ACCORD method on the same data (K\\H{o}v\\'ari et al. \\cite{AARNsgem}). ACCORD resulted also in anti-solar DR law, but with a stronger shear of $\\alpha=-0.021\\pm0.005$. From the latitudinal cross-correlation a coherent poleward flow was derived with a rate of $\\approx350$\\msec, which is in qualitative and also in quantitative agreement with the meridional flow component from LCT. The deviation could be originated from the different limitations of the two different approaches.  Sufficiently fast meridional flow can result in anti-solar type DR and the necessary angular momentum transport (i.e., the necessary rate of the meridional flow) can be estimated from Eq.~15 in Kitchatinov \\& R\\\"udiger (\\cite{kitrued}). Taking a turnover time of $\\tau=5.5\\times10^6$\\,s with a mixing length $l$ of $7\\times10^8$\\,m (cf. Gunn, Mitrou \\& Doyle \\cite{gumido98}, Patern\\`o et al. \\cite{pat02}) would yield $\\approx300$\\msec, which agrees with our results derived from two different methods. If the poleward surface migration on $\\sigma$\\,Gem is interpreted due to an underlying meridional circulation, our result could be regarded as a confirmation of the theoretical estimates in Kitchatinov \\& R\\\"udiger (\\cite{kitrued})."
    },
    "0711/0711.2297_arXiv.txt": {
        "abstract": "\\noindent We study spherically symmetric thin-shell wormholes in a string  cloud background  in (3+1)-dimensional spacetime. The  amount  of exotic matter required for the construction, the traversability  and the stability  under radial perturbations, are  analyzed as  functions of the parameters of the model. Besides, in the Appendices a non perturbative approach to the dynamics and a possible extension of the analysis to a related model are briefly discussed. ",
        "introduction": "After the well known leading paper by Morris and Thorne \\cite{motho}, considerable attention has been devoted to the study of  traversable Lorentzian wormholes  \\cite{visser}.  Such kind of geometries would  connect two regions of the same universe, or of two universes, by a traversable  throat.  If they actually exist, they would present some features of particular interest as, for example, the possibility of using them for  time travel \\cite{morris,novikov}. However, the  flare-out condition \\cite{hovis1} to be  satisfied at the throat  requires the presence of exotic matter, that is, matter which violates the null energy condition (NEC) \\cite{motho,visser,hovis1,hovis2}. However, it was shown in Ref. \\cite{viskardad}   that the amount of exotic matter necessary for the existence of a wormhole can be made infinitesimally small by  suitably choosing  the geometry. After this,  special attention has been devoted to quantifying the amount of exotic matter  \\cite{bavis,nandi1}; in particular, this  amount  has been pointed as an indicator of the physical viability of a traversable wormhole \\cite{nandi2}. Besides, for a wormhole to be considered traversable, the geometry must be such that the magnitude of tidal forces are admissible for an hypothetic traveller; this has also been analyzed for most physically meaningful configurations studied in the literature.  Of course, no unstable solution of the equations of gravitation could be of interest as a candidate for a traversable wormhole. Thus, besides the characterization of static solutions, their stability under perturbations must always be explored. In particular, this has been thoroughly studied for the case of small perturbations preserving the symmetry of the original configuration. A class of wormholes for which Poisson and Visser \\cite{poisson} developed a straightforward approach for analyzing this aspect  are thin-shell ones, that is, wormholes which are mathematically constructed by cutting and pasting two manifolds to obtain a geodesically complete new manifold \\cite{mvis}. In these wormhole configurations  the exotic matter lies in a shell placed at the joining surface; so the theoretical framework for dealing with them is the Darmois--Israel formalism, which leads to the Lanczos equations, that is, the  Einstein's equations  projected on the joining surface \\cite{daris,mus}. Once an equation of state for the exotic matter in the shell is provided, the solution of the Lanczos equations gives the dynamical evolution of the wormhole.  Such a procedure has been subsequently followed  to study the stability of more general spherically  symmetric configurations (see, for example, Refs. \\cite{eirom}).    According to present day theoretical developments, a scenario in which the fundamental building blocks of nature are extended objects instead of point objects should be considered quite seriously. In particular, 1--dimensional objects (strings)  are the most popular candidate for such fundamental objects. The study of the gravitational effects of matter in the form of clouds of both cosmic and fundamental strings has then deserved considerable attention; see for example Refs. \\cite{varios}. We are interested in the viability of wormholes; so within this framework any reasonable configuration including more parameters and thus allowing for improving its features as: amount of exotic matter, pressure, traversability and stability, deserves to be analyzed.   In the present work we  start from the metric proposed in the  leading paper by Letelier  \\cite{letelier} to address the study of thin-shell wormholes associated to a string cloud (and also  a global monopole; see below). We  study in detail the amount of exotic matter required for the construction, the traversability and also the stability of the configuration under perturbations preserving the original symmetry. As we shall see, comparing with the  Schwarzschild case, the string cloud allows for more freedom in the choice of the configurations to be stable, and also allows to reduce the amount of exotic matter without increasing the pressure. Because recently certain attention was devoted to exotic matter fulfilling the Chaplygin gas equation of state, in  Appendix A we explicitly impose it on the shell matter to obtain its time evolution beyond a perturbative approach. In Appendix B we discuss a possible extension of the analysis to a related geometry.  Throughout the paper we set units so that $c=G=1$. ",
        "conclusions": "We have built  traversable thin-shell wormholes applying the cut and paste procedure to the  geometry corresponding to a spherical cloud of strings. We have found that the amount of exotic matter --which is restricted to the throat-- can be reduced by a suitable choice of the parameters. Moreover, we have shown that for a fixed wormhole radius   $b_0\\gg r_{hor}$, there exists a range of values of the parameter $a$ such that, while the amount of exotic matter can be  reduced, the transverse pressure and the surface energy density can be kept under control. We have also studied the stability of the configuration under perturbations preserving the spherical symmetry, and we have  found that the stability regions --though shifted away--  become enlarged when the parameter $a$ is increased. Within the characterization of static configurations, we have considered the traversability  of large wormholes by evaluating the tidal force; we have found that it turns to be acceptable in the case of values of  $a$ and $b_0$ such that the wormhole is stable. Besides, in the Appendix A (see below) we briefly study the dynamics of the shell beyond the perturbative approach by considering the particular case of the Chaplygin equation of state for the shell matter. Also, the features of a related wormhole configuration are discussed in the Appendix B."
    },
    "0711/0711.2068_arXiv.txt": {
        "abstract": "We have obtained \\numptelobs{} near-infrared (NIR; \\jhk{}) measurements of \\numsneianewall{} Type Ia supernovae (\\sneia{}) using PAIRITEL, the 1.3-m Peters Automated InfraRed Imaging TELescope at Mount Hopkins, Arizona. This new set of observations nearly doubles the number of well-sampled NIR \\snia{} light curves. These data strengthen the evidence that \\sneia{} are excellent standard candles in the NIR, even without correction for optical light-curve shape. We construct fiducial NIR templates for normal \\sneia{} from our sample, excluding only the three known peculiar \\sneia{}: SN~2005bl, SN~2005hk, and SN~2005ke. The $H$-band absolute magnitudes in this sample of \\numsneianew{} \\sneia{} have an intrinsic RMS of only \\hptelrms{}~mag with no correction for light-curve shape. We found a relationship between the $H$-band extinction and optical color excess of $A_H=0.2 E(B-V)$. This variation is as small as the scatter in distance modulus measurements currently used for cosmology that are based on optical light curves after corrections for light-curve shape Combining the homogeneous PAIRITEL measurements with \\numsneialit{}~\\sneia{} from the literature, these \\numsneia{}~\\sneia{} have standard $H$-band magnitudes with an RMS scatter of \\hrms{}~mag. The good match of our sample with the literature sample suggests there are few systematic problems with the photometry. We present a nearby NIR Hubble diagram that shows no correlation of the residuals from the Hubble line with light-curve properties. Future samples that account for optical and NIR light-curve shapes, absorption, spectroscopic variation, or host-galaxy properties may reveal effective ways to improve the use of \\sneia{} as distance indicators.  Since systematic errors due to dust absorption in optical bands remain the leading difficulty in the cosmological use of supernovae, the good behavior of \\snia{} NIR light curves and their relative insensitivity to reddening make these objects attractive candidates for future cosmological work. ",
        "introduction": "Type Ia supernovae (\\sneia{}) observed at optical wavelengths in the rest frame have played a leading role in extragalactic astronomy and cosmology in the past decade.  \\sneia{} have been the key to measuring the Hubble constant \\citep{freedman01,jha99,riess05} and demonstrating cosmic acceleration \\citep{riess98,perlmutter99}. These discoveries rest on a foundation of photometric and spectroscopic similarities between high- and low-redshift \\sneia{} \\citep{hamuy96,riess99,goldhaber01,hook05,blondin06,jha06,conley06,garavini07,garg07,bronder07,foley07,ellis07}. The history of cosmic expansion has been traced out to redshift $z \\sim 1$ \\citep{tonry03,knop03,barris04}.  The predicted turnover to the matter-dominated era and the corresponding deceleration at $z>1$ have been seen by space-based work from {\\em HST}~\\citep{riess07}. Two current projects: ESSENCE (Equation of State: SupErNovae trace Cosmic Expansion; \\citealt{miknaitis07,wood-vasey07}), and SNLS (SuperNova Legacy Survey; \\citealt{astier06}) seek to constrain the nature of the dark energy responsible for cosmic acceleration by measuring the equation-of-state parameter $w$ with large, homogeneous samples of \\sneia{}. ESSENCE and SNLS data are both consistent with dark energy described by a cosmological constant, $w=-1.0$ to better than $\\pm0.1$.  These cosmological results are based on measurements of optical emission in the supernova rest frame, but recent work suggests that \\sneia{} may be superior distance indicators in the near infrared (NIR), with a narrow distribution of peak \\jhk{} magnitudes and 4--8 times less sensitivity to reddening \\citep{meikle00,krisciunas04a,krisciunas07}. Reddening and absorption of supernova light by dust in their host galaxies poses the single most vexing systematic question facing \\snia{} cosmology (\\citealt{wang06,conley07}).  Avoiding these complications would be very desirable, if the properties of the supernovae permit.  These questions can be approached by building a homogeneous and large set of NIR observations.  This paper is a step in that direction.  Early photometric NIR observations of \\sneia{} were made by \\citet{kirshner73,elias81,elias85,graham88,frogel87}, with more recent \\jhk{} light curves published by \\citet{jha99,hernandez00,valentini03,candia03,krisciunas00,krisciunas01,krisciunas03,benetti04,krisciunas04b,krisciunas04c,krisciunas05a,krisciunas06,elias-rosa06,elias-rosa07,krisciunas07,phillips06,pastorello07a,pastorello07b,prieto07,stritzinger07,wangx08,pignata08,taubenberger08}. Northern hemisphere work from the CfA Supernova Program uses observations from the robotic 1.3-m Peters Automated InfraRed Imaging TELescope (PAIRITEL; \\citealt{bloom06b}) at Mount Hopkins, Arizona.  The NIR observations are part of a systematic program of supernova observations that includes dense sampling of supernova spectra \\citep{matheson07} and concurrent $UBVri$ optical photometry (\\citealt{hicken08}, {\\it in prep.}).  This paper is a first report on the NIR data on \\sneia{} from PAIRITEL. Similar work is underway in the southern hemisphere by the Carnegie Supernova Project using observations at the Las Campanas Observatory in Chile \\citep{freedman05, hamuy06}. The prospect of large homogeneous data sets coupled with progress in modeling \\snia{} NIR light curves~\\citep{kasen06} raises hopes that \\sneia{}, especially in the rest-frame $H$ band, can be developed into the most precise and accurate of cosmological distance probes.  Recent work by \\citet{krisciunas04a} shows that \\sneia{} have a narrow range of luminosity in \\jhk{} at the time of $B$-band maximum light (\\tbmax{}) with smaller scatter than in the $B$ and $V$ bands.  \\citet{krisciunas04a} found no correlation between optical light-curve shape and intrinsic NIR luminosity. The NIR behavior is in sharp contrast with the optical light curves, where a variety of ingenious methods have been devised to reduce the scatter in distance estimates: the $\\Delta m_{15}$ method (\\citealt{hamuy96,phillips99,prieto06}), the multicolor light-curve shape method (MLCS/MLCS2k2; \\citealt{riess96,riess98,jha07}), the ``stretch'' method \\citep{perlmutter97,goldhaber01}, the color-magnitude intercept calibration method (CMAGIC; \\citealt{wang03}), the spectral adaptive template method (SALT/SALT2; \\citealt{guy05,astier06,guy07}), and SiFTO \\citep{conley08}. For a sample of 16 \\sneia{} observed in NIR bands, \\citet{krisciunas04a} found a RMS of $\\sigma_J=0.14$, $\\sigma_H=0.18$, and $\\sigma_{K_s}=0.12$~mag. We set out to construct an independent set of observations to test this remarkable result.  In this paper we present \\numsneianewall{}  \\sneia{} observed with PAIRITEL from 2005--2007. The \\numptelobs{} individual data points in this sample represent the largest homogeneous set of \\sneia{} to date and double the NIR \\snia{} observations in the literature. Data collection with PAIRITEL is discussed in \\S\\ref{sec:data}.  Data analysis, including the mosaic creation process and our photometry pipeline, is discussed in \\S\\ref{sec:analysis}. We construct NIR \\jhk{} templates from this new sample as detailed in \\S\\ref{sec:template} and fit these templates to the \\snia{} light curves to derive the \\jhk{} magnitudes at the time of $B$-band maximum light. These magnitudes are remarkably uniform, particularly in the $J$ and $H$ bands.  Distances are calculated in \\S\\ref{sec:redshifts}, and we discuss the standard magnitude of NIR \\sneia{} in \\S\\ref{sec:standard}.  We then compare our sample to that compiled in \\citet{krisciunas04a} as well as to some more recent \\sneia{}~\\citep{krisciunas05a,krisciunas06,krisciunas07,phillips06,pastorello07b,stritzinger07,wangx08,elias-rosa06,elias-rosa07,pignata08,taubenberger08} in \\S\\ref{sec:litsneia} and show that the $J$ and $H$ bands at maximum light, even from this heterogeneous sample, have a small range in absolute magnitude. Our conclusions are summarized in \\S\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  We have constructed improved NIR templates to find accurate luminosity distances to \\sneia{} using the $H$-band light curve from $[-10,+20]$~days after $B$-band maximum light. Within the photometric and local-flow uncertainties, the distribution about a best-fit \\lcdm{} model is consistent with no intrinsic dispersion of our $H$-band template fit magnitude values. Even without accounting for peculiar velocity uncertainties, the RMS scatter of \\hflowrms{}~mag from these {\\em uncorrected} inferred luminosity distance moduli is as small as {\\em corrected} distances from optical-based methods. No correction for light-curve shape has been made here and the treatment of dust extinction is relatively limited in its extrapolation over a factor of 2 in wavelength. It is quite plausible that using information on the NIR and optical light-curve shapes, dust absorption as fit from the optical to the infrared, spectroscopic variation, or host-galaxy properties will result in even smaller scatter and better distance determinations.  The \\citet{krisciunas04a} result was determined from an inhomogeneous sample of 16 light curves. This paper improves on that work by providing a homogeneous sample of \\numsneianew{} PAIRITEL \\snia{} light curves that more than doubles the number of independent photometric observations in the literature. With our ongoing program at Mt.~Hopkins adding $\\sim10$--$15$ \\sneia{} per year, PAIRITEL will continue building an extensive nearby training set of ground-based NIR \\sneia{}. While we observe all the \\sneia{} that meet our criteria, objects in the Hubble flow observed with adequate integration times to beat down the measurement errors will be the most helpful.  In this paper, we considered reddening values derived from optical data, and found a relationship of apparent $A_H = 0.06 A_V$ between the absolute $H$-band magnitude and the optical reddening. However, recent work on the highly extinguished SN~2002cv~\\citep{elias-rosa07} shows that reddening corrections will be necessary for some highly-absorbed \\sneia{} in the NIR. Reddening uncertainties derived from optical data alone currently represent the most significant systematic error affecting \\snia{} luminosity distance measurements. Simulations by \\citet{krisciunas07} demonstrate major improvements to be gained from the addition of NIR data.  \\citet{krisciunas07} show that simulated distance modulus errors are improved by factors of 2.7 and 3.5 by adding $J$, and \\jhk{} to $UBVRI$ data, respectively. Reducing systematics due to reddening are crucial to future space-based \\snia{} surveys which will be large enough to avoid limitations from sample size. In that case, improving the constraints on cosmological parameters will be limited by systematics. We will explore the utility of a full combination of optical and NIR data to more precisely measure and correct for reddening in \\citet{friedman08}.  Although ground-based NIR data can be obtained for low redshift objects, rest-frame NIR observations for high-redshift \\sneia{} will have to be done from space. Currently, rest-frame \\snia{} Hubble diagrams of high-redshift \\sneia{} have yet to be constructed beyond the $I$ band \\citep{freedman05,nobili05}, with limited studies of \\sneia{} and their host galaxies conducted in the mid-infrared with Spitzer \\citep{chary05,gerardy07}.  Because nearby \\sneia{} are excellent standard candles in the NIR, it may be worth careful consideration of the rest-frame NIR for the space missions using \\sneia{} for cosmology. For example, the SNAP~\\citep{aldering04} and DESTINY satellites \\citep{lauer05}, candidates for the NASA/DOE Joint Dark Energy Mission (JDEM) mission, are both currently designed with detectors sensitive out to $1.7$~$\\mu$m, which will only detect rest-frame $H$-band light ($1.6$~$\\mu$m) out to $z \\sim 0.1$.  Only a detector capable of observing rest-frame $H$-band at $z \\sim 0.5$--$2$ could take full advantage of Nature's gift to us of a superb standard candle in the rest frame $H$-band. Such a detector would require sensitivity from $2$--$5$~$\\mu$m. We can look forward to work of this type with JWST, just over the horizon. The astronomical community should discuss which aspects of a JDEM mission could best be carried out in the rest-frame NIR."
    },
    "0711/0711.3478_arXiv.txt": {
        "abstract": "Using a hydrodynamic adaptive mesh refinement code, we simulate the growth and evolution of a galaxy, which could potentially host a supermassive black hole, within a cosmological volume. Reaching a dynamical range in excess of 10 million, the simulation follows the evolution of the gas structure from super-galactic scales all the way down to the outer edge of the accretion disk. Here, we focus on global instabilities in the self-gravitating, cold, turbulence-supported, molecular gas disk at the center of the model galaxy, which provide a natural mechanism for angular momentum transport down to sub-pc scales. The gas density profile follows a power-law $\\propto r^{-8/3}$, consistent with an analytic description of turbulence in a quasi-stationary circumnuclear disk. We analyze the properties of the disk which contribute to the instabilities, and investigate the significance of instability for the galaxy's evolution and the growth of a supermassive black hole at the center. ",
        "introduction": "It is increasingly evident that the growth histories of different components of galaxies: stars, gas, and supermassive black holes, are intricately connected. Observations indicate a relationship between the masses of supermassive black holes (SMBHs) and various properties of their host galaxies, such as the spheroid mass \\citep{Magorrian98} and the velocity dispersion of stars in the bulge \\citep{FM00, Geb00, Trem02}. Recent simulations have shown that feedback from accreting SMBHs can regulate the growth of the black holes as well as the evolution of their host galaxies, making feedback a potentially important piece of the theory of SMBH-host galaxy co-evolution \\citep[e.g.][]{DiMatteo05, Springel05, DiMatteo07, Sijacki07}.  Building up the mass of a galaxy and driving AGN feedback requires a continuous replenishment of fuel in the center of the galaxy. A comprehensive understanding of such fueling is not possible without detailed knowledge about matter transport from large scales to the vicinity of the black hole. That transport helps determine, in particular, both the amount of material available for accretion, feedback, and star formation, as well as the strength and duration of the fueling. Large scale gravitational tidal fields created during major mergers of galaxies are thought to be effective at funneling matter toward the centers of galaxies. Indeed, simulations and semi-analytic modeling of hierarchical growth scenarios have shown that a combination of accretion and black hole mergers can effectively build-up the local black hole population \\citep{KaufHae00,YooME04,VolRees05,Malbonetal06,VolRees06,Lietal07}, as well as assemble the $\\sim 10^9$ M$_{\\sun}$ black holes already observed to be present in quasars at $z\\approx6$ \\citep{Fan03,Lietal07}.  Secular evolution of galaxies, particularly in the absence of major merger events, can drive fuel down to the center as well. Global bar instabilities are thought to be efficient at transporting angular momentum, allowing material in the disk to move toward the center of the galaxy \\citep{Robertsetal79,Simkinetal80,Noguchi88,barswinbars,KormKenn04,RegTeu04}. In the ``bars within bars'' scenario of \\citet{barswinbars}, a large scale galactic bar drives gas inward where it forms a self-gravitating disk. As this disk becomes unstable, a smaller secondary bar forms, driving material down to scales of order $10$ pc, at which point other physics can take over and transport the material the rest of the way toward the black hole \\citep[e.g.][]{ShlosRev90}.  Given the complexity of transporting matter from cosmological scales all the way down to the circumnuclear region of a galaxy, and ultimately to the vicinity of a SMBH, it is clear that a thorough understanding of the relationship between SMBHs and galaxy formation requires modeling of a wide range of scales. Recently, cosmological SPH simulations have been combined with smaller scale simulations of a SMBH host galaxy environment to study the effects of merger driven fueling and AGN feedback on supermassive black hole growth and demography \\citep{DiMatteo07, Sijacki07}. Reaching a large dynamic range in N-body+SPH simulations, \\citet{Mayeretal07} have studied galaxy dynamics during mergers resulting in a supermassive black hole binary. Small scale simulations have addressed gas dynamics in sub-galactic scale disks with high resolution \\citep{Fukuda00, Wada01, WadaNorman01,Escala07,WadaNorman07}, finding the development of a turbulent, multi-phase interstellar medium (ISM).  The study we present here follows the evolution of the circumnuclear region in a typical galaxy environment rather than a dramatic, rare event such as a major merger, in a self-consistent cosmological simulation. We have specifically chosen a simulated galaxy that will evolve into a typical $L_*$ galaxy at $z=0$. The simulation follows the galaxy during a phase of its evolution in which it has not undergone any major mergers within several dynamical time scales, in order to follow the development of instabilities in the circumnuclear disk, which may drive the transport of matter and angular momentum from large to small scales. The cosmological simulations use the adaptive mesh refinement (AMR) technique to self-consistently model the gas dynamics in a single galaxy at high resolution (sub-pc resolution in the center of the galaxy). A large dynamic range ($> 10$ million), achievable with AMR technique, allows us to bridge cosmological scales to scales relevant for molecular cloud formation (the birthplace of stars) and AGN fueling. After studying the physics in this basic model galaxy, we can begin to include physical processes that are directly relevant to the problem of SMBH growth in the context of galaxy evolution, such as AGN feedback.  It is a complex task to implement mergers, feedback and secular evolution in large, cosmological simulations all at once. Our approach is to split the problem into pieces to be addressed one at a time, ultimately building a more realistic simulation. While the present paper focuses on the structure of the galaxy, and the development of instabilities which may be responsible for driving matter and angular momentum transport within the galaxy, we plan to explore the time evolution of such transport, specifically the evolution of the accretion rates of mass and angular momentum, in a subsequent paper. The nature of the circumnuclear region and a quantitative analysis of accretion rates each have important consequences for galaxy evolution.  The organization of the paper is as follows. In Section \\ref{sec:sim}, we describe the details of the simulation and the adopted ``zoom-in'' method. In Section \\ref{sec:res}, we describe the features of the highly-resolved galaxy in a single zoom-in episode at $z=4$, including a stability analysis of the disk and the potential role of physics not included in the current simulation. Section \\ref{sec:miss} examines the potential role of physics missing from the simulations. Finally, Section \\ref{sec:disc} contains a summary of our conclusions and their interpretation. ",
        "conclusions": "\\label{sec:disc}  We have used a large-dynamic range cosmological simulation to study gas dynamics in the circumnuclear disk of typical mass galaxy at $z\\approx4$ (evolving into an $L_*$ galaxy at $z=0$), resolving the distribution of matter from megaparsec scales all the way down to sub-pc scales (with $20$ levels of refinement). The simulation reveals a cold, fully molecular, self-gravitating, and turbulent rotationally-supported gas disk, which is globally unstable but locally stable.  The global instability, operating on a range of scales comparable to the size of the system, generates turbulence down to the smallest scale resolved in the simulation. On small scales, the turbulence supports most of the disk against gravitational fragmentation and collapse. The disk, therefore, remains locally stable and reaches a quasi-stationary state. In that state, global instability drives bar-like and spiral-wave-like structures on time scales of the order of $100{,}000$ years at $100\\dim{pc}$ scales (and on proportionally shorter time scales at smaller radii), but on a $500{,}000$ year time-scale the structure of the disk remains quasi-steady.  The disk develops a power-law gas density profile with a well defined slope of $-5/3$ in surface density and $-8/3$ in spherically-averaged volume density. This slope is a natural consequence of the dynamical state of the disk: the turbulence, generated by the global instability, redistributes the angular momentum in the disk on a dynamical time scale, reducing the gradient of the viscous torque. The resulting (quasi-)steady state uniquely determines the gas density profile we find in the simulation and the distribution of the angular momentum with gas mass.  The turbulence in the disk drives the local outward transport of angular momentum and the inward flow of gas toward the supermassive black hole on time scales of at least 10 million years, which are too long to follow in a single zoom-in episode of the simulation. Thus, we only capture a single snapshot in the cosmological life of the supermassive black hole. In follow-up work, we plan to consider several snapshots taken at different cosmological times in order to describe the system on longer time scales.  The dynamical state of the disk that we find in our simulation appears to be consistent with the results of previous simulations of isolated circumnuclear disks in galaxies \\citep{Fukuda00, Wada01, WadaNorman01, Escala07}. A distinctive feature of our approach is that we follow the dynamics of the circumnuclear disk within cosmological simulations. While most of the volume of the simulation evolves little on time scales relevant for the dynamics of the circumnuclear disk, cosmological scales provide realistic boundary conditions for the dynamics on sub-kpc scales. Since we find that the circumnuclear disk rapidly reaches a quasi-stationary state, its evolution is entirely governed by the boundary conditions. For the same reason, the fact that the cosmological simulation that was used for the initial conditions did not resolve the scale of the circumnuclear disk, does not compromise our results: the quasi-stationary state of the disk does not depend on the initial conditions, and so the lack of power on scales below about $100\\dim{pc}$ in the initial cosmological simulation is not important.  The adopted approach of this work is comparable to the recent study by \\citet{Mayeretal07}, who used a cosmological simulation to model the coalescence of two supermassive black holes. While the detailed treatment of gas physics and the scientific questions answered by the two studies are different, many of our results are consistent with those of \\citet{Mayeretal07}. For example, they find a similar slope for the spherically averaged density profile, although they do not elaborate on the physical mechanisms for the formation of this profile.  The local stability of the disk, supported by highly supersonic turbulence, implies that the disk is capable of continuously feeding a central black hole, uninterrupted by catastrophic bursts of star formation, which could have consumed the available fuel were the disk locally unstable. This interesting dynamical state of the disk provides a potential solution to the problem of how AGN fueling is maintained by self-gravitating gas disks \\citep[for a related discussion, see e.g.][]{ShlosBeg89,Riceetal05,NayakshinKing07}.  The circumnuclear disk in the simulation extends all the way to the maximally resolved scale of $\\approx0.1\\dim{pc}$, which corresponds to the outer part of the black hole accretion disk. We find no toroidal-like structures on several parsec scales, which are commonly inferred to exist around AGN. A possible reason that the disk does not form an AGN torus on the appropriate scales (if it indeed, should) is the limitation imposed by our implementation of gas cooling in the code. The ART code, like all existing cosmological codes, assumes that the cooling radiation from cosmic gas escapes freely into the IGM. This assumption breaks down at the densities and temperatures reached by the simulation in the inner $10\\dim{pc}$. At this scale the disk becomes optically thick to its own cooling radiation from dust and molecular hydrogen. In that regime the disk may to heat up and acquire a substantial amount of pressure support, which may result in a puffier, more toroidal-like configuration for the inner several parsecs of the disk. A consistent treatment of a putative AGN torus will require simulations that incorporate optically-thick cooling."
    },
    "0711/0711.1966_arXiv.txt": {
        "abstract": "Photometric analysis of $BVR_C$ light curves of newly discovered eclipsing binary GSC~0008-00901 is presented. The orbital period is improved to 0.28948(11) days. Photometric parameters are determined, as well. The analysis yielded to conclusion that system is an over-contact binary of W UMa type with components not in thermal contact. The light curves from 2005 show the presence of a spot on the surface of one of the components, while light curves from 2006 are not affected by maculation. ",
        "introduction": "The variability of GCS 00008-00901 ($\\alpha_{2000}$= 00$^{\\mathrm h}$ 13$^{\\mathrm m}$ 22.689$^{\\mathrm s}$, $\\delta_{2000}$=+05$^{\\mathrm o}$ 40' 09.25''), located in the field of eclipsing binary DV Psc, was recently reported by Parimucha (2007). He analyzed $V$ and $R_C$ light curves and showed that the system is most probably contact binary with spotted regions on one of the components. Cutri et al. (2003) reported this star as suspected for variability during their 2MASS project. ",
        "conclusions": "Our light curve analysis of $BVR_C$ photometric observations of recently discovered eclipsing binary GSC 00008-00901, which is located in the field of another eclipsing binary DV~Psc showed, that this system is an overcontact binary of W~UMa type. Considering of its spectral type and orbital period, we concluded that the primary minimum is occultation (primary component is eclipsed by secondary one). So, this system belongs to an W-type of W UMa binaries. The temperature difference between components is about 350K.  In spite of problems with quality of the observations we showed that light curves from the both seasons 2005 and 2006 are different. The light curves from 2005 show different brightness in the both maxima, while 2006 light curves are symmetric. Asymmetry of the light curves from 2005 could be explained by the spot on one of the components, however, some conclusions on spot's temperature and/or its location (on primary or secondary component) are uncertain.  \\begin{figure} \\centering \\includegraphics[width=0.49\\textwidth]{star.eps} \\caption{3D model of GSC 00008-00901 with a cool spot on the secondary component} \\end{figure}"
    },
    "0711/0711.3963_arXiv.txt": {
        "abstract": " ",
        "introduction": "``Superhomogeneous\" \\cite{glasslike} or ``hyperuniform\" \\cite{To03} point patterns in $d$-dimensional Euclidean space $\\mathbb{R}^d$ are defined to be those in which infinite wavelength density fluctuations vanish. In other words, the structure factor (or power spectrum) $S({\\bf k})$ of the number density field at wave vector $\\bf k$ has the following behavior: \\begin{equation} \\lim_{\\bk \\rightarrow 0} S(\\bk) = 0. \\label{Equation4} \\end{equation} This defining characteristic of superhomogeneity (or hyperuniformity) is tantamount to saying that the usual mean-square particle-number fluctuations increases less rapidly than  $R^{d}$ for large $R$, where $R$ denotes the linear size of an observation window in $\\mathbb{R}^d$ \\cite{glasslike,To03}. Indeed, the magnitude of such local density fluctuations have been suggested as  a possible ``order metric\" to quantify the degree of order (disorder) of an arbitrary  point pattern \\cite{To03}. Any superhomogeneous point pattern can be seen as a typical configuration of a particular type of  ``critical\" point in that the {\\it direct correlation} function (defined through the Ornstein-Zernike relation) is long-ranged while the pair correlation function is short-ranged \\cite{To03}. Such remarkable behavior is diametric to that seen in usual thermal critical points in which the inverse is true, i.e., the pair correlation is long-ranged and the direct correlation function is short-ranged.   Although it is clear that any periodic point pattern is superhomogeneous, it is less obvious that statistically translationally and even rotationally invariant random point patterns in $\\mathbb{R}^d$ can have this property. We now know of a variety of intriguing translationally and rotationally invariant random point patterns that are superhomogeneous, including the ground state of liquid $^4\\mbox{He}$~\\cite{Fe54,Fe56,Re67}, maximally random jammed hard-sphere packings~\\cite{Do05}, certain one-component plasmas~\\cite{OCP-review,lebo_and_all,modOCP}, the matter distribution in the Universe~\\cite{glasslike, lebo_and_all}, and certain aperiodic tilings \\cite{lebo_and_all,To03,Ha06}.  An interesting application of superhomogeneous point patterns in cosmology is in the preparation of initial conditions for gravitational $N$-body simulations \\cite{book,lebo_and_all,modOCP,Ha06}. Superhomogeneous distributions also appear in cosmology in the context of the determination of bounds, first derived by Zeldovich \\cite{zeldovich-k4}, on the mass fluctuations at large scales generated by causal mechanisms (i.e. with physics respecting the causal constraints of cosmological models). Indeed, we note that in this context a simplified form of the analysis we develop here of the small-$k$ behavior of the structure factor is often used (see e.g. Ref.~\\cite{peebles}).  It is desirable to develop both theoretical and computational methods to generate a wide class of superhomogeneous random point patterns. Recently, a collective coordinate approach \\cite{Uc04,Uc06} has been employed to numerically generate translationally invariant superhomogeneous point processes. This procedure enables one to produce point patterns that completely suppress density fluctuations of modes for a positive range of wavenumbers around the origin.  In Ref.~\\cite{gabrielli_etal_04} an algorithm for generating discrete processes in one dimension with superhomogeneous mass fluctuations has been given (see also Ref.~\\cite{fratzl_etal1991}). An analytical methodology to relate superhomogeneous point processes and Voronoi tilings of space has recently been proposed and studied in Ref.~\\cite{Ga04}.  In this paper, we study the construction of superhomogeneous point patterns starting from generic tilings of Euclidean space $\\mathbb{R}^d$ with equal volume tiles. We show how to explicitly generate such point processes in which the structure factor for small wavenumbers has the power-law form $S(k) \\sim k^\\gamma$ for positive $\\gamma$, where $k\\equiv |{\\bf k}|$ is the wavenumber.  The constructions illustrate the very specific properties of these superhomogeneous point patterns in which the exponent $\\gamma$, characterizing the long wavelength fluctuation in $k$ space, is related to the detailed arrangement of the points on small scales. Our study also shows how the exponents of the small-$k$ behavior of the structure factor for these point processes encode properties of the tilings, and could thus possibly be used as a method for classifying them. In a related article by two of us \\cite{ag-mj-cloudprocesses} the two-point correlation properties of point processes generated by replacing each particle, in a point process with known two point properties, by a ``cloud''  of particles  are derived\\footnote{Results in this case are derived in \\cite{ag-mj-cloudprocesses} under the assumption that the stochastic process describing the generation of the ``clouds'' and the initial point process are independent. In the algorithm discussed here this is not the case, as the points are ascribed to each tile in a way which depends, in general, on the tile. For the particular case of a Bravais lattice tiling, however, both calculations are valid because of the equivalence of all tiles/points in such a  lattice. Indeed, in this case, the different general formulae derived in the present article and \\cite{ag-mj-cloudprocesses}  give the same result.}. One of us \\cite{Uc06} numerically generated disordered point distributions within a cubical box under periodic boundary conditions with $\\gamma > 4$. However, to our knowledge, prior to this paper and \\cite{ag-mj-cloudprocesses}, explicit analytical constructions of point processes with $\\gamma > 4$ have not been given previously in the literature.  It is instructive to recall qualitatively why tilings are a natural starting point for the construction of superhomogeneous point processes. A {\\it tiling} or {\\it tessellation} is a partition of Euclidean space $\\mathbb{R}^d$  into closed regions whose interiors are disjoint regions \\cite{To02}. Let us suppose we have a tiling of space by tiles which are (i) of equal volume $\\|T\\|$, and (ii) bounded, with maximal length  $\\Lambda$ in any direction. Let us now place one point in each tile and consider the number fluctuations in the point process so generated. If $N(R)$ is the number of points in a sphere of radius $R$, and of volume $V(R)$, it is simple to see that \\begin{equation} \\frac{V(R-\\Lambda)}{\\|T\\|} \\leq N(R) \\leq \\frac{V(R+\\Lambda)}{\\|T\\|}\\,. \\end{equation} The lower bound is the minimal number of tiles which can overlap the sphere of radius $R-\\Lambda$ [and all such tiles must contribute a point to $N(R)$], the upper bound is the maximal number of tiles which are fully enclosed in the sphere of radius $R+\\Lambda$ [and only such tiles can contribute to $N(R)$]. For $R \\rightarrow \\infty$ we have therefore \\begin{equation} |\\Delta N(R)| \\leq c R^{d-1} \\label{lim} \\end{equation} where $c$ is a constant, and $\\Delta N(R)= N(R) - \\overline{N}(R)$ with $\\overline{N}(R)=V(R)/\\|T\\|$. Averaging over configurations (or randomly placed centers for the spheres)  one anticipates that {\\em the slowest possible scaling of number fluctuations is} \\begin{equation} \\langle \\Delta N^2(R) \\rangle \\propto R^{d-1} \\end{equation} where $\\langle...\\rangle$ denotes the ensemble (or volume) average. This behavior of the variance, proportional to the surface, is a characteristic of superhomogeneous point processes. If there is appropriate long-range correlation in the tiling at arbitrarily large scale, the fluctuations could however, in principle, add coherently to give the more rapid growth up to \\begin{equation} \\langle \\Delta N^2(R) \\rangle \\propto R^{2(d-1)}\\,, \\end{equation} which corresponds to the limit equality in Eq.~(\\ref{lim}). While in $d=1$ this still corresponds to surface fluctuations\\footnote{Note that the case of $d=1$ is rather trivial given our assumptions: the only equal volume tiling is the lattice.}, for any $d \\geq 2$ it implies only the limiting small-$k$ behavior $S(k) \\propto k^\\gamma$ with $\\gamma \\ge -d+2$, which means that the point processes are not necessarily superhomogeneous for $d\\ge 2$. We will recover this result below, with the only difference that the bound is found to be  $\\gamma > -d+4$, which implies that even long-range correlations between tiles give superhomogeneous processes for $d\\leq 4$. The difference between this result and our naive estimate is simply due to the fact that below we constrain the particles to lie at the center of mass, rather than placing them randomly. In fact we will show here that by assigning more than one point in an appropriately constrained manner to each tile, we can  increase these bounds on the exponents without limit, and realize superhomogeneous processes with an arbitrary positive exponent, in any dimension. ",
        "conclusions": "\\label{conclusions}  We have studied a class of point processes generated by placing one or more points in each tile of a tiling of $\\mathbb{R}^d$. We have assumed that the tiles have equal volume as this is expected to lead to the suppression of fluctuations at large scales characteristic of superhomogeneous point processes. We have shown explicitly that one can build superhomogeneous point processes with an arbitrarily large exponent characterizing the small-$k$ behavior of the structure factor, and we have presented various examples. To our knowledge exact constructions of such point processes for the case $\\gamma > 4$ have not previously been given in the literature.  We have shown how in these algorithms the exponent $\\gamma$ depends (i) on the arrangements of the points ascribed to the tiles in the algorithm, and (ii) on the correlation properties of the shapes and orientations of the tiles. For the specific case of regular lattice tilings a non-trivial contribution to the small-$k$ behavior of $S(\\bk)$ arises only from the former, with the coefficients in the small-$k$ expansion depending explicitly only on the {\\it variance} of the (tensorial) mass moments of the points assigned to each cell. By arranging a sufficient number of points in a way which makes this variance zero for the first $n$ moments, one can obtain $\\gamma>2n$. In the case of an irregular tiling --- we have considered the example of pinwheel and quaquaversal tilings --- an identical arrangement of the points in each tile can be sufficient to produce continuous SF and translational and rotational invariant superhomogeneous point processes. The exponent $\\gamma$ then encodes information about the correlation properties of the shapes and orientations of the tiles. If these are short-range correlated, one obtains $\\gamma=4$ placing a single point at the center of mass of each tile, and $\\gamma=2(n+1)$ if one places a number of points with all mass moments up to the $n$-th equal to that of the elementary tile. If there is, on the hand, long-range correlation in the shapes and orientations of tiles, the exponent obtained will be modified in a way which depends on the nature of this correlation.  Our results shed light on the meaning of the exponent $\\gamma > 0$ characterizing a superhomogeneous point process. Up to the value $\\gamma=4$ previous explicit constructions of discrete processes (see, e.g., \\cite{andrea, gabrielli_etal_04}) have shown that the increase of $\\gamma$ can be associated with a suppression of fluctuations at large scales. Here we have seen that values $\\gamma>4$ correspond indeed to an increased order in the arrangement of the points, but now {\\it at small scales}: it is by changing how points are arranged within each tile, i.e., below a finite length scale (but subject always to the global constraints on fluctuations imposed by the tiling), that we can increase the exponent. Thus to ``undo'' the order represented by an exponent $\\gamma > 4$ with respect to a system with $\\gamma=4$ requires only the rearrangement of the system at small scales, while to ``undo'' that in a system with $\\gamma \\leq 4$ require a coherent rearrangement of points on arbitrarily large scales (i.e. on scales inverse to the wavenumber range in which the exponent is measured).  We have mentioned that our results are relevant in cosmology. Firstly the analysis given here makes more rigorous certain heuristic arguments used in this context regarding ``causal constraints'' on the generation of fluctuations from a uniform background \\cite{zeldovich-k4,peebles}. An algorithm like that described here, for the case of a single point placed at the center of mass of each tile, has been considered \\cite{peebles} as a toy model for the generation of fluctuations starting from an exactly uniform mass density, by a physical process which conserves mass and momentum {\\it locally}. The result $\\gamma=4$ is obtained by assuming that space is divided into finite cells whose positions are uncorrelated. The latter assumption is in fact not consistent: the division of space into equal volume cells implies that their positions are necessarily correlated. Our more rigorous analysis shows that this exponent $\\gamma=4$ does result generically, however, if the shapes and orientations of these cells (i.e. tiles) are short-range correlated, i.e., have integrable correlation functions.  Secondly, the generation of very uniform point processes is of relevance to the generation of initial conditions for numerical simulations of structure formation in the universe. In this context, to represent a given set of initial conditions, one must perturb appropriately  (see \\cite{discreteness1_mjbm} for a detailed discussion) a point distribution representing as well as possible the uniform (unperturbed) universe. To understand the effects coming from this chosen point distribution (which are non-physical) it is desirable to have different algorithms which can generate such configurations. It is for this reason that a special case of the algorithm we have studied here has been built explicitly and studied numerically in this context \\cite{Ha06}.  The analytical results we have given here complement these studies and give further algorithms for producing even more uniform point processes which may be useful in this context.  We note again in this respect that explicit algorithms for producing $\\gamma > 4$ have not previously been given. Such distributions in themselves provide interesting initial conditions (without any perturbation) for gravitational clustering, which have not previously been studied.  We conclude with some further remarks on our results and some other directions for further work:  \\begin{itemize}  \\item In our constructions of point processes we have always constrained the center of mass of the points in each tile to coincide with that of the tile. We have done so because our goal here has been to generate point processes which are as uniform as possible. It is a simple exercise to redo our calculation leading to Eqs.~(\\ref{Sk-calI}) and (\\ref{defn-calI}) when this constraint is relaxed, i.e., allowing the center of mass of the particles (or particle) in each tile to be displaced randomly from that of the tile. The result is that the leading term in Eqs.~(\\ref{Sk-calI}) is now at order $k^2$ rather than $k^4$. For short-range correlated tilings the leading behavior of the SF will then be proportional to $k^2$. This result will be valid if the displacements of the center of mass of the particles within a tile with respect to the center of mass of the associated tiles have a finite variance.  On the other hand, if the variance of these displacements diverges, the small-$k$ behavior of the SF is given by $S(\\bk) \\propto k^\\gamma$ where $0<\\gamma <2$, the value of the exponent depending on the precise behavior displacements PDF for large arguments. A detailed calculation of these cases for a randomly perturbed lattice can be found in Ref.~\\cite{andrea, book}.  \\item While we have shown analytically the existence of point processes with arbitrarily large exponents $\\gamma$, we have not done so for a case which is statistically translation and rotation invariant. In the latter case our results for the exponent are expressed in terms of the small-$k$ behavior of the ${\\tilde g}^{n,m} (\\bk)$ which encode, as we have explained, information about the correlation properties of the shapes and orientations of the tiles. For the one such example which has been numerically studied (in Ref.~\\cite{Ha06}, the quaquaversal tiling with a single point at the center of mass) the result indicates an asymptotic behavior close to (but, as noted in Ref.~\\cite{white_WDM}, slightly different to) that which would arise from a purely short-range correlation of the orientations of the tiles. Further numerical and analytical study of these points processes would clearly be of interest, in particular of the simpler pinwheel tiling.  \\item We have made in our derivations here an assumption of analyticity at $\\bk=0$ of the window function of the tiles, which corresponds to all moments of their mass distribution being finite. This is certainly valid if the tiles are of finite extent. It may, however, include other cases which might be of interest, e.g., in $d>1$ one may envisage that there is a non-trivial distribution of the shapes of the equal volume tiles, in which the extent of a tile is not limited. One could also consider relaxing the assumption that the volumes of the tiles are strictly equal, admitting a distribution of volumes with specified correlation properties. In analogy with what has been found in certain algorithms for $\\gamma \\leq 4$ \\cite{andrea, gabrielli_etal_04}, one would expect that such modifications would allow the generation of point processes with leading non-analytic behavior, and any value of the exponent $\\gamma$. Indeed one would expect non-analytic exponents to be related either to the divergence of moments of such a distribution of extent or volume or to the presence of long-range tile-tile correlations.   \\item While all our explicit examples have employed tilings which are congruent, our results for the small-$k$ behavior of the SF $S(\\bk)$ all apply only on the much weaker assumption of equal volume of the tiles. Thus for example we can apply these results to any tiling generated by a deformation of tiles which leaves their volume fixed, which could encompass a large range of systems of physical interest (cells, foams, etc.). We recall in this respect, as remarked above, that a point placed randomly in each cell, rather than at the center of mass, leads to the restoration of the order $k^2$ term in the expression given in Eqs.~(\\ref{Sk-calI}).  \\item It is useful to briefly remark on the physical realizability of superhomogeneous point distributions with arbitrary positive but bounded values of $\\gamma$. Such distributions are disordered to some degree and, although they are unusual, can be physically constructed. For example, the maximally random jammed (MRJ) state in three dimensions \\cite{To00} is a special disordered sphere packing that can be regarded to be a prototypical glass because it is perfectly rigid and yet is maximally disordered. It is a superhomogeneous point distribution characterized by an exponent $\\gamma=1$ \\cite{Do05}, but it is inherently a system out of equilibrium. While we have examples like the equilibrium one-component plasma that has $\\gamma=2$, can one devise equlibrium superhomogeneous point distributions in which $\\gamma$ is arbitrarily large? The answer is apparently in the affirmative but it requires more than just pair interactions, namely, two-, three- and four-body interactions as shown in Ref. \\cite{Uc04}.  \\end{itemize}   MJ and AG thank ST for hospitality at Princeton University during a visit in May 2006. MJ thanks B. Jancovici and J. Lebowitz for useful discussions. ST gratefully acknowledges the support of the Office of Basic Energy Sciences, DOE, under Grant No. DE-FG02-04ER46108.  \\appendix"
    },
    "0711/0711.4772_arXiv.txt": {
        "abstract": "Hot spots residing on the surface of an accretion disc have been considered as a model of short-term variability of active galactic nuclei. In this paper we apply the theory of random point processes to model the observed signal from an ensemble of randomly generated spots. The influence of general relativistic effects near a black hole is taken into account and it is shown that typical features of power spectral density can be reproduced. Connection among spots is also discussed in terms of Hawkes' process, which produces more power at low frequencies. We derive a semi-analytical way to approximate the resulting power-spectral density. ",
        "introduction": "Radiation from accreting black holes varies on different timescales \\citep{gas06}. In X-rays, the observed light-curve, $f\\,\\equiv\\,f(t)$, is a complicated noisy curve that can be represented by a broad-band power spectrum \\citep{2002MNRAS.332..231U}. It has been proposed \\citep{Abramowicz:1991:A&A,1992AIPC..254..251W} that `hot spots' are a possible contributor to this variability. These spots are supposed to occur on the surface of an accretion disc following its irradiation by coronal flares \\citep{Galeev:1979:APJ,2001MNRAS.328..958M,2004A&A...420....1C}. A model light-curve can be constructed as a sum of contributions from many point-like sources that are orbiting above an underlying accretion disc. The observed signal is modulated by relativistic effects as photons propagate towards a distant observer.  In order to characterise the light curves we need to introduce some appropriate estimator of the source variability. In a mathematical sense, one applies a functional: $f\\rightarrow \\mathcal{S}\\left[f\\right]$, where $\\mathcal{S}\\left[.\\right]$ is a map from functions defined on $\\mathbb{R}$ to functions on $\\mathbb{R}^k$ ($k\\geq 0$). The variability estimator can be a single number (for example, the mean flux or the `rms' characteristic), or function of one variable (power spectrum density or probability distribution) or of many variables (poly-spectra, rms--flux relation, etc). A signal of such a spotted accretion disc should be intrinsically stochastical. Hence, the variability estimator $\\mathcal{S}\\left[f\\right]$, derived from a piece of the light-curve, is a random value, too.  Various schemes have been proposed in which spots are mutually interconnected in some way \\citep{Poutanen:1999:MNRAS, Zycki:2005:MNRAS}. We want to investigate this type of models within a common mathematical basis. Other authors have developed different approaches  to the problem \\citep{1994PASJ...46...97M,1997MNRAS.292..679L}. ",
        "conclusions": "We have studied the properties of power spectral density within the model of accretion disc variability driven by orbiting spots. The origin and evolution of spots were described in terms of Poissonian and Hawkes' processes. The latter belongs to a category of avalanche models. We developed an analytical approximation of PSD and compared it with our numerical results from light-curve simulations. In this way we were able to demonstrate the precision of formulae (\\ref{PoissPSD}) and (\\ref{HawkPSD}). The analytical approximation evaluates very fast and provides the main trend of the PSD shape while avoiding the noisy form of the numerically simulated spectra. Our approach allows us to investigate the resulting PSD as a function of the assumed type of process, which describes creation of parent spots and the subsequent cascades of daughter spots. In particular, we can investigate the predicted PSD slope at different frequency ranges and we can locate the break frequency depending on the model parameters.  The resulting PSD can be approximated by a broken power-law. For every stationary process the quantities $S(0)$ and $\\int_0^\\infty S(\\omega){\\rm d}\\omega$ are finite. Therefore, the function $S(\\omega)$ flattens ($S(\\omega)\\approx \\omega^0$) near $\\omega=0$ and it must decrease faster than $1/\\omega$ at high frequencies. Power-spectra generated by the spot model behave in this way. The low-frequency limit is a constant, whereas the-high frequency behaviour depends mainly on the shape of the spot emission profile, $I(t,\\,\\mbox{\\boldmath $\\xi$}_j)$. In our calculations the emissivity was a decaying exponential and the slope was equal to $-2$ at high frequencies. The most interesting part of the spectra in between those two limits is influenced by both the emissivity profile and the underlying process.  \\ack We thank the Workshop participants for helpful comments and the Academy of Sciences for financial support (grant GAAV 300030510)."
    },
    "0711/0711.1361_arXiv.txt": {
        "abstract": "We derive a physical model for the observed relations between star formation rate (SFR) and molecular line (CO and HCN) emission in galaxies, and show how these observed relations are reflective of the underlying star formation law. We do this by combining 3D non-LTE radiative transfer calculations with hydrodynamic simulations of isolated disk galaxies and galaxy mergers. We demonstrate that the observed SFR-molecular line relations are driven by the relationship between molecular line emission and gas density, and anchored by the index of the underlying Schmidt law controlling the SFR in the galaxy. Lines with low critical densities (e.g. CO J=1-0) are typically thermalized and trace the gas density faithfully. In these cases, the SFR will be related to line luminosity with an index similar to the Schmidt law index. Lines with high critical densities greater than the mean density of most of the emitting clouds in a galaxy (e.g. CO J=3-2, HCN J=1-0) will have only a small amount of thermalized gas, and consequently a superlinear relationship between molecular line luminosity (\\lmol) and mean gas density ($\\bar{n}$). This results in a SFR-line luminosity index less than the Schmidt index for high critical density tracers. One observational consequence of this is a significant redistribution of light from the small pockets of dense, thermalized gas to diffuse gas along the line of sight, and prodigious emission from subthermally excited gas. At the highest star formation rates, the SFR-\\lmol \\ slope tends to the Schmidt index, regardless of the molecular transition.  The fundamental relation is the Kennicutt-Schmidt law, rather than the relation between SFR and molecular line luminosity. Our model for SFR-molecular line relations quantitatively reproduces the slopes of the observed SFR-CO (J=1-0), CO (J=3-2) and HCN (J=1-0) relations when a Schmidt law with index of $\\sim$1.5 describes the SFR.  We use these results to make imminently testable predictions for the SFR-molecular line relations of unobserved transitions. ",
        "introduction": "The rate at which stars form in galaxies has historically been parameterized in terms of ``laws'' relating the star formation rate (SFR) to the density of available gas. \\citet{sch59} originally proposed a power-law form for the SFR such that SFR $\\propto$ $\\rho^N$ (hereafter, referred to as a Schmidt Law).  Observed SFR relations typically come in two flavors. The first, relating surface SFR density to surface gas density takes the form: \\begin{equation} \\label{eq:ken} \\Sigma_{\\rm SFR} \\propto \\Sigma_{\\rm gas}^N \\, . \\end{equation} Observations of local galaxies have constrained the surface density SFR index, $N$, to roughly $N$=1.4$\\pm$0.15 \\citep[e.g.][and references therein]{ken98a,ken98b,ken07}.  The second varietal of SFR indicators relates the SFR to the mass of molecular gas above a given volumetric density. For example, local galaxy surveys have shown a relation exists between the SFR and molecular gas such that the SFR (as traced by the infrared luminosity - hereafter \\lir) is proportional to the $^{12}$CO (J=1-0) luminosity to the 1.4-1.6 power \\citep[][and references therein]{san91,sanmir96}. Because the J=1-0 transition of $^{12}$CO (hereafter CO) can be excited at relatively low densities ($n \\sim$10$^2$-10$^{3}$ \\cmthree), and lies a modest $\\sim$5 K above ground, it serves as a valuable tracer of total molecular \\htwo \\ gas content down to relatively low densities. These observed relations have been broadly interpreted as an increasing star formation efficiency (SFE; the SFR divided by \\mhtwo) as a function of molecular gas mass \\citep{gao04a,gao04b,san91}.   Recent observations of HCN have suggested that a perhaps more fundamental volumetric SFR relation for galaxies exists in terms of the {\\it dense} molecular gas. Because HCN (J=1-0) has a relatively high critical density ($n_{\\rm crit} \\sim$10$^5$ \\cmthree) compared to that of CO (J=1-0), HCN is typically only thermalized in the dense cores of molecular clouds. Thus, in the limit that the bulk of HCN luminosity originates from thermalized gas, HCN serves as a good tracer of the dense molecular gas that is actively involved in the star formation process. This is in contrast to CO (J=1-0) which tends to emit from both dense cores as well as diffuse molecular filaments and cloud atmospheres.  Pioneering millimeter-wave observations by \\citet{gao04a,gao04b} uncovered a tight linear correlation between the infrared luminosity and HCN luminosity in a sample of local galaxies, ranging from quiescent spirals to luminous and ultraluminous infrared galaxies (LIRGs and ULIRGs, respectively). This relationship was found by \\citet{wu05} to hold in individual star forming cores in the Milky Way as well.  Observations of a roughly linear correlation between \\lir \\ and CO (J=3-2) emission (with critical density $n_{\\rm crit} \\sim$10$^4$ \\cmthree) in a similar sample of galaxies further corroborated this result, and provided evidence against HCN-related chemistry driving the observed \\citet{gao04a,gao04b} relations \\citep{nar05,yao03}.  Despite a plethora of observations, a consensus physical interpretation of the seemingly disparate \\lir-\\lhcnjone, \\lir-\\lcojone \\ and \\lir-\\lcojthree \\ relations has yet to be borne out of the literature. \\citet{gao04a,gao04b} suggest that the linear relationship between \\lir-\\lhcnjone \\ reflects a scenario in which the SFR in galaxies is directly proportional to the amount of dense molecular gas, and similarly, a constant star formation efficiency in terms of dense molecular gas. In this picture, LIRGs and ULIRGs simply have a higher fraction of their molecular gas in a dense gas phase which is reflected in the linear \\lir-\\lhcnjone \\ slope\\footnote{Throughout this paper we use the term slope interchangeably with 'index' (as in the exponent in a given SFR relation - e.g. Equation~\\ref{eq:ken}). We do this as we are considering all relations in log-log space.}  and non-linear \\lir-\\lcojone \\ slope. \\citet{nar05} suggest a similar interpretation for their observed relation of \\lir $\\propto$ \\lcojthree$^{0.92}$.  \\citet{wu05} extended these interpretations to include the observed linear \\lir-\\lhcnjone \\ relation within the Galaxy. These observations suggested that dense molecular cores may represent a fundamental unit of star formation. In this case, the only difference between a ULIRG like Arp 220 and a Galactic star forming region is the number of individual star forming units present. This leads to a natural linear relationship between the SFR and amount of dense molecular gas.  From the theoretical side, \\citet{krumthom07} made headway towards understanding the observed relations by providing a model which quantitatively reproduces the observed \\lir-\\lcojone \\ and \\lir-\\lhcnjone \\ relations for star forming clouds. By combining 1D non-local thermodynamic equilibrium (LTE) radiative transfer calculations with models of giant molecular clouds (GMCs), these authors found that the relationship between the SFR and molecular line luminosity in star forming clouds depends on how the critical density of the molecule compares to the mean density of the observed source. Lines with critical densities which are typically below the mean density in a galaxy (e.g. CO J=1-0) probe the total molecular gas commonly from galaxy to galaxy. In this scenario, the SFR-molecular line exponent can be represented by the index that relates SFR to total volumetric gas density (i.e. the Schmidt index).  Conversely, observations of molecules which typically trace densities well above the mean density of the galaxy (e.g. HCN J=1-0) trace similar conditions from galaxy to galaxy - i.e. the peaks in the density spectrum. In these cases, the molecular line luminosity rises faster than linearly with increasing gas density and the corresponding relation between SFR and line luminosity is near linear. A key direction forward beyond these models is understanding the effects of emitting GMCs on a galaxy-wide scale, with the potential effects of molecular line radiative transfer.  In this paper, we build on the body of observational and theoretical work by providing a physical model for the origin of the observed molecular SFR indicators on galaxy-wide scales, and relating them to observed relations. We additionally make model-distinguishing testable predictions for how the SFR in galaxies relates to as yet unobserved transitions in CO and HCN. We do this by combining 3D non-LTE molecular line radiative transfer codes \\citep{nar06a,nar06b,nar07a,nar07b} with smoothed particle hydrodynamic (SPH) simulations of both isolated star forming galaxies and galaxy mergers. Our methodology includes the effects of both collisional and radiative molecular excitation and de-excitation, a multi-phase ISM, and star formation \\citep{spr05a}. We additionally include a methodology for black hole growth and the winds associated with both black hole growth and star formation, though note that these processes play a minimal role on the SFR-\\lmol \\ relation in our simulations.  The paper is outlined as follows. In \\S~\\ref{section:methods} we describe our hydrodynamic and radiative transfer simulations. In \\S~\\ref{section:origins} we quantitatively describe the origin for the observed SFR-CO and SFR-HCN slopes and follow by making testable predictions for submillimeter-wave telescopes in \\S~\\ref{section:predictions}. In \\S~\\ref{section:observations} we compare our results to the body of observational work in this field, and in \\S~\\ref{section:othertheory}, relate these simulations to other models and interpretations for molecular SFR indicators. In \\S~\\ref{section:conclusions}, we conclude with a summary.  Throughout the work we assume a $\\Lambda$CDM cosmology with $h$=0.7, $\\Omega_\\Lambda$=0.7, $\\Omega_M$=0.3. ",
        "conclusions": "\\label{section:conclusions}  We have utilized a combination of 3D non-LTE radiative transfer calculations with hydrodynamic simulations of isolated disk galaxies and galaxy mergers to derive a physical model for the observed SFR-molecular line relations. We specifically focus on the examples of the SFR-CO (J=1-0), CO (J=3-2) and HCN (J=1-0) relations as they are the best constrained observationally, and show that our model quantitatively reproduces the observed relations when a Schmidt index of $\\sim$1.5 is assumed.   While the linear relationship between SFR and high critical density tracers such as HCN (J=1-0) and CO (J=3-2) in galaxies have been standardly interpreted as a fundamental SFR law owing to dense gas being the formation site of massive stars, our models suggest that this is not the entire story.  The linear relations of SFR and HCN (J=1-0) and CO (J=3-2) arise as a consequence of a superlinear relation between line luminosity and mean gas density in galaxies. This owes to small amounts of thermalized gas in high critical density tracers. The fundamental relation is instead the underlying Schmidt law which sets the way in which observed transitions trace the molecular gas.   Our model makes the prediction that for CO lines with J$_{\\rm upper} >$3 and HCN lines with J$_{\\rm upper} >$2, the SFR-line luminosity relationships will be sublinear (Figure~\\ref{figure:mvtransition}). These predictions can directly be tested with existing submillimeter-wave technology, as well as with ALMA.  Our models additionally provide specific interpretation regarding the existing observed SFR-molecular line relations:  \\begin{enumerate}  \\item The slope in a given SFR-molecular line luminosity relation is dependent on both the underlying Schmidt law controlling the SFR for the galaxy and the relationship between molecular line luminosity and density of emitting gas (the details of which are outlined in \\S~\\ref{section:origins}). When line luminosity traces gas density linearly, the resultant SFR-line luminosity index is similar to the Schmidt law index. In cases where the line luminosity increases with gas density superlinearly, the SFR is related to line luminosity with an index less than the Schmidt index. The relationship between line luminosity and gas density depends on how the critical density of the line compares with the mean density of the bulk of the emitting clouds. This is similar to the explanations posited by \\citet{krumthom07} who utilized models for turbulence-regulated GMCs. This directly affects the observed relations between SFR and CO (J=1-0), CO (J=3-2) and HCN (J=1-0) in the following way:  \\begin{enumerate} \\item Owing to its low critical density, the CO (J=1-0) line is roughly thermalized throughout most regions of the galaxies in our simulation sample. This results in a linear rise in CO (J=1-0) luminosity with increasing gas density, and a consequent SFR-CO (J=1-0) relationship with index similar to the Schmidt index. For the case of a Schmidt index of $\\sim$1.5, our simulations reproduce the observed relation between SFR and CO (J=1-0) luminosity.   \\item The critical densities of CO (J=3=2) and HCN (J=1-0) are much higher than the mean density of the bulk of the clouds in our simulated galaxies. Because of this, the relationship between \\lmol \\ and $\\bar{n}$ is superlinear, causing an SFR-\\lmol \\ relation which has index less than the Schmidt index.  An important observational consequence of this is that significant amounts of emission can arise from subthermally excited diffuse gas in the vicinity of denser regions.  For the case of a Schmidt index of $\\sim$1.5, the observed relations between SFR and HCN (J=1-0) and CO (J=3-2) emission are recovered.   \\end{enumerate}      \\item The emission processes driving the line luminosity-density relations (and consequently SFR-line luminosity relations) are variable such that some galaxies exhibit mostly thermalized gas for high critical density tracers whereas others are largely subthermally excited. Generally, the galaxies with higher SFRs have more of their gas thermalized which may drive the upturn in the \\lir-HCN (J=1-0) relationship for extremely high luminosity sources observed by \\citet{gao07}. This result is additionally recovered by models for GMCs by \\citet{krumthom07}.  \\item The physical basis for these models has a similar reasoning as those posited by \\citet{krumthom07} for the SFR-\\lmol \\ relation in turbulence regulated GMCs. In this sense, two completely different modeling techniques arrive at similar physical motivations for the observed SFR-\\lmol \\ relation in galaxies.  \\item We reemphasize that these models can be {\\it directly tested} via observations of high critical density tracers (see Figure~\\ref{figure:mvtransition}). In particular, these models suggest that observations of HCN lines higher than J=1-0 and CO lines higher than CO J=3-2 should bear sublinear relations between SFR and \\lmol.  \\item A natural consequence of this model is that at the highest SFRs, the SFR-\\lmol relation will have an index similar to that of the underlying Schmidt index. This is a generic feature of our models, and happens regardless of the inclusion of AGN.  \\end{enumerate}"
    },
    "0711/0711.3014_arXiv.txt": {
        "abstract": "We study the formation of disk-dominated galaxies in a $\\Lambda$CDM universe.  Their existence is considered to be a challenge for the $\\Lambda$CDM cosmology, because galaxy mergers isotropize stellar disks and trigger angular momentum transport in gas disks, thus fostering the formation of central stellar spheroids. Here, we postulate that the formation of stellar spheroids from gas-rich disks is controlled by two parameters that characterize galaxy mergers, the mass ratio of merging dark matter halos, and the virial velocity of the larger merging halo.  We utilize merger histories generated from realizations of the cosmological density field to calculate the fraction of dark matter halos that have avoided spheroid formation, and compare the derived statistics with the spheroid occupation fractions in surveys of nearby galaxies.  We find, for example, that the survival rate of disk-dominated galaxies in $\\Lambda$CDM is just high enough to explain the observed fractional representation of disk-dominated galaxies in the universe if the only mergers which lead to central spheroid formation are those with mass ratios $M_2/M_1>0.3$ and virial velocities $V_{\\rm vir,1}>55\\textrm{ km s}^{-1}$. We discuss the physical origin of this criterion. ",
        "introduction": "\\label{sec:intro}  The existence of disk-dominated galaxies, with little or no bulge, is frequently cited as a challenge to the $\\Lambda$CDM cosmology \\citep[e.g.,][]{Kautsch:06,Kormendy:07}. Apart from the question of whether or not the theory of galaxy formation in $\\Lambda$CDM can succeed in making these observed rotationally-supported disk galaxies in the first place \\citep{D'Onghia:04, Abadi:03a}, the survival of such disks, once formed, is our focus here.  The mergers that every galaxy-hosting dark matter halo experiences can trigger angular momentum transport in the interstellar medium of the merger remnant. If a substantial amount of gas is transported into the central kiloparsec of the remnant, the gas can fuel a starburst forming a central stellar system---a ``classical'' bulge (i.e., self-gravitating, baryon-dominated stellar system supported by random motions rather than rotation). By contrast, ``pseudobulges'' can also arise in some disk galaxies if they have \\textit{not} experienced recent major mergers, by the secular transport of angular momentum \\citep[e.g., by galactic bars,][]{Jogee:05}.  These pseudobulges (sometimes called ``disky bulges'') are supported more by rotation than random motion, however \\citep[e.g.,][and references therein]{Kormendy:04}.\\footnote{ Pseudobulges are to a larger degree supported by rotation than classical bulges and can have rotational velocity-to-1D velocity dispersion ratios $V_{\\rm max}/\\sigma > 1$.  Their velocity dispersions are smaller than expected from the Faber-Jackson relation.  Pseudobulges tend to have flattened isophotes and surface brightness profiles close to exponential. The classical bulges and pseudobulges can be distinguished from a third class, the boxy or peanut-shaped bulges, which are bars seen edge-on \\citep{Bureau:99,Athanassoula:05}.}  About one third of all disk galaxies in the local universe do not contain bulges or pseudobulges \\citep{Kautsch:06}\\footnote{The classification of galaxies by their bulge content is usually done using simple bulge-disk or concentration models, which may not be accurate in extreme disk-dominated galaxies \\citep[e.g.,][]{Boeker:03}.}  and another third contain only pseudobulges.  \\citet{Allen:06} carried out a S\\'ersic spheroid and exponential disk decomposition on a large sample of galaxies and find that $30\\%$ of exponential disks have small bulge-to-total ratio $B/T<0.2$.  \\citet{Barazza:07} report that $20\\%$ of disk galaxies can be visually classified as bulgeless.\\footnote{ Pure disk galaxies also contain nuclear star clusters \\citep[e.g.,][]{Boeker:02,Walcher:05} which could be products of secular gas transport \\citep{Milosavljevic:04}, but these star clusters are tiny by comparison with the bulges described above.}  For our purposes here, we shall assume that the survival of disk-dominance means that no classical bulge is formed.  Although the precise characteristics of mergers that form classical bulges remain unknown, disk-dominated galaxies must have avoided major mergers during and after formation. Just how large the mass ratio of the merging galactic halos must be in order to induce bulge formation is somewhat uncertain, however. Mergers of similar-mass galaxies have been shown to trigger starbursts and result in elliptical galaxies, by gas dynamical and $N$-body simulations for galaxy halo masses $M \\gtrsim 10^{11} M_\\odot$, merging at relative velocities of the order of the virial velocity \\citep{Mihos:96, Cox:07}. In that case, a bulge forms from the momentum-exchange and compression of gas in merger shocks and the outward angular momentum transport induced by merger torques.  Pure $N$-body simulations of such mergers find that the pre-existing stellar disks are mixed and destroyed \\citep[e.g.,][]{Naab:03}, leaving an elliptical galaxy. For minor mergers, $N$-body and gas dynamical simulations with halo masses $M \\gtrsim 10^{11} M_\\odot $ indicate that disks survive but bulges can also grow \\citep[][see also \\citealt{D'Onghia:06}]{Mihos:94, Cox:07, Eliche-Moral:06}. However pure $N$-body simulations of this process find that these bulges are pseudobulge-like (i.e., partially supported by rotation), rather than classical, for mass ratios $0.1\\lesssim M_2/M_1\\lesssim 0.25$ \\citep{Bournaud:05}.  Bulge formation by the merging of lower mass ($V_{\\rm max} \\lesssim 70\\textrm{ km s}^{-1}$) and gas-rich galaxies has not yet been simulated.    Semianalytic models of galaxy formation \\citep[e.g.,][]{Kauffmann:93b, Baugh:96} assume they can track the morphological type of galaxies by converting a disk component into a spheroidal component in mergers with mass ratios $\\mu\\equiv M_2/M_1$ greater than some threshold. These models successfully reproduce the distribution of earlier morphological types by tuning the critical mass ratio for disk destruction, and adopting a critical bulge-to-total mass ratio that discriminates broadly between disk and elliptical galaxies.  Recent semianalytic models employing merger trees extracted from hydrodynamic $N$-body simulations yield a similar result \\citep{Maller:06}. However, those studies which focus on the survival of disks generally do not address the abundance of {\\it disk-dominated} galaxies. Existing ab initio cosmological simulations have yielded disks with bulges \\citep[e.g.,][]{Abadi:03a}, but currently lack the dynamic range to explore a large enough volume to sample the statistics of galaxy morphology while simultaneously resolving the formation and mergers of individual galaxies.   The purpose of this work is to compare the predicted disk survival probabilities during hierarchical merging in a $\\Lambda$CDM universe with the observed statistics of galaxies at the end of the Hubble sequence.  This comparison is then used to place constraints on the physics of bulge-forming mergers.  In \\S~\\ref{sec:mergers}, we discuss the role of mergers in bulge formation.  In \\S~\\ref{sec:jeans-mass}, we describe the effect of cosmic reionization and Jeans-mass filtering on bulge formation. In \\S~\\ref{sec:trees}, we present a calculation of bulge formation and disk survival probabilities derived from galactic halo merger trees generated from realizations of cosmological density fields.  In \\S~\\ref{sec:results}, we compare observed disk galaxy statistics with these merger tree results and place constraints on the properties of bulge-forming mergers.  In \\S~\\ref{sec:discussion}, we summarize our main conclusions.  Standard cosmological parameters consistent with the {\\it Wilkinson Microwave Anisotropy Probe} \\citep{Spergel:07} are assumed throughout. ",
        "conclusions": "\\label{sec:discussion} In order to explain the observed space density and fraction of disk-dominated galaxies within the $\\Lambda$CDM cosmology, we propose a bulge-forming criterion such that only those halo mergers with mass ratio greater than $\\mu_{\\rm crit} \\sim 0.3$ and halo virial velocity (of the larger halo) above $V_{\\rm vir,crit} \\sim 55 \\textrm{ km s}^{-1}$ formed classical bulges, while other mergers did not. This criterion has some degeneracy between $\\mu_{\\rm crit}$ and $V_{\\rm vir,crit}$, but $V_{\\rm vir,crit}$ cannot be larger than about $65 \\textrm{ km s}^{-1}$ without underproducing the galaxy fraction with bulges, or much smaller than $\\sim 60 \\textrm{ km s}^{-1}$ since Jeans-mass filtering after reionization inhibits such small-mass halos from acquiring and retaining baryons or forming stars.  This bulge-forming criterion also gives a reasonable dimensionless condition, $\\mu_{\\rm crit} \\beta \\mathcal{M}_{\\rm mer,crit} \\sim 1$, for the impact of merger on the gas disk from the point of view of angular momentum transport. The validity of this bulge formation criterion needs to be confirmed by further analytic calculation or hydrodynamical simulations of mergers in the halo mass range $V_{\\rm vir,crit} \\sim 60 \\textrm{ km s}^{-1}$."
    },
    "0711/0711.2823_arXiv.txt": {
        "abstract": "Sterile neutrinos with keV masses can constitute all or part of the cosmological dark matter.  The electroweak-singlet fermions, which are usually introduced to explain the masses of active neutrinos, need not be heavier than the electroweak scale; if one of them has a keV-scale mass, it can be the dark-matter particle, and it can also explain the observed pulsar kicks. The relic sterile neutrinos could be produced by several different mechanisms. If they originate primarily from the Higgs decays at temperatures of the order of 100 GeV, the resulting dark matter is much ``colder'' than the warm dark matter produced in neutrino oscillations. The signature of this form of dark matter is the spectral line from the two-body decay, which can be detected by the X-ray telescopes.  The same X-rays can have other observable manifestations, in particular, though their effects on the formation of the first stars. ",
        "introduction": "Most discoveries in particle physics amount to either a measurement of some parameter related to a known particle, or a detection of some new degrees of freedom,  new particles.  The discovery of the neutrino mass\\cite{review} is both.  Not only is it a measurement of the non-zero mass, but it also implies the existence of some additional, SU(2) singlet fermions, ``right-handed'' neutrinos.  The corresponding particles can be made very heavy even for small masses of the active neutrinos (the seesaw mechanism~\\cite{seesaw}), but they can also be light, in which case they are called sterile neutrinos.  The name {\\em sterile neutrino} was coined by Bruno~Pontecorvo, who hypothesized the existence of the right-handed neutrinos in a seminal paper~\\cite{pontecorvo}, in which he also considered vacuum neutrino oscillations in the laboratory and in astrophysics, the lepton number violation, the neutrinoless double beta decay, some rare processes, such as $\\mu \\rightarrow e \\gamma$, and several other questions that have dominated the neutrino physics for the next four decades.   Most models of the neutrino masses introduce sterile (or right-handed) neutrinos to generate the masses of the ordinary neutrinos via the seesaw mechanism~\\cite{seesaw}.  The seesaw lagrangian \\begin{equation} {\\cal L} = {\\cal L_{\\rm SM}}+\\bar N_{a} \\left(i \\gamma^\\mu \\partial_\\mu \\right ) N_{a} - y_{\\alpha a} H \\,  \\bar L_\\alpha N_{a} - \\frac{M_{a}}{2} \\; \\bar {N}_{a}^c N_{a} + h.c. \\,, \\label{lagrangianM} \\end{equation} where ${\\cal L_{\\rm SM}}$ is the lagrangian of the Standard Model, includes some number $n$  of singlet neutrinos $N_a$ ($a=1,...,n$) with Yukawa couplings $ y_{\\alpha a}$.  Here $H$ is the Higgs doublet and $L_\\alpha$ ($\\alpha=e,\\mu,\\tau$) are the lepton doublets. Theoretical considerations do not constrain the number $n$ of sterile neutrinos. In particular, there is no constraint based on the anomaly cancellation because the sterile fermions do not couple to the gauge fields. The experimental limits exist only for the larger mixing angles~\\cite{sterile_constraints}. To explain the neutrino masses inferred from the atmospheric and solar neutrino experiments, $n=2$ singlets are sufficient~\\cite{2right-handed}, but a greater number is required if the lagrangian (\\ref{lagrangianM}) is to explain the r-process nucleosynthesis~\\cite{r}, the pulsar kicks~\\cite{ks97,Kusenko:1998yy} and the strength of the supernova explosion~\\cite{Fryer,Hidaka:2006sg}, as well as dark matter~\\cite{dw,Fuller,shi_fuller,smirnov,nuMSM}.  The same particle can play an important role in the formation of the first stars~\\cite{reion} and other astrophysical phenomena~\\cite{biermann_munyaneza}.  The scale of the right-handed Majorana masses $M_{a}$ is unknown; it can be much greater than the electroweak scale~\\cite{seesaw}, or it may be as low as a few eV~\\cite{deGouvea:2005er}. The seesaw mechanism~\\cite{seesaw} can explain the smallness of the neutrino masses in the presence of the Yukawa couplings of order one if the Majorana masses $M_a$ are much larger than the electroweak scale. Indeed, in this case the masses of the lightest neutrinos are suppressed by the ratios $ \\langle H \\rangle/M_a$. However, the origin of the Yukawa couplings remains unknown, and there is no experimental evidence to suggest that these couplings must be of order 1. In fact, the Yukawa couplings of the charged leptons are much smaller than 1. For example, the Yukawa coupling of the electron is as small as $(10^{-6})$.  The Majorana mass can arise from the Higgs mechanism\\cite{Chikashige:1980ht}. For example, let us consider the following modification of the lagrangian (\\ref{lagrangianM}): \\begin{equation} {\\cal L} =   {\\cal L}_{0}+\\bar N_{a} \\left(i \\gamma^\\mu \\partial_\\mu \\right ) N_{a}  - y_{\\alpha a} H \\,  \\bar L_\\alpha N_{a}  - \\frac{h_a}{2} \\, S \\, \\bar {N}_{a}^c N_{a} +V(H,S) + h.c. \\,, \\label{lagrangianS} \\end{equation} where $ {\\cal L}_{0}$ includes the gauge and kinetic terms of the Standard Model, $H$ is the Higgs doublet, $S$ is the real boson which is SU(2)-singlet, $L_\\alpha$ ($\\alpha=e,\\mu,\\tau$) are the lepton doublets, and $ N_{a}$ ($a=1,...,n$) are the additional singlet neutrinos. After the symmetry breaking, the Higgs doublet and singlet fields each develop a VEV, $\\langle H\\rangle= v_0=247$~GeV, $\\langle S\\rangle= v_1$, and the singlet neutrinos acquire the Majorana masses $ M_a = h_a v_1$.   As discussed below, this model is suitable for generating dark matter in the form of sterile neutrinos. ",
        "conclusions": "The fundamental physics responsible for the neutrino masses is likely to involve some additional SU(2)-singlet fermions, or sterile neutrinos.  The Majorana masses of these states can range from a few eV to some values well above the electroweak scale.  A sterile neutrino with a keV mass is a viable dark matter candidate.  One can discover the relic sterile neutrinos using the X-ray observations.   The same neutrinos can be produced in the supernova explosions, and the anisotropy in their emission can explain the observed pulsar velocities.    The X-rays from the sterile neutrino decays can play an important role in the production of molecular hydrogen, which is necessary for the formation of the first stars.  This work was supported in part by the DOE grant DE-FG03-91ER40662 and by the NASA ATP grants NAG~5-10842 and NAG~5-13399.  The author appreciates the hospitality of Aspen Center for Physics, where part of this work was done."
    },
    "0711/0711.2362_arXiv.txt": {
        "abstract": "Single-star stellar population (ssSSP) models are usually used for spectral stellar population studies. However, more than 50\\% of stars are in binaries and evolve differently from single stars. This suggests that the effects of binary interactions should be considered when modeling the stellar populations of galaxies and star clusters. Via a rapid spectral stellar population synthesis ($RPS$) model, we give detailed studies of the effects of binary interactions on the Lick indices and colours of stellar populations, and on the determination of the stellar ages and metallicities of populations.  Our results show that binary interactions make stellar populations less luminous, bluer, with larger age-sensitive Lick index (H$\\beta$) and smaller metallicity-sensitive indices (e.g., Mgb, Fe5270 and Fe5335) compared to ssSSPs. It also shows that when ssSSP models are used to determine the ages and metallicities of stellar populations, smaller ages or metallicities will be obtained, when using two line indices (H$\\beta$ and [MgFe]) and two colours (e.g., $u-R$ and $R-K$), respectively. Some relations for linking the stellar-population parameters obtained by ssSSPs to those obtained by binary-star stellar populations (bsSSPs) are presented in the work. This can help us to get some absolute values for stellar-population parameters and is useful for absolute studies. However, it is found that the relative luminosity-weighted stellar ages and metallicities obtained via ssSSPs and bsSSPs are similar. This suggests that ssSSPs can be used for most spectral stellar population studies, except in some special cases. ",
        "introduction": "Stellar population synthesis is a powerful technique to study the stellar contents of galaxies and star clusters (see, e.g., \\citealt{Yungelson:1997}, \\citealt{Tout:1997}, \\citealt{Pols:1998}, \\citealt{Hurley:2007}). It is also an important method to study the formation and evolution of galaxies. Simple stellar population (SSP) models that do not take binary interactions into account are usually used for spectral stellar population studies, as most models are ssSSP models (e.g., \\citealt{Bruzual:2003}, \\citealt{Vazdekis:1999}, \\citealt{Fioc:1997}, \\citealt{Worthey:1994}). However, as pointed by, e.g., \\cite{Duquennoy:1991}, \\cite{Pinfield:2003}, and \\cite{Lodieu:2007}, about 50\\% of stars are in binaries and they evolve differently from single stars. We can see this when comparing the isochrone of an ssSSP to that of a bsSSP (Fig.1). In fact, bsSSPs better fit the colour-magnitude diagrams (CMDs) of star clusters than ssSSPs \\citep{Li:2007database}. This suggests that binary interactions can affect stellar population synthesis studies significantly and it is, therefore, necessary to consider binary interactions. This is also supported by some observational results, e.g., the Far-UV excess of elliptical galaxies \\citep{Han:2007} and blue stragglers in star clusters (e.g., \\citealt{Davies:2004}; \\citealt{Tian:2006}; \\citealt{Xin:2007}). These phenomena can be naturally explained via stellar populations with binary interactions, without any special assumptions.  A few works have tried to model populations via binary stars and have presented some results on the effects of binary interactions on spectral stellar population synthesis. For example, \\citet{Zhang:2004, Zhang:2005} showed that binary interactions can make bsSSPs bluer than ssSSPs. However, there is not a more detailed investigation about how binary interactions affect the Lick indices and colours of stellar populations, and on the determination of stellar ages and metallicities. One of our previous works, \\citet{Li:2007database}, compared bsSSPs to ssSSPs, but various stellar population models are used. This makes it difficult to understand the effects of the changes of Lick indices and colours of populations that result only from binary interactions. Furthermore, it did not show how binary interactions affect the determinations of stellar ages and metallicities. In this case, we have no clear picture for the differences between the predictions of bsSSPs and ssSSPs, and do not well know the differences between luminosity-weighted stellar-population parameters (age and metallicity) determined by bsSSPs and ssSSPs. Because all galaxies contain some binaries, detailed studies of the effects of binary interactions on the Lick indices and colours of populations are important, as are the determinations of stellar-population parameters. In this work, we perform a detailed study of the effects of binary interactions on spectral stellar population synthesis studies, via the rapid spectral population synthesis ($RPS$) model of \\citet{Li:2007database}.  The paper is organized as follows. In Sect. 2, we briefly introduce the $RPS$ model. In Sect. 3, we study the effects of binary interactions on the isochrones, spectral energy distributions (SEDs), Lick Observatory Image Dissector Scanner absorption line indices (Lick indices), and colours of stellar populations. In Sect. 4, we investigate the differences between stellar ages and metallicities fitted by ssSSPs and bsSSPs. Finally, in Sect. 5, we give our discussions and conclusions. ",
        "conclusions": "We investigated the effects of binary interactions on the isochrones, SEDs, Lick indices and colours of simple stellar populations, and then on the determination of light-weighted stellar ages and metallicities. The results showed that binary interactions can affect stellar population synthesis studies significantly. In detail, binary interactions make stellar populations less luminous and bluer, while making the H$\\beta$ index larger and metal line indices smaller compared to ssSSPs. The colour changes (2$\\sim$5\\%) caused by binary interactions are smaller than the systematic errors (about 6\\%) of the $RPS$ model while similar to observational errors (4$\\sim$7\\%). Note that the systematic error of 6\\% did not take the uncertainties due to the free parameters of the star model into account (see Sect. 2). The changes (3$\\sim$6\\%) of Lick indices caused by binary interactions are somewhat smaller than the systematic errors (about 6\\%) of stellar population synthesis model but larger than observational errors (1$\\sim$4\\%). Therefore, if we measure luminosity-weighted stellar-population parameters (metallicity and age) via bsSSPs instead of ssSSPs, higher (0.0010 on average) metallicities and significantly larger ages will be obtained via a Lick-index method, and significantly higher (about 0.0030) metallicities and similar ages will be obtained via a photometric method. Because simple stellar population models are usually used for studying the populations of early-type galaxies or globular clusters, which possibly have old ($>$ 7\\,Gyr) and relative metal-poor populations, the changes ($\\sim$ 1.8\\,Gyr in age and 0.0030 in metallicity) caused by binary interactions in stellar ages and metallicities are larger than the typical uncertainties. In particular, we found that the relative results of stellar population studies obtained by ssSSPs and bsSSPs are similar. The bs-fitted stellar-population parameters can be calculated from the ss-fitted ones, via equations presented by the paper.  The relations between bs-fitted ages and ss-fitted stellar-population parameters are useful for some special investigations. For example, when studying the age of the universe via the stellar ages of some distant globular clusters, we can estimate the absolute age of star clusters using ss-fitted results. Although the results shown in this paper can help us to give some estimates for the absolute stellar ages and metallicities of galaxies, it is far from getting accurate values because of the large uncertainties in stellar population models (see, e.g., \\citealt{Yi:2003}). In addition, different stellar population models usually give different absolute results for stellar population studies. Note that the results obtained by the Lick-index method are affected slightly by $\\alpha$-enhancement, according to the work of \\citet{Thomas:2003}, but this is not the case for the results obtained by photometric methods.   In this work, all bsSSPs contains about 50\\% binaries with orbital periods less than 100 yr (the typical value of the Galaxy). If the binary (with orbital periods less than 100 yr) fraction of galaxies are different from 50\\%, the results shown in the paper will change. The higher the fraction of binaries, the larger the difference between ss- and bs-fitted stellar-population parameters. Thus the results obtained by in this paper may not be proper for investigating galaxies or star clusters with binary fractions obviously different from 50\\%. Furthermore, when building bsSSPs, we assumed that the masses of the two components of a binary are correlated \\citep{Li:2007database}, according to previous works. We did not try to take other distributions for secondary mass and binary period in this work, because it is limited by our present computing ability. We will give further studies in the future. The differences between the Lick indices and colours of bsSSPs and ssSSPs do not evolve smoothly with age. This possibly relates to the method used to calculate the integrated features of stellar populations. In fact, the Monte Carlo method usually leads to some scatter (about 2\\%) in the integrated features of populations. The analytic fits and the binary algorithm used by Hurley code can also lead to some scatter (about 5\\%).  We investigated the effects of binary interactions via only some simple stellar populations, but the real populations of galaxies and star clusters are usually not so simple. In other words, the populations of galaxies and star clusters seem to be composite stellar populations including populations with various ages and metallicities (e.g., \\citealt{Yi:2005}). It seems that the effects of binary interactions and population-mixing are degenerate. This is a complicated subject, which requires further study."
    },
    "0711/0711.0684_arXiv.txt": {
        "abstract": "We present 21 cm HI line observations of 5$\\times$1 square degrees centered on the local Abell cluster 1367 obtained as part of the Arecibo Galaxy Environment Survey. One hundred sources are detected (79 new HI measurements and 50 new redshifts), more than half belonging to the cluster core and its infalling region. Combining the HI data with SDSS optical imaging we show that our HI selected sample follows scaling relations similar to the ones usually observed in optically selected samples. Interestingly all galaxies in our sample appear to have nearly the same baryon fraction independently of their size, surface brightness and luminosity. The most striking difference between HI and optically selected samples resides in their large scale distribution: whereas optical and X-ray observations trace the cluster core very well, in HI there is almost no evidence of the presence of the cluster. Some implications on the determination of the cluster luminosity function and HI distribution for samples selected at different wavelength are also discussed. ",
        "introduction": "With the advent of wide field surveys our knowledge of galaxy properties has significantly improved. The large amount of high quality data of nearby and high redshift objects obtained so far is allowing astronomers to shed light on how galaxies formed and evolved at different redshifts and in different environments (e.g. \\citealp{cowie96,phenomen,kauff03}). However our picture of the Universe still remains limited since it is mostly based on the study of rest-frame optical selected objects. First of all optically-selected galaxies already contain, by definition, a large amount of stars which may not be the case in younger or less evolved galaxies. Secondly, surface brightness selection is very dramatic and nearly all catalogued objects lie in an extremely narrow range of surface brightnesses \\citep{freeman70}.\\\\ Since each wavelength has its own selection effects and traces different baryonic components, the combination of blind surveys at various frequencies is of vital importance. Alternative methods of baryon selection will allow us to have a less biased view of our Universe and enable us to correctly reconstruct the evolutionary history of galaxies (e.g. \\citealp{vero06}). In particular, the most interesting results can come from the comparison of optical surveys with blind surveys tracing components not directly related to stellar emission like, for example, neutral and molecular hydrogen. For molecular hydrogen we will have to wait a couple of years for the first results from ALMA, but with the advent of multi-beam instruments on large single dish radio telescopes it has become possible during the last decade to carry out fully sampled blind HI surveys of the sky (e.g. \\citealp{dualbeam,henning00,barnes01,lang03,minchin03,davies04}). Neutral hydrogen is extremely important because it represents the fuel for the future star formation activity of a galaxy. It is also one of the galactic components that can most easily be affected by environmental mechanism (e.g. ram pressure, tidal interaction) influencing its star formation history. Therefore HI line observations of galaxies can provide us with some of the most powerful diagnostics on the role of the environment in regulating the evolution of galaxies \\citep{haynes,giova85}. What we lack is deep and high spatial resolution HI blind surveys of cluster of galaxies and their outskirts. The HI Parkes all Sky Survey (HIPASS) has represented a real breakthrough for HI blind surveys, but its low spatial resolution ($\\sim$15 arcmin) has made it very difficult to carry out detailed studies of high density regions. Fortunately, with the refurbishing of the 305m Arecibo radio telescope, blind HI surveys with much improved sensitivity, spatial  (3.5 arcmin $\\rm beam^{-1}$) and  velocity (5 $\\rm km~s^{-1}$) resolution have become possible thanks to the installation of the ALFA multi beam instrument.  The Abell cluster 1367, due to its proximity (z$\\sim$0.0216, V$\\sim$6500 \\kms) and to the fact that it is currently forming at the intersection of two filaments in the Great Wall \\citep{COGA04}, represents an ideal target for HI surveys (e.g. \\citealp{sullivan81,chinca83,gav87,solanes01,peppo06}). It is ideal for the study of the properties of HI selected galaxies and environmental effects. For these reasons, Abell 1367 is one of the region observed as a part of the Arecibo Galaxy Environment Survey (AGES, \\citealp{auld06}), one of the new HI blind survey carried out with the ALFA multi-beam system\\footnote{The other three main extragalactic ALFA surveys are ALFALFA \\citep{alfaalfa05}, AUDS \\citep{auds} and ZOA \\citep{zoa}.}. AGES aims to study the atomic hydrogen properties of different galactic environments to faint sensitivity limits; low HI masses (6$\\times10^{8}~\\rm M_{\\odot}$, assuming a 200 \\kms velocity width at 92.8 Mpc, the distance of the A1367 cluster) and column densities of $\\sim$ 3$\\times10^{18}~\\rm cm^{2}$ (for a source that fills the beam). The environments AGES will survey range from apparent voids in the large-scale structure of galaxies, via isolated spiral galaxies and their haloes, to galaxy-rich regions associated with galaxy clusters and filamentary structures. AGES plans to cover $5\\times4$ square degrees centered on Abell 1367. Here we report on the early results obtained from observations of a subset of the whole AGES-A1367 region: a $5\\times1$ square degree area centered on the cluster core. ",
        "conclusions": "We have presented the results of a HI blind survey of 5$\\times$1 square degrees centered on the core of Abell 1367 as part of the Arecibo Galaxy Environment Survey. One hundred HI galaxies have been detected (79 new measurements and 50 new redshifts), half of which belong to the Abell 1367 volume. Comparing the properties of  the HI selected galaxies with an optically selected sample extracted from SDSS-DR5 we have shown that the large scale distribution of galaxies is strongly wavelength dependent: in HI the cluster core almost completely disappears and HI selected galaxies are homogeneously distributed in the cluster volume. The difference between the two samples is even more evident when we compare their luminosity and HI distribution. Our optically selected sample misses a considerable number of low luminosity, low surface brightness galaxies underestimating by almost a factor 2 the number of HI rich galaxies present in the cluster volume. On the contrary our HI sample contains a factor 3 less bright galaxies than the optical one. These are mainly ellipticals or early type, red spirals but also a not negligible fraction of the blue sequence cluster spirals is missed. All these blue gas-poor objects are located near the edge of the X-ray cluster contours suggesting that their gas has been recently ($\\sim$ 100 Myr) stripped by ram pressure so that the galaxies are already gas poor whereas their stellar populations are still young.\\\\ Although the spatial and number distribution of samples selected at different wavelengths are significantly different, the internal properties of galaxies seem not to be significantly wavelengths dependent. HI galaxies seem to follow the same scaling relations observed in optical and bigger galaxies are redder, have a higher surface brightness and a lower gas content than lower luminosity galaxies. The most interesting result of our analysis is that the baryon fraction of our sample appears to be almost constant, suggesting the difference in the star to gas ratio observed between bright and faint galaxies is only due to a different star formation efficiency (i.e. the time scale necessary to convert gas into stars). This is the first time that such a result is obtained from an HI selected sample and even if consistent with models of galaxy formation it suggests that during their evolutionary history less massive galaxies do not lose a larger fraction of their baryons than bigger objects.  The analysis presented in this paper points out the importance of HI blind surveys like AGES in order to have an unbiased view of the properties of galaxies in the local Universe. Once the survey is at a more advanced stage and the size of our sample is significantly bigger than the one presented here it will be possible to carry out more detailed quantitative studies and comparisons between optically and HI selected samples, gaining more insight into the evolutionary history of nearby galaxies."
    },
    "0711/0711.2154_arXiv.txt": {
        "abstract": "We present APEX observations of HCN (9-8) and (4-3) lines toward a sample of hot cores. The spectral shapes of the main transitions are asymmetric and self-absorbed, as expected for high optical depth in a possibly infalling envelope. For spherical symmetry, the large column densities of these sources would mean that the central region is obscured by dust above a certain frequency. However, we detected the vibrationally excited satellite lines ($v_2=1$; $J=$9-8) at 797 GHz, which  originate from the inner regions. This indicates that high-frequency ALMA observations of hot core centers will be feasible. ",
        "introduction": "\\label{intro} The early stages of star formation are known to occur deep inside interstellar clouds of dust and molecular gas. This is particularly true when massive protostars are produced: the hydrogen-burning reactions in a high-mass protostar are ignited while the object is still accreting and deeply embedded in its parent cloud.  Hence, due to its much higher energy output, it stirs and heats up the natal core in a much more effective way than a low-mass object, giving rise to a hot molecular core.  \\begin{figure} \\includegraphics[width=0.45\\textwidth]{dust.eps} \\caption{In the top panel, the region with high dust opacity (dark brown) is very small, while the region with a density high enough to excite molecules (light brown) is large, one sees emission spectra.  As one goes higher in frequency (and excitation energy), the dust photosphere grows while the critical density region shrinks, and absorption components toward the opaque dust core appear (center panel), while at very high energies and frequencies the region where the molecule can emit becomes embedded in the opaque dust, and no line radiation can escape to be detected by us (bottom panel).} \\label{fig:dust}       % \\end{figure}   \\begin{figure*} \\includegraphics[angle=-90, width=0.85\\textwidth]{spec798.eps} \\caption{The observed spectra around 798/810 GHz are shown for two sources, together with preliminary line identifications} \\label{fig:spec798}       % \\end{figure*}  This stage of high-mass star formation is characterized by processes that occur in the dense and heated (to $>$ 100 K) surroundings of the protostar: Ice mantles around dust grains evaporate and the frozen molecules are thereby released into the gas phase. Due to the high temperature and density, the molecules are easily excited, and due to the large column densities, they can readily be observed. Since hot cores display a plethora of molecular lines (see Figure~\\ref{fig:spec798}), they can be studied very well using molecular lines, particularly of highly excited transitions (for a review, see e.g. van der Tak 2004 \\cite{vdTak04}). ",
        "conclusions": "\\label{results} We found evidence for the (9-8, $v_2=1$) line toward all sour- ces where the (4-3, $v_2=1$) was detected, indicating that, at least at this wavelength and for this source sample, the central core is still visible.  Hence ALMA observations of the central engines of hot cores even at the highest frequencies appear to be feasible, although the cores may shut down at higher frequencies relevant for Herschel/HIFI.  The ground state lines often display self-absorption features. Their asymmetric line shapes with the blue part being the stronger one could be indicative of infall motion. To extract the information on the source structure from the line strengths and shapes, we plan to analyze the data with detailed radiative transfer modeling (Hogerheijde \\& van der Tak 2000 \\cite{Hogerheijde00}). Preliminary results indicate that modelling the source as a spherical cloud cannot reproduce all observed lines, indicating deviations from this simple structure. Investigations of effects of such deviations are under way."
    },
    "0711/0711.1843.txt": {
        "abstract": "{%% In this work we present the results of a systematic search for diffuse bands (DBs, hereafter) in the circumstellar envelopes of a carefully selected sample of post-AGB stars. We concentrated on the analysis of 9 of the DBs most commonly found in the interstellar medium. The strength of these features is determined using high resolution optical spectroscopy and the results obtained are compared with literature data on field stars affected only by interstellar reddening. Based on the weak features observed in the subsample of post-AGB stars dominated by circumstellar reddening we conclude that the carrier(s) of these DBs must not be present in the circumstellar environment of these sources, or at least not under the excitation conditions in which DBs are formed. The conclusion is applicable to all the post-AGB stars studied, irrespective of the dominant chemistry or the spectral type of the star considered. A detailed radial velocity analysis of the features observed in individual sources confirms this result, as the Doppler shifts measured are found to be consistent with an interstellar origin. ",
        "introduction": "%% The diffuse interstellar bands (DIBs) are absorption features, showing a broad range of widths and strengths, which appear over-imposed on the spectra of bright stars whose lines of sight probe (extra)galactic diffuse to dense interstellar clouds. Currently, more than 200 DIBs have been identified and catalogued in the spectral range from 3600 to 10200~\\AA\\ \\citep{Jenniskens94b,Cox05}, the most studied ones being those found at 4430, 5780, 5797 and 6284~\\AA. Since their discovery \\citep{Heger22}, they have been associated to the interstellar medium (ISM), because their strengths show a positive relationship with the observed extinction \\citep{Merrill36} as well as to the neutral sodium column density \\citep{Herbig93}. Many carriers have been proposed, however, no unambiguous identification has yet been made and it is debated whether they arise from the dust or the gas component of the ISM (see reviews by \\citealt{Herbig95} and \\citealt{Sarre06}). There is increasing observational evidence that the DIB carriers constitute a set of carbonaceous gas phase molecules as evidenced from substructures resembling rotational contours in some bands (\\citealt{Sarre95}, \\citealt{Ehrenfreund96}). In particular, photo-UV-resistant organic molecules, such as carbon chains \\citep{Douglas77}, PAHs \\citep{Salama99,Allamandola99}, fullerenes \\citep{Foing94,Iglesias-Groth07} and / or buckyonions \\citep{Iglesias-Groth04} are promising candidates. The local interstellar environmental conditions set the balance of local formation and destruction of the carriers as well as their level of ionization and hydrogenation. The interstellar radiation field is one of the most important factors in this \\citep{Ruiterkamp05}.  There is a possible link between the DIB carriers and the carriers of the unidentified (aromatic) infrared bands (UIBs), the so-called PAH-DIB hypothesis (\\citealt{Crawford85, LegerDHendecours85, VanderZwet-Allamandola85}).  Thus, although PAHs are thought to reside and to be processed (ionisation, dehydrogenation, destruction) in the diffuse ISM, this does not exclude the scenario that these molecules (or their parent species) are produced elsewhere. Since circumstellar shells are sources of replenishment of the ISM, it has been argued that DIBs (and/or parent structures) may have a circumstellar origin, either in dense stellar winds or circumstellar shells, thus somehow contravening the name they were initially given. The suspected connection between DIB carriers and some carbon-rich compounds can be investigated attending to the usually known chemistry and physical properties of these circumstellar shells.  Observationally, the detection of diffuse bands (DBs, hereafter) around evolved stars is hampered by the fact that most of them are mass-losing stars, usually strongly variable, and surrounded by very cool extended atmospheres where molecules are the dominant source of opacity. These stars are very difficult to model and DBs are hardly detected (in absorption) against the forest of features attributed to molecular transitions which appear over-imposed on the stellar continuum. This has hampered the systematic search for DBs in evolved stars in the past. Furthermore, if detected, it is difficult to determine whether the DBs are originating from the interstellar or the circumstellar environment, or even both.  In face of these difficulties \\citet{SnowWallerstein73} and \\citet{Snow73} searched for circumstellar diffuse bands (at 4430, 5780, 5797 and 6614~\\AA) in 26 stars with suspected circumstellar dust shells / envelopes but found no evidence for their presence. Several other authors have since searched for and commented on the presence of diffuse bands and their interstellar or circumstellar nature, in spectra observed towards planetary nebulae (\\mbox{\\object{NGC 6210}}, \\mbox{\\object{NGC 7027}}, \\mbox{\\object{IC 351}} and \\mbox{\\object{AFGL 2688}} by \\citealt{PritchetGrillmair84}; \\mbox{\\object{IRAS 21282+5050}} by \\citealt{CohenJones87}; \\mbox{\\object{NGC 7027}} and \\mbox{\\object{IRAS 21282+5050}} by \\citealt{LeBertreLequeux92}), a post-AGB star (\\mbox{\\object{HR 4049}}; \\citealt{WatersEtal89}), and a carbon star (\\mbox{\\object{IRAS 07270-1921}} or \\mbox{CGCS 1732}; \\citealt{LeBertre90}). \\citet{Lebertre93} studied a new sample consisting of carbon-, oxygen- or nitrogen-rich mass-losing sources, such as (pre)planetary nebulae (\\mbox{\\object{BD+30 3639}}, \\mbox{\\object{CPD-56 8032}}, \\mbox{Hen 104}), a carbon rich RV Tauri star (\\mbox{\\object{AC Her}}), wolf-rayet stars (\\mbox{\\object{WR137}}, \\mbox{\\object{WR140}}) and post-AGB stars (\\mbox{\\object{HR 4049}}, \\mbox{\\object{HD213985}}), revisiting several sources studied in the past (\\emph{e.g.} \\mbox{CS 776}, \\mbox{\\object{NCG 7027}}, \\mbox{\\object{HR 4049}}, \\mbox{\\object{IRAS 21282+5050}}). These authors did not find any evidence for circumstellar DBs and refuted previous claims and thus concluded that DBs are depleted in circumstellar environments. Notably, these authors did not detect any bands in the spectra of sources showing strong PAH emission (UIB) at mid-infrared wavelengths. This suggested that carrier molecules, if present, in circumstellar envelopes are in a different state of ionization / hydrogenation than in the ISM. Strong DBs were detected toward carbon-rich sources that do not show PAH emission, as well as toward most of their oxygen-rich and nitrogen-rich sources in the sample, although in all cases the observed DBs could be attributed to the interstellar material in the lines of sight. For unexplained reasons enhanced DBs were detected toward WN stars and LBVs. Exceptionally, narrow emission features possibly related to DIBs as well have been observed toward the \\object{Red Rectangle} \\citep{Scarrott92}, although their identification and nature remains controversial.  A largely unexplored alternative exists. This concerns the so-called post-AGB stars that are in a short-lived transition phase between the Asymptotic Giant Branch (AGB) and the Planetary Nebula (PN) stage, evolving very rapidly in the H-R diagram while they are still surrounded by the remnant of the AGB circumstellar shell. Post-AGB stars show all possible spectral types from M to B in what probably represents an evolutionary sequence of increasing effective temperature in their way to become PNe \\citep{Garcia-Lario97b}. This means that in these stars we should easily be able to detect DBs formed in the remnant AGB shell over-imposed on the intermediate or early-type spectrum of the central star without the confusion originating from the presence of molecular bands in AGB stars. Interestingly, while many of these DBs are common to those observed in the ISM, the relative ratios are sometimes found to be very different \\citep{Garcia-Lario99}. These circumstellar DBs could form and survive for some time under conditions which might be substantially different to those found in the ISM in terms of density, UV radiation field, etc. and could hold the key to understand and solve this long-standing problem. Another advantage is the fact that the chemical composition of the gas and dust in these shells can easily be determined from observations in the optical, infrared, mm / sub-mm or radio wavelengths. In addition, post-AGB stars are located in many cases at relative high galactic latitudes, and are as such affected only by little interstellar reddening. This facilitates the attribution of a circumstellar origin to the features observed.  The potential formation of DBs around post-AGB stars has, however, been explored so far only occasionally for a limited number of sources. Nevertheless, the presence of strong DBs has been reported in the optical spectra of a few post-AGB stars (\\citealt{Lebertre93,Garcia-Lario99,Zacs99,Zacs01,Klochkova99,Klochkova00,Kendall02}) and some carbon rich (barium) stars \\citep{Zacs03}. In several cases tentative claims have been put forward of DBs detected at radial velocities coinciding with the photospheric absorption lines or shell/envelope expansion velocity (\\mbox{\\object{IRAS 04296+3429}} by \\citealt{Klochkova99}; and \\object{HD 179821} by \\citealt{Zacs99}). All other studies listed above gave non-conclusive results.  Another method recently employed utilised nearby background stars to probe the circumstellar environments of the carbon star \\mbox{\\object{IRC+10216}} \\citep{Kendall02} and the Helix PN \\citep{MauronKendall04}. Again, no circumstellar DBs were detected, confirming the lack of DIB carriers in these environments.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  \\begin{table*}[ht!] \\caption{Observation log.}\\centering \\label{tb:observations} \\resizebox{\\textwidth}{!}{%% \\begin{tabular}{ l c c c c c c c }\\hline\\hline Name                  &      Other Names          & Observing Date&   Observatory &   Telescope   & Instrument& Spectral range (\\AA) & Resolution\\\\ \\hline \\object{IRAS 01005+7910}  &               &   12.09.03    &   La Palma    &   TNG (3.58 m)&   SARG    &   4960$-$10110    &   50000   \\\\ \\object{IRAS Z02229+6208} &               &   12.09.03    &   La Palma    &   TNG (3.58 m)&   SARG    &   4960$-$10110    &   50000   \\\\ \\object{IRAS 04296+3429}  &               &   22.08.94    &   La Palma    &   WHT (4.2 m) &   UES     &   5570$-$10200    &   55000   \\\\ \\object{IRAS 05113+1347}  &               &   23.02.94    &   La Palma    &   WHT (4.2 m) &   UES     &   4480$-$10200    &   55000   \\\\ \\object{IRAS 05251-1244}  & \\object{IC 418}       &   12.09.03    &   La Palma    &   TNG (3.58 m)&   SARG    &   4960$-$10110    &   50000   \\\\ \\object{IRAS 05341+0852}  &               &   30.09.98    &   ESO-La Silla&   NTT (3.58 m)&   EMMI    &   5570$-$10040    &   65000   \\\\ \\object{IRAS 06530-0213}  &               &   12.01.01    &   ESO-Paranal &   VLT-U2 (8 m)&   UVES    &   4790$-$6810     &   100000  \\\\ \\object{IRAS 07134+1005}  & \\object{HD 56126}         &   20.12.93    &   La Palma    &   WHT (4.2 m) &   UES     &   3980$-$10400    &   55000   \\\\ \\object{IRAS 08005-2356}  &               &   02.03.94    &   La Palma    &   WHT (4.2 m) &   UES     &   3690$-$11050    &   55000   \\\\ \\object{IRAS 08143-4406}  &               &   16.01.01    &   ESO-Paranal &   VLT-U2 (8 m)&   UVES    &   3750$-$10520    &   100000  \\\\ \\object{IRAS 08544-4431}  &               &   26.01.99    &   ESO-La Silla&   ESO-1.52    &   FEROS   &   3700$-$8860     &   50000   \\\\ \\object{IRAS 12175-5338}  & \\object{SAO 239853}       &   24.03.00    &   ESO-La Silla&   ESO-1.52    &   FEROS   &   3730$-$8850     &   50000   \\\\ \\object{IRAS 16594-4656}  &               &   20.05.00    &   ESO-Paranal &   VLT-U2 (8 m)&   UVES    &   3750$-$10520    &   100000  \\\\ \\object{IRAS 17086-2403}  &               &   14.07.01    &   La Palma    &   WHT (4.2 m) &   UES     &   4300$-$9000     &   50000   \\\\ \\object{IRAS 17097-3210}  & \\object{HD 155448}        &   07.07.97    &   ESO-La Silla&   NTT (3.58 m)&   EMMI    &   5800$-$10430    &   60000   \\\\ \\object{IRAS 17150-3224}  & \\object{RAFGL 6815}       &   10.06.00    &   ESO-Paranal &   VLT-U2 (8 m)&   UVES    &   3750$-$10520    &   100000  \\\\ \\object{IRAS 17245-3951}  &               &   21.05.00    &   ESO-Paranal &   VLT-U2 (8 m)&   UVES    &   3750$-$10520    &   100000  \\\\ \\object{IRAS 17395-0841}  &               &   13.07.01    &   La Palma    &   WHT (4.2 m) &   UES     &   4300$-$9000     &   55000   \\\\ \\object{IRAS 17423-1755}  & \\object{Hen 3-1475}       &   13.07.01    &   La Palma    &   WHT (4.2 m) &   UES     &   4300$-$9000     &   55000   \\\\ \\object{IRAS 17436+5003}  & \\object{HD 161796}        &   01.09.94    &   La Palma    &   WHT (4.2 m) &   UES     &   3650$-$10200    &   55000   \\\\ \\object{IRAS 18025-3906}  &               &   30.09.98    &   ESO-La Silla&   NTT (3.58 m)&   EMMI    &   3980$-$10430    &   65000   \\\\ \\object{IRAS 18062+2410}  & \\object{HD 341617}        &   13.07.01    &   La Palma    &   WHT (4.2 m) &   UES     &   4300$-$9000     &   50000   \\\\ & \\object{HD 172324}        &   08.08.95    &   ESO-La Silla&   NTT (3.58 m)&   EMMI    &   3650$-$10240   &   55000   \\\\ \\object{IRAS 19114+0002}  & \\object{HD 179821}        &   07.07.97    &   ESO-La Silla&   NTT (3.58 m)&   EMMI    &   5800$-$10430    &   60000   \\\\ \\object{IRAS 19200+3457}  & \\object{LS II +34 1}      &   11.09.03    &   La Palma    &   TNG (3.58 m)&   SARG    &   4960$-$10110    &   50000   \\\\ \\object{IRAS 19386+0155}  &               &   30.09.98    &   ESO-La Silla&   NTT (3.58 m)&   EMMI    &   5980$-$8320     &   65000   \\\\ \\object{IRAS 19500-1709}  & \\object{HD 187885}        &   07.08.95    &   ESO-La Silla&   NTT (3.58 m)&   EMMI    &   3650$-$10040    &   55000   \\\\ \\object{IRAS 20000+3239}  &               &   23.08.94    &   La Palma    &   WHT (4.2 m) &   UES     &   4440$-$10040    &   50000   \\\\ \\object{IRAS 20462+3416}  & \\object{LS II +34 26}     &   17.08.96    &   La Palma    &   WHT (4.2 m) &   UES     &   5300$-$9380     &   55000   \\\\ \\object{IRAS 22023+5249}  & \\object{LS III +52 24}    &   14.07.01    &   La Palma    &   WHT (4.2 m) &   UES     &   4300$-$9000     &   55000   \\\\ \\object{IRAS 22223+4327}  & \\object{BD+42 4388}       &   24.08.94    &   La Palma    &   WHT (4.2 m) &   UES     &   4440$-$10040    &   50000   \\\\ \\object{IRAS 22272+5435}  & \\object{HD 235858}        &   23.08.94    &   La Palma    &   WHT (4.2 m) &   UES     &   4440$-$10220    &   55000   \\\\ \\object{IRAS 23304+6147}  &               &   23.08.94    &   La Palma    &   WHT (4.2 m) &   UES     &   4440$-$10040    &   55000   \\\\ \\hline \\end{tabular} } \\end{table*}  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  \\begin{figure}[th!] \\begin{center} \\includegraphics[width=\\columnwidth]{5282fig1.eps} \\caption{%% Removal of telluric lines at 6284, 6993 and 7224~\\AA. A few examples are shown before (left) and after correction (right). In the middle panel we show the stellar spectrum of \\object{HD 172324} (B9Ib), one of the sample stars, which was used as telluric divisor (see text).} \\label{fig:mezcla} \\end{center} \\end{figure}  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  \\begin{table*}[ht!] \\caption{Main characteristics of the post-AGB stars selected for the analysis} \\centering \\label{tb:referencias} \\begin{tabular}{l r@{~$\\pm$~}l c c c c c r r }\\hline\\hline IRAS Name    & \\multicolumn {2}{c}{$E(B-V)$}    &    Ref.      & Chemistry & Ref.&  Sp.Type   &   Ref.   &   GLON.   &   GLAT.    \\\\ \\hline 01005+7910   & 0.2  & 0.1   & (25)      &  C   &   (25)   &   B2       &   (25)   &   123.57  &   +16.59   \\\\ 02229+6208   & 1.67 & 0.09  & (17)      &  C   &   (37)   &   G9Ia     &   (37)   &   133.73  &   +1.50    \\\\ 04296+3429   & 1.3  & 0.1   & (7, 15, 36)  &  C   &    (27)   &   G0Ia     &   (3)    &   166.24  &   $-$9.05  \\\\ 05113+1347   & 1.1  & 0.2   & (26, 36, 40) &  C   &    (26)   &   G5I      &   (40)   &   188.86  &   $-$14.29 \\\\ 05251-1244   & 0.23 & 0.09  & (35)      &  C   &   (42)   &   PN/07f   &   (14)   &   215.21  &   $-$24.28 \\\\ 05341+0852   & 1.65 & 0.09  & (7, 17, 40)  &  C   &    (26)   &   F5I      &   (15)   &   196.19  &   $-$12.14 \\\\ 06530-0213   & 1.7  & 0.2   & (36, 40)     &  C   &    (18)   &   F5I      &   (18)   &   215.44  &   $-$0.13  \\\\ 07134+1005   & 0.4  & 0.1   & (15, 40)     &  C   &    (26)   &   F7 Ie    &   (40)   &   206.75  &   +9.99    \\\\ 08005-2356   & 0.7  & 0.3   &  (15, 40)     &  O   &   (5)    &   F5I      &   (3)    &   242.36  &   +3.58    \\\\ 08143-4406   & 0.8  & 0.1   &  (36, 40)     &  C   &   (39)   &   F8I      &   (36)   &   260.83  &   $-$5.07  \\\\ 08544-4431   & 1.45 & 0.09  &  (29)         &  O   &   (29)   &   F3       &   (29)   &   265.50  &   +0.39    \\\\ 12175-5338   & 0.25 & 0.09  &  (19, 38)     &  C   &   (46)   &   A9Iab    &   (32)   &   298.30  &   +8.67    \\\\ 16594-4656   & 2.2  & 0.3   &  (20, 40, 44) &  C   &   (9)    &   B7       &   (44)   &   340.39  &   $-$3.29  \\\\ 17086-2403   & 0.86 & 0.09  &  (40)         &  C?  &   (10)   &   PN+G5?   &   (34)   &   359.84  &   +8.99    \\\\ 17097-3210   & 0.06 & 0.05  &  (30)         &  C   &   (45)   &   B9       &   (30)   &   353.36  &   +4.03    \\\\ 17150-3224   & 0.68 & 0.09  &  (23)         &  O   &   (23)   &   G2I      &   (23)   &   353.84  &   +2.98    \\\\ 17245-3951   & 1.0  & 0.1   &  (40)         &  O   &   (20)   &   F6I      &   (40)   &   348.81  &   $-$2.84  \\\\ 17395-0841   & 1.1  & 0.2   &  (11, 40)     &  O?  &   (10)   &   PN+G     &   (11)   &   17.02   &   +11.10   \\\\ 17423-1755   & 1.13 & 0.09  &  (11, 40)     &  O   &   (41)   &   Be       &   (11)   &   9.36    &   +5.78    \\\\ 17436+5003   & 0.24 & 0.09  & (19)      &  O   &   (4)    &   F3 Ib    &   (48)   &   77.13   &   +30.87   \\\\ 18025-3906   & 1.15 & 0.09  & (22, 40)     &   O   &   (24)   &   G1I      &   (40)   &   353.27  &   $-$8.72  \\\\ 18062+2410   & 0.6  & 0.2   & (33, 40)     &   O   &   (33)   &   B1 I     &   (33)   &   50.67   &   +19.79   \\\\ HD 172324    & 0.03 & 0.02  & (6)       &  O   &   (1)    &   B9Ib     &   (1)    &   66.18   &   +18.58   \\\\ 19114+0002   & 0.60 & 0.05  & (19, 40)     &   O   &   (4)    &   G5 Ia    &   (3)    &   35.62   &   $-$4.96  \\\\ 19200+3457   & 0.3  & 0.1   & (11)      &  O?  &   (11)   &   B        &   (11)   &   67.57   &   +9.51    \\\\ 19386+0155   & 1.05 & 0.09  & (2)       &  O   &   (28)   &   F5I      &   (31)   &   40.51   &   $-$10.09 \\\\ 19500-1709   & 0.37 & 0.09  & (19)      &  C   &   (47)   &   F2-6 Ia  &   (3)    &   23.98   &   $-$21.04 \\\\ 20000+3239   & 1.6  & 0.1   & (15, 26)     &   C   &   (21)   &   G8Ia     &   (21)   &   69.68   &   +1.16    \\\\ 20462+3416   & 0.38 & 0.09  & (8,  43)     &   O   &   (8)    &   B1.5     &   (8)    &   76.60   &   $-$5.75  \\\\ 22023+5249   & 0.52 & 0.09  & (12)      &  O   &   (13)   &   B        &   (12)   &   99.30   &   $-$1.96  \\\\ 22223+4327   & 0.2  & 0.1   & (15, 26)     &   C   &   (26)   &   G0 Ia    &   (15)   &   96.75   &   $-$11.56 \\\\ 22272+5435   & 0.9  & 0.2   & (15)      &  C   &   (16)   &   G5 Ia    &   (15)   &   103.35  &   $-$2.52  \\\\ 23304+6147   & 1.4  & 0.2   & (15, 36)     &   C   &   (15)   &   G2 Ia    &   (15)   &   113.86  &   +0.59    \\\\ \\hline \\end{tabular} %} \\resizebox{1\\textwidth}{!}{%% \\begin{tabular}{l l l l}\\\\ (1) \\citet{Arellano01}      &   (13) \\citet{Gauba04}    & (25) \\citet{Klochkova02}&     (37) \\citet{Reddy99}  \\\\ (2) \\citet{Arkhipova00}     &   (14) \\citet{Heap87}     & (26) \\citet{Kwok95}     &     (38) \\citet{Reed96}  \\\\ (3) \\citet{Bakker97}        &   (15) \\citet{Hrivnak95}  & (27) \\citet{Kwok99}     &     (39) \\citet{Reyniers04} \\\\ (4) \\citet{Bujarrabal92}    &   (16) \\citet{Hrivnak91}  & (28) \\citet{Lewis00}    &     (40) \\citet{Suarez06} \\\\ (5) \\citet{Desmurs02}       &  (17) \\citet{Hrivnak99a}      & (29) \\citet{Maast03}    &     (41) \\citet{Te_Lintel_Hekkert91} \\\\ (6) \\citet{Fernie83}        &  (18) \\citet{Hrivnak03}   & (30) \\citet{Malfait98}  &         (42) \\citet{Torres-Peimbert80} \\\\ (7) \\citet{Fuji02}      &  (19) \\citet{Hrivnak89}   & (31) \\citet{Meixner99}  &     (43) \\citet{Turner84} \\\\ (8) \\citet{Garcia-Lario97}  &  (20) \\citet{Hrivnak99b}  & (32) \\citet{Oudmaijer92}&     (44) \\citet{Steene03} \\\\ (9) \\citet{Garcia-Lario99}  &  (21) \\citet{Hrivnak00}   & (33) \\citet{Parthasarathy00a} &   (45) \\citet{Van_der_Veen89} \\\\ (10) Garc\\'\\i a-Lario (priv.comm.)& (22) \\citet{Hu93a}      & (34) \\citet{Parthasarathy00b} &   (46) \\citet{Van_Winckel97} \\\\ (11) \\citet{Gauba03a}       &   (23) \\citet{Hu93b}      & (35) \\citet{Pottasch04} &     (47) \\citet{Van_Winckel00} \\\\ (12) \\citet{Gauba03b}       &    (24) \\citet{Hu94}      & (36) \\citet{Reddy96} &        (48) \\citet{Volk89} \\\\ \\end{tabular} } \\end{table*}   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%   In this work we present the first systematic survey to detect the presence of circumstellar DBs (DCBs) in a carefully selected sample of galactic post-AGB stars\\footnote{The sample also includes three very young planetary nebulae, which are here also considered post-AGB stars in a broad sense.}. The goal is to perform a detailed analysis of the differential properties observed in the DBs associated to post-AGB stars in comparison with the standard DBs observed towards reddened, early-type field stars (where these bands are expected to be essentially of interstellar origin).  To perform this task we have studied the intensity of 9 of the strongest absorption features identified as DBs in the spectral range 4000--10000~\\AA\\ using high resolution optical spectroscopy. The comparison of the properties observed in carbon-rich and oxygen-rich post-AGB stars is used to test the carbon-rich nature of the DB carrier(s) and determine which of the DBs detected are most probably of circumstellar origin (if any).  In Section~\\ref{section:observation} we describe the observations made and the data reduction process. The strategy followed in our analysis is presented in Section~\\ref{section:relation}. The main results are discussed in Section 4 as a function of various observational parameters. Finally, the conclusions are presented in Section~\\ref{section:conclusions}.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\begin{figure}[t!] \\begin{center} \\includegraphics[width=0.9\\columnwidth]{5282fig2.eps} \\caption{%% Equivalent width measurements taken from the literature for field stars, plotted as a function of $E(B-V)$. Table~\\ref{tb:caracdbs} gives slopes $EW$/$E(B-V)$ and correlations \\emph{r} for the linear fits.} \\label{fig:thor_02} \\end{center} \\end{figure} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{section:conclusions}  The equivalent widths of 9 DBs commonly found in the ISM have been determined for a representative sample of galactic post-AGB stars displaying a wide variety of observational properties. We present here the results of our extensive survey to look for DBs in envelopes of evolved stars.  We have carefully disentangled the observed extinction by assessing the expected interstellar extinction for each of the observed targets. This allowed us to select a sub-sample of targets whose line of sight reddenings are dominated ($>50\\%$) by circumstellar dust.  In general, the strengths of the DBs are found to follow the same correlation with $E(B-V)$ observed in field stars only in those sources showing little circumstellar contribution to the overall reddening. In contrast, DBs are weak or absent in sources dominated by circumstellar reddening, irrespective of the dominant chemistry and spectral type of the central star, although our conclusions should be taken with caution due to the relatively small sample size.   The results obtained suggest that the carrier(s) of the DBs do not form or at least they are not ``available'' to produce any detectable spectral feature during the post-AGB phase. The carriers, if present in the circumstellar envelope of these stars are not found under the environmental conditions needed to excite the transitions which we identify as DBs in the ISM.  The radial velocity analysis of the features observed in individual sources confirm this result, as the Doppler shifts measured are always found to be consistent with an interstellar origin for the bands observed.  DB carriers may be carbonaceous species or radicals attached to large organic molecules, trapped in lattice or more complex structures, or constituents of the mantle of circumstellar dust grains which are liberated to the ISM only after strong UV irradiation (either UV photons from the central star or from the more energetic interstellar UV field).  In this sense, the identification of the carriers as strongly ionised PAHs and/or radicals liberated from carbonaceous species as a consequence of photo-evaporation of dust grains in the ISM looks tempting and would be consistent with our observations.  However, we do not find any evidence of the carbonaceous nature of the carrier(s) in our sample stars, something generally accepted in the literature, nor any correlation with the presence of PAHs in the mid-infrared spectrum of these sources, as it has been claimed by several authors in the past.  If DBs are connected with PAHs or with any other carbonaceous species such as the ones suggested in the introduction of this paper, their carrier(s) must form at a later stage, probably under different excitation conditions, once the envelope of the post-AGB star is totally diluted in the interstellar medium as a result of the expansion of the shell.   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "0711/0711.1375_arXiv.txt": {
        "abstract": "{}{We present multi-band light curves and % distances for five type Ia supernovae at intermediate redshifts, $0.18<z<0.27$.}{Three telescopes on the Canary Island of La~Palma, \\inttel, \\nottel, and \\jkt, were used for discovery and follow-up of type Ia supernovae in the $g'$ and $r'$ filters. Supernova fluxes were measured by simultaneously fitting a supernova and host galaxy model to the data, and then calibrated using star catalogues from the Sloan Digital Sky Survey.}{The light curve shape and colour corrected peak luminosities are consistent with the expectations of a flat $\\Lambda$CDM universe at the $1.5\\sigma$ level. One supernova in the sample, SN\\sndr, shows surprisingly large reddening, considering both that it is located at a significant distance from the core of its host ($\\sim4$ times the fitted exponential radius) and that the galaxy can be spectroscopically classified as early-type with no signs of ongoing star formation. }{} ",
        "introduction": "Measurements of type Ia supernovae (\\snia) have had a dra\\-ma\\-tic im\\-pact on the field of cosmology, and lead to the discovery of the acceleration of the Universe (\\cite{1998AJ....116.1009R,1999ApJ...517..565P}). Since then the number of well measured \\snia:s has increased considerably through the joint effort of several independent research groups. However, the majority of the search campaigns have either been looking for high redshift \\snia (\\cite{knop03,2006A&A...447...31A,riess07,astro-ph/0701041} and references therein) where the difference in the relative luminosity distance between different cosmological models increases, as shown by Goobar \\& Perlmutter (1995), or low redshift \\snia (\\cite{1996AJ....112.2438H,1999AJ....117..707R,jha06}) which are essential for both anchoring the Hubble diagram as well as learning more about \\snia properties. With the exception of the ongoing three year Sloan Digital Sky Survey \\sn campaign initiated in 2005 (\\cite{astro-ph/0708.2749,astro-ph/0708.2750}), the redshift region around $z\\sim0.2$ in the Hubble diagram has not been heavily explored.  This missing link in the \\snia Hubble diagram is of great importance as it corresponds to the epoch of dark energy domination. Thus it presents interesting opportunities for testing for a possible time evolution of the dark energy density with data-sets probably less plagued by systematic uncertainties. Effects such as \\snia brightness evolution or gravitational lensing are expected to be much less than at higher redshifts.  The current best estimates of cosmological parameters from \\snls (\\cite{2006A&A...447...31A}) and \\essence (\\cite{astro-ph/0701041}) are approaching the systematic uncertainties for these projects. This further increases the need of a better understanding of \\snia properties with the aim of calibrating them to even better precision.  Also, for these redshifts the rest frame ultraviolet region of the \\snia spectrum has been redshifted into the visible wavelength area, illustrated in Figure~\\ref{fig:intzfilters}, where the detector efficiency is higher and the atmosphere is transparent and less variable. \\begin{figure*}[!tb] \\centering \\includegraphics[height=\\textwidth,angle=-90]{figures/intzfilt.eps} \\caption{Overlap between rest frame $UBV$ (dashed) and observed, effective, \\inttel $g'$ and $r'$ (solid) filters for the mean redshift, $z=0.22$, of the \\sn sample presented here. The dotted curve shows the normalised spectral template of a type~Ia supernova (\\cite{2002PASP..114..803N}) at maximum. \\label{fig:intzfilters}} \\end{figure*} This is a part of the type~Ia spectrum that contains a large fraction of the total flux, and it is therefore very important to understand the intrinsic UV-properties for interpreting high redshift data.  Furthermore, the UV part of the spectrum is of particular interest as it has been suggested that a redshift evolution of the \\snia properties due to differences in the progenitor metallicity, if present, may be detectable in this wavelength region (\\cite{1998ApJ...495..617H,2000ApJ...530..966L}). Intermediate redshift \\sne are bright enough to allow high precision spectroscopy (\\cite{2008ApJ...674...51E}), and a large number of well observed \\snia at these redshifts may very well turn out to be essential for understanding rest frame $U$-band properties.  To address these issues, a pilot campaign for photometric and spectroscopic follow-up of intermediate redshift \\snia was launched in 1999. Here we present photometric follow-up of five \\snia from the 1999 campaign. The spectra of these \\sne have already been published in Balland~\\etal~(2006).  In \\secref{sec:searchfollow} we describe the search and the photometric follow-up. \\secref{sec:redphot} covers the reduction and photometric measurements, followed by a description of the light curve analysis and possible systematic errors in~\\secref{sec:analysis}. In \\secref{sec:discussion} the results are summarised and discussed. ",
        "conclusions": "We have presented $g'$ and $r'$-band light curves of five \\sne\\ at intermediate redshift $z\\sim0.2$ obtained with the \\inttel, \\nottel and \\jkt telescopes. The photometry was extracted by simultaneously fitting the \\sn flux and the background galaxy. Four \\sne were discovered and have data in at least one filter before rest frame $B$-band maximum. This constrains the fitted light curve and the light curve shape and colour corrected peak $B$-band magnitude could be estimated to $\\sim10\\,\\%$.  One \\sn of our sample (\\sndr) shows an unusual red colour, given that the \\sn is located far away from the core of its early-type host galaxy. Unfortunately the light curve sampling around its peak brightness is too poor to draw any extensive conclusions about the intrinsic colour properties of this object.  After stretch and colour correcting the peak $B$-band magnitudes, the data is consistent with the $\\Lambda$CDM cosmology at the $\\sim1.5\\sigma$ level.  \\begin{acknowledgement} The authors would like to thank David Rubin at UC Berkeley for assisting in parts of the analysis. Ariel Goobar and Vallery Stanishev would like to thank the G\\\"{o}ran Gustafsson Foundation for financial support. \\end{acknowledgement}  \\begin{appendix}"
    },
    "0711/0711.1696_arXiv.txt": {
        "abstract": "Non-linear evolution of density fluctuations yields secondary anisotropies in the cosmic microwave background (CMB), which are correlated with the same density fluctuations that can be measured by weak lensing (WL) surveys. We study the CMB-WL correlation using analytical models as well as $N$-body simulations. We show that an analytical model based upon the time derivative of non-linear matter power spectrum agrees with simulations. All-sky cosmic-variance limited CMB and WL surveys allow us to measure the correlation from non-linearity with high significance ($50\\sigma$) for $l_{\\rm max}=10^4$, whereas the forthcoming missions such as Planck and LSST are expected to yield only marginal detections. The CMB-WL correlation is sensitive to the time derivative of structure growth. We study how this property may be used to constrain the nature of dark energy. While the expected constraints are not very strong, they may provide a cross check of results from other observations. ",
        "introduction": "\\label{sec:intro} Secondary temperature anisotropies of the cosmic microwave background (CMB) provide invaluable information on the structure formation in the universe. The sources of anisotropies include the Sunyaev-Zel'dovich (SZ) effects by galaxy clusters as well as by reionization, the Rees-Sciama (RS) effect, and CMB lensing.  The forthcoming Planck satellite and ground-based observations such as Atacama Cosmology Telescope \\citep[ACT;][]{Kosowsky:2003} and South-Pole Telescope \\citep[SPT;][]{SPT:2004} are designed to measure temperature fluctuations at arc-minute scales, and thus are expected to detect some of the secondary anisotropies.  The Rees-Sciama effect is, in principle, a unique probe of the {\\it time-variation} of gravitational potential, as it is given by \\citep{SachsWolfe:1967, ReesSciama:1968} \\begin{equation} \\frac{\\Delta T(\\hat{\\bs{n}})}{T} = 2\\int_0^{r_*} dr \\Phi'(\\hat{\\bs{n}}r,r), \\label{eq:1} \\end{equation} where $\\Phi$ is the Newtonian potential, $\\hat{\\bs{n}}$ is a unit direction vector, $r$ is the conformal {\\it lookback} time, and $r_*$ is $r$ out to the photon decoupling epoch. The prime denotes $\\partial/\\partial r$, which is equal to $-\\partial/\\partial\\eta$, where $\\eta$ is the conformal time.  In the linear regime, $\\Phi'$ vanishes in a matter-dominated universe; however, when either curvature or dark energy dominates, $\\Phi$ decays in time and thus $|\\Phi|' > 0$. In the {\\it non-linear} regime, on the other hand, $\\Phi$ {\\it grows} in time, $|\\Phi|' < 0$. Therefore, one expects $\\Phi\\Phi'>0$ on large scales where density fluctuations are still linear, and $\\Phi\\Phi'<0$ on small scales where fluctuations are non-linear.  The auto-correlation of the RS effect, both linear and non-linear, has been studied \\citep{Seljak:1996,Tuluie.etal:1996,Cooray:2002a}. The cross-correlation between the {\\it linear} RS effect and large-scale structure traced by galaxies has also been studied \\citep[e.g.,][]{Crittenden:1995,Peiris:2000,Cooray:2002b,Afshordi:2004}, and detected for various galaxy surveys \\citep[e.g.,][]{Boughn:2003,Nolta:2003,Afshordi.etal:2004}.  Much less attention has been given to the correlation of the {\\it non-linear} RS effect and large-scale structure. As this effect changes the sign at the linear-to-non-linear transition scale, the signature of non-linearity is distinct.  The RS effect is not the only thing that is correlated with the large-scale structure. The SZ effects, as well as radio and infrared point sources, also trace the large-scale structure. Fortunately, multi-frequency data enable us to separate these contributions that have unique and specific frequency dependence.  In this paper, we calculate the cross-correlation of the RS effect and large-scale structure traced by weak lensing (WL) surveys. In particular, we simulate this directly using a $N$-body code for the first time.  We adopt the standard $\\Lambda$CDM model with $(\\Omega_{m0},\\Omega_{\\Lambda0}, \\sigma_8, h, n_s)$ $=(0.26, 0.74, 0.76, 0.7, 1)$, which is consistent with \\citet{spergel.etal:2006}. ",
        "conclusions": "\\label{sec:summary} We have studied the cross-correlation between CMB (the RS effect) and large-scale structure traced by WL. We have developed a simple analytical model based upon the time derivative of non-linear matter power spectrum, and tested its validity analytically with 3rd-order PT as well as numerically with $N$-body simulations.  We have shown that all-sky cosmic-variance limited CMB and deep (shallow) WL surveys can yield a $50\\sigma$ ($10\\sigma$) detection of the {\\it non-linear} CMB-WL correlation for $l_{\\rm max}=10^4$. The forthcoming surveys are not expected to yield significant detections. We expect $\\sim 1.5\\sigma$ from Planck+LSST and $0.1\\sigma$ from ACT+CFHTLS.  The change of the sign of $C_l$ at the cross-over, $l\\sim 800$, offers a unique probe of the nature of DE. Sensitivity of the CMB-WL correlation to DE turns out to be not very strong; however, as the direction of degeneracy on $w_0-w_1$ is different from that of the baryon oscillations and Type Ia supernovae, it may provide an independent cross-check of the results from these observations."
    },
    "0711/0711.3693_arXiv.txt": {
        "abstract": "% We investigate, by means of numerical simulations, the kinematics of elliptical-spiral merger remnants. Counterrotation can appear both in coplanar and in non-coplanar retrograde mergers, and it is mostly associated to the presence of a disk component, which preserves part of its initial spin. In turn, the external regions of the two interacting galaxies acquire part of the orbital angular momentum, due to the action of tidal forces. ",
        "introduction": "% The simulations analyzed are a small subset of the simulations realized in the framework of the GalMer Project \\citep{dimatteo07a, dimatteo07b}. % ",
        "conclusions": ""
    },
    "0711/0711.5006_arXiv.txt": {
        "abstract": "Black hole formation may occur if the spectrum of the curvature perturbation $\\zeta$ increases strongly as the scale decreases.  As no such increase is observed on cosmological scales, black hole formation requires  strongly positive running $n'$ of the spectral index $n$, though the running might only kick in  below the `cosmological  scales' probed by the CMB anisotropy and galaxy surveys. A concrete and well--motivated way of producing this running is through the running mass model of slow roll inflation. We obtain a new observational  bound $n'<0.026$ on the running provided by this   model, improving an earlier result by a factor two. We also discuss black hole production in more general scenarios. We show that the usual conditions $\\epsilon \\ll 1$ and $| \\eta | \\ll 1$ are enough to derive the spectrum ${\\cal P}_{\\zeta}(k)$, the introduction of higher order parameters $\\xi^{2}$ etc. being optional. ",
        "introduction": "The primordial curvature perturbation $\\zeta$ is only of order $10\\mfour$ on cosmological scales, but it might be of order 1 on smaller scales. Primordial black holes will then form as those scales enter the horizon, with possibly observable consequences. The purpose of this paper is to see to what extent the value of order 1 is reasonable, taking into account observational constraints and current thinking about the origin of $\\zeta$.  In Section II we see what is required for black hole formation, in terms of the spectral index $n(\\ln k)\\equiv d\\ln\\calpz/d\\ln k$, which specifies the scale-dependence of the spectrum of $\\zeta$. Averaged over the whole range of scales we need strongly increasing $n$ corresponding to running $n'\\sim 10\\mtwo$. Up to this point we assumed nothing about the origin of $\\zeta$.  In Section III we introduce the assumption that it  originates from the inflaton perturbation during slow-roll inflation (the standard paradigm). Within this paradigm, the only extant model giving the required running $n'$ is the running mass model, which typically makes $n'$ roughly constant hence  requiring  $n'\\sim 10\\mtwo$ on cosmological scales.  In Section IV we ask whether such a large value of the running is still permitted by current data, thereby updating an earlier work.  We find that it is.  The question then arises, whether black hole formation can still be achieved if $n'$ is negligible on cosmological scales, as might be required by future data.   In Sections V to VIII we show that black hole formation can indeed be achieved within the standard paradigm. Along the way, we are led to consider the standard paradigm in more detail than before.  In Section IX we depart from the standard paradigm, by allowing a curvaton-type mechanism to contribute to the curvature perturbation. We show that black hole formation can occur if there is a switch from the standard paradigm to a curvaton-type paradigm as we go up in scale, or vice-versa.  We conclude in Section X. ",
        "conclusions": "\\label{sconc}  The possibility of primordial black hole formation at the end of inflation has a long history, which was long overdue for an update. The update is needed partly because observation now requires {\\em on cosmological scales} a tilt far below $0.3$ (and with negative sign) and not too much running. It is also needed because the original paradigm, that the inflaton perturbation is entirely responsible for the curvature perturbation, is now only one possibility.  According to the standard paradigm, the curvature perturbation is generated during slow roll inflation from the vacuum fluctuation of the inflaton field. Within this  paradigm, the running-mass model provides a well-motivated way of achieving black hole formation. To form black holes, the  model  probably requires strong running on cosmological scales, $n'\\sim 0.01$.  Such  running is allowed by the data. If a value $n' \\sim 10\\mtwo$ on cosmological scales  is ruled out in the future,  $n'$ will still have to increase to $\\gsim 10\\mtwo$ in order to  form black holes. We saw that this may be achieved within the standard paradigm by a suitable potential. Alternatively, it might be achieved by a switch from the standard paradigm  to a curvaton-type paradigm,  or by a switch from the curvaton-type paradigm to the standard paradigm.  {}For observers, we would like to re-iterate and earlier conclusion \\cite{paris}, that an upper bound $n'<10\\mthree$, or detection, would have important implications and is a very worthwhile goal."
    },
    "0711/0711.1139_arXiv.txt": {
        "abstract": "We have analysed deep optical spectra of the `born-again' planetary nebula Abell~58 and its hydrogen-deficient knot, surrounding V605~Aql, which underwent a nova-like eruption in 1919.  Our analysis shows that the extinction towards the central knot is much higher than previously thought, with $c({\\rm H}\\beta)$=2.0.  The outer nebula is less reddened, with $c({\\rm H}\\beta)$=1.04.  We find that the outer nebula has a Ne/O ratio higher than the average PN value.  The electron temperature we derive for the central knot varies widely depending on the diagnostic used.  The [O~{\\sc iii}] nebular-to-auroral transition ratio gives a temperature of 20\\,800\\,K, while the ratio of the [N~{\\sc ii}] nebular and auroral lines give T$_{\\rm e}$=15\\,200\\,K.  The helium line ratios $\\lambda$5876/$\\lambda$4471 and $\\lambda$6678/$\\lambda$4471 imply temperatures of 350\\,K and 550\\,K respectively.  Weakly temperature-sensitive O~{\\sc ii} recombination line ratios imply similarly low electron temperatures.  Abundances derived from recombination lines are vastly higher than those found from collisionally excited lines, with the abundance discrepancy factor (adf) for O$^{2+}$ reaching 89 -- the second highest known value after that found for the hydrogen deficient knots in Abell~30.  The observed temperature diagnostics and abundances support the idea that, like Abell~30, the knot of Abell~58 contains some very cold ionised material.  Although the central star is carbon-rich (C/O$>$1), the knot is found to be oxygen-rich, a situation not predicted by the single-star `born again' theory of its formation.  We compare the known properties of Abell~58 to those of Abell~30, Sakurai's Object and several novae and nova remnants.  We argue that the abundances in the ejecta observed in A\\,30 and A\\,58 have more in common with neon novae than with Sakurai's Object, which is believed to have undergone a final helium flash.  In particular, the C/O ratio of less than unity and presence of substantial quantities of neon in the ejecta of both Abell~30 and Abell~58 are not predicted by very late thermal pulse models. ",
        "introduction": "\\begin{figure} \\epsfig{file=figures/a58_3.eps, width=8cm, height=8cm, angle=0} \\caption{Hubble Space Telescope image of Abell 58 (WFPC2, F658N filter isolating [N~{\\sc ii}] emission), showing the faint outer nebulosity and bright central knot.  North is up and east is to the left, and the field of view is approximately 70 arcsec on each side.  The position marks indicate the central knot; the point source within the nebula to the north-east is a field star.} \\label{A58fig} \\end{figure}  Abell~58 (V605~Aql) consists of a large (44 $\\times$ 36 arcsec) faint shell, with a brighter knot at its geometric centre (Figure~\\ref{A58fig}).  The knot is assumed to have formed in a nova-like outburst which peaked in 1919, by which time the central star had increased in brightness by five magnitudes.  A\\,58 is one of a very small number of objects which have been seen to undergo rapid stellar evolution, and is considered an `older twin' of Sakurai's object (V4334 Sgr).  Since the outburst, the central star has become heavily obscured, and its spectrum is now only visible in scattered light (Clayton et al. 2006).  The central knot is extremely hydrogen-deficient (Seitter 1985, Guerrero \\& Manchado 1996), and high-resolution spectroscopy by Pollacco et al. (1992) showed that its emission was blue-shifted relative to the main nebula by $\\sim$100\\,km\\,s$^{-1}$.  They proposed that the knot was one side of a bipolar collimated flow, the other side being obscured by dust.  The spectrum of A\\,58's central knot has been studied by Guerrero \\& Manchado (1996, hereafter GM96), who found He/H to be 1.24 and O/H to be 1.25$\\times$10$^{-2}$, by number.  However, they considered this value of O/H to be unreliably high due to probable shock excitation.  V605~Aql itself cannot be seen directly, as it is hidden by a thick dust torus, but stellar emission lines are seen via scattered light.  Clayton et al. (2006) derived surface abundances for the star from observations of its scattered light, finding them to be typical of Wolf-Rayet central stars of PNe, with $\\sim$55\\% helium and $\\sim$40\\% carbon by mass.  They estimated the star's surface temperature to be 95\\,000\\,K.  Abell 30 is a nebula containing similar but much older hydrogen-deficient knots, ejected from its central star $\\sim$1000 years ago (Reay, Atherton \\& Taylor 1983).  The knots of Abell\\,30 have been studied in detail by Wesson, Liu \\& Barlow (2003) and Ercolano et al (2003).  Empirical analysis and subsequent photoionisation modelling of the polar knots showed that they contain extremely cold ($\\lesssim$2000\\,K) ionised material.  The existence of knots of cold ionised material in `normal' planetary nebulae has been proposed as a way of reconciling the ubiquitous discrepancy whereby abundances derived from optical recombination lines (ORLs) are higher than those derived from collisionally excited lines (CELs) (Liu et al. 2000).  The `normal' nebula which shows the most extreme abundance discrepancy factor (that is, the ratio of the abundance of an element derived from recombination lines to that derived from collisionally excited lines) is Hf\\,2--2, for which the oxygen abundance derived from ORLs is 68 times higher than the value derived from ORLs (Liu et al. 2006).   This is intermediate between the typical discrepancy for oxygen of 3-5 (Wesson et al 2005) and the extreme value of $\\sim$700 seen in A\\,30 (Wesson et al. 2003).  The generally accepted explanation of the origin of the knots in A\\,30 and A\\,58 is that the central star, after the formation of its surrounding planetary nebula, underwent a very late thermal pulse (VLTP), which ejected freshly processed stellar material into the centre of the nebula (Iben et al. 1983, Herwig 2001).  However, the analysis of Abell~30's polar knots showed that they were oxygen-rich (Wesson et al. 2003), the opposite of the situation predicted by the born-again theory.  The observed C/O ratios are similar to those seen in PNe with large abundance discrepancies, such as NGC\\,6153 (Liu et al. 2000).  This paper presents an analysis of long-slit spectra of A\\,58's central knot.  We find that its properties are very similar to those found for A\\,30's knots and consistent with the Ercolano et al. (2003) model whereby a cold, extremely hydrogen-deficient core contributes essentially all the strong ORL emission, and a hot outer shell emits CELs.  As with A\\,30, the knots are found to be oxygen-rich.  We discuss the implications of these findings for the born-again theory, and for planetary nebulae in general. ",
        "conclusions": "Of the five known hydrogen-deficient planetary nebulae (A\\,30, A\\,58, A\\,78, IRAS 15154-5258 and IRAS 18333-2357), we have now carried out detailed ORL/CEL abundance studies of two of them.  In both cases, we find that the VLTP born-again scenario commonly invoked to account for the production of H-deficient material within an old planetary nebula, and which predicts C/O$>$1, cannot account for the abundance ratios observed, while the collimated outflows seen in both objects seem inconsistent with a single star at the centre.  The presence of knots of cold hydrogen-deficient material has commonly been invoked to account for the observed discrepancy in planetary nebulae whereby ORLs give much higher chemical abundances than CELs for heavy elements (e.g. Liu et al. 2000, Tsamis et al. 2004, Wesson et al. 2005).  The origin of this H-deficient material in normal nebulae is as yet unknown.  Given the uncertainty at the moment about how the knots in A\\,30 and A\\,58 have been produced, it is difficult to say if the proposed knots present in `normal' nebulae could have a similar origin.  However, we note that it is difficult for the evolution of a single star to explain the morphology and abundances in A\\,30 and A\\,58, and that recent work has suggested that most or all central stars of planetary nebulae could be binary systems (Moe and De Marco 2006).  However, spatially resolved spectroscopy and high resolution imaging of nebulae with high abundance discrepancy factors, such as NGC 6153, seem to suggest that knots in `normal' nebulae must be very small, $<$60\\,AU across, numerous, and have a smooth distribution, as no clumps are seen in HST images, nor spikes in long slit spectra, but rather a smooth decline of adf from centre to edge, for example in NGC 6153 (Liu et al. 2000).  This is in contrast to the highly clumpy distribution of H-deficient material seen in Abell 30, Abell 58 and Abell 78."
    },
    "0711/0711.4693_arXiv.txt": {
        "abstract": "We have conducted a wide-field photometric survey in a single 52$'$$\\times$52$'$ field towards the Lupus Galactic Plane in an effort to detect transiting Hot Jupiter planets. The planet Lupus-TR-3b was identified from this work. The dataset also led to the detection of 494 field variables, all of which are new discoveries. This paper presents an overview of the project, along with the total catalog of variables, which comprises 190 eclipsing binaries (of contact, semi-contact and detached configurations), 51 miscellaneous pulsators of various types, 237 long period variables (P$\\ge$2d), 11 $\\delta$ Scuti stars, 4 field RR Lyrae (3 disk and 1 halo) and 1 irregular variable. Our survey provides a complete catalog of W UMa eclipsing binaries in the field to V$=$18.8, which display a Gaussian period distribution of 0.277$\\pm$0.036d. Several binary systems are likely composed of equal mass M-dwarf components and others display evidence of mass transfer. We find 17 candidate blue stragglers and one binary that has the shortest period known, 0.2009d (V$=$20.9). The frequency of eclipsing binaries (all types) is found to be 1.7$\\pm$0.4$\\times$10$^{-3}$ per star, substantially higher (by a factor of 3-10) than previously determined in the haloes of the globular clusters 47 Tuc and $\\omega$ Cen. This indicates that cluster dynamics aids mass segregation and binary destruction. ",
        "introduction": "The discovery rate of field variable stars has dramatically increased in recent years due to the advent of large-scale photometric surveys with wide fields of view. Many such surveys are dedicated to the detection of transiting short period planets or microlensing events (for example \\citet{Woz2004} and references there-in). They are perfectly suited to the detection of variable stars \\citep{Sos2006,Alc2003}, as they incorporate long temporal baselines, high resolution imagery, and achieve simultaneous high precision photometry for tens of thousands of stars.  Variable stars provide important information on the frequency, nature and evolution of stellar variability (of all  types) in various regions of the Galaxy. Binary stars provide information on stellar mass-radius relationships across many orders of magnitude, particularly for the little-understood relation for the lowest mass stars. Pulsating variables (such as $\\delta$ Scuti stars) offer insights into the internal structure and evolution of main sequence and post-main sequence objects. RR Lyrae and Cepheids are often used as distance indicators and their identification (particularly in stellar clusters) provides additional distance information for comparing stellar properties with that provided by theoretical isochrones.  Unlike microlensing surveys, transit surveys have temporal resolution of several minutes, permitting the detection of variability on very short time-scales, as well as longer period variability depending on dataset properties. \\citet{Kane2005} present an example of a variable star catalog from a wide angle transit survey in the general field, containing the types of variability commony found. Similarly, \\citet{W2004,W2007a} present the variables identified during a transit survey of the haloes of 47 Tucanae (47 Tuc) and $\\omega$ Centauri ($\\omega$ Cen). Open clusters have also been targeted (\\citealt{PB2006} and references there-in) and the Permanent All Sky Survey \\citep{Deeg2004} has the goal of permanently tracking variable stars in the whole sky with high temporal resolution.  This paper presents the total catalog of 494 variable stars identified during a deep, wide single-field survey for transiting Hot Jupiter planets ($\\sim$Jupiter-mass planets with orbital periods of a few days) towards the Lupus Galactic Plane. In addition to searching for transiting planets, this survey acts as a control field for our previous 47 Tuc and $\\omega$ Cen transit surveys, both of which produced significant null results. One object, Lupus-TR-3b, that is almost certainly a transiting low-mass Hot Jupiter has been identified in our survey \\citep{W2007b}. The survey also serves as an excellent test into the techniques and strategies for use in the soon-to-begin 5.7deg$^{2}$ SkyMapper transit survey \\citep{BS2007}. SkyMapper has the capacity to discover dozens of new transiting planets and thousands of variable stars in the near future.  Here we present the properties of all 494 variables, and illustrative V+R lightcurves for the majority of the catalog. We detail a preliminary analysis on their occurrence rates and likely nature, particularly in comparison to previous work in the fields of the globular clusters 47 Tuc and $\\omega$ Cen \\citep{W2004,W2007a,W2007c}. We find a significant difference in the occurrence rate for binary stars, with our field having three times the observed binary frequency of $\\omega$ Cen and ten times that of 47 Tuc. This supports mass segregation, particularly for 47 Tuc, and overall binary destruction in globular clusters due to cluster dynamics. Several of our binary systems are likely composed of M dwarf components, of interest in the study of low mass stars, and several other binaries display evidence of mass transfer, important for studies into binary system evolution and stellar interactions. We have also identified 17 blue straggler candidates and a binary which very likely contains a $\\delta$ Scuti component.  Section 2 of this paper details the observational strategy and data reduction techniques. Section 3 describes the production of the stellar time-series, the removal of dataset systematics, resulting photometric accuracy, and stellar colors and astrometry.  Section 4 details the variability search methods and dataset completeness. Section 5 presents the variable catalog, detailing each main type of variable and provides the derived binary frequency in our field, with a comparison to the globular clusters and distance estimates for our four RR Lyrae stars. We conclude in Section 6. ",
        "conclusions": "We present a variable star catalog, along with a preliminary analysis of the stars, detected during a long-baseline (53 nights: 26 contiguous nights in June 2005 and 27 in June 2006) high temporal resolution (6-7 minutes) photometric survey to detect transiting short period planets towards Lupus. A total of 494 variable stars were detected in our 0.66 deg$^{2}$ field, all of which are new discoveries. The catalog comprises 190 eclipsing binaries (of contact, semi-detached and detached configurations), 51 miscellaneous pulsators, 237 long period variables, 11 $\\delta$ Scuti stars, 4 field RR Lyrae and a single irregular variable.  We determine a period distribution of 0.277$\\pm$0.035d for our detected short period W UMa binaries (to V$=$18.8). Several binaries appear to be composed of equal-mass M-dwarf components, and others display evidence of mass transfer. We detected 17 candidate blue straggler stars and we also detected a binary with period 0.2009d (V=20.9), which constitutes the shortest period eclipsing binary known. Our brightest binary likely harbors a $\\delta$ Scuti component.  We have used this catalog to determine the occurrence frequencies of the various types of variable in our field, and compared the overall frequency of detected binary stars to that previously determined for the haloes of the globular clusters 47 Tuc and $\\omega$ Cen. We find that the frequency of binary stars is significantly larger in our field than that found in either cluster (ten times for 47 Tuc and three times for $\\omega$ Cen). This favors the scenario in which binaries are readily mass segregated into the cluster cores and/or destroyed by cluster dynamical processes."
    },
    "0711/0711.4977_arXiv.txt": {
        "abstract": "We examine the accuracy of strong gravitational lensing determinations of the mass of galaxy clusters by comparing the conventional approach with the numerical integration of the fully relativistic null geodesic equations in the case of weak gravitational perturbations on Robertson-Walker metrics.  In particular, we study spherically-symmetric, three-dimensional singular isothermal sphere models and the three-dimensional matter distribution of \\citet{NFW}, which are both commonly used in gravitational lensing studies.  In both cases we study two different methods for mass-density truncation along the line of sight: hard truncation and conventional (no truncation). We find that the relative error introduced in the total mass by the thin lens approximation alone is less than $0.3\\%$ in the singular isothermal sphere model, and less than $2\\%$ in the model of \\citet{NFW}.  The removal of hard truncation introduces an additional error of the same order of magnitude in the best case, and up to an order of magnitude larger in the worst case studied. Our results ensure that the future generation of precision cosmology experiments based on lensing studies will not require the removal of the thin-lens assumption, but they may require a careful handling of truncation. ",
        "introduction": "\\label{intro:sec}  In the current age of precision cosmology, the fundamental parameters of the favored cosmological model can be measured with high accuracy.  Strong and weak gravitational lensing studies are particularly important in this endeavor as an independent and relatively assumption-free measure of mass morphology in clusters of galaxies.  Conventional gravitational lensing studies continue to utilize the standard assumption of thin lenses.  This allows for the actual path of photons from source to observer to be approximated by its asymptotics, with a sharp bending at the plane where the lens is located.  There is, clearly, some error involved in this approximation, but the size of this error has not been rigorously examined, to our knowledge.  A question that will be important as the accuracy increases is at what point the error in the conventional approach becomes comparable to the uncertainty in the cosmological parameters.  At that point, accurate modeling of lensing events  will require a more accurate version of the lens equation, or an efficient algorithm to calculate the photon paths. In this respect, seminal proposals that improve on the conventional approach exist, dating from the 1990's.  In the work of \\citet{pyne}, higher-order lens equations in cosmological backgrounds are developed.  A different style of approximation which correctly utilizes the thin-lens paths as the zeroth-order path is introduced in \\citet{knp1,knp2}.  The same motivation sparked some interest in revisiting the foundations of gravitational lensing theory with the aim of developing the concept of the lens equation without reference to a lens plane, as in \\citet{efn,fn} and \\citet{perlick}.  These works develop a consistent theory of strong lensing entirely on the basis of the null geodesics in an arbitrary spacetime. Image distortion and weak gravitational lensing were considered from this perspective in \\citet{fkn2,fkn3}. An application of this no-lens-plane approach in (non-cosmological) spacetimes representing spherically symmetric non-singular matter distributions is presented in \\citet{kn}. The case of Schwarzschild black holes is developed in \\citet{fkn1} and the resulting lensing predictions are compared to \\citet{ve}, where a lens-plane approximation is applied to lightrays that undergo large bending.  Large bending angles by static and spinning black holes are also treated with an approximation scheme in  \\citet{bz,bz1}, with an application to Sag A$^*$.  The purpose of this paper is to test the accuracy of the conventional approach in determining the total mass of clusters. For this purpose we integrate the null geodesic equations in cosmological lensing spacetimes. We work with actual or realistic strong-lensing scenarios in the standard cosmology, assuming matter distributions of either a three-dimensional singular isothermal sphere (SIS) model or the model of \\citet{NFW} (NFW). Both these models are unrealistic to the extent that the mass density does not fall off fast enough to allow for a bounded mass in all space. In principle, some way of truncating the mass density is implicit in the model, although the details of the truncation mechanism are not considered of fundamental relevance to the predicted observables.  Often in the application of gravitational lensing to astrophysical data, truncation is imposed only in the directions transverse to the line of sight. This practice is justified on the basis that the projected mass density required for the thin-lens approximation is actually well defined in the model, as it involves only a one dimensional integration of the mass density along the line of sight, and therefore, truncation of the model along the line of sight is not necessary. Naturally, this practice introduces further error in addition to the error involved in the use of the thin-lens assumption.  The accuracy of this conventional model (thin-lens plus lack of truncation along the line of sight) is the issue in our current study.  Our method consists of comparing the observables predicted by the conventional approach to those predicted by the theory with no approximations. The theory is provided by the bending of light in the presence of mass as described by general relativity. The physical scenario common to both problems is the mass-density model. This is used differently in both approaches.  In the conventional approach, the mass density is only used after integration along the line of sight, with no truncation.  In the relativistic approach, which here we consider to be the exact treatment of the problem, the mass density is used as a source of the Newtonian potential that determines the components of the metric whose null geodesics are sought in exact form. The general relativistic (or exact) approach does not make sense in the absence of truncation along the line of sight.  Therefore, our study ends up comparing the results of a non-truncated thin-lens model to a properly truncated general-relativistic numerical integration, both of which represent the same physical situation.  For a proper interpretation of our results, it is important to emphasize that the non-truncated thin-lens models are conventionally intended to represent, with some accuracy, the physical situation of a properly truncated mass density, and for this reason the comparison is entirely justified.  Clearly, however, the comparison of the conventional approach with the relativistic treatment provides a measure of the total error incurred by two separate mechanisms: the thin lens approximation, and the lack of proper truncation along the line of sight.  The error introduced by the truncation mechanism, or the lack of it, into the thin-lens scheme can be analyzed as an end in itself. In this respect, progress has been made recently in introducing better physical truncation schemes that improve over hard truncation, including \\citet{baltz} and \\citet{takada}. Of importance, \\citet{baltz} considers differences between truncated and non-truncated lensing observables, giving a measure of the inaccuracy introduced by the truncation mechanism into the thin-lens scheme.  These studies take the properly truncated thin-lens scheme as the exact model of the physical situation, thus they fail to properly account for the actual inaccuracy with respect to the general relativistic theory. The size of the truncation error by itself within the thin-lens scheme has little meaning if the actual error introduced by the thin-lens assumption is not known.  In order to give an idea of the sizes of both errors (thin-lens and lack of truncation) with respect to the relativistic theory, we also compare our integration of the null geodesics with thin-lens models truncated along the line of sight.  In this case, the inaccuracy of the model is entirely due to a single assumption: that of thin lenses. Our mechanism of truncation is as simple as possible: a hard cut off of the mass density after an arbitrarily chosen radius.  In Section~\\ref{back:sec}, we outline the ways the Robertson-Walker (RW) background geometry enters our calculation, and in Section~\\ref{eqns:sec}, we outline the equations of motion.  The SIS and NFW gravitational potentials and thin-lens models are developed in section~\\ref{models:sec}.  Section~\\ref{sim:sec} describes our experimental set-up, numerical approach, and numerical accuracy. Our main results comparing the accuracy of the thin-lens models are presented in section~\\ref{compare:sec}.  Section~\\ref{cosmo:sec} presents the variation of thin-lens accuracy as a function of cosmology parameters for a range of flat cosmologies.  We discuss and explain our results in section~\\ref{discussion:sec} before concluding with some comments about our results and the their",
        "conclusions": "\\label{conclusion:sec}  The principal finding of this paper is that, in the context of the current generation of precision cosmology experiments, the use of the conventional approach in strong gravitational lensing is justified. Our figures and tables demonstrate that in a wide variety of strong lensing scenarios, and for two widely used mass models, the conventional approach works well in predicting cluster masses. Except in cases where the truncation radius is on the order of 1 to 2 Mpc, the errors in the SIS predicted mass are generally less than $2\\%$. The error in concentration parameter for the NFW model is also small except for very high redshift clusters.  This shows that concerns over the possibility of artificially high estimates of dark matter content in clusters as a result of the use of the thin-lens approximation are generally unfounded.  For the first time, to our knowlegde, we have been able to show that the thin lens assumption and the removal of the truncation radius along the line of sight contribute to the error in very different measures.  Interestingly, the two approximations introduce errors in opposite directions: the thin-lens assumption underestimates the mass, whereas the removal of the truncation radius overestimates it. The error introduced by removing the truncation along the line of sight is of the same order of magnitude as that of the thin-lens assumption in the best case, and up to an order of magnitude larger in the worst case within our study.   The truncation error is thus the dominant source of  error in the conventional approach.  This is a strong indication that, as precision increases in strong lensing observations, in order to gain a significant increase in accuracy it will be necessary, and perhaps even sufficient at first, simply to recover the truncation radius as an additional parameter in thin-lens models.  Truncation issues aside, we find that the error introduced by the thin lens assumpion alone for a given truncation method (in our case, hard truncation) varies between $0.03\\%$ and $0.3\\%$ in the SIS case, and between $0.3\\%$ and $2\\%$ in the NFW case for lenses at redshift less than 1.  Knowledge of the size of this error should be taken into account when considering additional approximations such as proper truncation methods.  Errors of this size may become significant in the next 20 years as the observational precision increases.  Perhaps more important than the specific strong lensing finding in this paper is the possible implication of this finding for weak lensing surveys of galaxy clusters.  Strong lensing depends of the first derivative of the gravitational potential (the Christoffel symbols of the underlying metric), but weak lensing depends on the second derivative (the curvature).  One may advance the conjecture that the higher order (derivative) measurements of weak lensing might be less sensitive to the thin-lens approximation, and therefore more accurate in the determination of the cluster mass.  At the time at which the fundamental parameters are known to $1\\%$ or better, this paper develops the appropriate steps for pursuing strong gravitational lensing without lens planes.  In particular, we have demonstrated gravitational lensing by weak gravitational potentials in the appropriate cosmological backgrounds by numerical integration of the null geodesic equations of general relativity to be a feasible task.  A long-standing argument for the use of the thin-lens approximation and against integrating the null geodesic equations has always been the high computational cost of numerical integration compared with the low cost of the (essentially algebraic) thin-lens approach. Ten years ago this was essentially true. However, the current speed and memory capacity of even small workstations removes this argument against using the general relativistic equations instead of the thin-lens method or other higher order approximations such as those proposed by \\citet{pyne} or \\citet{knp1}."
    },
    "0711/0711.2027_arXiv.txt": {
        "abstract": "{The onset of the thermonuclear runaway in a Chandrasekhar-mass white dwarf, leading to the explosion as a type Ia supernova, is studied with hydrodynamical simulations. We investigate the evolution of temperature fluctuations (``bubbles'') in the WD's convective core by means of 2D numerical simulations. We show how the occurrence of the thermonuclear runaway depends on various bubble parameters. The relevance of the progenitor's composition for the ignition process is also discussed. ",
        "introduction": "\\label{intro}  The transition from the hydrostatic to the explosive C-burning in a carbon-oxygen white dwarf (CO WD), which accretes mass in a binary system until it approaches the Chan\\-dra\\-se\\-khar mass, marks the start of the explosion of this progenitor star as a type Ia supernova (SN Ia). From a computational point of view, the modelling of this phase is extraordinarily challenging because of the span of involved length and time scales at ignition.  Nonetheless there are several studies on the evolution of the progenitor to the ignition of SN Ia explosion. \\citet{lht06} study the accretion phase of the progenitor and find interesting links between binary population synthesis and the physical conditions of the WD at ignition. On shorter time intervals, the last stage before the thermonuclear runaway has been explored with analytical models \\citep{wwk04,ww04} and numerical simulations \\citep{hs02,kwg06}. Though widely different in their numerical approach, these studies agree in identifying the convective flow prior to the runaway as a crucial issue for the ignition process and the early stage of the explosion.  A complementary way to address the ignition problem is given by small-scale (length scales of the order of $1$\\ km or less, whereas the WD diameter is about $2000$\\ km) investigations of the evolution to runaway (e.g.~\\citealt{gsb05}, \\citealt{zd07}). According to the models of \\citet{gsw95} and \\citet{ww04}, the thermonuclear runaway is triggered by temperature perturbations (``bubbles'') on the km-scale, generated in the WD core by the turbulent convection. The study of the bubble features is therefore a powerful tool to better understand the ignition process.  In this contribution, the numerical simulations (Sect.~\\ref{setup} and \\ref{simulations}) of \\citet{ibh06} (henceforth I06) are reviewed (we refer to that paper for further details). In Sect.~\\ref{discussion} we extend the discussion of the results to assess the role of the WD composition and the relevance of new nuclear measurements \\citep{srr07}. ",
        "conclusions": "\\label{discussion}  \\subsection{The parameter study and its implications for SN Ia ignition}  The role of the bubble parameters on the physics of temperature perturbations has been explored with the performed parameter study. The main results are reported here:  \\begin{itemize} \\item The bubble diameter $D$ affects the buoyant rise time because $\\tau_{\\mathrm{brt}} \\propto D^{1/2}$ ($v_{\\mathrm{b}} \\propto D^{1/2}$ as well, as in \\citealt{g02}). Larger bubbles have longer $\\tau_{\\mathrm{brt}}$ and are therefore favoured to go to thermonuclear runaway.  \\item The bubble temperature $T$ affects the nuclear timescale, because of the high sensitivity of the burning to $T$. Hotter bubbles are favoured to go to runaway. However, to some extent, also the buoyant rise time decreases with increasing temperature, because a larger temperature contrast leads to an increase of  $g_{\\mathrm{eff}}$, producing a more effective bubble dispersion.  \\item The increase of the central distance $R$ increases $g_{\\mathrm{eff}}$ and thus decreases the buoyant rise time, because both the bubble temperature contrast and $g(R)$ increase with increasing $R$. Bubbles at smaller distance from the WD centre are favoured to go to runaway. \\end{itemize}  The evolution of 2D bubbles is a highly simplified model for the small-scale ignition, nonetheless the results can be used to provide some clues on the ignition process in SNe Ia at larger length scales. First, we observe that the estimated convective velocity before the runaway ($50 - 100\\ \\mathrm{km\\ s^{-1}}$) is much larger than the bubble buoyant velocity. The features of the convective flow in the progenitor prior to runaway are therefore needed to model the ignition properly. Comparing convective velocities and timescales for bubble evolution, we estimate that the ignition is initiated at a central distance of about $150$\\ km, with a dispersion of the order of $100$\\ km, which agrees well with previous results \\citep{wwk04,ww04}.  The bubbles are generated in the WD's core with some probability distribution function of temperature fluctuations. Our bubble analysis indicates that the (few) hotter bubbles are more effectively dispersed than the (many) milder temperature fluctuations, because of their shorter bubble rise time. It suggests, indirectly, that the multi-point ignition scenario is a viable model for the onset of the SN Ia explosion.  \\subsection{Role of other parameters}  The study of the diversity in explosion simulations (e.g.~\\citealt{rgr06}), caused by the physical parameters of the WD and linked with the observed diversity in SNe Ia, has received comparatively more attention than the exploration of the diversity in the ignition phase. The bubble model described here has intrinsic and numerical limitations in its predictive power of the ignition properties of SNe Ia, but it can be profitably used as a ``probe'' of the influence on the ignition process of other physical parameters.  In \\citet{rgr06} the role of the WD composition is explored by varying the carbon/oxygen ratio and $^{22}\\mathrm{Ne}$ mass fraction (representative of the progenitor's metallicity, \\citealt{tbt03}). \\citet{dt06} found that the ignition timescale depends strongly on the $^{22}\\mathrm{Ne}$ abundance. We run some test simulations which confirm their results in the framework of the bubble physics. We suggest that the importance of both the C/O ratio and $X(^{22}\\mathrm{Ne})$ for the {\\em ignition} process should be carefully scrutinised in future works.  Recently \\citet{srr07} presented new measurements of the cross sections of the reactions $^{12}\\mathrm{C}(^{12}\\mathrm{C},\\alpha)^{20}\\mathrm{Ne}$ and $^{12}\\mathrm{C}(^{12}\\mathrm{C},\\mathrm{p})^{23}\\mathrm{Na}$ from $E = 2.10$\\ to $4.75$\\ MeV. The authors claim that the new reaction rates have interesting astrophysical applications. The improvement of the experimental determination of these reaction rates in the considered energy range (or at lower energies) could significantly influence the hydrostatic carbon burning in the progenitor and is worth being explored. The resonant screening effects should also be further investigated \\citep{itw03}."
    },
    "0711/0711.1815_arXiv.txt": {
        "abstract": "We analyze the Type II Plateau supernovae (SN II-P) 2005cs and 2006bp with the non-LTE model atmosphere code CMFGEN.  We fit 13 spectra in the first month for SN 2005cs and 18 for SN 2006bp.  {\\sl Swift} ultraviolet photometry and ground-based optical photometry calibrate each spectrum. Our analysis shows both objects were discovered less than 3 days after they exploded, making these the earliest SN II-P spectra ever studied. They reveal broad and very weak lines from highly-ionized fast ejecta with an extremely steep density profile. We identify He{\\,\\sc ii}\\,4686\\AA\\ emission in the SN 2006bp ejecta. Days later, the spectra resemble the prototypical Type II-P SN 1999em, which had a supergiant-like photospheric composition. Despite the association of SN 2005cs with possible X-ray emission, the emergent UV and optical light comes from the photosphere, not from circumstellar emission.  We surmise that the very steep density fall-off we infer at early times may be a fossil of the combined actions of the shock wave passage and radiation driving at shock breakout. Based on tailored CMFGEN models, the direct-fitting technique and the Expanding Photosphere Method both yield distances and explosion times that agree within a few percent. We derive a distance to NGC 5194, the host of SN 2005cs, of 8.9$\\pm$0.5\\,Mpc and 17.5$\\pm$0.8\\,Mpc for SN 2006bp in NGC 3953. The luminosity of SN 2006bp is 1.5 times that of SN 1999em, and 6 times that of SN 2005cs.  Reliable distances to Type II-P supernovae that do not depend on a small range in luminosity provide an independent route to the Hubble Constant and improved constraints on other cosmological parameters. ",
        "introduction": "Although radiation represents a negligible fraction of the energy in core-collapse supernova (SN) explosions, most of what we know about these events is inferred from the analysis of spectra and light curves. The gravitational binding-energy of the newly-formed protoneutron star is on the order of 10$^{53}$\\,erg. Within milliseconds after core bounce, a few $\\times$ 10$^{52}$\\,erg of this energy is used to photodissociate the infalling outer iron core. Starting with an electron-neutrino burst when the core reaches nuclear densities, the radiation of neutrinos of all flavors operates over a few tens of seconds after the bounce as the protoneutron star cools, and carries away a few $\\times$ 10$^{52}$\\,erg. These neutrinos are believed to deposit $\\sim$10$^{51}$\\,erg (1\\% of the total) into the infalling progenitor mantle, reversing the accretion, and to provide the internal and kinetic energy for the SN ejecta (Woosley \\& Janka 2005). Only $\\sim$10$^{49}$\\,erg (0.01\\% of the total) is eventually processed into light, that gets radiated with a rate equivalent to a few $\\times$ 10$^8$\\lsun\\, sustained for three months.  Depending on the progenitor radius, the shock emerges at the surface between a few hours and a day after core bounce, with a radiative precursor that is expected to heat and accelerate the surface layers, producing a soft X-ray flash (Chevalier 1982; Ensman \\& Burrows 1992; Matzner \\& McKee 1999; Blinnikov et al. 2000). Cooling, due to adiabatic expansion and radiative losses, is moderated as time progresses by recombination of hydrogen during the photospheric phase of Type II SN and by non-thermal excitation by unstable isotopes.  Although UV emission is significant for the first few days after shock breakout, the spectral energy distribution (SED) subsequently peaks further and further into the red, first in the visual and then in the near-IR, a few months after explosion (Kirshner et al. 1973,1975; Mitchell et al. 2002; Leonard et al. 2002a; Brown et al. 2007).  Core-collapse SN explosions constitute an excellent laboratory for inferring the properties of massive stars at the end of their lives and, indirectly, their pre-SN evolution. As the material expands, the photosphere recedes to deeper and deeper layers of the star's former envelope, allowing the observer to probe all the way from the surface at shock breakout to the innermost layers just outside the compact remnant. The nebular phase, the epoch when the ejecta become optically thin throughout ($\\sim$3-4 months after explosion for Type II SNe), permits the inspection of the regions where the blast originated, offering a means to constrain the mechanism and the morphology of the explosion as well as the nucleosynthetic yields.  Detailed quantitative analyzes focus on the early, photospheric, evolution (Ho\\\"{e}flich 1988; Eastman \\& Kirshner 1989; Schmutz et al. 1990; Baron et al 1996, 2000, 2003, 2004, 2007; Mitchell et al. 2002; Dessart \\& Hillier 2005a,2006ab), because the ejecta are best suited for ``standard'' model atmosphere computations: There is an optically-thick base where the radiation thermalizes; there is no contribution from non-thermal excitation by isotopes like $^{56}$Ni and its daughter product $^{56}$Co; electron scattering dominates the opacity; and line and continuum formation is spatially confined. Advances in computer technology permit better handling of non-Local-Thermodynamic-Equilibrium (non-LTE) effects and the critical role played by metals is now systematically accounted for (Baron et al. 2004; Dessart \\& Hillier 2005a). Non-LTE effects are important even at early times, for example, to reproduce the He{\\,\\sc i} lines observed in optical spectra.  Here, we present a quantitative spectroscopic analysis of the early photospheric-phase evolution of the Type II SNe 2005cs in NGC 5194 and SN 2006bp in NGC 3953, both with plateau-type light curves as shown in Fig.~\\ref{fig_lc_obs}. We employ the non-LTE model atmosphere code CMFGEN (Hillier \\& Miller 1998; Dessart \\& Hillier 2005a), in its steady-state configuration, following a similar approach to Dessart \\& Hillier (2006a; hereafter DH06) and their analysis of SN 1999em. A novel component of our study is the use of {\\sl Swift} (Gehrels et al. 2004) Ultra-Violet/Optical Telescope (UVOT; Roming et al. 2005) ultraviolet photometry that places important constraints on the SED in the UV, the early cooling of the photosphere, and the reddening. We infer from our models that SN 2005cs and SN 2006bp, were discovered only $\\sim$2 days after explosion, the earliest detections after explosion since SN 1987A (Type II peculiar) and SN 1993J (Type IIb; see, e.g., Matheson et al. 2000). As a result of this prompt discovery, we have the opportunity to observe a bright UV spectrum and to see high-excitation lines that have not previously been observed in SN II-P.  Finally, the fine time sampling of our spectra during the first month, along with dense photometric monitoring allows us to make an accurate distance determination. This is a good demonstration of the potential of Type II-P SNe, when coupled with detailed models, to provide accurate extragalactic distances.  Unlike the standard approach that depend on the Cepheid period-luminosity relation in the LMC and other nearby galaxies, our method is based on physical models fitted directly to observations, and the distances it provides are independent of all the lower rungs of the distance ladder.  We aim to assess the accuracy of distance determinations based on Type II-P SNe, and to refine the value of the Hubble constant $H_0$, in the spirit of the earlier efforts of Schmidt et al. (1994). Measurements of the total matter and baryonic content of the Universe via CMB experiments or large-scale structure correlations depend on the square of $H_0$ (Spergel et al. 2007; Eisenstein et al. 2005). In concordance solutions, these uncertainties affect measurements of the Dark Energy density and its associated equation-of-state parameter $w$ (Garnavich et al. 1998; Riess et al 2005; see Macri et al. 2006 for an illustration of the effect an improved accuracy on $H_0$ has on $w$). A small improvement in our knowledge of $H_0$ will lead to a significant improvement in the understanding the energy content of the Universe.  \\begin{figure*} \\epsscale{1.1} \\plottwo{f1a.ps}{f1b.ps} \\caption{ Observed light curves for SN 2005cs (left; optical data include the observations of T06, P06, and those of the CfA) and for SN 2006bp (right). Besides optical magnitudes in $UBVRI$ ($i'$ is shown instead of $I$ for SN 2006bp), we also include the {\\it Swift} UVOT magnitudes in the filters $uvw2$, $uwm2$, and $uvw1$ (corresponding values and photometric errors are given in Table~\\ref{tab_phot_05cs} \\& \\ref{table_swift_05cs} for SN 2005cs and Table~\\ref{tab_phot_06bp} \\& \\ref{table_swift_06bp} for SN 2006bp, for each filter and date). [See the electronic edition of the Journal for a color version of this figure, and see \\S\\ref{sect_obs_phot} and \\S\\ref{swift_obs} for discussion.\\label{fig_lc_obs}] } \\end{figure*} This paper is structured as follows. In the next section, we present the photometric and spectroscopic datasets. In \\S\\ref{sect_mod_pres}, we present the model atmosphere code CMFGEN and describe our methods. In \\S\\ref{sect_05cs}, we present a quantitative spectroscopic analysis at 13 photospheric-phase epochs for SN 2005cs, and in \\S\\ref{sect_06bp}, analyze 18 epochs for SN 2006bp. In \\S\\ref{sect_dist}, we apply the Expanding Photosphere Method (EPM) as well as a direct-fitting method (in the spirit of the Spectral-fitting Expanding Atmosphere Method, SEAM, of Baron et al. 2000) to infer the distance and the explosion date for both SN events. In \\S\\ref{sect_discussion}, we discuss the implications of our results, and in \\S\\ref{sect_conclusion}, we conclude. \\begin{deluxetable*}{lccccccccc} \\tablewidth{0pt} \\tabletypesize{\\scriptsize} \\tablecaption{Journal of photometric observations of SN 2005\\lowercase{cs}\\label{tab_phot_05cs}} \\tablehead{ \\colhead{UT Date} & \\colhead{JD} & \\colhead{U} & \\colhead{B} & \\colhead{V} & \\colhead{R} & \\colhead{I} & \\colhead{r'} & \\colhead{i'} & \\colhead{Source} \\\\ \\colhead{} & \\colhead{(+2,453,000)} & \\colhead{} & \\colhead{} & \\colhead{} & \\colhead{} & \\colhead{} & \\colhead{} & \\colhead{} & \\colhead{} } \\startdata 2005-06-30  &552.4 &13.48 $\\pm$  0.05  &14.36 $\\pm$  0.05  &14.48 $\\pm$  0.02  &14.46 $\\pm$  0.04 &14.44  $\\pm$ 0.04  &..               &..                 &  P06 \\\\ 2005-07-01  &553.2 &..                 &14.42 $\\pm$  0.02  &14.42 $\\pm$  0.02  &..                &..                 &14.68 $\\pm$ 0.03 &14.34  $\\pm$ 0.03  &  CfA \\\\ 2005-07-01  &553.3 &13.53 $\\pm$  0.06  &14.36 $\\pm$  0.05  &14.46 $\\pm$  0.05  &14.40 $\\pm$  0.02 &14.34  $\\pm$ 0.06  &..               &..                 &  P06 \\\\ 2005-07-02  &554.2 &13.50 $\\pm$  0.02  &14.47 $\\pm$  0.01  &14.41 $\\pm$  0.01  &..                &..                 &14.64 $\\pm$ 0.02 &14.32  $\\pm$ 0.02  &  CfA \\\\ 2005-07-02  &554.5 &13.69 $\\pm$  0.10  &14.45 $\\pm$  0.04  &14.51 $\\pm$  0.03  &14.47 $\\pm$  0.03 &14.29  $\\pm$ 0.10  &..               &..                 &  P06 \\\\ 2005-07-03  &555.2 &..                 &14.52 $\\pm$  0.02  &14.46 $\\pm$  0.02  &..                &..                 &..               &14.35  $\\pm$ 0.02  &  CfA \\\\ 2005-07-03  &555.3 &..                 &14.45 $\\pm$  0.02  &14.45 $\\pm$  0.02  &14.30 $\\pm$  0.02 &14.37  $\\pm$ 0.05  &..               &..                 &  T06 \\\\ 2005-07-04  &556.2 &..                 &14.55 $\\pm$  0.01  &14.52 $\\pm$  0.02  &..                &..                 &14.66 $\\pm$ 0.02 &14.36  $\\pm$ 0.02  &  CfA \\\\ 2005-07-05  &557.4 &..                 &14.52 $\\pm$  0.04  &14.54 $\\pm$  0.04  &14.42 $\\pm$  0.05 &14.29  $\\pm$ 0.08  &..               &..                 &  P06 \\\\ 2005-07-06  &557.8 &14.02 $\\pm$  0.06  &14.60 $\\pm$  0.05  &14.55 $\\pm$  0.03  &14.36 $\\pm$  0.03 &14.40  $\\pm$ 0.03  &..               &..                 &  P06 \\\\ 2005-07-06  &558.2 &13.87 $\\pm$  0.02  &14.61 $\\pm$  0.02  &14.52 $\\pm$  0.02  &..                &..                 &14.60 $\\pm$ 0.03  &14.33  $\\pm$ 0.03  &  CfA \\\\ 2005-07-06  &558.3 &..                 &14.57 $\\pm$  0.02  &14.54 $\\pm$  0.02  &14.34 $\\pm$  0.02 &14.29  $\\pm$ 0.03  &..               &..                 &  T06 \\\\ 2005-07-07  &559.2 &..                 &14.58 $\\pm$  0.03  &14.53 $\\pm$  0.01  &..                &..                 &14.59 $\\pm$ 0.02 &14.33  $\\pm$ 0.01  &  CfA \\\\ 2005-07-07  &559.4 &..                 &14.65 $\\pm$  0.04  &14.56 $\\pm$  0.09  &14.41 $\\pm$  0.09 &14.35  $\\pm$ 0.11  &..               &..                 &  P06 \\\\ 2005-07-08  &560.2 &..                 &..                 &14.54 $\\pm$  0.01  &..                &..                 &..               &14.23  $\\pm$ 0.01  &  CfA \\\\ 2005-07-08  &560.3 &..                 &14.72 $\\pm$  0.01  &..                 &..                &..                 &..               &..                 &  CfA \\\\ 2005-07-09  &561.3 &..                 &14.74 $\\pm$  0.02  &14.50 $\\pm$  0.02  &14.26 $\\pm$  0.02 &14.25  $\\pm$ 0.03  &..               &..                 &  T06 \\\\ 2005-07-10  &562.2 &14.64 $\\pm$  0.03  &14.84 $\\pm$  0.02  &14.52 $\\pm$  0.02  &..                &..                 &..               &..                 &  CfA \\\\ 2005-07-11  &563.1 &..                 &14.93 $\\pm$  0.01  &14.50 $\\pm$  0.02  &..                &..                 &..               &..                 &  CfA \\\\ 2005-07-11  &563.3 &..                 &..                 &14.50 $\\pm$  0.02  &..                &..                 &..               &..                 &  T06 \\\\ 2005-07-11  &563.4 &14.84 $\\pm$  0.05  &14.90 $\\pm$  0.04  &14.54 $\\pm$  0.04  &14.37 $\\pm$  0.02 &14.24  $\\pm$ 0.05  &..               &..                 &  P06 \\\\ 2005-07-12  &564.3 &..                 &14.92 $\\pm$  0.02  &14.50 $\\pm$  0.02  &14.32 $\\pm$  0.02 &14.20  $\\pm$ 0.03  &..               &..                 &  T06 \\\\ 2005-07-13  &565.4 &..                 &15.08 $\\pm$  0.05  &14.65 $\\pm$  0.04  &14.42 $\\pm$  0.04 &14.29  $\\pm$ 0.06  &..               &..                 &  T06 \\\\ 2005-07-13  &565.4 &15.27 $\\pm$  0.06  &15.04 $\\pm$  0.05  &14.64 $\\pm$  0.02  &14.37 $\\pm$  0.03 &14.29  $\\pm$ 0.05  &..               &..                 &  P06 \\\\ 2005-07-14  &566.3 &15.36 $\\pm$  0.06  &15.19 $\\pm$  0.02  &14.63 $\\pm$  0.02  &14.39 $\\pm$  0.02 &14.32  $\\pm$ 0.03  &..               &..                 &  T06 \\\\ 2005-07-14  &566.4 &15.29 $\\pm$  0.05  &15.10 $\\pm$  0.05  &14.67 $\\pm$  0.02  &14.42 $\\pm$  0.04 &14.27  $\\pm$ 0.02  &..               &..                 &  P06 \\\\ 2005-07-15  &567.3 &..                 &15.13 $\\pm$  0.03  &14.60 $\\pm$  0.02  &14.34 $\\pm$  0.02 &14.23  $\\pm$ 0.03  &..               &..                 &  T06 \\\\ 2005-07-17  &569.4 &15.92 $\\pm$  0.07  &15.31 $\\pm$  0.05  &14.67 $\\pm$  0.03  &14.40 $\\pm$  0.03 &14.27  $\\pm$ 0.04  &..               &..                 &  P06 \\\\ 2005-07-19  &571.3 &..                 &15.40 $\\pm$  0.02  &14.72 $\\pm$  0.02  &14.43 $\\pm$  0.02 &14.27  $\\pm$ 0.03  &..               &..                 &  T06 \\\\ 2005-07-19  &571.4 &..                 &15.34 $\\pm$  0.02  &14.72 $\\pm$  0.02  &14.38 $\\pm$  0.04 &14.22  $\\pm$ 0.04  &..               &..                 &  T06 \\\\ 2005-07-19  &571.4 &..                 &15.36 $\\pm$  0.12  &14.72 $\\pm$  0.05  &14.45 $\\pm$  0.09 &14.26  $\\pm$ 0.07  &..               &..                 &  P06 \\\\ 2005-07-20  &572.4 &..                 &15.39 $\\pm$  0.07  &14.71 $\\pm$  0.03  &14.45 $\\pm$  0.03 &14.27  $\\pm$ 0.03  &..               &..                 &  P06 \\\\ 2005-07-25  &577.4 &..                 &15.46 $\\pm$  0.09  &14.72 $\\pm$  0.04  &14.36 $\\pm$  0.06 &14.24  $\\pm$ 0.06  &..               &..                 &  P06 \\\\ 2005-07-27  &579.4 &..                 &15.60 $\\pm$  0.07  &14.73 $\\pm$  0.04  &14.39 $\\pm$  0.04 &14.16  $\\pm$ 0.04  &..               &..                 &  P06 \\\\ 2005-07-28  &580.3 &..                 &15.69 $\\pm$  0.05  &14.73 $\\pm$  0.03  &14.42 $\\pm$  0.03 &14.20  $\\pm$ 0.04  &..               &..                 &  T06 \\\\ \\enddata \\tablecomments{Sources for photometric observations are CfA, P06 and Tsvetkov et al. (2006; T06). For the distance determinations performed in \\S\\ref{sect_dist_05cs}, we use the photometry corresponding to the dates on which spectroscopic observations were obtained. If there is no photometry on the date of a spectrum, we linearly interpolate the slowly varying magnitude obtained at bracketing dates.} \\end{deluxetable*}    \\begin{deluxetable*}{lcccccc} \\tablewidth{18cm} \\tabletypesize{\\scriptsize} \\tablecaption{Journal of photometric observations of SN 2006\\lowercase{bp}\\label{tab_phot_06bp}} \\tablehead{ \\colhead{UT Date} & \\colhead{JD} & \\colhead{U} & \\colhead{B} & \\colhead{V} & \\colhead{r'} & \\colhead{i'} \\\\ \\colhead{} & \\colhead{(+2,453,000)} & \\colhead{} & \\colhead{} & \\colhead{} & \\colhead{} & \\colhead{} } \\startdata 2006-04-10 &836.27  &   ..                 & 15.59 $\\pm$  0.02    &   15.44 $\\pm$ 0.02  &   15.42 $\\pm$ 0.02  &     15.39 $\\pm$ 0.02 \\\\ 2006-04-11 &837.20  &   14.97 $\\pm$ 0.03   & 15.47 $\\pm$  0.02    &   15.31 $\\pm$ 0.02  &   15.24 $\\pm$ 0.02  &     15.23 $\\pm$ 0.02 \\\\ 2006-04-13 &839.20  &   14.92 $\\pm$ 0.03   & 15.44 $\\pm$  0.02    &   15.16 $\\pm$ 0.02  &   15.03 $\\pm$ 0.02  &     14.94 $\\pm$ 0.02 \\\\ 2006-04-14 &840.23  &   ..                 & 15.43 $\\pm$  0.02    &   15.12 $\\pm$ 0.02  &   14.97 $\\pm$ 0.02  &     14.91 $\\pm$ 0.01 \\\\ 2006-04-17 &843.24  &   15.04 $\\pm$ 0.03   & 15.46 $\\pm$  0.02    &   15.11 $\\pm$ 0.02  &   14.94 $\\pm$ 0.02  &     14.87 $\\pm$ 0.02 \\\\ 2006-04-18 &844.16  &   15.09 $\\pm$ 0.03   & 15.47 $\\pm$  0.02    &   15.12 $\\pm$ 0.02  &   14.93 $\\pm$ 0.02  &     14.86 $\\pm$ 0.02 \\\\ 2006-04-20 &846.17  &   15.20 $\\pm$ 0.02   & 15.52 $\\pm$  0.02    &   15.15 $\\pm$ 0.02  &   14.93 $\\pm$ 0.02  &     14.87 $\\pm$ 0.02 \\\\ 2006-04-22 &848.22  &   15.33 $\\pm$ 0.03   & 15.59 $\\pm$  0.02    &   15.18 $\\pm$ 0.02  &   14.94 $\\pm$ 0.02  &     14.90 $\\pm$ 0.02 \\\\ 2006-04-23 &849.23  &   15.42 $\\pm$ 0.03   & 15.63 $\\pm$  0.02    &   15.19 $\\pm$ 0.02  &    ..               &     14.91 $\\pm$ 0.02 \\\\ 2006-05-02 &858.27  &   16.44 $\\pm$ 0.03   & 16.01 $\\pm$  0.02    &   15.22 $\\pm$ 0.02  &   14.93 $\\pm$ 0.02  &      ..              \\\\ 2006-05-03 &859.19  &   16.55 $\\pm$ 0.04   & 16.07 $\\pm$  0.02    &   15.24 $\\pm$ 0.02  &   14.93 $\\pm$ 0.02  &     14.89 $\\pm$ 0.02 \\\\ 2006-05-05 &861.17  &   16.74 $\\pm$ 0.04   & 16.16 $\\pm$  0.02    &   15.27 $\\pm$ 0.02  &   14.95 $\\pm$ 0.02  &     14.88 $\\pm$ 0.02 \\\\ 2006-05-06 &862.16  &   16.79 $\\pm$ 0.05   & 16.21 $\\pm$  0.02    &   15.26 $\\pm$ 0.02  &   14.96 $\\pm$ 0.02  &     14.88 $\\pm$ 0.02 \\\\ 2006-05-11 &867.19  &   17.12 $\\pm$ 0.06   & 16.39 $\\pm$  0.03    &   15.32 $\\pm$ 0.02  &   15.00 $\\pm$ 0.02  &     14.89 $\\pm$ 0.02 \\\\ 2006-05-14 &870.20  &   ..                 & 16.48 $\\pm$  0.04    &   15.37 $\\pm$ 0.02  &   15.05 $\\pm$ 0.02  &     14.92 $\\pm$ 0.02 \\\\ 2006-05-30 &886.23  &   18.36 $\\pm$ 0.07   & 16.85 $\\pm$  0.02    &   15.52 $\\pm$ 0.02  &   15.11 $\\pm$ 0.02  &     14.98 $\\pm$ 0.02 \\\\ \\enddata \\tablecomments{ Log of CfA photometric observations. For the distance determination done in \\S\\ref{sect_dist_06bp}, we only use the $B$, $V$, and $i'$ magnitudes. We use  photometric data on the dates of the spectra, or, if there is no photometry on the day of a spectrum, linearly interpolate between magnitudes obtained at bracketing dates.} \\end{deluxetable*}   \\begin{deluxetable*}{ccccccccc} \\tablewidth{16cm} \\tabletypesize{\\scriptsize} \\tablecaption{Journal of spectroscopic observations of SNe 2005\\lowercase{cs} and 2006\\lowercase{bp} with FLWO 1.5\\,m+FAST\\label{Table:speclog}} \\tablehead{ \\colhead{UT Date{\\tablenotemark{a}}} & \\colhead{JD{\\tablenotemark{a}}} & \\colhead{Epoch{\\tablenotemark{b}}} & \\colhead{$\\Delta {\\rm PA}${\\tablenotemark{c}}} & \\colhead{Airmass} & \\colhead{Flux stds.{\\tablenotemark{d}}} & \\colhead{Seeing{\\tablenotemark{e}}} & \\colhead{Exp. time} & \\colhead{Observer{\\tablenotemark{f}}} \\\\ \\colhead{} & \\colhead{(+2,453,000)} & \\colhead{(d)} & \\colhead{($^\\circ$)} & \\colhead{} & \\colhead{} & \\colhead{($\\arcsec$)} & \\colhead{(s)} & \\colhead{} } \\startdata \\multicolumn{9}{c}{SN 2005cs} \\\\ \\hline 2005-06-30.23 & 551.73 &   4 &  1.1 & 1.31 & BD28, BD17 & \\nodata & 1200 & RH \\\\ 2005-07-01.27 & 552.77 &   5 &  5.0 & 1.53 & BD28, BD17 & 4   &  900 & RH \\\\ 2005-07-02.23 & 553.73 &   6 &  3.1 & 1.34 & BD28, BD26 & 2   & 1200 & RH \\\\ 2005-07-03.15 & 554.65 &   7 & 32.7 & 1.09 & F34, BD33 & 2   &  720 & RH \\\\ 2005-07-04.16{\\tablenotemark{g}} & 555.66 & 8 & 85.4 & 1.11 & BD28,BD17 & 1 & 1200 & JG \\\\ 2005-07-04.18{\\tablenotemark{g}} & 555.68 & 8 & 78.4 & 1.15 & BD28,BD17 & 1 & 1200 & JG \\\\ 2005-07-05.18 & 556.68 &   9 &  8.9 & 1.16 & F34, BD17 & 1   & 1200 & JG \\\\ 2005-07-06.17 & 557.67 &  10 &  5.2 & 1.15 & BD28, BD17 & 1   & 1200 & JG \\\\ 2005-07-07.15 & 558.65 &  12 &  1.8 & 1.11 & BD28, BD17 & 1   & 1200 & EF \\\\ 2005-07-09.19 & 560.69 &  13 &  0.7 & 1.24 & BD28, BD17 & 1-2 & 1200 & PB \\\\ 2005-07-10.18 & 561.68 &  14 &  1.1 & 1.22 & BD28, BD17 & 1-2 & 1200 & PB \\\\ 2005-07-11.18 & 562.68 &  15 &  2.9 & 1.23 & BD28, BD17 & 1-2 & 1200 & MC \\\\ 2005-07-12.16 & 563.66 &  16 &  2.8 & 1.17 & BD28, BD17 & 1-2 & 1200 & MC \\\\ 2005-07-28.17 & 579.67 &  32 &  1.0 & 1.38 & BD28, BD17 & 1-2 & 1200 & PB \\\\ 2005-07-29.15 & 580.65 &  33 &  2.1 & 1.30 & BD28, BD17 & 2-3 &  960 & MC \\\\ 2005-12-25.56 & 730.04 & 182 &  5.0 & 1.11 & BD28, BD17 & 1   & 1800 & MC \\\\ \\hline \\multicolumn{9}{c}{SN 2006bp} \\\\ \\hline 2006-04-20.30 & 845.80 & 12 & 12.3 & 1.13 & F34, HD84 & 1.5   & 1200 & WP \\\\ 2006-04-21.21 & 846.71 & 13 &  3.5 & 1.07 & F34, HD84 & 1.5   & 1200 & WP \\\\ 2006-04-24.19 & 849.69 & 16 & 67.9 & 1.08 & F34, BD26 & 1.5   & 1500 & JH \\\\ 2006-05-01.35 & 856.85 & 23 &  0.8 & 1.41 & F34, HD84 & 1.5   &  780 & WB \\\\ 2006-05-02.14 & 857.64 & 24 &  3.8 & 1.10 & F34, HD84 & 1.5-2 & 1200 & WB \\\\ 2006-05-03.22 & 858.72 & 25 & 64.4 & 1.08 & F34, HD84 & 1.5-2 & 1200 & TG \\\\ 2006-05-04.15 & 859.65 & 26 &  2.8 & 1.09 & F34, HD84 & 1.5-2 & 1200 & TG \\\\ 2006-05-07.19 & 862.69 & 29 & 12.2 & 1.07 & F34, BD17 & 1.5-2 & 1200 & WP \\\\ \\enddata  \\tablecomments{All the above spectra were obtained with the 1.5\\,m Tillinghast telescope using the FAST spectrograph (Fabricant et al. 1998). The grating used yields a resolution of $\\sim 7$\\AA. The observed wavelength range roughly spans 3500-7400\\,\\AA. (See \\S\\ref{sect_obs_spec} for discussion).}  \\tablenotetext{a}{UT Date and Julian Date at midpoint of observations.} \\tablenotetext{b}{Epoch of spectrum relative to the estimated explosion date.} \\tablenotetext{c}{Absolute difference between the observed position angle and average parallactic angle over the course of the observation.} \\tablenotetext{d}{Standard star pairs used for flux calibration: BD17 $=$ BD+17$^\\circ$4708; BD26 $=$ BD+26$^\\circ$2606; BD28 $=$ BD+28$^\\circ$4211; BD33 $=$ BD+33$^\\circ$2642; F34 $=$ Feige 34; HD84 $=$ HD 84937.} \\tablenotetext{e}{Seeing is based upon estimates by the observers.} \\tablenotetext{f}{Observers: EF $=$ E.~Falco; JG $=$ J.~Gallagher; JH $=$ J.~Hernandez; MC $=$ M.~Calkins; PB $=$ P.~Berlind; RH $=$ R.~Hutchins; TG $=$ T.~Groner; WB $=$ W.~Brown; WP $=$ W.~Peters.} \\tablenotetext{g}{Both spectra of SN~2005cs on UT 2005-07-04 were combined into a single spectrum for plotting and modeling purposes.}  \\end{deluxetable*} ",
        "conclusions": "\\label{sect_conclusion}  We have presented a multi-epoch multi-wavelength quantitative spectroscopic analysis of the two Type II-P SN events 2005cs and 2006bp, over the first two months after explosion, and have reached similar conclusions on the ejecta properties to those for SN 1999em (DH06). The progenitor stars have a composition compatible with a BSG/RSG progenitor. The spectral evolution shows a reddening of the SED with time, with a flux peaking in the UV at very early times, but progressively shifting to longer wavelengths. The photosphere recedes to deeper layers with time, reaching mass shells moving at about 2000-3000\\,\\kms about 6 weeks after explosion and having enhanced carbon and oxygen, and reduced nitrogen, abundances, compared to the surface layers. We find a density distribution that is very steep at early times, and we surmise that this could be an imprint of the shock breakout epoch. From about a week until six weeks after explosion, we find a more standard and constant density exponent of ten for the corresponding power law distribution, compared to a value of 50 (for SN 2006bp) and 20 (for SN 2005cs) used to model the first observation. {\\sl Swift} UVOT photometry and ground-based optical spectroscopy and photometry support a common origin for the UV and the optical light, so that we anticipate that if CSM interaction does occur, it does not contribute significantly to the observed UV and optical fluxes.  Using our analysis, we infer the distance and the explosion date of each SN. We find, in that order, JD 2453547.6$\\pm$0.5 days (2.8 days prior to discovery) and 8.9$\\pm$0.5\\,Mpc for SN 2005cs, and JD 2453833.4$\\pm$0.4\\,days (1.7 days prior to discovery) and 17.5$\\pm$0.8\\,Mpc for SN 2006bp. The large number of observations, the quality of our fits, the agreement between the two methods followed suggests a very high level of internal consistency of the distances determined here with Type II SNe. The actual level of accuracy on such determined distances, which we estimate to be on the order of 10\\% (related to the spectral fitting accuracy), would be best assessed by confrontation with alternate methods (Cepheids for example, although all alternate methods have their own inaccuracies, claimed and otherwise), the determination of distances to a galaxy host with well-recorded multiple Type II SN events (the typical Type II SN rate of 2 per century per galaxy makes this proposition somewhat unrealistic), a Galactic SN that could also be resolved with an interferometer or a large-aperture telescope. Our perspective on the uncertainty of our distance measurements is that, were we to increase the number of Type II SNe suitable for analysis in M 51 to 10 or 100, the scatter in the distance determined with our technique for that sample would not be larger than 10\\%. But this would still be no proof that the resulting average distance is indeed the right one since enlarging the sample would only allow for an internal consistency check on our method. In the short term, a Cepheid distance to the galaxy host of 2005cs and 2006bp is highly desirable, and would provide an independent check on the distance computed with Type II-P SNe, as was done with success in the past with SN 1999em (Leonard et al. 2003; Baron et al. 2004; DH06), giving further support to the potential use of such objects for distance determinations in the Universe.  This work fits within an ambitious project to build a general understanding of SN ejecta, to infer properties of the progenitor star (with ramifications for massive star evolution) and the explosion mechanisms at the origin of the event. A major ongoing theoretical development is the treatment of time-dependent effects both in the radiation field and in the level populations to allow for a direct modeling of the evolution over months after shock breakout, through the photospheric and into the nebular phases, and based on a range of hydrodynamical models of the explosion. In parallel, high-quality multi-epoch multi-wavelength observations for a substantial number of nearby SNe, irrespective of the type, galaxy host properties, covering from as soon as possible after explosion until the nebular phase, is eagerly sought."
    },
    "0711/0711.1212_arXiv.txt": {
        "abstract": "{We present $UBVR_{\\rm \\scriptstyle C}I_{\\rm \\scriptstyle C}$ magnitudes of 49 comparison stars in the fields of the Seyfert galaxies Mrk~335, Mrk~79, Mrk~279, Mrk~506, 3C~382, 3C~390.3, NGC~6814, Mrk~304, Ark~564, and NGC~7469 in order to facilitate the photometric monitoring of these objects; 36 of the stars have not been calibrated before. The comparison stars are situated in $5 \\times 5$ arcmin fields centred on the Seyfert galaxies, their $V$ band flux ranges from 11.7 to 18.2 mag with a median value of 16.3 mag, and their $B-V$ colour index ranges from 0.4 to 1.6 mag with a median value of 0.8 mag. The median errors of the calibrated $UBVR_{\\rm \\scriptstyle C}I_{\\rm \\scriptstyle C}$ magnitudes are 0.08, 0.04, 0.03, 0.04, and 0.06 mag, respectively. Comparison stars were calibrated for the first time in three of the fields (Mrk~506, 3C~382, and Mrk~304). The comparison sequences in the other fields were improved in various aspects. Extra stars were calibrated in four fields (Mrk~335, Mrk~79, NGC~6814, and NGC~7469)~-- most of these stars are fainter and are situated closer to the Seyfert galaxies compared to the existing comparison stars. The passband coverage of the sequences in five fields (Mrk~335, Mrk~79, Mrk~279, NGC~6814, and Ark~564) was complemented with $U$ band.} ",
        "introduction": "\\label{intro}  Variability is a common property of Active Galactic Nuclei (AGNs). Monitoring programmes undertaken till now provided a lot of information about variability characteristics of AGNs that was successfully used in studying the nuclear activity of galaxies (e.g. Schramm et~al. \\cite{schramm93}; Dietrich et~al. \\cite{dietrich98}; Shemmer et~al. \\cite{shemmer01}; Bachev \\& Strigachev \\cite{bachev04}; Bachev, Strigachev \\& Semkov \\cite{bachev05}; Villata et~al. \\cite{villata06}).  \\begin{table*}[t] \\caption{A list of the selected Seyfert galaxies. Equatorial coordinates, other names, redshifts, and Seyfert types of activity were taken from NASA/IPAC Extragalactic Database (NED). The number of the different observing nights, $N$, at which the data were taken is specified in the last column; the $U$ band data of Mrk~335, Mrk~79, Mrk~506, NGC~6814, and Mrk~304 were taken at a single epoch. The superscript $n$ means the number of nights used in the internal check for variability of the newly calibrated stars (see Sect.~\\ref{genchar}); the missing superscript means no new comparison stars were added to the corresponding fields.} \\label{gals} \\centering \\begin{tabular}{@{}llrrrrl@{}} \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} $\\rm Galaxy$ & $\\rm Other$ $\\rm names$ & $\\rm RA\\,(J2000)$ & $\\rm Dec\\,(J2000)$ & $z$ & $\\rm Sy$ & $N^n$ \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} $\\rm Mrk~335  $ &                                                    & $00$:$06$:$19.521$ & $ 20$:$12$:$10.49$ & $0.02578$ & $1.2$ & $2^3$ \\\\ $\\rm Mrk~79   $ & $\\rm UGC\\,03973,\\,CGCG\\,235-030,\\,MCG\\,+08-14-033$ & $07$:$42$:$32.797$ & $ 49$:$48$:$34.75$ & $0.02219$ & $1.2$ & $2^3$ \\\\ $\\rm Mrk~279  $ & $\\rm UGC\\,08823,\\,CGCG\\,336-028,\\,MCG\\,+12-13-022$ & $13$:$53$:$03.447$ & $ 69$:$18$:$29.57$ & $0.03045$ & $1.5$ & $1  $ \\\\ $\\rm Mrk~506  $ & $\\rm CGCG\\,170-020,\\,MCG\\,+05-41-012             $ & $17$:$22$:$39.899$ & $ 30$:$52$:$53.01$ & $0.04303$ & $1.5$ & $2^3$ \\\\ $\\rm 3C~382   $ & $\\rm CGCG\\,173-014                               $ & $18$:$35$:$03.391$ & $ 32$:$41$:$46.82$ & $0.05787$ & $1.0$ & $2^3$ \\\\ $\\rm 3C~390.3 $ & $\\rm VII\\,Zw\\,838                                $ & $18$:$42$:$08.990$ & $ 79$:$46$:$17.13$ & $0.05610$ & $1.0$ & $2  $ \\\\ $\\rm NGC~6814 $ & $\\rm MCG\\,-02-50-001                             $ & $19$:$42$:$40.644$ & $-10$:$19$:$24.57$ & $0.00521$ & $1.5$ & $3^4$ \\\\ $\\rm Mrk~304  $ & $\\rm II\\,Zw\\,175,\\,CGCG\\,428-065                 $ & $22$:$17$:$12.260$ & $ 14$:$14$:$20.90$ & $0.06576$ & $1.0$ & $2^4$ \\\\ $\\rm Ark~564  $ & $\\rm UGC\\,12163,\\,CGCG\\,495-018,\\,MCG\\,+05-53-012$ & $22$:$42$:$39.345$ & $ 29$:$43$:$31.31$ & $0.02468$ & $1.8$ & $2  $ \\\\ $\\rm NGC~7469 $ & $\\rm UGC\\,12332,\\,CGCG\\,405-026,\\,MCG\\,+01-58-025$ & $23$:$03$:$15.623$ & $ 08$:$52$:$26.39$ & $0.01632$ & $1.2$ & $2^4$ \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\end{tabular} \\end{table*}  \\begin{table*} \\caption{Apparent Johnson-Cousins $UBVR_{\\rm \\scriptstyle C}I_{\\rm \\scriptstyle C}$ magnitudes of the comparison stars, their errors (in parentheses), and the number of the individual estimates used to obtain the final magnitude. We list the other designations of the comparison stars that have been calibrated by other authors as well; the superscripts to these designations refer to the corresponding literature sources. The stars without other designations have been calibrated for the first time by us. Dots mean we have not been able to obtain the corresponding magnitudes (see Sect.~\\ref{comm} for comments).} \\label{naomag} \\centering \\begin{tabular}{@{}l@{\\hspace{3.2mm}}r@{\\hspace{3.2mm}}r@{\\hspace{3.2mm}}r@{\\hspace{3.2mm}}r@{\\hspace{3.2mm}}r@{\\hspace{3.2mm}}r@{\\hspace{3.2mm}}r@{}} \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} $\\rm Star$ & $\\rm RA\\,(J2000)$ & $\\rm Dec\\,(J2000)$ & $U$ & $B$ & $V$ & $R_{\\rm \\scriptstyle C}$ & $I_{\\rm \\scriptstyle C}$ \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\multicolumn{8}{c}{$\\rm \\textbf {Mrk~335}$} \\\\ $\\rm \\textbf {A}$:$\\;\\rm 1^{\\rm b},\\,B^{\\rm c},\\,4^{\\rm e} $ & $00$:$06$:$20.107$ & $20$:$10$:$50.59$ & $16.17\\;(0.08)\\;1$ & $15.42\\;(0.03)\\;2$ & $14.25\\;(0.03)\\;2$ & $13.70\\;(0.03)\\;2$ & $13.12\\;(0.06)\\;2$ \\\\ $\\rm \\textbf {B}                                           $ & $00$:$06$:$14.808$ & $20$:$11$:$34.20$ & $17.31\\;(0.08)\\;1$ & $17.38\\;(0.03)\\;2$ & $16.85\\;(0.03)\\;2$ & $16.55\\;(0.03)\\;2$ & $16.22\\;(0.05)\\;2$ \\\\ $\\rm \\textbf {C}$:$\\;\\rm 2^{\\rm b},\\,C^{\\rm c},\\,6^{\\rm e} $ & $00$:$06$:$17.971$ & $20$:$13$:$17.60$ & $15.40\\;(0.08)\\;1$ & $15.43\\;(0.03)\\;2$ & $15.00\\;(0.03)\\;2$ & $14.78\\;(0.03)\\;2$ & $14.55\\;(0.06)\\;2$ \\\\ $\\rm \\textbf {D}                                           $ & $00$:$06$:$17.135$ & $20$:$14$:$08.72$ & $16.28\\;(0.09)\\;1$ & $16.14\\;(0.04)\\;2$ & $15.38\\;(0.03)\\;2$ & $15.00\\;(0.04)\\;2$ & $14.59\\;(0.06)\\;2$ \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\multicolumn{8}{c}{$\\rm \\textbf {Mrk~79}$} \\\\ $\\rm \\textbf {A}                                     $ & $07$:$42$:$35.026$ & $49$:$47$:$01.59$ & $19.18\\;(0.09)\\;1$ & $18.84\\;(0.04)\\;2$ & $17.94\\;(0.06)\\;2$ & $17.48\\;(0.05)\\;2$ & $17.05\\;(0.11)\\;2$ \\\\ $\\rm \\textbf {B}                                     $ & $07$:$42$:$25.704$ & $49$:$48$:$07.52$ & $17.62\\;(0.09)\\;1$ & $17.60\\;(0.05)\\;2$ & $16.88\\;(0.05)\\;2$ & $16.53\\;(0.05)\\;2$ & $16.13\\;(0.08)\\;2$ \\\\ $\\rm \\textbf {C}$:$\\;4^{\\rm e}                       $ & $07$:$42$:$22.633$ & $49$:$48$:$17.49$ & $15.02\\;(0.10)\\;1$ & $14.98\\;(0.06)\\;1$ & $14.30\\;(0.05)\\;1$ & $13.98\\;(0.06)\\;1$ & $13.61\\;(0.07)\\;1$ \\\\ $\\rm \\textbf {D}                                     $ & $07$:$42$:$42.113$ & $49$:$49$:$46.54$ & $20.08\\;(0.09)\\;1$ & $18.66\\;(0.04)\\;2$ & $17.42\\;(0.06)\\;2$ & $16.70\\;(0.07)\\;2$ & $16.04\\;(0.07)\\;2$ \\\\ $\\rm \\textbf {E}                                     $ & $07$:$42$:$24.547$ & $49$:$49$:$54.38$ & $19.46\\;(0.09)\\;1$ & $18.39\\;(0.08)\\;2$ & $17.01\\;(0.04)\\;2$ & $16.12\\;(0.08)\\;2$ & $15.22\\;(0.06)\\;2$ \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\multicolumn{8}{c}{$\\rm \\textbf {Mrk~279}$} \\\\ $\\rm \\textbf {A}$:$\\;\\rm D^{\\rm c,f}                 $ & $13$:$53$:$21.781$ & $69$:$16$:$30.34$ & $16.37\\;(0.06)\\;1$ & $16.23\\;(0.05)\\;1$ & $15.53\\;(0.03)\\;1$ & $15.16\\;(0.05)\\;1$ & $14.82\\;(0.08)\\;1$ \\\\ $\\rm \\textbf {B}$:$\\;\\rm A^{\\rm c,f}                 $ & $13$:$53$:$21.294$ & $69$:$18$:$07.97$ & $14.07\\;(0.05)\\;1$ & $13.11\\;(0.05)\\;1$ & $12.06\\;(0.03)\\;1$ & $          \\ldots$ & $          \\ldots$ \\\\ $\\rm \\textbf {C}$:$\\;\\rm C^{\\rm c,f}                 $ & $13$:$52$:$54.690$ & $69$:$20$:$14.75$ & $16.69\\;(0.07)\\;1$ & $15.79\\;(0.05)\\;1$ & $14.78\\;(0.03)\\;1$ & $14.23\\;(0.06)\\;1$ & $13.81\\;(0.09)\\;1$ \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\multicolumn{8}{c}{$\\rm \\textbf {Mrk~506}$} \\\\ $\\rm \\textbf {A}                                     $ & $17$:$22$:$37.829$ & $30$:$50$:$10.71$ & $20.01\\;(0.06)\\;1$ & $19.13\\;(0.03)\\;2$ & $17.72\\;(0.02)\\;2$ & $16.90\\;(0.04)\\;2$ & $16.03\\;(0.05)\\;2$ \\\\ $\\rm \\textbf {B}                                     $ & $17$:$22$:$32.619$ & $30$:$51$:$23.33$ & $19.02\\;(0.05)\\;1$ & $18.23\\;(0.03)\\;2$ & $17.12\\;(0.02)\\;2$ & $16.49\\;(0.04)\\;2$ & $15.87\\;(0.04)\\;2$ \\\\ $\\rm \\textbf {C}                                     $ & $17$:$22$:$43.293$ & $30$:$52$:$00.91$ & $20.65\\;(0.05)\\;1$ & $19.49\\;(0.03)\\;2$ & $17.92\\;(0.02)\\;2$ & $16.86\\;(0.04)\\;2$ & $15.45\\;(0.04)\\;2$ \\\\ $\\rm \\textbf {D}                                     $ & $17$:$22$:$33.955$ & $30$:$52$:$02.37$ & $17.85\\;(0.05)\\;1$ & $16.62\\;(0.03)\\;2$ & $15.33\\;(0.02)\\;2$ & $14.59\\;(0.04)\\;2$ & $13.92\\;(0.04)\\;2$ \\\\ $\\rm \\textbf {E}                                     $ & $17$:$22$:$47.853$ & $30$:$52$:$12.09$ & $17.86\\;(0.06)\\;1$ & $18.03\\;(0.03)\\;2$ & $17.46\\;(0.02)\\;2$ & $17.14\\;(0.05)\\;2$ & $16.80\\;(0.05)\\;2$ \\\\ $\\rm \\textbf {F}                                     $ & $17$:$22$:$42.842$ & $30$:$52$:$59.42$ & $14.92\\;(0.05)\\;1$ & $15.05\\;(0.03)\\;2$ & $14.61\\;(0.02)\\;2$ & $14.38\\;(0.03)\\;2$ & $14.13\\;(0.04)\\;2$ \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\multicolumn{8}{c}{$\\rm \\textbf {3C~382}$} \\\\ $\\rm \\textbf {A}                                     $ & $18$:$35$:$06.547$ & $32$:$40$:$02.03$ & $18.32\\;(0.09)\\;2$ & $17.46\\;(0.05)\\;1$ & $16.42\\;(0.03)\\;2$ & $15.82\\;(0.04)\\;2$ & $15.37\\;(0.05)\\;2$ \\\\ $\\rm \\textbf {B}                                     $ & $18$:$34$:$57.385$ & $32$:$41$:$17.26$ & $17.37\\;(0.12)\\;2$ & $17.00\\;(0.05)\\;1$ & $16.27\\;(0.02)\\;2$ & $15.86\\;(0.03)\\;2$ & $15.57\\;(0.04)\\;2$ \\\\ $\\rm \\textbf {C}                                     $ & $18$:$35$:$01.184$ & $32$:$42$:$43.29$ & $17.86\\;(0.07)\\;2$ & $17.99\\;(0.05)\\;1$ & $17.41\\;(0.02)\\;2$ & $17.03\\;(0.02)\\;2$ & $16.77\\;(0.05)\\;2$ \\\\ $\\rm \\textbf {D}                                     $ & $18$:$35$:$05.136$ & $32$:$42$:$44.90$ & $17.72\\;(0.21)\\;2$ & $16.80\\;(0.05)\\;1$ & $15.76\\;(0.03)\\;2$ & $15.18\\;(0.02)\\;2$ & $14.75\\;(0.04)\\;2$ \\\\ $\\rm \\textbf {E}                                     $ & $18$:$35$:$01.718$ & $32$:$42$:$58.12$ & $17.34\\;(0.13)\\;2$ & $17.16\\;(0.05)\\;1$ & $16.51\\;(0.02)\\;2$ & $16.10\\;(0.02)\\;2$ & $15.80\\;(0.04)\\;2$ \\\\ $\\rm \\textbf {F}                                     $ & $18$:$35$:$07.086$ & $32$:$43$:$27.46$ & $16.30\\;(0.12)\\;2$ & $15.83\\;(0.05)\\;1$ & $15.00\\;(0.02)\\;2$ & $14.52\\;(0.02)\\;2$ & $14.16\\;(0.04)\\;2$ \\\\ $\\rm \\textbf {G}                                     $ & $18$:$34$:$57.009$ & $32$:$43$:$33.51$ & $15.95\\;(0.04)\\;2$ & $15.99\\;(0.05)\\;1$ & $15.55\\;(0.02)\\;2$ & $15.15\\;(0.02)\\;2$ & $14.87\\;(0.05)\\;2$ \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\multicolumn{8}{c}{$\\rm \\textbf {3C~390.3}$} \\\\ $\\rm \\textbf {A}$:$\\;\\rm A^{\\rm a},\\,1^{\\rm b}        $ & $18$:$42$:$23.706$ & $79$:$45$:$39.09$ & $13.65\\;(0.08)\\;2$ & $12.74\\;(0.04)\\;2$ & $11.73\\;(0.08)\\;2$ & $11.11\\;(0.05)\\;1$ & $10.66\\;(0.07)\\;1$ \\\\ $\\rm \\textbf {B}$:$\\;\\rm B^{\\rm a,f},\\,2^{\\rm b}      $ & $18$:$42$:$00.991$ & $79$:$47$:$35.25$ & $15.12\\;(0.05)\\;2$ & $15.02\\;(0.04)\\;2$ & $14.31\\;(0.03)\\;2$ & $13.94\\;(0.06)\\;2$ & $13.56\\;(0.05)\\;2$ \\\\ $\\rm \\textbf {C}$:$\\;\\rm D^{\\rm a,f},\\,3^{\\rm b}      $ & $18$:$41$:$29.706$ & $79$:$47$:$59.64$ & $15.72\\;(0.06)\\;2$ & $15.45\\;(0.05)\\;2$ & $14.66\\;(0.05)\\;2$ & $14.29\\;(0.08)\\;2$ & $13.92\\;(0.07)\\;2$ \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\multicolumn{8}{c}{$\\rm \\textbf {NGC~6814}$} \\\\ $\\rm \\textbf {A}                                     $ & $19$:$42$:$34.053$ & $-10$:$21$:$03.98$ & $          \\ldots$ & $19.43\\;(0.08)\\;2$ & $17.91\\;(0.02)\\;3$ & $17.16\\;(0.03)\\;3$ & $16.47\\;(0.06)\\;2$ \\\\ $\\rm \\textbf {B}                                     $ & $19$:$42$:$32.218$ & $-10$:$18$:$55.05$ & $17.75\\;(0.08)\\;1$ & $17.28\\;(0.03)\\;2$ & $16.31\\;(0.02)\\;3$ & $15.84\\;(0.03)\\;3$ & $15.40\\;(0.05)\\;2$ \\\\ $\\rm \\textbf {C}                                     $ & $19$:$42$:$38.269$ & $-10$:$17$:$46.90$ & $16.71\\;(0.06)\\;1$ & $15.98\\;(0.03)\\;2$ & $14.90\\;(0.02)\\;3$ & $14.40\\;(0.03)\\;3$ & $13.91\\;(0.04)\\;2$ \\\\ $\\rm \\textbf {D}                                     $ & $19$:$42$:$47.459$ & $-10$:$17$:$42.35$ & $18.06\\;(0.08)\\;1$ & $17.42\\;(0.03)\\;2$ & $16.38\\;(0.02)\\;3$ & $15.88\\;(0.03)\\;3$ & $15.41\\;(0.05)\\;2$ \\\\ $\\rm \\textbf {E}                                     $ & $19$:$42$:$45.795$ & $-10$:$17$:$40.85$ & $17.81\\;(0.07)\\;1$ & $16.33\\;(0.03)\\;2$ & $14.93\\;(0.03)\\;3$ & $14.28\\;(0.03)\\;3$ & $13.68\\;(0.05)\\;2$ \\\\ $\\rm \\textbf {F}                                     $ & $19$:$42$:$36.006$ & $-10$:$17$:$22.74$ & $17.52\\;(0.07)\\;1$ & $16.92\\;(0.03)\\;2$ & $15.93\\;(0.02)\\;3$ & $15.45\\;(0.03)\\;3$ & $15.01\\;(0.05)\\;2$ \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\multicolumn{8}{c}{$\\rm \\textbf {Mrk~304}$} \\\\ $\\rm \\textbf {A}                                     $ & $22$:$17$:$17.819$ & $14$:$13$:$28.11$ & $15.69\\;(0.06)\\;1$ & $15.42\\;(0.04)\\;2$ & $14.65\\;(0.03)\\;2$ & $14.23\\;(0.04)\\;2$ & $13.84\\;(0.06)\\;2$ \\\\ $\\rm \\textbf {B}                                     $ & $22$:$17$:$04.362$ & $14$:$14$:$37.28$ & $17.86\\;(0.06)\\;1$ & $17.82\\;(0.04)\\;2$ & $17.09\\;(0.03)\\;2$ & $16.68\\;(0.04)\\;2$ & $16.31\\;(0.06)\\;2$ \\\\ $\\rm \\textbf {C}                                     $ & $22$:$17$:$14.323$ & $14$:$15$:$08.79$ & $17.58\\;(0.05)\\;1$ & $17.66\\;(0.04)\\;2$ & $16.98\\;(0.03)\\;2$ & $16.55\\;(0.04)\\;2$ & $16.15\\;(0.06)\\;2$ \\\\ $\\rm \\textbf {D}                                     $ & $22$:$17$:$08.602$ & $14$:$15$:$25.62$ & $19.32\\;(0.05)\\;1$ & $18.03\\;(0.04)\\;2$ & $16.85\\;(0.03)\\;2$ & $16.16\\;(0.04)\\;2$ & $15.59\\;(0.06)\\;2$ \\\\ $\\rm \\textbf {E}                                     $ & $22$:$17$:$05.070$ & $14$:$15$:$29.10$ & $17.02\\;(0.06)\\;1$ & $16.31\\;(0.04)\\;2$ & $15.33\\;(0.03)\\;2$ & $14.78\\;(0.04)\\;2$ & $14.27\\;(0.06)\\;2$ \\\\ $\\rm \\textbf {F}                                     $ & $22$:$17$:$20.088$ & $14$:$15$:$43.74$ & $18.47\\;(0.07)\\;1$ & $18.32\\;(0.05)\\;2$ & $17.59\\;(0.03)\\;2$ & $17.11\\;(0.05)\\;2$ & $16.70\\;(0.07)\\;2$ \\\\ $\\rm \\textbf {G}                                     $ & $22$:$17$:$19.755$ & $14$:$16$:$06.40$ & $17.44\\;(0.07)\\;1$ & $17.38\\;(0.05)\\;2$ & $16.65\\;(0.03)\\;2$ & $16.21\\;(0.05)\\;2$ & $15.81\\;(0.07)\\;2$ \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\end{tabular} \\end{table*}  \\setcounter{table}{1}  \\begin{table*}[t] \\caption{Continued.} \\centering \\begin{tabular}{@{}l@{\\hspace{4.0mm}}r@{\\hspace{4.0mm}}r@{\\hspace{4.0mm}}r@{\\hspace{4.0mm}}r@{\\hspace{4.0mm}}r@{\\hspace{4.0mm}}r@{\\hspace{4.0mm}}r@{}} \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} $\\rm Star$ & $\\rm RA\\,(J2000)$ & $\\rm Dec\\,(J2000)$ & $U$ & $B$ & $V$ & $R_{\\rm \\scriptstyle C}$ & $I_{\\rm \\scriptstyle C}$ \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\multicolumn{8}{c}{$\\rm \\textbf {Ark~564}$} \\\\ $\\rm \\textbf {A}$:$\\;3^{\\rm d}                       $ & $22$:$42$:$39.265$ & $29$:$44$:$21.03$ & $14.16\\;(0.09)\\;2$ & $14.17\\;(0.07)\\;2$ & $13.58\\;(0.05)\\;2$ & $13.27\\;(0.07)\\;2$ & $12.97\\;(0.06)\\;2$ \\\\ $\\rm \\textbf {B}$:$\\;1^{\\rm f}                       $ & $22$:$42$:$48.379$ & $29$:$45$:$22.30$ & $13.90\\;(0.13)\\;2$ & $13.97\\;(0.09)\\;2$ & $13.44\\;(0.07)\\;2$ & $13.17\\;(0.09)\\;2$ & $12.93\\;(0.08)\\;2$ \\\\ $\\rm \\textbf {C}$:$\\;2^{\\rm d},\\,13^{\\rm f}          $ & $22$:$42$:$32.084$ & $29$:$45$:$27.02$ & $15.85\\;(0.11)\\;2$ & $15.77\\;(0.08)\\;2$ & $15.07\\;(0.07)\\;2$ & $14.74\\;(0.08)\\;2$ & $14.43\\;(0.08)\\;2$ \\\\ $\\rm \\textbf {D}$:$\\;1^{\\rm d},\\,12^{\\rm f}          $ & $22$:$42$:$35.122$ & $29$:$45$:$37.96$ & $15.56\\;(0.09)\\;2$ & $15.42\\;(0.08)\\;2$ & $14.73\\;(0.07)\\;2$ & $14.41\\;(0.09)\\;2$ & $14.12\\;(0.08)\\;2$ \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\multicolumn{8}{c}{$\\rm \\textbf {NGC~7469}$} \\\\ $\\rm \\textbf {A}                                     $ & $23$:$03$:$10.785$ & $08$:$50$:$43.14$ & $18.59\\;(0.17)\\;2$ & $18.81\\;(0.05)\\;2$ & $18.20\\;(0.04)\\;2$ & $17.93\\;(0.04)\\;2$ & $17.55\\;(0.06)\\;2$ \\\\ $\\rm \\textbf {B}                                     $ & $23$:$03$:$07.717$ & $08$:$51$:$43.43$ & $17.22\\;(0.10)\\;2$ & $17.38\\;(0.05)\\;2$ & $16.73\\;(0.05)\\;2$ & $16.43\\;(0.04)\\;2$ & $16.02\\;(0.06)\\;2$ \\\\ $\\rm \\textbf {C}                                     $ & $23$:$03$:$12.216$ & $08$:$52$:$12.93$ & $17.46\\;(0.10)\\;2$ & $17.65\\;(0.04)\\;2$ & $17.06\\;(0.04)\\;2$ & $16.75\\;(0.04)\\;2$ & $16.32\\;(0.05)\\;2$ \\\\ $\\rm \\textbf {D}                                     $ & $23$:$03$:$23.233$ & $08$:$53$:$21.83$ & $18.34\\;(0.13)\\;2$ & $18.62\\;(0.06)\\;2$ & $17.96\\;(0.05)\\;2$ & $17.67\\;(0.04)\\;2$ & $17.29\\;(0.06)\\;2$ \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\end{tabular} \\flushleft{$^{\\rm a}$\\,Penston et~al. (\\cite{penston71}), $^{\\rm b}$\\,Curry et~al. (\\cite{curry98}), $^{\\rm c}$\\,Bachev et~al. (\\cite{bachev00}), $^{\\rm d}$\\,Shemmer et~al. (\\cite{shemmer01}), $^{\\rm e,\\,f}$\\,Doroshenko et~al. (\\cite{doroshenko05a}, \\cite{doroshenko05b}).} \\end{table*}  Calibration of optical monitoring light curves is important as AGN monitoring programmes typically combine data from different observatories~-- to get dense temporal coverage of the resulting light curves over a certain period of time~-- and from different frequency domains~-- to study the AGN emission across the electromagnetic spectrum. The preferable photometric technique used to calibrate AGN light curves is relative photometry against local comparison stars. Compared to multi-airmass calibration, it is less time-consuming and makes possible the use of non-photometric nights, too. Therefore, the presence of comparison sequences in the fields of target AGNs is of importance for monitoring programmes. Aside from photometric monitoring of AGNs, the comparison sequences could also be used for surface photometry calibration of the corresponding active galaxies.  Johnson-Cousins comparison sequences in the fields of Seyfert galaxies were presented in the papers of Penston, Penston \\& Sandage (\\cite{penston71}), Miller (\\cite{miller81},\\cite{miller86}), Hamuy \\& Maza (\\cite{hamuy89}), Curry et~al. (\\cite{curry98}), Bachev, Strigachev \\& Dimitrov (\\cite{bachev00}), Gonz\\'alez-P\\'erez, Kidger \\& Mart\\'in-Luis (\\cite{gonzalez01}), and Doroshenko et~al. (\\cite{doroshenko05a}, \\cite{doroshenko05b}). Further improvement of the existing standard sequences along with establishing new ones is highly recommended. The improvement of a comparison star sequence could comprise: \\begin{itemize} \\item[--] Addition of new comparison stars in order to extend the magnitude and the colour index ranges covered by the comparison sequence and/or to optimize the distribution of the standards in the field of the target object. The larger number of comparison stars would also increase the accuracy of the zero-point magnitude derived by them; \\item[--] Addition of new flux measurements for the existing comparison stars in order to increase the accuracy of their magnitudes. Further check of the comparison stars for variability could also be done; \\item[--] Extension of the passband coverage of the calibrated magnitudes. \\end{itemize}  Since 1997 $UBVR_{\\rm \\scriptstyle C}I_{\\rm \\scriptstyle C}$ observations of a number of Seyfert galaxies have been obtained in the course of active galaxy surface photometry programme at the Rozhen National Astronomical Observatory (NAO), Bulgaria, by the authors; about a dozen of the objects have been observed two or more times. We were not able to do a multi-airmass calibration of the data taken during a few of the observing nights due to weather or technical problems. A part of these non-calibrated data was transformed to the standard Johnson-Cousins system using comparison sequences established in the corresponding object fields, whereas for the rest this approach failed because of the following reasons: (1) the comparison stars lay outside our field of view, (2) the standards were saturated in our frames, (3) the standard sequence did not cover some of the photometric bands (usually $U$), and (4) there was no comparison sequence established in the field. A possible approach in these cases is matching the non-calibrated host galaxy profiles or galaxy multi-aperture magnitudes to calibrated ones (e.g. Kotilainen, Ward \\& Williger \\cite{kotilainen93}). However, this type of calibration is of lower accuracy compared to both multi-airmass and relative calibration.  In order to transform the non-calibrated data of ours to the standard system and to facilitate future AGN monitoring, we re-examined our observational data with the purpose to find objects suitable (1) for establishing new comparison sequences in their fields and (2) for improving the standard sequences already existing. A total of ten Seyfert galaxies were selected as a result of this search; some information about these objects is listed in Table~\\ref{gals}. We present the results of the $UBVR_{\\rm \\scriptstyle C}I_{\\rm \\scriptstyle C}$ calibration of comparison stars in the fields of the selected galaxies in this paper.  The paper is organized as follows. The observations and data reduction are presented in Sect.~\\ref{obsred}. The flux measurement and magnitude calibration are described in Sect.~\\ref{fluxmc}. The general characteristics of the calibrated stars are discussed in Sect.~\\ref{genchar}. Comments on the individual fields are presented in Sect.~\\ref{comm} and we summarize the main results of this paper in Sect.~\\ref{summ}. ",
        "conclusions": "\\label{summ}  We have calibrated Johnson-Cousins $UBVR_{\\rm \\scriptstyle C}I_{\\rm \\scriptstyle C}$ magnitudes of 49 stars in the fields of the Seyfert galaxies Mrk~335, Mrk~79, Mrk~279, Mrk~506, 3C~382, 3C~390.3, NGC~6814, Mrk~304, Ark~564, and NGC~7469 in order to facilitate photometric monitoring of these objects; 36 stars have been measured for the first time. Our magnitudes are in good agreement with the published ones and we have found no signs of variability of the calibrated comparison stars at least with amplitudes down to the estimated magnitude errors.  We have calibrated in Johnson-Cousins $UBVR_{\\rm \\scriptstyle C}I_{\\rm \\scriptstyle C}$ system the comparison stars in three of the fields, Mrk~506, 3C~382, and Mrk~304, for the first time (to our knowledge) and we have improved the existing standard sequences in the other fields. New comparison stars have been added to the fields of Mrk~335 (two stars), Mrk~79 (four stars), NGC~6814 (six stars), and NGC~7469 (four stars)~-- most of the newly added stars are fainter and are situated closer to the Seyfert galaxies compared to the existing standards. The passband coverage of the comparison sequences in the fields of Mrk~335, Mrk~79, Mrk~279, NGC~6814, and Ark~564 have been complemented with the $U$ band.  Future observations of the presented comparison sequences are welcome in order to improve them further."
    },
    "0711/0711.1162_arXiv.txt": {
        "abstract": "Ultra compact dwarf galaxies (UCDs) have similar properties as massive globular clusters or the nuclei of nucleated galaxies. Recent observations suggesting a high dark matter content and a steep spatial distribution within groups and clusters provide new clues as to their origins. We perform high-resolution $N$-body / smoothed particle hydrodynamics simulations designed to elucidate two possible formation mechanisms for these systems: the merging of globular clusters in the centre of a dark matter halo, or the massively stripped remnant of a nucleated galaxy. Both models produce density profiles as well as the half light radii that can fit the observational constraints. However, we show that the first scenario results to UCDs that are underluminous and contain no dark matter. This is because the sinking process ejects most of the dark matter particles from the halo centre. Stripped nuclei give a more promising explanation, especially if the nuclei form via the sinking of gas, funneled down inner galactic bars, since this process enhances the central dark matter content. Even when the entire disk is tidally stripped away, the nucleus stays intact and can remain dark matter dominated even after severe stripping. Total galaxy disruption beyond the nuclei only occurs on certain orbits and depends on the amount of dissipation during nuclei formation. By comparing the total disruption of CDM subhaloes in a cluster potential we demonstrate that this model also leads to the observed spatial distribution of UCDs which can be tested in more detail with larger data sets. ",
        "introduction": "A new population of subluminous and extremely compact objects have been recently discovered in cluster and group environments (Drinkwater et al. 2000; Phillips et al. 2001) These ultra compact dwarf galaxies (hereafter UCDs) are dynamically distinct systems having intrinsic sizes $\\lsim$ 100 pc and absolute magnitudes in the $B$-band in the range from $-13$ to $-11$ placing them in the lower range of dwarf galaxy luminosities (Mateo 1998). A number of different scenarios have been proposed for origin of UCDs (e.g. Evstigneeva et al. 2007 and references within): (a)\\,They are simply very big and luminous globular clusters, (b)\\,they are nucleated dwarf galaxies, (c)\\,they are the resulting objects of the coalescence of several globular clusters, (d)\\,they are the remnants of stripped disk galaxies. In this work we compare the last two scenarios.  Oh and Lin (2000), Oh, Lin and Richer (2000) as well as Fellhauer and Kroupa (2002) investigated the third formation scenario. The latter authors performed $N$-body simulations of the merging of several globular clusters and argued that the resulting object resembled a UCD. Their simulations initiated with a supercluster, an accumulation of up to fifty globular clusters, orbiting inside a host galaxy on fairly elliptical orbits and with apo-galactic distances around 20\\,kpc. The globular clusters merged within the supercluster giving rise to an object with values for surface brightness and absolute bolometric luminosity comparable to UCDs.  The last formation scenario has been discussed in Oh, Lin \\& Aarseth (1995) as well as Bekki et al. (2001, 2003). The latter authors performed numerical simulations of the dynamical evolution of nucleated dwarf galaxies orbiting inside NGC 1399 and the Fornax cluster. Adopting a plausible scaling relation for dwarf galaxies, Bekki et al. found that the outer stellar envelope of a nucleated dwarf was totally removed by tidally stripping over the course of several passages from the central region of their host. The nucleus is so dense that it will always survive the tidal field of a group or cluster potential. By construction in the initial conditions, the size and luminosity of the remnant were similar to those observed for UCDs and the host galaxy was a Plummer model halo which has a constant density core and therefore easy to tidally disrupt leaving behind the central nucleus which would be associated with a UCD galaxy.  The main goal of the present study is to shed light into the formation of UCDs investigating the last two formation scenarios in more detail and with more realistic initial conditions. First, we adopt the Fellhauer and Kroupa (2002) model and combine it with the cold dark matter model (CDM) paradigm. We assume that globular clusters form and subsequently orbit around dark matter haloes having masses comparable to that of the Fornax dwarf spheroidal and containing no other baryonic matter. Due to dynamical friction, these globular cluster would spiral in towards the centre of the halo. We perform collisionless $N$-body simulations of this process and show that the object resulting from the coalescence of the globular clusters at the halo centre resembles a UCD. As we demonstrate below, the above model suffers from two major drawbacks which rule out its applicability for UCD formation.  The first issue is related to the observed high M/L ratio (between 6 and 9 in solar units) recently reported for UCDs \\cite{ACS}. One has to note here that the M/L ratio of UCDs is still debated (Micheal Drinkwater, private communication, compare Evstigneeva et al. 2007 who quote a M/L for UCDs 3 and 5 in solar units). Such high values can probably not be achieved by the above mechanism. This is because sinking globular clusters will expel most of the dark matter particles from the halo centre \\cite{amr,merritt,goerdt}. If there is no dark matter in UCDs, Fellhauer and Kroupa (2006) describe a possible way to enhance the mass-to-light ratios of UCDs through tidal interactions. The second problem with this scenario is related to the total luminosity of an UCD. At least today, dark matter halos with a sufficient number of globular clusters to produce such bright objects are very rare \\cite{sharina,boeker1}.  Second, we examine the scenario of Bekki et al. (2001, 2003) which is based on the hypothesis that UCDs are remnants of stripped nucleated galaxies. This mechanism has also been discussed by Kazantzidis et al. (2003). We test this model using SPH simulations of a low mass galaxy which forms a nucleus via gas inflow to the inner $\\approx 100$ parsecs. Once this galaxy is placed on a critical disrupting orbit within a cluster potential, the surviving nucleus is in excellent agreement with the latest observational constraints for UCDs, including their dark matter content and spatial distribution.  The paper is organised as follows. In section 2 we describe the globular cluster numerical simulations and compare the properties of the resulting object with those of UCDs. Section 3 contains results from the simulations of tidal stripping of disk galaxies inside a host cluster potential. Finally, in Section 4 we summarise our main conclusions. ",
        "conclusions": "UCDs have similar properties as massive globular clusters or the nuclei of nucleated galaxies. Recent observations of a high dark matter content and their steep spatial distribution within groups and clusters give us new clues as to their origins. We perform $N$-body simulations and compare two possible mechanisms for their formation: the merging of globular clusters in the centre of a dark matter halo, or the massively stripped remnant of a nucleated galaxy. Our simulations reveal, that a swarm of ten as well as one of fifty globular clusters born in a Fornax dwarf spheroidal size dark matter halo will normally evolve to an object, which has the same density profile as an UCD. We performed a second series of simulations in which a disk galaxy experiences mass loss processes inside a cluster environment. The disk is entirely stripped and the remaining nucleus exhibits all of the observed UCD properties. Both models produce density profiles as well as the half light radii that can fit the observational constraints very well.  However, we show that the first scenario produces UCDs that are underluminous and contain no dark matter - the sinking process ejects most of the dark matter particles from the halo centre. The stars, which consist the globulars, would replace the dark matter particles and expel them from the centre of the halo \\cite{amr,merritt} lowering the M/L ratio beyond observational constraints \\cite{ACS}. The other drawbacks of this model is that it is very unlikely to have so many globular cluster in a halo of that size. It becomes apparent in Fig. \\ref{figsm} that it is not sufficient to merge only ten globulars: approximately fifty are necessary. However such a halo may not exist \\cite{sharina}.  Stripped nuclei give a more promising explanation, especially if the nuclei form via the sinking of gas, funneled down inner galactic bars, since this process enhances the central dark matter content. Even when the entire disk is stripped away, the nucleus remains intact and can be dark matter dominated. The total disruption of the galaxy beyond the nuclei only occurs on certain orbits and depends on the amount of dissipation during nuclei formation. By comparing the total disruption of CDM subhaloes in a cluster potential we show that this model also leads to the observed spatial distribution of UCDs. This scenario can be tested in more detail with larger data sets."
    },
    "0711/0711.0744_arXiv.txt": {
        "abstract": "We establish an indirect link between relic neutrinos and the dark energy sector which originates from the vacuum energy contributions of the neutrino quantum fields. Via renormalization group effects they induce a running of the cosmological constant with time which dynamically influences the evolution of the cosmic neutrino background. We demonstrate that the resulting reduction of the relic neutrino abundance allows to largely evade current cosmological neutrino mass bounds and discuss how the scenario might be probed by the help of future large scale structure surveys and Planck data. ",
        "introduction": "Two of the outstanding advances of the last decades were the confirmation of neutrino oscillations and the observation of a late-time acceleration of the cosmological expansion attributed to a form of {\\it dark energy}, its simplest origin being a Cosmological Constant (CC). In this work, we discuss the implications for cosmological neutrino mass bounds arising from an interaction between the CC and relic neutrinos.  According to Big Bang theory the Universe was merely about one second old, when cosmic neutrinos fell out of equilibrium due to the freeze-out of the weak interactions. From the standard presumption that the neutrinos have been non-interacting ever since decoupling, it follows that their distribution is frozen into a freely expanding Fermi-Dirac momentum distribution. Consequently, these neutrinos are assumed to homogenousely permeate the universe as the cosmic neutrino background (\\CnuB) with a predicted average number density of $\\sum n_{\\nu_{0}}=\\sum n_{\\bar{\\nu}_{0}}=168\\,{\\rm{cm}}^{-3}$ today. This substantial relic abundance only falls second to the photons of the cosmic microwave background (CMB), the exact analog of the \\CnuB. However, owing to the feebleness of the weak interaction, all attempts to directly probe the \\CnuB in a laboratory setting have so far been spoiled~\\cite{Hagmann:1999kf,Ringwald:2004np,Gelmini:2004hg,Ringwald:2005zf}. Yet, evidence for its existence and knowledge of its properties can indirectly be inferred from cosmological measurements sensitive to its presence like the abundance of light elements as well as CMB and Large Scale Structure (LSS) measurements (see e.g.\\ Ref.~\\cite{Lesgourgues:2006nd} for a recent review). For example, LSS data can reflect kinematical signatures of relic neutrinos which arise, since on scales below their free streaming scale massive neutrinos cannot contribute to the gravitational clustering of matter. Through the metric source term in the perturbed Einstein equations, this lack of neutrino clustering translates into a small scale suppression of the matter power spectrum which depends on the fractional energy density provided by neutrinos to the total matter density. Hence, within the standard cosmological model, LSS data allows to infer bounds on the absolute neutrino mass scale $\\sum m_{\\nu}=0.2 \\div 1$ eV (2$\\sigma$) depending on the employed data sets (see e.g.~\\cite{Goobar:2006xz,Seljak:2006bg,Feng:2006zj,Cirelli:2006kt,Hannestad:2006mi,Fogli:2006yq,Spergel:2006hy,Tegmark:2006az,Zunckel:2006mt,Kristiansen:2006ky}). However, these bounds are in tension with the claim of part of the Heidelberg-Moscow collaboration of a $>4\\sigma$ evidence for neutrinoless double beta decay translated into a total neutrino mass of $\\sum m_{\\nu}>1.2$ eV at $95\\%$ c.l.~\\cite{KlapdorKleingrothaus:2004wj}. While this claim is still considered as controversial (see e.g.~\\cite{Elliott:2004hr}), at present other independent laboratory experiment are not sensitive enough to verify it. However, considering that cosmological neutrino mass bounds strongly rely on theoretical assumptions on the \\CnuB properties such an apparent tension might hint at new physics beyond the standard model. Namely, it might imply the presence of non-standard neutrinos interactions which reduce the relic neutrino abundance and thus relax cosmological neutrino mass bounds (see Refs.~\\cite{Lesgourgues:2006nd,Beacom:2004yd} and references therein). Hence, the improvement of current mass limit from tritium $\\beta$ decay experiments, $\\sum m_{\\nu}<6.6$ eV ($2\\sigma$)~\\cite{Weinheimer:2003fj,Lobashev:2001uu}, by the approved tritium experiment KATRIN with a projected sensitivity of $0.2$ eV~\\cite{Osipowicz:2001sq,KATRIN}, offers the exciting possibility to probe such new neutrino interactions in the near future. What kind of new interactions could neutrinos have?  Recently, it has been argued that a natural origin for new neutrino interactions would be the sector responsible for dark energy, a smooth energy component, which is assumed to drive the observed accelerated expansion of the universe~\\cite{Gu:2003er,Fardon:2003eh,Amendola:2007yx,Wetterich:2007kr}. A scenario has been proposed~\\cite{Fardon:2003eh}, in which the homogenousely distributed relic neutrinos are promoted to a natural dark energy candidate, if they interact through a new scalar force. As a clear and testable signature of this so-called Mass Varying Neutrino (MaVaN) Scenario the new neutrino interaction gives rise to a variation of neutrino masses with time. The resulting phenomenological consequences have been explored by many authors, in particular the implications for cosmology and astrophysics have been considered in Refs.~\\cite{Li:2004tq,Brookfield:2005td,Brookfield:2005bz,Afshordi:2005ym,Ringwald:2006ks,Schrempp:2006mk,Friedland:2007vv,Bjaelde:2007ki,Bean:2007ny}.  In this paper we propose an alternative approach for the realization of a dynamical influence of the dark sector on relic neutrinos (and vice versa), which is complementary to the avenue in the MaVaN Scenario in the following sense. While we also start from the requirement of energy-momentum conservation of the coupled two-component system, we, however, require the neutrino masses to be constant. As a direct consequence, the energy exchange with the dark sector demands a non-conservation of the number of neutrinos and accordingly its variation with time. Such a set-up naturally arises on the basis of Einstein's theory of general relativity in the appropriate framework provided by quantum field theory on curved space time~\\cite{Birrell:1982ix}. Namely, this description requires the running of the vacuum (or zero point) energy or equivalently the CC as governed by the renormalization group~\\cite{Elizalde:1993ew,Elizalde:1994av,Shapiro:1999zt, Shapiro:2000dz, Guberina:2002wt, Grande:2007wj}, which implies an energy exchange between the CC and matter fields~\\cite{Horvat:2005ua, Carneiro:2006yv, SilvaeCosta:2007je, Bilic:2007gr}. Within a suitable renormalization scheme, we find the relic neutrinos to govern the running of the CC at late times, when they turn non-relativistic. We analyse the consequences of the energy-momentum exchange between the relic neutrinos and the CC both for the late-time dynamics of the running CC and for the phase-space distribution and the abundance of relic neutrinos. Interestingly, independent of the absolute neutrino mass scale realized in nature, in the redshift range accessible by current cosmological probes we find the effective equation of state of the running CC to be extremely close to $-1$ as suggested by recent data (e.g.~\\cite{Astier:2005qq}). However, we demonstrate that within the considered scenario the effects of the \\CnuB on cosmological measurements are drastically reduced. As a consequence, current neutrino mass limits arising from cosmology can be largely evaded. Furthermore, we discuss how the considered scenario can be tested in the near future by the help of cosmological and laboratory based measurements.  The outline of the paper is as follows. In Sec.~\\ref{sec:RGE} we provide the theoretical framework for the considered scenario, and in Sec.~\\ref{sec:Evolution} we analyse in detail the time evolution of the neutrinos and the CC arising from their interaction. Sec.~\\ref{sec:Pheno} deals with the phenomenological consequences of the new neutrino interaction, in particular, with the implications for the inference of neutrino mass bounds from current LSS and CMB data. In Sec.~\\ref{sec:Probing} we provide an outlook how the considered scenario can be tested both by future cosmological probes and by experiments sensitive to Pauli-Blocking effects arising from the presence of the \\CnuB. ",
        "conclusions": "The claimed evidence for neutrinoless double beta decay translates into a neutrino mass bound of $\\sum m_{\\nu}>1.2$ eV at $95\\%$~\\cite{KlapdorKleingrothaus:2004wj} which is in tension with current neutrino mass bounds derived from cosmology~\\cite{Goobar:2006xz,Seljak:2006bg,Feng:2006zj,Cirelli:2006kt,Hannestad:2006mi,Fogli:2006yq,Spergel:2006hy,Tegmark:2006az,Zunckel:2006mt,Kristiansen:2006ky}. In this work, we have shown that cosmological neutrino mass bounds can be relaxed and brought into agreement with the Heidelberg-Moscow claim, if a newly proposed neutrino interaction is taken into account which can be tested by future CMB and LSS measurements. It acts between relic neutrinos and the CC and arises from the zero-point energy contributions of the neutrino quantum fields. Since they induce a scale-dependence of the vacuum energy via renormalization group effects, the CC becomes time-dependent as long as the renormalization scale runs. In this case, an energy-momentum exchange between the relic neutrinos and the CC is implied through the Bianchi identity.  We have studied in detail the time-evolution of the coupled system, in particular its consequences for the dynamics of the CC, the spectrum and the abundance of relic neutrinos and its impact on the interpretation of cosmological measurements. We have found that owing to the relative smallness of neutrino masses, the interaction becomes of dynamical influence at late times when neutrinos turn non-relativistic, $T_{\\nu} \\simeq m_{\\nu}/3$. Moreover, it effectively ceases after a relatively short period, at a temperature also set by the neutrino mass, $T_{\\nu} \\simeq m_{\\nu}/9$. This decoupling behavior can be attributed to our particle specific choice for the renormalization scale. Since it was taken to be the average available neutrino energy which asymptotically runs to a value slightly above the neutrino mass, its evolution and thus the running of the CC effectively cease in the non-relativistic neutrino regime. By taking into account other fermions of higher masses, it was also shown that the corresponding vacuum contributions are becoming more and more important at late times, whereas they are completely sub-dominant in the early universe. During the cosmological evolution, the CC is therefore approaching a tracking regime at late times.  Accordingly, in the relativistic neutrino regime after decoupling we have found the neutrino distribution to essentially maintain its standard Fermi-Dirac form, while the CC runs only logarithmically with time, but with a smaller value than today. However, when turning non-relativistic, the neutrinos were found to efficiently transfer energy-momentum to the CC, thus driving its energy density to its value as measured today. Since as a result, high neutrino momenta turned out to be removed from the neutrino spectrum while low momenta got enhanced, remarkably, the interaction was found to deform the neutrino momentum distribution into a maximally degenerate form. In comparison to Fermi-Dirac distributed neutrinos of equal mass, the mean neutrino momentum decreased by one order of magnitude implying a reduction of the neutrino energy density to $1/30$ of its standard value after the interaction has ceased. Interpreted in terms of the standard relation between the neutrino energy density and mass, the non-relativistic neutrinos redshift as CDM, but appear lighter $m_{\\rm app}=1/30\\,m_{\\nu}$. Accordingly, the presence of the non-standard neutrino interaction was found to considerably relax cosmological neutrino mass bounds. In particular, within the considered scenario current galaxy redshift surveys were shown not to be sensitive at all to neutrino masses below the upper bound $m_{\\nu}=2.2$ eV from tritium experiments~\\cite{Weinheimer:2003fj,Lobashev:2001uu} and could thus be evaded. In addition, we have argued that the present comparatively mild neutrino mass bound from WMAP3 year data of $\\sum 2\\div 3$ eV at $2\\sigma$ can be relaxed in the presence of the proposed neutrino interaction with the CC. The reason primarily is that in comparison to standard $\\Lambda$CDM cosmology, the characteristic integrated effects on the expansion rate turned out to be smaller than in the case of a $\\Lambda$MDM model.  However, we have proposed possible tests for the non-standard neutrino interactions with the CC by future cosmological probes as well as by laboratory based experiments. In the former case, the explanatory power on possible new physics turned out to increase with the neutrino mass. Since future weak lensing surveys probe a broad range of comoving wavenumbers at different redshifts, they provide unbiased, tomographic information about the matter power spectrum. Combined with Planck data the most ambitious projects promise a $2\\sigma$ sensitivity to the minimal value in case of a normal hierarchy ($\\sum m_{\\nu}\\sim0.06$ eV)~\\cite{Lesgourgues:2006nd} such that a non-observation of neutrino free streaming signatures would provide a hint for non-standard neutrino interactions. Interpreted in terms of the considered scenario this would mean that the interaction has reduced the neutrino energy density below the sensitivity of these future probes implying the sum of neutrino masses to be in the sub-eV range.  In this case it might still be feasible for proposed laboratory based experiments to reveal the unique signatures of a possible interaction with the CC by directly probing the neutrino distribution function through Pauli-blocking effects~\\cite{Cocco:2007za,Takahashi:2007ec}.  If, on the other hand, the sum of neutrino masses is larger than $\\sum m_{\\nu}>1.2$ eV as suggested by part of the Heidelberg-Moscow collaboration, we found future LSS probes also to be able to unambiguously reveal the characteristic free streaming signatures of neutrinos having interacted with the CC. Namely, Fermi-Dirac distributed neutrinos can free stream larger distances the lighter they are. However, as a unique signature of the energy-exchange with the CC, neutrinos become slower than in the absence of the interaction and thus free stream a shorter distance, while at the same time they appear lighter due to their reduced abundance.  In addition, it seems likely that future Planck data alone could provide a complementary test for the Heidelberg-Moscow result within the considered scenario, since in this case the interacting phase overlaps with the recombination epoch. Considering that the projected sensitivity in the absence of the neutrino-interaction is $\\sigma(\\sum m_{\\nu})\\lwig 0.48$ eV (see e.g. Ref.~\\cite{Lesgourgues:2006nd} and references therein), it thus seems feasible by Planck data to trace the characteristic modifications to the expansion rate resulting from the neutrino -- CC interaction.  Let us in the following give an outlook on possible extensions of the considered scenario. Since theories beyond the standard model predict other relics as warm or cold dark matter candidates, e.g.\\ axions, gravitinos or supersymmetric partners of standard model fermions, their effects on the running of the CC relevant at energy scales of the order of their masses, could consistently be studied within the proposed renormalization scheme. Such an extension of the considered scenario is especially interesting, because the masses of bosonic species enter the RGE for the CC with an opposite sign as fermions. This might indicate that while the interaction with the CC generically causes the fermionic species to lose energy-momentum, in turn, the bosonic species would gain energy-momentum. Thus, presumably this would tend to increase the momenta of the bosonic distribution, making the bosonic dark matter candidate hotter, while the energy density in the CC would decrease. However, in order to be consistent with observations, within this scenario strong restriction on the particle masses are expected to arise or their existence could turn out to be excluded after all.  After the posting of the first version of this paper, Ref.~\\cite{Foot:2007wn} appeared. The corresponding authors claim that in the work on the RGE running of the CC~\\cite{Shapiro:1999zt,Shapiro:2000dz} part of the one-loop corrections to the CC have been missed out, namely the explicit $\\mu$ dependent part of the full one-loop effective action. It is furthermore argued that on including the missing part the CC does not formally run with the renormalisation scale $\\mu$. In our opinion this claim is misleading, since it is precisely this allegedly missed out one-loop correction which gives rise to the non-zero beta-function for the CC in Eq.~(\\ref{eq:QF-RGE}) (determined by the {\\it partial} $\\mu-$derivative). Moreover, in the framework of an RG improvement within the effective field theory approach taken in our work, it is as usual the (non-zero) beta-function which describes the RG evolution. Therefore, we do not see the appropriateness of our RG analysis of the CC invalidated by the above paper."
    },
    "0711/0711.2607_arXiv.txt": {
        "abstract": "If active galaxies are defined as extragalactic objects with appreciably non thermal spectra then a continuity exists in redshift from the highest redshift quasars to low redshift Seyferts, AGNs and allied galaxies.  Evidence is discussed for this sequence to be an evolutionary track with objects evolving from high to low intrinsic redshift with time. At the end of this evolution the objects are nearly the same age as our own galaxy and they come to rest on the traditional Hubble relation. ",
        "introduction": "The standard Hubble line in Fig. 1 is delineated by E galaxies which are the brightest in their clusters. The objects which fall above this line can be loosely referred to as ``active''. We will use in the following discussion the fact that their underlying spectra comprise non thermal radiation which cuts off further to the red as they age. Observationally they span the extremes between the strong synchrotron continuum of the quasars to the far red cut off of the radio galaxies and finally to the oldest E galaxies which define the accepted Hubble line in Figs.1 and 2.  At this point we invoke the analyses of individual galaxy - quasar associations which establish that the quasars start out their life at very small luminosity and very high redshift (Arp 1998a, 2003). The general solution of the field equations then show that the particles near zero mass in recently created proto quasars increase as $t^2$ (t being the age of the created particles in the variable mass hypothesis; Narlikar 1977; Narlikar and Das 1980). In order to conserve momentum the particles slow down, lowering the temperature and helping the growing gravity to condense the plasma into a coherent proto galaxy (Actually the effect for which ``dark matter'' was invented.) After this rapid wringing out of the synchrotron energy, low mass atoms form and, because of increasing mass of the electrons in energy level transitions, the intrinsic wavelengths of emissiion and absorption, the redshifts, decrease. (Narlikar and Arp 1993.)  The resultant evolution in the z - m diagram is to brighten apparent magnitude initially followed by a deep decline of intrinsic redshift. Then a slower brightening of the object as star formation sets in. Looking at the distribution of objects in Fig. 2 would suggest a fast evolution in luminosity from about z = 6 to 4, then a long decline in intrinsic redshift from z= 5 to 0.1 and finally a slow approach to the Hubble line for high mass AGNs and somewhat below the Hubble line for low mass progenitors.  As they near the end point of their evolution (when their intrinsic redshift has stabilized near zero) they have ended up somewhat under the standard Hubble line (because most of them finish as intermediate bright galaxies). If one were to shut off the galaxy creating mechanism one would see the active galaxies disappearing with time leaving only the standard Hubble line.  The key to this is the solution of the generalized field equations based on the Hoyle Narlikar Machian gravitation theory as described by Narlikar (1977). For a constant mass approximation the theory reduces to the standard Einstein theory. However, the input from Mach's principle, suggests that the inertia of a newly created particle starts off as zero and grows with its age as it begins to get Machian contributions from more and more remote matter in the universe. The typical wavelengths emitted by a particle (such as the electron in a hydrogen atom) would reduce as its mass grows. Thus newly created matter would exhibit high intrinsic redshift. In such a framework, the intrinsic redshift of all galaxies created at the same time as our own will always give a perfect Hubble line because the look back time to a distant galaxy will always reveal it at a younger age when its intrinsic redshift was exactly that predicted by the Hubble law, cz = d x Ho.   In other words the Hubble line is the the line of repose, the evolutionary end for all galaxies of the age of our own galaxy. Younger galaxies have higher intrinsic redshifts which do not signal their velocity or distance but only their age and their luminosity at that time  which can be much fainter than our own, contrary to the tenants of current astronomy. ",
        "conclusions": ""
    },
    "0711/0711.0602_arXiv.txt": {
        "abstract": "% A number of X-ray binaries (neutron stars or black holes accreting from a companion star) have such short orbital periods that ordinary, hydrogen rich, stars do not fit in. Instead the mass-losing star must be a compact, evolved star, leading to the transfer of hydrogen deficient material to the neutron star. I discuss the current knowledge of these objects, with focus on optical spectroscopy. ",
        "introduction": "In low-mass X-ray binaries (LMXBs) a compact object (either a neutron star or a black hole) accretes material from a companion star. In the process, copious X-rays are produced, making these objects visible to large (many kpc) distances. Distinction based on the mass of the donor is made between high-mass and low-mass X-ray binaries \\citep[see various chapters in][]{lk06}. Combining Kepler's law and approximations to the maximum volume a star can occupy before starting to transfer mass to a companion (the Roche lobe), the following equality can be derived for stars that exactly fill their Roche lobe \\citep[e.g.][]{vh95} \\begin{equation} P_{\\rm orbital} \\approx 9 \\mbox{hr} \\left(\\frac{\\rho}{\\rho_{\\odot}}\\right)^{-1/2} \\end{equation} which together with the fact that the maximum density of hydrogen rich stars (around the hydrogen burning limit) is about 100 gr/cc, implies orbital periods for LMXBs with hydrogen rich donors larger than about 1 hour. Reversing the argument: any LMXB with a shorter orbital period necessarily has a donor star that is more compact that a main sequence star and thus (in view of the evolution of stellar cores) is hydrogen deficient \\citep[see][]{nrj86}. Therefore these short period LMXBs have a special status and are called Ultra-compact X-ray binaries (UCXBs). For a short review see \\citet{nj06}. The current known sample consists of 27 systems, 8 with known periods, 4 with tentative periods and 15 candidate systems \\citep[see][]{zjm+07}. ",
        "conclusions": "Ultra-compact X-ray binaries are puzzling end products of binary evolution and despite substantial progress in recent years both in observational as well as theoretical work it is not yet clear how the observed systems are formed. The (detailed) chemical composition of the transferred material will be a key observable to distinguish different formation channels. In principle optical and X-ray spectroscopy in combination with detailed modelling of the donor stars and the disc spectra can provide this information. However of the currently 27 (candidate) systems that are know, for only 4 or 5 we know their (rough) chemical composition: 3 (or 4) C/O discs and 1 (or 2) He discs."
    },
    "0711/0711.5027_arXiv.txt": {
        "abstract": "{We study the merger histories of galaxy dark matter halos using a high resolution $\\Lambda$CDM $N$-body simulation. Our merger trees follow $\\sim 17,000$ halos with masses $M_0 = (10^{11} - 10^{13}) \\Msun$ at $z = 0$ and track accretion events involving objects as small as $m \\simeq 10^{10} \\Msun$.  We find that mass assembly is remarkably self-similar in $m/M_0$, and dominated by mergers that are $\\sim 10 \\%$ of the final halo mass. While very large mergers, $m \\gtrsim 0.4 \\, M_0$, are quite rare, sizeable accretion events, $m \\sim 0.1 \\, M_0$, are common.  Over the last $\\sim 10$ Gyr, an overwhelming majority ($\\sim 95\\%$) of Milky Way-sized halos with $M_0 = 10^{12} \\Msun$ have accreted at least one object with greater total mass than the Milky Way disk ($m > 5 \\times 10^{10} \\Msun$), and approximately $70 \\%$ have accreted an object with more than twice that mass ($m > 10^{11} \\Msun$). Our results raise serious concerns about the survival of thin-disk dominated galaxies within the current paradigm for galaxy formation in a $\\Lambda$CDM universe.  In order to achieve a $\\sim 70 \\%$ disk-dominated fraction in Milky Way-sized $\\Lambda$CDM halos, mergers involving $m \\simeq 2 \\times 10^{11} \\Msun$ objects must not destroy disks.  Considering that most thick disks and bulges contain old stellar populations, the situation is even more restrictive: these mergers must not heat disks or drive gas into their centers to create young bulges.} ",
        "introduction": "\\label{Introduction}  In the cold dark matter (CDM) model of structure formation, dark matter halos form via the continuous accretion of smaller systems \\citep{Peebles82,Blumenthal84,Davis85,FakhouriMa07,NeisteinDekel08,Cole08}. Mergers of the type predicted can help explain many properties of the observed universe.  Major mergers are believed to play an important role in shaping the Hubble sequence \\cite[][]{tt72,barnes88,hernquist93,nb03,kb05,cox06a,robertson06b,robertson06a,r06,maller06,jesseit07,bournaud07} and triggering star formation and AGN activity \\cite[][]{mihos96,kolatt99,cox06,woods06,barton07}.  Minor mergers may help explain the origin of thick disks \\citep{quinn93,walker_etal96,abadi_03,brook04,dalcanton_etal05,k07,HayashiChiba06}, cause anti-truncation \\citep{younger07}, and produce extended diffuse light components around galaxies \\citep{johnston96,helmi99,bkw01,bj05,purcell07,bell07}.  However, there is lingering concern that mergers are too common in CDM cosmologies for thin disk-dominated systems to survive \\citep[][]{toth_ostriker92,wyse01,kormendy05,kautsch06}.  In this paper we present the merger statistics necessary for addressing this issue.  The formation of disk galaxies within hydrodynamic simulations in hierarchical CDM cosmologies has proven problematic \\citep[e.g.][]{ns00}.  While there have been some successes in forming galaxies with disks in cosmological simulations \\citep[][]{abadi_03,sommer_larsen_etal03,brook04,robertson04,kaufmann07a,governato_etal07}, the general problem is far from resolved. The resultant disks are often fairly thick and accompanied by large bulges, and the systems that form disks tend to have special merger histories.  The resultant thick disk and bulge stars also tend towards a broad range of stellar ages, instead of being dominated by predominantly old stars. Moreover, the successes depend strongly on effective models that describe physics on scales far below the simulation resolution, which are poorly understood. Given the  current difficulties in  understanding {\\it ab  initio} disk formation, one can consider a  less ambitious, but more well-posed question. Even if disk galaxies can form within   CDM halos, can they survive the predicted merger histories?   Unfortunately, the prevalence of thin-disk or even disk-dominated galaxies in the universe is difficult to quantify with large observational samples.  Some promising approaches use asymmetry vs. concentration to define morphological type \\citep[e.g.,][]{Ilbert06alt}, and some use a combination color and concentration indicators \\cite[e.g.][]{Choi07, Park07}.  Despite the wide range of definitions, the general consensus in the literature is that $\\sim 70 \\%$ of $\\sim 10^{12} \\Msun$ halos host disk-dominated, late-type galaxies \\citep[e.g.][]{Weinmann06,vandenBosch06,Ilbert06alt,Choi07,Park07}. We adopt this number for the sake of discussion in what follows, but none of our primary results on merger statistics depend on this number.  Also relevant to the discussion of galaxy merger histories is the prevalence of pure disk galaxies in the universe.  \\cite{kautsch06} have compiled a statistically meaningful sample of edge-on disk galaxies and found that $\\sim 16 \\%$ of these objects are ``simple disks'' with no observable bulge component.  In principle, this statistic places strong constraints on the merger histories of galaxies.  Moreover, a large fraction of disk galaxies with bulges contain pseudo-bulges, which may be the products of secular processes and not the remnants of an early merger event \\citep[e.g.][]{kk04,carollo07}.  These cases provide further motivation to quantify the predicted merger histories of galaxy halos in the favored CDM cosmology.   Here we use a large dissipationless cosmological $\\Lambda$CDM N-body simulation to track the merger histories of an ensemble of $\\sim 17,000$ halos with $z=0$ masses $M_0 = 10^{11} - 10^{13} \\Msun$.  We focus on halos of fixed mass at $z=0$, and concentrate specifically on Milky Way-sized systems, $M_0 \\simeq 10^{12} \\Msun$.  We categorize the accretion of objects as small as $m \\simeq 10^{10} \\Msun$ and focus on the infall statistics into {\\em main progenitors} of $z=0$ halos as a function of lookback time.  As discussed below, the {\\em main progenitor} is defined to be the most massive progenitor of a $z=0$ halo tracked continuously back in time.  A {\\em merger} is defined here to occur when an infalling halo first crosses within the virial radius of the main progenitor.  In most cases we do not track subhalo evolution after accretion.  We have chosen to track mergers in this way in order to provide a robust prediction.  An understanding of an accreted halo's subsequent orbital evolution and impact with the central disk region is essential for any complete understanding of {\\em galaxy} merger statistics.  However, this evolution will be sensitive to the baryonic distribution within both the main progenitor halo and the satellites themselves.  The {\\em halo} merger rate we present is a relatively clean measure that can be used as a starting point for more detailed investigations of galaxy--galaxy encounters.  Still, it is worth pointing out that for most of the mergers we consider, impacts with the central disk region should occur relatively shortly after accretion.  As we show in the Appendix, events with $m \\gtrsim 0.1 M_0$ typically happen at a redshift $z$ when the main progenitor mass, $M_z$, is significantly smaller than $M_0$, such that the merger ratio is fairly large $m/M_z \\gtrsim 0.2$.  Therefore, even ignoring the enhanced orbital decay that will be caused by a central disk potential, the dynamical friction decay times are expected to be short for these events, with central impacts occurring within $\\tau \\lesssim 3$ Gyr for typical orbital parameters \\citep{bk07,z05,zb03}.  As discussed in conjunction with Figure 5 in \\S 3, destruction times of $\\sim 3$ Gyr are consistent with our measurements of subhalo evolution.  The outline of this paper is as follows.  In \\S \\ref{Simulation} we discuss the numerical simulations used and the method of merger tree construction.  In \\S \\ref{Results} we present our principle results, which characterize the accretion mass functions of halos and the fraction of halos with mergers as a function of lookback time.  In \\S \\ref{Discussion} we discuss these results in reference to the problem of disk survival in a hierarchical universe, and we summarize our main conclusions in \\S \\ref{Conclusion}.    \\begin{figure*}[t!] \\begin{center} \\subfigure{\\includegraphics[width=.45\\textwidth]{tree810.eps}} % \\subfigure{\\includegraphics[width=.45\\textwidth]{tree853.eps}} % \\caption{Sample merger trees, for halos with $M_{0} \\simeq 10^{12.5} \\Msun$.   Time progresses downward,  with the redshift $z$ printed on the left hand side. The  bold, vertical line at the center corresponds to   the main progenitor,  with  filled  circles proportional to the radius of  each halo.  The minimum  mass halo shown in  this diagram has   $m =   10^{9.9}  \\Msun$. Solid (black) and    dashed (red) lines and    circles correspond   to  isolated field  halos,  or  subhalos, respectively. The dashed (red) lines that do not merge with main progenitor represent surviving subhalos at $z=0$. \\emph{Left:} A ``typical'' merger history, with a merger of mass $m \\simeq 0.1 M_0  \\simeq 0.5 M_z$ at $z = 0.51$. \\emph{Right:} A halo that experiences an unusually large merger $m \\simeq 0.4 M_0 \\simeq 1.0 M_z$ at $z = 0.65$.} \\label{mergertrees} \\end{center} \\end{figure*} ",
        "conclusions": "\\label{Conclusion}  We have used a high-resolution $\\Lambda$CDM $N$-body simulation to investigate the merger histories of $\\sim 17,000$ galaxy dark matter halos with masses $M_0 = 10^{11-13} \\Msun$ at $z=0$.  Mergers with objects as small as $m = 10^{10} \\Msun$ were tracked. The principle goal has been to present the raw statistics necessary for tackling the issue of thin disk survival in $\\Lambda$CDM and for providing a cosmological context for more focused simulations aimed at understanding the role of mergers for processes like morphological transformation, star formation triggering, and AGN fueling.  Our main results may be summarized as follows:  \\begin{enumerate}  \\item  Mass accretion into halos of mass $M_0$ at $z=0$ is dominated by mergers with objects of mass $m  \\simeq (0.03 - 0.3) M_0$ (Figure \\ref{dfdm}). Typically, $\\sim 1-4$ mergers of this  size occur over a halo's history (Figure \\ref{NgtM}).   Because these mergers  tend  to occur when  the main  progenitor's mass, $M_z$,  was somewhat smaller  than $M_0$, these dominant events have fairly large {\\em merger ratios},  $m/M_z \\simeq 0.1-0.6$ (see Appendix).  \\item The mass accretion function, $n(m)$, of mergers larger than $m$ accreted over a halo's history is well-described, on average, by a simple analytic form, $n(x \\equiv m/M_0) = A \\, x^{-\\alpha} \\, (0.5-x)^{\\beta}$, with $\\alpha = 0.61$ and $\\beta = 2.3$. The normalization increases as a function of the halo's mass at $z=0$, $M_0$, as $A \\simeq 0.47 \\log_{10}(M_0)-3.2$.  By extrapolating this fit, we find that the total mass fraction accreted in objects of any mass ($m>0$) does not asymptote to 1.0, but rather increases with $M_0$ from $\\sim 50 \\%$ in $M_0 = 10^{11.5} \\Msun$ halos to $\\sim 70 \\%$ in $M_0 = 10^{13} \\Msun$ halos.  This suggests that a non-zero fraction of a halos mass may be accreted as truly ``diffuse'' material.   \\item An overwhelming majority ($95 \\%$) of Milky Way-sized halos with $M_0 \\simeq  10^{12} \\Msun$ have accreted  an object  larger than the Milky Way's disk ($m \\gtrsim 5 \\times  10^{10} \\Msun$) in the last 10 Gyr.  Approximately  $70 \\%$ have had  accretions with $m > 10^{11} \\Msun$ objects over the same  period, and  $40 \\%$  have had $m > 2 \\times 10^{11} \\Msun$ events (Figure \\ref{MMfrac}).  \\item Halo merger histories are approximately self-similar  in $m/M_0$ for  halos with  masses in  the range $M_0 = 10^{11}-10^{13}  \\Msun$ (Figure \\ref{LMsincet}).   This suggests that the empirical trend  for late-type  galaxies to   be more common   in smaller halos is not governed by differences in merger histories, but rather is associated with baryon physics.    \\item Typically, a small fraction, $\\sim 20-30  \\%$, of a halo's final mass $M_0$ is accreted {\\em after} the most recent large merger with $m > (0.1-0.2) \\, M_0$ objects (Figure \\ref{Macc}).  This suggests that the ``regrowth'' of a disk from newly accreted material after a large merger is unlikely.  Note that this does not rule out the possibility that a disk reforms from gaseous material involved in the merger itself.   \\end{enumerate}  The relatively high fraction of halos with large $m \\sim 0.1 M_0$ merger events raises concerns about the survival of thin disk galaxies within the current paradigm for galaxy formation in a $\\Lambda$CDM universe.  If we naively match percentages using Figure \\ref{MMfrac}, we find that in order to achieve a $\\sim 70 \\%$ disk-dominated fraction in $M_0 = 10^{12} \\Msun$ halos, then $m \\simeq 0.2 \\, M_0 \\simeq 2 \\times 10^{11} \\Msun \\simeq 3 \\times 10^{11} M_\\odot$ objects must not (always) destroy disks. Furthermore, since stars in the local thick disk and bulge are predominantly older than $\\sim 10$ Gyr, this suggests that these mergers in the past $\\sim 10$ Gyr must not drive gas towards the bulge or significantly thicken the disk. Note that the total mass in such an accreted object is $\\sim 10$ times that of the Milky Way disk itself. Moreover halos typically do not accrete a significant fraction of their final mass after these mergers ($\\sim 20 \\%$ on average). Finally, as noted in the Appendix, $m \\sim 0.2 \\, M_0$ events typically have merger {\\em ratios} of $m/M_z \\simeq 0.4$ at the time of the merger.  These numbers do not seem encouraging for disk survival, and may point to a serious problem with our current understanding of galaxy formation in a $\\Lambda$CDM universe.  Our basic conclusion is unlikely to be sensitive to uncertain cosmological parameters.  Note that the simulations considered here have a fairly high $\\sigma_8 = 0.9$.  At a fixed $\\Omega_m$, a lower $\\sigma_8$ will systematically produce slightly {\\em more} recent merger events \\citep[e.g.][]{zb03}.  However, given that the merger fractions measured using $m/M_0$ are approximately self-similar in $M_0$ (and therefore in $M_0/M_\\star$), we expect that the overall merger fractions will be fairly insensitive to power-spectrum normalization.  As discussed in the introduction, a complete investigation into the issue of disk survival will require an understanding of the orbital evolution of objects once they have fallen within the main progenitor halo's virial radius and on the subsequent impact of interacting galaxies.  Both of these outcomes will depend sensitively on the baryonic components in the main halo and in the smaller merging object.  For this reason, the present work has focused on {\\em halo mergers}, defined to occur when an infalling halo first crosses within the virial radius of the main progenitor halo.  The merger statistics presented here are relatively devoid of uncertainties and can be used as a starting point for direct simulations of galaxy-galaxy encounters.  Simulations of this kind will be essential to fully address the broader implications of these frequent, large mergers, which seem to pose a serious challenge to disk survival."
    },
    "0711/0711.4810_arXiv.txt": {
        "abstract": "Based on dramatic observations of the cosmic microwave background radiation with WMAP and of Type Ia supernovae with the Hubble Space Telescope and ground-based facilities, it is now generally believed that the Universe's expansion is accelerating. Within the context of standard cosmology, this type of evolution leads to the supposition that the Universe must contain a third `dark' component of energy, beyond matter and radiation. However, the current data are still deemed insufficient to distinguish between an evolving dark energy component and the simplest model of a time-independent cosmological constant. In this paper, we examine the role played by our cosmic horizon $R_0$ in our interrogation of the data, and reach the rather firm conclusion that the existence of a cosmological constant is untenable. The observations are telling us that $R_0\\approx ct_0$, where $t_0$ is the perceived current age of the Universe, yet a cosmological constant would drive $R_0$ towards $ct$ (where $t$ is the cosmic time) only once, and that would have to occur right now. In contrast, scaling solutions simultaneously eliminate several conundrums in the standard model, including the `coincidence' and `flatness' problems, and account very well for the fact that $R_0\\approx ct_0$. We show in this paper that for such dynamical dark energy models, either $R_0=ct$ for all time (thus eliminating the apparent coincidence altogether), or that what we believe to be the current age of the universe is actually the horizon time $t_h\\equiv R_0/c$, which is always shorter than $t_0$. Our best fit to the Type Ia supernova data indicates that $t_0$ would then have to be $\\approx 16.9$ billion years. Though surprising at first, an older universe such as this would actually eliminate several other long-standing problems in cosmology, including the (too) early appearance of supermassive black holes (at a redshift $>6$) and the glaring deficit of dwarf halos in the local group. ",
        "introduction": "Over the past decade, Type Ia supernovae have been used successfully as standard candles to facilitate the acquisition of several important cosmological parameters. On the basis of this work, it is now widely believed that the Universe's expansion is accelerating (Riess et al. 1998; Perlmutter et al. 1999). In standard cosmology, built on the assumption of spatial homogeneity and isotropy, such an expansion requires the existence of a third form of energy, beyond the basic admixture of (visible and dark) matter and radiation.  One may see this directly from the (cosmological) Friedman-Robertson-Walker (FRW) differential equations of motion, usually written as \\begin{equation} H^2\\equiv\\left(\\frac{\\dot a}{a}\\right)^2=\\frac{8\\pi G}{3c^2}\\rho-\\frac{kc^2}{a^2}\\;, \\end{equation} \\begin{equation} \\frac{\\ddot a}{a}=-\\frac{4\\pi G}{3c^2}(\\rho+3p)\\;, \\end{equation} \\begin{equation} \\dot\\rho=-3H(\\rho+p)\\;, \\end{equation} in which an overdot denotes a derivative with respect to cosmic time $t$, and $\\rho$ and $p$ represent, respectively, the total energy density and total pressure. In these expressions, $a(t)$ is the expansion factor, and $(r,\\theta,\\phi)$ are the coordinates in the comoving frame. The constant $k$ is $+1$ for a closed universe, $0$ for a flat (or open) universe, and $-1$ for an open universe.  Following convention, we write the equation of state as $p=\\omega\\rho$. A quick inspection of equation (2) shows that an accelerated expansion ($\\ddot a>0$) requires $\\omega<-1/3$. Thus, neither radiation ($\\rho_r$, with $\\omega_r=1/3$), nor (visible and dark) matter ($\\rho_m$, with $\\omega_m\\approx 0$) can satisfy this condition, leading to the supposition that a third `dark' component $\\rho_d$ (with $\\omega_d<-1/3$) of the energy density $\\rho$ must be present. In principle, each of these contributions to $\\rho$ may evolve according to its own dependence on $a(t)$.  Over the past few years, complementary measurements (Spergel et al. 2003) of the cosmic microwave background (CMB) radiation have indicated that the Universe is flat (i.e., $k=0$), so $\\rho$ is at (or very near) the ``critical\" density $\\rho_c\\equiv 3c^2H^2/8\\pi G$. But among the many peculiarities of the standard model is the inference, based on current observations, that $\\rho_d$ must itself be of order $\\rho_c$. Dark energy is often thought to be the manifestation of a cosmological constant, $\\Lambda$, though no reasonable explanation has yet been offered as to why such a fixed, universal density ought to exist at this scale. It is well known that if $\\Lambda$ is associated with the energy of the vacuum in quantum theory, it should have a scale representative of phase transitions in the early Universe---many, many orders of magnitude larger than $\\rho_c$.  Many authors have attempted to circumvent these difficulties by proposing alternative forms of dark energy, including Quintessence (Wetterich 1988; Ratra and Peebles 1988), which represents an evolving canonical scalar field with an inflation-inducing potential, a Chameleon field (Khoury and Weltman 2004; Brax et al. 2004) in which the scalar field couples to the baryon energy density and varies from solar system to cosmological scales, and modified gravity, arising out of both string motivated (Davli et al. 2000), or General Relativity modified (Capozziello et al. 2003; Carroll et al. 2004) actions, which introduce large length scale corrections modifying the late time evolution of the Universe. The actual number of suggested remedies is far greater than this small, illustrative sample.  Nonetheless, though many in the cosmology community suspect that some sort of dynamics is responsible for the appearance of dark energy, until now the sensitivity of current observations has been deemed insufficient to distinguish between an evolving dark energy component and the simplest model of a time-independent cosmological constant $\\Lambda$ (see, e.g., Corasaniti et al. 2004). This conclusion, however, appears to be premature, given that the role of our cosmic horizon has not yet been fully folded into the interrogation of current observations.  The purpose of this paper is to demonstrate that a closer scrutiny of the available data, if proven to be reliable, can in fact already delineate between evolving and constant dark energy theories, and that a simple cosmological constant $\\Lambda$, characterized by a fixed $\\omega_d= \\omega_\\Lambda=-1$, is almost certainly ruled out. In \\S~2 of this paper, we will introduce the cosmic horizon and discuss its evolution in time, demonstrating how measurements of the Hubble constant $H$ may be used to provide strict constraints on $\\omega$, independent of Type Ia supernova data. We will compare these results with predictions of the various classes of dark energy models in \\S\\S~3 and 4, and conclude with a discussion of the consequences of these comparisons in \\S~5. ",
        "conclusions": "Our main goal in this study has been to examine what we can learn about the nature of dark energy from a consideration of the cosmic horizon $R_0$ and its evolution with cosmic time. A principal outcome of this work is the realization that the so-called ``coincidence\" problem in the standard model is actually more severe than previously thought. We have found that in a $\\Lambda$CDM universe, $R_0\\rightarrow ct_0$ only once, and according to the observations, this must be happening right now. Coupled to the fact that the basic $\\Lambda$CDM model does not adequately account for the Type Ia supernova data without introducing new parameters and patching together phases of deceleration and acceleration that must somehow blend together only in the current epoch, this is a strong indication that dark energy almost certainly is not due to a cosmological constant. Other issues that have already been discussed extensively in the literature, such as the fact that the vacuum energy in quantum theory should greatly exceed the required value of $\\Lambda$, only make this argument even more compelling. Of course, this rejection of the cosmological-constant hypothesis then intensifies interest into the question of why we don't see any vacuum energy at all, but this is beyond the scope of the present paper.  Many alternatives to a cosmological constant have been proposed over the past decade but, for the sake of simplicity, we have chosen in this paper to focus our attention on scaling solutions. The existence of such cosmologies has been discussed extensively in the literature, within the context of standard General Relativity, braneworlds (Randall-Sundrum and Gauss-Bonnett), and Cardassian scenarios, among others (see, e.g., Maeda 2001; Freese and Lewis 2002; and Fujimoto and He 2004).  Our study has shown that scaling solutions not only fit the Type Ia supernova data much better than the basic $\\Lambda$CDM cosmology, but they apparently simultaneously solve several conundrums with the standard model. As long as the time-averaged value of $\\omega$ is less than $-1/3$, they eliminate both the coincidence and flatness problems, possibly even obviating the need for a period of rapid inflation in the early universe (see, e.g., Guth 1981; Linde 1982).  But most importantly, as far as this study is concerned, scaling solutions account very well for the observed fact that $R_0\\approx ct_0$. If $\\langle\\omega\\rangle=-1/3$ exactly, then $R_0(t)=ct$ for all cosmic time, and therefore the fact that we see this condition in the present Universe is no coincidence at all. On the other hand, if $\\langle\\omega\\rangle<-1/3$, scaling solutions fit the Type Ia supernova data even better, but then we have to accept the conclusion that the Universe is older than the horizon time $t_h\\equiv R_0/c$. According to our calculations, which produce a best fit to the supernova data for $\\langle\\omega\\rangle=-0.47$ (corresponding to a dark energy equation of state with $\\omega_d=-2/3$), the age of the Universe should then be $t_0\\approx 16.9$ billion years. This may be surprising at first, perhaps even unbelievable to some, but the fact of the matter is that such an age actually solves other major problems in cosmology, including the (too) early appearance of supermassive black holes, and the glaring deficit of dwarf halos in the local group of galaxies.  When thinking about a dynamical dark energy, it is worth recalling that scalar fields arise frequently in particle physics, including string theory, and any of these may be appropriate candidates for this mysterious new component of $\\rho$. Actually, though we have restricted our discussion to equations of state with $\\omega_d\\ge -1$, it may even turn out that a dark energy with $\\omega_d<-1$ is providing the Universe's acceleration. Such a field is usually referred to as a Phantom or a ghost. The simplest explanation for this form of dark energy would be a scalar field with a negative kinetic energy (Caldwell 2002). However, Phantom fields are plagued by severe quantum instabilities, since their energy density is unbounded from below and the vacuum may acquire normal, positive energy fields (Carroll et al. 2003). We have therefore not included theories with $\\omega_d<-1$ in our analysis here, though a further consideration of their viability may be warranted as the data continue to improve.  On the observational front, the prospects for confirming or rejecting some of the ideas presented in this paper look very promising indeed. An eagerly anticipated mission, SNAP (Rhodes et al. 2004), will constrain the nature of dark energy in two ways. First, it will observe deeper Type Ia supernovae. Second, it will attempt to use weak gravitational lensing to probe foreground mass structures. If selected, SNAP should be launched by 2020. An already funded mission, the Planck CMB satellite, probably won't have the sensitivity to measure any evolution in $\\omega_d$, but it may be able to tell us whether or not $\\omega_d=-1$. An ESA mission, Planck is scheduled for launch in mid-2008.  Finally, we may be on the verge of uncovering a class of sources other than Type Ia supernovae to use for dark-energy exploration. Type Ia supernovae have greatly enhanced our ability to study the Universe's expansion out to a redshift $\\sim 2$. But this new class of sources may possibly extend this range to values as high as $5-10$. According to Hooper and Dodelson (2007), Gamma Ray Bursts (GRBs) have the potential to detect dark energy with a reasonable significance, particularly if there was an appreciable amount of it at early times, as suggested by scaling solutions. It is still too early to tell if GRBs are good standard candles, but since differences between $\\Lambda$CDM and dynamical dark energy scenarios are more pronounced at early times (see figures~3, 6, and 9), GRBs may in the long run turn out to be even more important than Type Ia supernovae in helping us learn about the true nature of this unexpected ``third\" form of energy."
    },
    "0711/0711.1118_arXiv.txt": {
        "abstract": "We report on an imaging survey with the {\\it Spitzer Space Telescope } of 62 brightest cluster galaxies with optical line emission. These galaxies are located in the cores of X-ray luminous clusters selected from the ROSAT All-Sky Survey. We find that about half of these sources have a sign of excess infrared emission;  22 objects out of 62 are detected at 70 $\\mu$m, 18 have 8 to 5.8 $\\mu$m flux ratios above 1.0 and 28 have 24 to 8 $\\mu$m flux ratios above 1.0.  Altogether 35 of 62 objects in our survey exhibit at least one of these signs of infrared excess. Four galaxies with infrared excesses have a  4.5/3.6 $\\mu$m flux ratio indicating the presence of hot dust, and/or an unresolved nucleus at 8 $\\mu$m. Three of these have high measured $[$OIII$]$(5007\\AA)/H$\\beta$ flux ratios suggesting that these four, Abell 1068, Abell 2146, and Zwicky 2089, and R0821+07, host dusty active galactic nuclei (AGNs). 9 objects (including the four hosting dusty AGNs) have infrared luminosities greater than $10^{11} L_\\odot$ and so can be classified as luminous infrared galaxies (LIRGs). Excluding the four systems hosting dusty AGNs, the excess mid-infrared emission in the remaining brightest cluster galaxies is likely related to star formation. ",
        "introduction": "X-ray observations of many galaxy clusters show a density increase and a temperature decrease toward the centers of these systems, implying that gas should be cooling at rates of a few to 1000~$M_\\odot$~yr$^{-1}$ (e.g., \\citealt{allen00}). A number of studies have found evidence for cooled gas and star formation in galaxy clusters, but at a rate below that corresponding to the cooling rate predicted from the X-ray observations (e.g., \\citealt{mcnamara89,cardiel98,crawford99,bregman06}). Although mechanisms to counteract radiative cooling of the intracluster medium have been identified (e.g., \\citealt{birzan04,rafferty06,dunn06,peterson06,mcnamara07}), these energy sources do not seem to be sufficiently efficient to shut off cooling altogether.  Up to $10^{11} M_\\odot$ molecular gas is detected in BCGs with the largest mass deposition rates (e.g., \\citealt{edge02,salome03}). Most cooling flow systems exhibit optical emission-lines such as H$\\alpha$ (\\citealt{heckman81,Hu85}). Emission line spectroscopy suggests that some of the emission is due to photoionization and heating by hot stars \\citep{voit97}. The gulf between the spectral limits on the condensation rates and the sink of cooling material, the primary objection to cooling flows, has narrowed dramatically.  In some systems, star formation rates are close to the X-ray and far UV estimates of the rates of gas condensation \\citep{hicks05,rafferty06,mcnamara04}.  To probe the efficiency of cooling and star formation in cluster galaxies, we aim to provide more direct measurements of the star formation rate in central  cluster galaxies using broad band images obtained with the {\\it Spitzer Space Telescope} (SST). Recent work using SST observations of a few brightest cluster galaxies (BCGs) includes the study by \\citet{egami06} finding that the majority of BCGs are not infrared luminous.  This study of 11 galaxies suggested that star formation was responsible for non-AGN excess infrared emission in 3 BCGs and that star formation was taking place primarily in objects with high X-ray luminosity.  Two of the objects studied by \\citet{egami06} have estimated infrared luminosities classifying them as Luminous Infrared Galaxies (LIRGs), $10^{11} < L_{\\rm IR} < 10^{12} L_\\odot$. Star formation rates for the infrared luminous objects were approximately 10 times lower than the mass deposition rates. While the galaxies studied by \\citet{egami06} are moderately distant ($0.15< z<0.5$), \\citet{donahue07} found that the BCG in the nearby ($z=0.0852$) cluster Abell 2597 can also be classified as a LIRG.  In this paper we build on the previous studies of a small number of BCGs \\citep{egami06,donahue07} with an imaging survey of 62 BCGs using the Spitzer Space Telescope (SST). In this paper we describe the sample and survey data and briefly discuss the morphology seen in the images.  Using aperture photometry we assemble spectral energy distributions (SEDs) of the central regions of the galaxies and identify the sources that have infrared excesses and estimate their infrared luminosities. Using infrared colors and morphology (whether a central source was red and unresolved) and optical emission line ratios we differentiate between sources with infrared emission consistent with contribution from a dusty AGN and those likely dominated by star formation. Comparison data are assembled to search for correlations between infrared luminosities and other cluster properties; however, the actual comparisons are done in a companion paper (paper II; \\citealt{odea07}). In this paper all luminosities have been corrected or computed to be consistent with a Hubble constant $H_0 = 70$~Mpc$^{-1}$~km~s$^{-1}$ and a concordant cosmology ($\\Omega_M=0.3$ and flat). ",
        "conclusions": "We have presented Spitzer IRAC and MIPS photometry of a sample of 62 BCGs selected on cluster X-ray flux and BCG H$\\alpha$ or [NII] flux to favor objects in cool cluster cores. We find that one half to one-third of the sample show an infrared excess above that expected for the old stellar population. These results confirm the previous results of \\citet{egami06}, based on a sample of 11 BCGs in X-ray luminous clusters, that some BCGs in cool core clusters show an infrared excess due to star formation.  An interesting difference between our survey and previous studies \\citep{egami06,donahue07} is the discovery of a small number of objects that host dusty AGNs, thanks to our much larger sample. Three galaxies exhibit red 4.5/3.6$\\mu$m flux ratios and unresolved nuclei seen in IRAC color maps indicating the presence of hot dust. These three, Abell 1068, Abell 2146, and Zwicky 2089, have measured high $[$OIII$]$(5007)/H$\\beta$ flux ratios suggesting that they host a dusty AGN. An additional BCG, R0821+07, has a red unresolved nucleus at 8$\\mu$m and a high $[$OIII$]$(5007)/H$\\beta$ ratio suggesting that it too hosts a dusty AGN.  These four BCGs have infrared luminosities greater than $10^{11} L_\\odot$ and so can be classified as LIRGs. A comparison between their infrared luminosity and 1.4 GHz radio luminosity yields ratios of $10^6 - 10^7$ suggesting that they can be considered radio quiet. The only similar previously known object is the more distant ($z=0.44$) and hyperluminous infrared galaxy IRAS 09104+4109 \\citep{fabian95}, also a BCG in an X-ray luminous cluster.  In addition to the 4 hosting AGNs, there are 5 other BCGs with infrared luminosities greater than $10^{11} L_\\odot$ that can be classified as luminous infrared galaxies (LIRGs). Excluding the AGNs, the remaining brightest cluster galaxies with infrared excesses are likely bright in the mid-infrared because of star formation as discussed in our companion paper \\citep{odea07}. In this work we have discussed the identifications and broad properties of this sample. Our companion paper (paper II; \\citealt{odea07}) searches for correlations between star formation rates, radio, H$\\alpha$, CO and X-ray luminosities and mass deposition rates estimated from the X-ray observations. Planned future work includes fitting model spectral energy distributions to the observations."
    },
    "0711/0711.1432_arXiv.txt": {
        "abstract": "Many extrasolar planets orbit closely to their parent star. Their existence raises the fundamental problem of loss and gain in their mass. For exoplanet HD209458b, reports on an unusually extended hydrogen corona and a hot layer in the lower atmosphere seem to support the scenario of atmospheric inflation by the strong stellar irradiation. However, difficulties in reconciling evaporation models with observations call for a reassessment of the problem.  Here, we use HST archive data to report a new absorption rate of $\\sim 8.9\\%\\pm 2.1$\\% by atomic hydrogen during the HD209458b transit, and show that no sign of evaporation could be detected for the exoplanet. We also report evidence of time variability in the HD209458 \\lya\\ flux, a variability that was not accounted for in previous studies, which corrupted their diagnostics. Mass loss rates thus far proposed in the literature in the range $5\\times 10^{10}-10^{11}\\, {\\rm g\\, s^{-1}}$ must induce a spectral signature in the Lyman-$\\alpha$ line profile of HD209458 that cannot be found in the present analysis. Either an unknown compensation effect is hiding the expected spectral feature or else the mass loss rate of neutrals from HD209458 is modest. ",
        "introduction": "Of all the planets discovered outside the realm of our solar system, some of the most dramatic new classes of objects are those in which the planet is a gas giant orbiting at merely a few stellar radii ($\\sim 0.02 {\\rm AU}$) from its parent star. These close-in extrasolar planets are Jupiter-like giants that are exposed to strong fluxes, magnetic fields, and plasma winds---a very harsh and active stellar environment. Because of their stars' proximity, gravity, through tidal effects, distorts the shape of their atmosphere while the continuous extreme ultraviolet (UV) energy deposition inflates it (\\cite{lam03}; \\cite{bar04}; \\cite{lec04}; \\cite{yel04}; \\cite{jar05}; \\cite{tia05}; \\cite{mun07}; \\cite{vil07}). Unfortunately, little is known about those regions that separate extrasolar giant planets from their stars, particularly the immediate environment of the planet.  One of the most extensively studied extrasolar systems is HD209458. For reference, HD209458b was first discovered transiting its parent star and covering $1.5$\\% of its disk (\\cite{cha00}; \\cite{hen00}).  Some of the first attempts to learn about the immediate environment of the planet were \\lya\\  ($121.6$ nm) observations of the system using the Space Telescope Imaging Spectrometer (STIS) onboard the Hubble Space Telescope (HST).  A first program of observation, obtained with the STIS/G140M grating and the 52\"x0.1\" slit, was implemented in $2001$ during HD209458 planetary transit but no conclusions were reported (see Table~\\ref{tbl-1}). Soon after, a second program visited the target during three transits (\\cite{vid03}).  An initial analysis of this data set concluded that a huge cloud of hydrogen is covering $15\\%\\pm 4$\\% of the stellar disk (\\cite{vid03}); it also claimed that spectral absorption during transit is deeper on the blue side of the stellar line.  Accordingly, the hydrogen cloud was required to extend beyond the planetary Roche limit where an intense escape of $\\sim~ 10^{10}\\, {\\rm g\\, s^{-1}}$ of hydrogen is a priori operating. These results, along with other far UV low-resolution observations of heavy constituents, led to the conclusion that the upper atmosphere of HD209458b should be in a hydrodynamic blow-off state (\\cite{vid03, vid04}).  Numerous studies then followed on different mechanisms for hydrogen loss from hot exoplanets closely orbiting their stars (\\cite{lam03}; \\cite{bar04}; \\cite{jar05}; \\cite{lec04}; \\cite{yel04, yel06}; \\cite{tia05}; \\cite{mun07}). As noted by \\cite{mun07}, all loss rates thus far proposed by theoretical models in the range $5\\times 10^{10}-10^{11}\\, {\\rm g\\, s^{-1}} $ exceed the lower limit provided by \\cite{vid03}. Unfortunately,  most studies neglected to quantitatively translate their loss rate to a spectral absorption in the \\lya\\ line profile that could be tested with the HST observation. Independently, pointing out that the observed mass function distribution of extrasolar giant planets (EGPs) follows a trend ${\\rm M^{-1}}$ for mass range $\\sim 0.2-5\\, {\\rm M_{J}}$, where ${\\rm M_{J}}$ is the Jovian mass, \\cite{hub07} derived the same mass function distribution for highly irradiated EGPs orbiting at distances smaller than $\\sim 0.07$ AU. Accordingly, \\cite{hub07} rejected substantial mass loss during EGPs' migration to smaller distances from their star, unless the loss mechanism is compensated by an unknown process. When combined with the unusual scales derived for the hydrogen extent and escape, all these studies then call for a careful reassessment of the HST \\lya\\ observations thus far obtained on HD209458, at least in order to provide validated constraints on theoretical models. ",
        "conclusions": "We use HST archive observations of the \\lya\\ emission of HD209458 to report an absorption rate of $\\sim 8.9\\% \\pm 2.1$\\% by atomic hydrogen during the transit of the planetary companion. If the planet is sketched as a compact blocking body, our analysis requires an H cloud effective extent that does not exceed $\\sim 2.5$ R$_{\\rm P}$---a size that falls short of the Roche limit $\\sim 4.08$ R$_{\\rm P}$  of HD209458b. In addition, time variability of the stellar flux is evidenced, but no sign of extra or Doppler-shifted absorption could be detected during transit. This absence of extra absorption during transit and the relatively small size of the effective area of the hydrogen cloud around the exoplanet make it difficult to conceive of significant atmospheric evaporation from the planet. Of course, we cannot rule out that a complex atmospheric distribution, related to a particular planet-star interaction scenario, may hide or compensate the loss signature during the observing time. Future HST (when repaired) FUV observation of the system during a full planetary orbit should help to disentangle the different processes in play."
    },
    "0711/0711.3601_arXiv.txt": {
        "abstract": "{}{We study the interaction of a low-mass planet with a protoplanetary disk with a realistic treatment of the energy balance by doing radiation-hydrodynamical simulations. We look at accretion and migration rates and compare them to isothermal studies.} {We used a three-dimensional version of the hydrodynamical method RODEO, together with radiative transport in the flux-limited diffusion approach.} {The accretion rate, as well as the torque on the planet, depend critically on the ability of the disk to cool efficiently. For densities appropriate to 5 AU in the solar nebula, the accretion rate drops by more than an order of magnitude compared to isothermal models, while at the same time the torque on the planet is \\emph{positive}, indicating outward migration. It is necessary to lower the density by a factor of 2 to recover inward migration and more than 2 orders of magnitude to recover the usual Type I migration. The torque appears to be proportional to the radial entropy gradient in the unperturbed disk. These findings are critical for the survival of protoplanets, and they should ultimately find their way into population synthesis models. }{} ",
        "introduction": "Planet formation takes place in protoplanetary disks that are commonly found around young stars \\citep{1996Natur.383..139B}. These disks play a critical role in structuring the forming planetary system in terms of masses and orbits, and planet-disk interaction is therefore a very important process to study in order to understand planet formation.  Planet-disk interaction affects a forming planet in three ways. First of all, the mass that the disk is able to supply to the young planet is limited, which leads to a maximum gas mass that a planet can achieve \\citep{1999ApJ...526.1001L,1999MNRAS.303..696K,2003ApJ...586..540D, 2003MNRAS.341..213B}. This limiting mass is of the order of a few Jupiter masses (\\mjup). But also for planets of lower mass the disk regulates the mass accretion, which becomes especially apparent in three-dimensional numerical simulations \\citep{2003ApJ...586..540D}.  Second, the disk is able to change the orbital radius of the planet through tidal interaction \\citep{1980ApJ...241..425G}. Depending on the mass of the planet, two modes of migration can be distinguished \\citep{1997Icar..126..261W}: type I migration for low-mass planets that generate a linear disk response, and type II for high-mass, gap-opening planets. Recently a third type of migration was put forward, in which strong corotational torques force a very fast mode of migration \\citep{2003ApJ...588..494M,2004ASPC..324...39A}. This Type III migration is very sensitive to local density gradients, in terms of migration speed and even migration direction, and therefore the structure of the gas disk determines the outcome of the migration process. This is unlike Type I and Type II migration, for which all planets move inward and only the time scale varies with disk properties.  Finally, the eccentricity of the planet's orbit may be affected by the disk. Usually, eccentricity is damped \\citep{1993ApJ...419..166A}, but high-mass planets ($M_\\mathrm{p} > 5 $ \\mjup) that open deep gaps may experience eccentricity growth \\citep{2001A&A...366..263P,2006A&A...447..369K}. However, due to possible saturation of corotation resonances planets of $M_\\mathrm{p} \\approx$ \\mjups may pick up significant eccentricity during the gap formation process \\citep{2004ApJ...606L..77S}, although this requires a slightly eccentric orbit to start with \\citep{2003ApJ...585.1024G}.  Observations of the masses and orbits of extrasolar planets provide insight in the result of planet-disk interaction. One of the most striking results that were obtained is the discovery of a new class of planets, the so-called Hot Jupiters. They are gas giant planets orbiting very close to their central star. It is unlikely that they have formed in situ \\citep{1999ApJ...521..823P} and inward migration through planet-disk tidal interaction has been put forward as the explanation for their existence.  The main problem for this scenario is that the time scales for migration as obtained from analytical \\citep{2002ApJ...565.1257T} and numerical \\citep{2003MNRAS.341..213B,2003ApJ...586..540D} arguments are usually smaller than or comparable to the total lifetime of the disk. This means that theory essentially predicts that all planets should fall onto the central star, which poses a problem for planet formation theory. If it is already hard to form a single giant planet within the lifetime of the disk \\citep{1996Icar..124...62P}, how are we going to form a huge amount of them, of which only a tiny fraction survives? Clearly a stopping mechanism is required that prevents planets from migrating all the way towards the star.  One possibility is that the planet encounters an inner hole in the disk, possibly created by the magnetic field of the central star. When there is no disk around to interact with, the planet will stop migrating. Another possibility is that the planet undergoes Type III outward migration at some point \\citep{2003ApJ...588..494M}. Finally, direct interaction of the disk with a magnetic field may stop the planet, either by a toroidal magnetic field \\citep{2003MNRAS.341.1157T} or, for a low-mass planet, its migration may be slowed down through magnetic turbulence \\citep{2004MNRAS.350..849N}.  However, most numerical hydrodynamical work on planet-disk interaction lacks proper treatment of the energy balance, with only a few notable exceptions \\citep{2003ApJ...599..548D,2006A&A...445..747K,2003MNRAS.346..915M}. The usually adopted locally isothermal equation of state assumes that all excess energy generated by compression, viscous dissipation or shocks can be radiated away efficiently, thereby keeping the temperature profile fixed. However, the ability of the disk to cool is strongly linked to the opacity, and therefore to the density.  The effects of temperature structure on planet migration in realistic disks were studied by \\cite{2003ApJ...593.1116J,2004ApJ...608..497J, 2005ApJ...619.1123J} who used a detailed disk model including radiative transfer to calculate the torque on the planet in an analytical way. They found that temperature perturbations resulting from shadowing and illumination at the surface of the disk can decrease the migration rate by a factor of 2. \\cite{2004ApJ...606..520M} also find reduced migration rates in realistic T Tauri $\\alpha$-disks, and they point out the importance of sudden changes in the opacity, for example near the snow line.  In this study, we aim at relaxing the isothermal assumption in hydrodynamical simulations of planet-disk interaction. The first results on planet migration were presented in \\cite{radlett}, where it was shown that low-mass planets may migrate \\emph{outward} in disks with realistic temperature structures. In this paper, we continue our study of migration of low-mass objects, with more eye for detail than in \\cite{radlett}, and include results on accretion.  In contrast with \\cite{2006A&A...445..747K}, we focus on low-mass planets that do not open a gap in the disk. For these deeply embedded planets temperature effects are the most important. However, due to their low mass their gravitational sphere of influence is small, and one needs very high resolution locally to resolve the atmosphere of the planet. Furthermore, because radiation is the main cooling agent, we need to perform radiation-hydrodynamical simulations. This is a major leap from isothermal simulations, because we are not only adding an equation for the gas energy but also the evolution of the radiation field needs to be followed.  The plan of this paper is as follows. We briefly review the relevant cooling time scale in Sect. \\ref{3DsecCool}, effects of radiation in Sect. \\ref{3DsecRad} and the condition for convection in Sect. \\ref{3DsecConv}. In Sect. \\ref{3DsecModel} we present the adopted disk model, and in Sect. \\ref{3DsecNum} we discuss the numerical method. In Sect. \\ref{3DsecRes} we present the results, from simple isothermal models to the full radiation-hydrodynamical models. We give a discussion on the results in Sect. \\ref{3DsecDisc}, and we conclude in Sect. \\ref{3DsecCon}. ",
        "conclusions": "\\label{3DsecCon} In this paper we have presented the first global radiation-hydrodynamical simulations of the interaction between low-mass planets and protoplanetary disks. Specifically, we looked at accretion and migration rates in the mass regime of linear disk response.  We have found that accretion and migration rates depend sensitively on the ability of the disk to radiate away energy generated by compression in the various flow regions around the planet. Compression deep within the planetary envelope leads to a decrease in accretion rate onto the planet of more than an order of magnitude. Compression within the tidal waves generated by the planet leads to a decreased Lindblad torque, which makes the planet move inward but at a slower rate than predicted by analytical models for isothermal Type I migration. Compression in the horseshoe region, finally, leads to a large positive corotation torque that is able to change the direction of migration, depending on the radial entropy gradient.  All these effects depend critically on the cooling properties of the disk. For a density appropriate for the Minimum Mass Solar Nebula at 5 AU we find outward migration, while a ten times lower density gives rise to inward migration. Lowering the density even more we are able to reproduce the analytical Type I torque. This means that inward Type I migration is restricted to the outer parts of protoplanetary disks, or to older disks that have a lower density.  Because low-mass planet migration is driven by the decrease of density in the disk, it proceeds on the viscous time scale, just as Type II migration for gap-opening planets. This means that corotational heating is able to solve the problem of protoplanetary cores that move inward too fast to form giant planets. All planets within the region where the cooling time scale is much larger than the dynamical time scale will experience outward migration, depending on the local entropy gradient.  Future studies should be aimed at finding the most appropriate recipe to simulate gas accretion onto the planet, and at including irradiation from the central star to calculate the equilibrium disk temperature profile self-consistently."
    },
    "0711/0711.4351_arXiv.txt": {
        "abstract": "We present a new determination of the metallicity gradient in M33, based on Keck/LRIS measurements of oxygen abundances using the temperature-sensitive emission line [\\ion{O}{3}]~$\\lambda$4363~\\AA\\ in 61 \\hii\\ regions.  These data approximately triple the sample of direct oxygen abundances in M33.  We find a central abundance of 12 + log(O/H) = $8.36 \\pm 0.04$ and a slope of $-0.027 \\pm 0.012$~dex~kpc$^{-1}$, in agreement with infrared measurements of the neon abundance gradient but much shallower than most previous oxygen gradient measurements.  There is substantial intrinsic scatter of 0.11~dex in the metallicity at any given radius in M33, which imposes a fundamental limit on the accuracy of gradient measurements that rely on small samples of objects.  We also show that the ionization state of neon does not follow the ionization state of oxygen as is commonly assumed, suggesting that neon abundance measurements from optical emission lines require careful treatment of the ionization corrections. ",
        "introduction": "\\label{intro}  M33 was one of the first galaxies in which a radial abundance gradient was recognized \\citep{searle71}.  Beginning with \\citet{smith75}, a series of studies showed that \\hii\\ region abundances in M33 decline by approximately an order of magnitude over the visible extent of the galaxy \\citep{ka81,vilchez88,garnett92}, with a simple exponential providing an accurate fit to the gradient.  A typical value for the gradient from this work is $-0.12 \\pm 0.02$~dex~kpc$^{-1}$ \\citep[][hereafter V88]{vilchez88}.  More recently, \\citet[][hereafter C06]{crockett06} presented new metallicity measurements for a small sample of \\hii\\ regions that indicated nearly constant oxygen abundances across M33, with a best-fit gradient of only $-0.012 \\pm 0.011$~dex~kpc$^{-1}$.  If confirmed, this very shallow gradient would call into question a variety of results concerning, e.g., the metallicity dependence of the Cepheid period-luminosity relation \\citep{lee02} and the CO-H$_{2}$ conversion factor \\citep{eros03} that depend on the abundance gradient in M33.  To complicate matters further, infrared measurements of the neon abundance gradient, which should be less sensitive to systematic effects such as extinction and the temperature of the \\hii\\ region, yield an intermediate gradient of $-0.034 \\pm 0.015$~dex~kpc$^{-1}$ that is not in very good agreement with any of the published oxygen gradients from optical data \\citep[][hereafter WNP02]{wnp02}.  In this paper, we present new, direct metallicity measurements based on detections of the temperature-sensitive \\oiii\\ line in 61 \\hii\\ regions in M33, approximately tripling the sample available in the literature for M33 (or any other galaxy).  We use these data to provide an improved measurement of the abundance gradient in M33.  In \\S \\ref{observations} we describe the observations, data reduction, and derivation of metallicities.  In \\S \\ref{results} we determine the oxygen abundance gradient in M33, attempt to derive neon abundances, and discuss some of the implications of our results.  We briefly summarize the paper in \\S \\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  Using data from the M33 Metallicity Project, we have obtained direct oxygen abundance measurements based on detections of the [\\ion{O}{3}] $\\lambda 4363$~\\AA\\ line for 61 \\ion{H}{2} regions in the southwest half of M33.  This data set approximately triples the sample of M33 \\hii\\ region metallicities in the literature and makes M33 by far the best-studied external galaxy in terms of its ISM abundances.  We presented a refined oxygen gradient for M33 and found an exponential abundance profile with a gradient of $-0.027 \\pm 0.012$~dex~kpc$^{-1}$ and a central abundance of $12+\\log(\\mathrm{O/H})=8.36\\pm 0.04$. There is an intrinsic scatter of 0.11 dex around the trend, which complicates the measurement of the shallow gradient in the galaxy. When this scatter is accounted for, previous studies (which generally derived steeper gradients) were shown to be consistent with our measurements.  We also noted that there is a significant correlation of the Ne$^{+2}$/O$^{+2}$ ratio with the fraction of oxygen found in the doubly ionized state, which suggests that neon determinations based solely on optical lines require more sophisticated ionization correction factors than are typically assumed."
    },
    "0711/0711.3753_arXiv.txt": {
        "abstract": "The origins of Gamma-ray Burst prompt emission are currently not well understood and in this context long, well-observed events are particularly important to study. We present the case of GRB\\,070616, analysing the exceptionally long-duration multipeaked prompt emission, and later afterglow, captured by all the instruments on-board {\\it Swift} and by {\\it Suzaku} WAM. The high energy light curve remained generally flat for several hundred seconds before going into a steep decline. Spectral evolution from hard to soft is clearly taking place throughout the prompt emission, beginning at 285 s after the trigger and extending to 1200 s. We track the movement of the spectral peak energy, whilst observing a softening of the low energy spectral slope. The steep decline in flux may be caused by a combination of this strong spectral evolution and the curvature effect. We investigate origins for the spectral evolution, ruling out a superposition of two power laws and considering instead an additional component dominant during the late prompt emission. We also discuss origins for the early optical emission and the physics of the afterglow. The case of GRB\\,070616 clearly demonstrates that both broadband coverage and good time resolution are crucial to pin down the origins of the complex prompt emission in GRBs. ",
        "introduction": "The prompt emission mechanism for Gamma-ray Bursts (GRBs) is commonly attributed to internal shocks due to collisions of shells of different Lorentz factors ejected from the vicinity of a compact object \\citep{Rees}. Later afterglow emission comes from an external shock as the GRB blastwave decelerates through interaction with the surrounding medium \\citep{Meszaros2}. Though the general picture appears applicable to most GRBs, the details are far from understood and other models, in particular magnetised flows, have also been proposed \\citep[e.g.][]{Meszaros3,Usov,Kumar1,Zhang1}.  The advent of the {\\it Swift} mission \\citep{Gehrels1} has revealed additional features, for example steep decays and X-ray flares \\citep{Burrows2}, whose properties are consistent with an internal origin \\citep[e.g.][]{Tagliaferri,Nousek,O'Brien,Chincarini,Falcone}. The steeply decaying phases that directly follow both prompt emission and X-ray flares are usually interpreted as due to the curvature effect \\citep[e.g.][]{Kumar2} where high latitude emission is delayed with respect to that on-axis. However, significant spectral evolution is not expected in this model, and new mechanisms must be invoked to explain those observations that do show spectral evolution during the steep decay phase (Zhang, Liang \\& Zhang 2007).  Of particular interest are the longest duration GRBs in which the relationship between various possible early emission components can be studied. Observationally, very few GRBs are detected in $\\gamma$-rays for more than 400 seconds as quantified using the T$_{\\rm 90}$ parameter. For example, approximately 0.5\\% of the BATSE sample meet this criteria \\citep{Paciesas}. Some of the very long duration GRBs include those with a pre-cursor or a late flare \\citep[e.g.][]{Price,Nicastro}. Interestingly, the very long GRBs also include several FRED-like (fast rise, exponential decay) events which can be spectrally quite soft \\citep{Giblin,Zand}.  In the {\\it Swift} era it remains true that very few GRBs have a T$_{\\rm 90} >400$ s. The longest is GRB\\,060218, which is an unusually long, low luminosity, spectrally soft GRB \\citep{Campana2}. In the first {\\it Swift} BAT catalogue \\citep{Sakamoto} there are only three others: GRB\\,060929, which has a spectrally soft, late giant flare \\citep{Palmer}; GRB\\,070129, for which BAT triggered on a precursor \\citep{Krimm}; and GRB\\,070616.  Here we present data for the case of GRB\\,070616, in which the prompt emission shows a very complex multipeaked structure, leading to one of the longest prompt emission durations ever recorded. We take advantage of extensive coverage of such a long burst by the {\\it Swift} Burst Alert Telescope \\citep[BAT;][]{Barthelmy}, the X-Ray Telescope \\citep[XRT;][]{Burrows1} and the UV/Optical Telescope \\citep[UVOT;][]{Roming}. Combining data from {\\it Swift} and {\\it Suzaku} \\citep{Mitsuda} we study the evolution of the prompt emission spectrum, following the temporal variability of the peak energy, and study the development of the afterglow component.  In Section 2 we describe the {\\it Swift} and {\\it Suzaku} observations. In Section 3 we outline the X- and $\\gamma$-ray temporal characteristics. In Section 4 we model the X-ray to $\\gamma$-ray spectra and investigate spectral evolution. We discuss the discovery of an optical transient in the UVOT images in Section 5 and go on to model the spectral energy distributions in Section 6. In Section 7 we present possible interpretations for our findings, and discuss them in the context of the blastwave model and in comparison with other GRBs, with summary and conclusions in Section 8. ",
        "conclusions": "GRB\\,070616 has afforded us the rare opportunity to track the detailed spectral and temporal evolution from $\\gamma$-ray to optical wavelengths throughout much of the prompt emission phase, owing to {\\it Swift's} fast slew capability, the broadband coverage provided by the combination of {\\it Swift} with {\\it Suzaku} WAM, and the long prompt emission duration. The high energy light curve remains generally flat for several hundred seconds before going into a steep decline.  Spectral evolution from hard to soft is clearly taking place through the prompt emission (flat and decaying phases) of GRB\\,070616, beginning at 285 s out to at least 1000 s. The high energy spectrum from 0.3--800 keV is well modelled by a Band function. We track the spectral peak energy moving from 135 keV down to $\\sim$4 keV and measure a softening of the low energy spectral slope from $\\Gamma$$\\sim$1.1--2.3. The curvature effect, whilst contributing to the spectral shape, is not dominating during the steep decay at $\\sim$1000 s.  The presence of an additional component, perhaps not present in most GRBs, is favoured by the inability of a two-component fit to model the light curve. This unusual light curve shape of an initial constant flux followed by a sharp drop, plus the strong spectral evolution, is most prominent in GRB\\,070616, but we note a similar structure can be found in a small number of other GRBs which are also not dominated by the curvature effect and not well fit by the two-component model.  We find that the optical data during the prompt phase are not consistent with a simple extrapolation of the high energy spectrum, requiring either significant intrinsic extinction or some extinction plus a break in the spectrum between the X-ray and optical bands.  The afterglow is consistent with the synchrotron model, with the cooling break lying in between optical and X-ray bands at a few thousand seconds after the trigger.  The prompt emission of GRB\\,070616 comprises a component well fitted with a Band function and a possible further component. It is clear that both broadband coverage and good time resolution are crucial to pinning down the origins of the complex prompt emission in GRBs.   \\begin{figure*} \\begin{center} \\includegraphics[width=18cm, angle=0]{Sidewaystable.ps} \\label{replacement} \\end{center} \\end{figure*}"
    },
    "0711/0711.3565_arXiv.txt": {
        "abstract": "We present continuum data from the Submillimetre Common-User Bolometer Array (SCUBA) on the James Clerk Maxwell Telescope (JCMT), and the Mid-Infrared Photometer for Spitzer (MIPS) on the Spitzer Space Telescope, at submillimetre and infrared wavelengths respectively. We study the Taurus molecular cloud 1 (TMC1), and in particular the region of the Taurus Molecular Ring (TMR). In the continuum data we see no real evidence for a ring, but rather we see one side of it only, appearing as a filament. We name the filament `the bull's tail'. The filament is seen in emission at 850, 450 and 160~$\\mu$m, and in absorption at 70~$\\mu$m. We compare the data with archive data from the Infra-Red Astronomical Satellite (IRAS) at 12, 25, 60, 100~$\\mu$m, in which the filament is also seen in absorption. We find that the emission from the filament consists of two components: a narrow, cold ($\\sim$8~K), central core; and a broader, slightly warmer ($\\sim$12~K), shoulder of emission. We use a radiative transfer code to model the filament's appearance, either in emission or absorption, simultaneously at each of the different wavelengths. Our best fit model uses a Plummer-like density profile and a homogeneous interstellar dust grain population. Unlike previous work on a similar, but different filament in Taurus, we require no grain coagulation to explain our data. ",
        "introduction": "The Taurus molecular cloud is one of the nearest star forming regions, at a distance of $\\sim$140pc \\citep{1978ApJ...224..857E}. As such it has often been the target of observational studies into the processes of molecular cloud evolution and star formation. The large scale structure of the molecular cloud has been mapped using the emission from CO \\citep{1987ApJS...63..645U,1996ApJ...465..815O,2005prpl.conf.8268G}, and also the extinction of starlight \\citep{1999A&A...345..965C,2002ApJ...580L..57P,2005PASJ...57S...1D}.  The cloud has been mapped to study the formation mechanism of molecular clouds \\citep{1999ApJ...527..285B}, and the nature of the turbulence that pervades them \\citep[see review by][]{2004ARA&A..42..211E}. On smaller scales, detailed studies of the densest regions of the cloud have greatly increased our knowledge of the star-formation process itself \\citep{2000ApJ...537L.135W,2001A&A...365..440M,2002ApJ...575..950O, 2004A&A...427..651D,2005MNRAS.360.1506K}.  Taurus differs from other nearby sites of star formation, in that it is forming stars in a quiescent manner. There are no newly formed, massive stars injecting energy into the cloud; and compared to more dynamic star-formation regions like Orion or $\\rho$-Ophiuchus, each of the star-forming cores are widely separated \\citep[e.g.][]{1998A&A...336..150M,1999MNRAS.305..143W}. The cloud also appears to have a core mass function which differs from that found in other regions \\citep{2002ApJ...575..950O,2004A&A...419..543G}.  \\begin{figure*} \\includegraphics[angle=0,width=175mm]{figs/fig1.eps} \\caption{The Taurus molecular cloud seen in extinction \\citep{2005PASJ...57S...1D}. The Av=1 and Av=5 contours are shown. TMC1 is marked, as are a number of Lynds dark clouds \\citep{1962ApJS....7....1L}. The outlines of the regions mapped by MIPS at 160$\\mu$m (large box), and SCUBA at 850$\\mu$m (small box), are shown. The SCUBA data are centred around TMC1.} \\label{taurus_overview} \\end{figure*}  \\begin{figure*} \\begin{minipage}[t]{87mm} \\vspace{0pt} \\includegraphics[angle=0,width=80mm]{./figs/fig2a.eps} \\end{minipage} \\begin{minipage}[t]{87mm} \\vspace{0pt} \\includegraphics[angle=0,width=80mm]{./figs/fig2b.eps} \\end{minipage} \\caption{A close-up of the TMC1 region at (a) 160~$\\mu$m, and (b) 850~$\\mu$m. The bright components of TMC1 and HCL2 are marked (TMC1-CP and TMC1-NH3 refer to the cyclopolyyne and ammonia peaks respectively -- \\citealp{1988A&A...196..194O}). Some bright IRAS sources are also marked. The approximate locus of the TMR is shown as a dashed line. The ring structure is not very clear in the continuum images. Instead we see a filament coincident with the eastern portion of the ring, and simply a collection of point sources and clumps in the western half. We here name the filament `the bull's tail'.} \\label{taurus_overview_zoom} \\end{figure*}  \\begin{figure*} \\begin{minipage}[t]{87mm} \\includegraphics[angle=0,width=87mm,height=94mm]{./figs/fig3a.eps} \\end{minipage} \\begin{minipage}[t]{87mm} \\includegraphics[angle=0,width=75.8mm,height=94mm]{./figs/fig3b.eps} \\end{minipage} \\begin{minipage}[t]{87mm} \\vspace{0pt} \\includegraphics[angle=0,width=87mm,height=101mm]{./figs/fig3c.eps} \\end{minipage} \\begin{minipage}[t]{87mm} \\vspace{0pt} \\includegraphics[angle=0,width=75.8mm,height=101mm]{./figs/fig3d.eps} \\end{minipage} \\caption{An even closer view of the bull's tail filament at (a) 850~$\\mu$m, (b) 450~$\\mu$m, (c) 160~$\\mu$m, and (d) 70~$\\mu$m. The filament is seen in emission at the three long wavelengths, and in absorption at 70~$\\mu$m. Each of the maps has been smoothed to match the 45-arcsec angular resolution of the MIPS 160-$\\mu$m data. The remaining differences in the width of the filament are therefore intrinsic to the source itself. The filament is narrowest at 850 and 70~$\\mu$m, where there is good agreement between the emission at 850 and absorption at 70~$\\mu$m. The filament appears broader at the other two wavelengths.} \\label{taurus_maps1} \\end{figure*}  \\begin{figure*} \\includegraphics[angle=270,width=175mm]{figs/fig4.ps} \\caption{One-dimensional cuts through the data perpendicular to the filament at each of the eight wavelengths studied. Each cut has been normalised to its own peak.} \\label{slices} \\end{figure*}  The far-infrared and submillimetre continuum represent an efficient means of detecting cold star-forming cores in molecular clouds, even in energetic regions such as Orion \\citep{1995MNRAS.274.1219W,1996A&A...314..477B}. Figure~\\ref{taurus_overview} shows the Taurus molecular cloud as seen in extinction \\citep{2005PASJ...57S...1D}, with the contours showing the visual extinctions of Av=1 and Av=5. A number of the Lynds clouds which are associated with the Taurus molecular cloud are labelled.  The majority of the cloud has been mapped in the far-infrared by Spitzer \\citep[][for more details, see Section \\ref{spitzer}]{2007prpl.conf..329G}. The region mapped by Spitzer is outlined in white on Figure~\\ref{taurus_overview}. A somewhat smaller region, centred on Taurus molecular cloud 1 (TMC1), has been mapped by SCUBA on the JCMT (see Section \\ref{scuba}). This smaller region is also outlined in white on Figure~\\ref{taurus_overview}. The densest part of the cloud is TMC1. Together with Heiles Cloud~2 \\citep[HCL2;][]{1968ApJ...151..919H} this makes up the Taurus Molecular Ring (TMR) \\citep[see][]{2004A&A...420..533T}.  Figures~\\ref{taurus_overview_zoom}(a) and (b) show the region around TMC1 at 160 and 850~$\\mu$m respectively, with the brighter components of Heiles Cloud 2 and TMC1 marked. In addition, the approximate locus of the TMR is also marked. For more details about these data, see Section~\\ref{observations}. However, we note here that the molecular ring is not really seen in the continuum. Instead, we see a filament in the eastern half and a series of point-like and extended sources scattered roughly around the western half. We therefore challenge the existence of the ring as a single coherent entity, as that is not what is seen in these data. We here name the filament `the bull's tail'.  In this paper, we present a study of a portion of TMC1 using infrared and submillimetre data between 12 and 850~$\\mu$m. This range of wavelengths straddles the peak of the spectral energy distribution (SED) for the cold dust that is found in the densest regions of molecular clouds. Infrared and sub-mm data are therefore a powerful tool for breaking the degeneracy of temperature and density that is present when only sub-mm data are considered. ",
        "conclusions": "We have presented submillimetre data from SCUBA on the JCMT and far-infrared data taken by Spitzer and IRAS, for a filament in TMC1 which we have named the bull's tail. TMC1 is the densest part of the Taurus molecular cloud region. The filament is seen in emission at all wavelengths longer than 100~$\\mu$m, and in absorption at all shorter wavelengths.  The filament is significantly narrower at 850~$\\mu$m, which we attribute to a narrow, dense inner filament, running through the centre of the broader and warmer main filament. We have shown intensity cuts through the filament at wavelengths of 850, 450, 160, 100, 70, 60, 25 and 12~$\\mu$m, which are consistent with this scenario.  We have fitted SEDs to both the inner and outer components of the filament and find that the temperature decreases from 12~K in the outer part to 8~K at the centre. We then used a radiative transfer code to fit the profile of the filament. We modelled the filament using a Plummer-like density profile with constant grain parameters throughout, and find good fits to all wavelengths."
    },
    "0711/0711.4173_arXiv.txt": {
        "abstract": "We study the non-thermal emission from old shell-type supernova remnants (SNRs) on the frame of a time-dependent model. In this model, the time-dependent non-thermal spectra of both primary electrons and protons as well as secondary electron/positron ($e^{\\pm}$) pairs can be calculated numerically by taking into account the evolution of the secondary $e^{\\pm}$ pairs produced from proton-proton (p-p) interactions due to the accelerated protons collide with the ambient matter in an SNR. The multi-wavelength photon spectrum for a given SNR can be produced through leptonic processes such as electron/positron synchrotron radiation, bremsstrahlung and inverse Compton scattering as well as hadronic interaction. Our results indicate that the non-thermal emission of the secondary $e^{\\pm}$ pairs is becoming more and more prominent when the SNR ages in the radiative phase because the source of the primary electrons has been cut off and the electron synchrotron energy loss is significant for a radiative SNR, whereas the secondary $e^{\\pm}$ pairs can be produced continuously for a long time in the phase due to the large energy loss time for the p-p interaction. We apply the model to two old SNRs, G8.7$-$0.1 and G23.3$-$0.3, and the predicted results can explain the observed multi-wavelength photon spectra for the two sources. ",
        "introduction": "\\label{sec:intro}  In recent years, various observational evidences indicate that particles in SNRs can be accelerated up to relativistic energies. For example, X-ray observations show that electrons can be accelerated to about 10 TeV in the shell of an SNR (Koyama et al. 1995; Reynolds 1996; Bamba et al. 2003, 2005a, 2005b). TeV $\\gamma$-ray observations indicate that both electrons and protons can be accelerated to several tens TeV in young shell-type SNRs such as RX~J1713.7$-$3946 (Enomoto et al. 2002; Aharonian et al. 2006a) and RX~J0852.0$-$4622 (Katagiri et al. 2005; Aharonian et al. 2005). More recently, a survey of the inner part of the Galactic Plane in VHE $\\gamma$-ray band has been performed with \\textit{HESS} Cherenkov telescope system. Aharonian et al. (2006b) presented detailed spectral and morphological information of some new sources discovered in the \\textit{HESS} survey along with some discussions on possible counterparts in other wavelength bands. Some sources discovered in the survey can be associated with old SNRs. For example, HESS J1804$-$216 has a flux of nearly 25\\% of the flux of the Crab Nebula with a photon index of about 2.72. The source can be associated with either the shell-type SNR G8.7$-$2.1 or the young Vela-like pulsar PSR J1803$-$2137 (Aharonian et al. 2006b). Furthermore, Fatuzzo, Melia, \\& Crocker (2006) found that the SNR G8.7$-$2.1 can plausibly account for all known radiative characteristics of HESS J1804$-$216. Another case is HESS~J1834$-$087, which is an extended TeV source discovered in the \\textit{HESS} survey of the inner Galaxy in VHE $\\gamma$-ray band. The source, spatially coincident with the SNR G23.3$-$0.3, has a $\\gamma$-ray flux of 8\\% of the flux of the Crab Nebula above 200 GeV, and the $\\gamma$-ray spectrum can be expressed by a power law with a photon index of $2.45\\pm0.16$ (Aharonian et al. 2006b). The new \\xmm\\ observation reveals diffuse X-ray emission within the \\textit{HESS} source and suggests an association between the diffuse X-rays and the VHE $\\gamma$-rays, moreover, G23.3$-$0.3 is estimated to be an old SNR with a distance of $4\\pm0.2$ kpc and an age of $\\sim10^5$ yr (Tian et al. 2007).  Theoretically, it is believed that the non-thermal photon spectrum for a given SNR is produced through electron/positron synchrotron radiation, bremsstrahlung and inverse Compton scattering as well as $\\pi^0$ decay from the p-p interactions. The question is, however, whether the origin of the non-thermal photons for young SNRs is different from that for old ones. Sturner et al. (1997) presented a temporal evolution model to calculate the non-thermal particle and photon spectra from shell-type SNRs, and Zhang \\& Fang (2007) modified the model to investigate the possible hadronic contributions to TeV $\\gamma$-ray emission from young SNRs. For the young SNRs, the secondary particles produced via the p-p interaction have a negligible contribution to the non-thermal photon spectrum (e.g. Zhang \\& Fang 2007). On the other hand, Yamazaki et al. (2006) studied the emission from old SNRs and found that proton acceleration could be efficient enough to emit TeV $\\gamma$-rays both at the shock of the SNR and at the giant molecular cloud. They presented that the energy flux ratio $\\rm{R}_{\\rm{TeV/X}}$ can be more than $10^2$ for some SNRs, and such sources may be the origins of the recently discovered unidentified TeV sources. For the old SNRs, the energies of the primary electrons accelerated by the shock wave are limited by the synchrotron cooling, then the roll-off energy of the synchrotron emission of the primary electrons is much smaller than that for the young SNRs (Sturner et al. 1997; Yamazaki et al. 2006; Zhang \\& Fang 2007). However, the secondary $e^{\\pm}$ pairs can be produced from the p-p collisions in an SNR for a long time due to the large energy loss time for the p-p interaction. Therefore, the contributions of secondary electrons and positrons to the non-thermal emission can become prominent for the old SNRs.  Motivated by above discussions, we study the non-thermal photon spectra from old SNRs based on the time-dependent model in which the contribution of the secondary $e^{\\pm}$ pairs is included (see \\S 2.2). It should be noted that our model is different from that given by Yamazaki et al. (2006), the latter assumes a steady particle distribution. This paper is organized as follows. The temporally evolving model of the non-thermal particle and photon spectra are discussed in Section \\ref{sec:model}, including the SNR evolution, shock acceleration, and various photon production mechanisms involved in the model. The results from the applications of the model to two old SNRs, G8.7$-$0.1 and G23.3$-$0.3, are shown in Section \\ref{sec:application}. The main conclusions and some discussions are given in Section \\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion} In this paper, by taking into account the evolution of secondary $e^{\\pm}$ pairs produced via the p-p interaction due to high-energy protons collide with the ambient matter in an SNR, we developed a time-dependant model of multi-waveband non-thermal particle and photon spectra both for young and for old shell-type SNRs. The primary electrons accelerated directly by the shock wave usually dominate the secondary $e^{\\pm}$ pairs produced from the p-p collisions for a young SNR, however, the emission from the secondary $e^{\\pm}$ pairs becomes more and more important and eventually surpasses the radiation from the primary electrons as the SNR grows old because the source of the primary electrons is cut off and the synchrotron radiation loss is significant for electrons when the SNR is in the radiative phase, whereas the secondary $e^{\\pm}$ pairs can be produced for a long time due to the small cross section for the p-p interaction (see Fig. \\ref{fig:dtpu}). $K_{\\rm{ep}}$, which can not be directly deduced from the observations now and is usually limited by comparison of the model results with the multi-band observations for a given SNR, is also a crucial parameter to determine whether the contribution to the final non-thermal emission of the secondary $e^{\\pm}$ pairs dominates that of the primary electrons. For an old SNR with small $K_{\\rm{ep}}$, the non-thermal photons with energies $<10\\,\\rm{MeV}$ are usually produced mainly via the radiations of the secondary $e^{\\pm}$ pairs (see Fig. \\ref{fig:dtpu}).  The model is applied to two old shell-type SNRs, G8.7$-$0.1 and G23.3$-$0.3, whose VHE $\\gamma$-rays are detected in the \\textit{HESS} survey of the inner Galaxy. G8.7$-$0.1 is a shell-type SNR with an age of $\\sim15000$ yr. The SNR must be in the radiative phase to make the calculated results consistent with the \\textit{Chandra} X-ray observations. Finally, a set of parameters, a distance of 6 kpc, an age of 18000 yr, $M_{\\rm{ej}}=1.5\\,M_{\\odot}$, $n_{\\rm{ISM}}=6\\,\\rm{cm}^{-3}$, $K_{\\rm{ep}}=0.15$, and $B_{\\rm{ISM}}=6\\,\\mu\\rm{G}$ is used to model the multi-wavelength non-thermal emission from the SNR G8.7$-$0.1. It can be concluded that (1) the radio emission from the SNR is mainly from the synchrotron radiation of the primary electrons, whereas the non-thermal photons with energies around 0.1 keV primarily come from the synchrotron radiation of the secondary $e^{\\pm}$ pairs produced from the p-p collisions; (2) the observed X-ray photons with energies ranging from 2 to 10 keV are produced mainly via the bremsstrahlung of the primary electrons; (3) the TeV photons with energies $<1$ TeV are primarily from both the bremsstrahlung of the primary electrons and the p-p collisions of the primary protons, however, those with higher energies are produced mainly through the p-p interaction (see Fig. \\ref{fig:216}). Another old SNR is G23.3$-$0.3, which is an asymmetric shell-type SNR and spatially coincide with HESS~J1834$-$087. Tian et al. (2007) reported new H I observations from the VLA Galactic Plane System and a new \\xmm\\ observation for HESS~J1834$-$087, and argued that G23.3$-$0.3 is an old SNR with a distance of $4\\pm0.2$ kpc and an age of $\\sim10^5$ yr, moreover, a density of 6 cm$^{-3}$ was estimated around the SNR from the VGPs H I line emission associated with the SNR. We find that, using $n_{\\rm{ISM}}=6\\,\\rm{cm}^{-3}$ and an age of $\\sim10^5$ yr with other appropriate parameters, the model result agrees with the radio and VHE $\\gamma$-ray observations, but the calculated X-ray flux is significantly lower than the observed one given by \\xmm\\ (see Fig. \\ref{fig:dy087}). A smaller age and a smaller $n_{\\rm{ISM}}$ are needed to make the calculation result consistent with the multi-band observations. Finally, we use parameters a distance of 4 kpc, $M_{\\rm{ej}}=1.5\\,M_{\\odot}$, an age of $2.98\\times10^4$ yr, $n_{\\rm{ISM}}=1.5\\,\\rm{cm}^{-3}$, $B_{\\rm{ISM}}=8\\,\\mu\\rm{G}$, and $K_{\\rm{ep}}=0.05$ to model the multi-band non-thermal emission from the SNR G23.3$-$0.3. The modeling results can be consistent with both the radio, VHE $\\gamma$-ray observations and the X-ray flux with these parameters. Moreover, the emission from the secondary $e^{\\pm}$ pairs is negligible and the multi-wavelength photons with those parameters for the SNR mainly come from the radiation of the primary particles, and the TeV photons primarily have hadronic origin (see Fig. \\ref{fig:087}).  The model in this paper includes the contribution of the secondary $e^{\\pm}$ pairs produced from the p-p collisions to the non-thermal emission from SNRs. The secondary $e^{\\pm}$ pairs can be produced for a long time and experience energy losses and diffusion, and the contribution of the secondary $e^{\\pm}$ pairs to the final non-thermal emission from an SNR can become more and more important and eventually surpasses that of the primary electrons when the SNR ages in the radiative phase. Whether the contribution of the secondary $e^{\\pm}$ pairs to the final non-thermal emission is prominent or not depends heavily on the SNR age, $K_{\\rm{ep}}$ and $n_{\\rm{ISM}}$. These parameters can be limited by the comparison of the model result with the multi-wavelength observations for a given SNR, however, the shortage of precise observations and uncertainties such as the initial eject mass and the initial velocity of the shock lessen the reliability of the calculated results. On the other hand, for an old SNR with observed TeV emission, if the TeV photons primarily have hadronic origin, the emission from the secondary $e^{\\pm}$ pairs is usually prominent. The morphology of the resulting spectrum for an SNR in the energy range between 10 MeV to 1 TeV directly relates to the primary origin of the VHE photons. Fortunately, the Gamma Ray Large Area Space Telescope (\\textit{GLAST}) have a good sensitivity and resolution in the energy range between 30 MeV and 300 GeV, and the observations with \\textit{GLAST} will give a direct evidence on whether the TeV photons from an SNR mainly have hadronic origin or leptonic origin. Obviously, \\textit{GLAST} will give more limits on the parameters in the modeling process with the model in this paper for an SNR, and then we can get more insights on the acceleration and photon production processes involved in SNRs."
    },
    "0711/0711.1083_arXiv.txt": {
        "abstract": "We present in this paper an analysis of the faint and red near-infrared selected galaxy population found in near-infrared imaging from the Palomar Observatory Wide-Field Infrared Survey. This survey covers 1.53 deg$^{2}$ to 5 $\\sigma$ detection limits of  K$_{\\rm vega}$ = 20.5-21 and J$_{\\rm vega}$ = 22.5, and overlaps with the DEEP2 spectroscopic redshift survey. We discuss the details of this NIR survey, including our J and K band counts. We show that the K-band galaxy population has a redshift distribution that varies with K-magnitude, with most K $< 17$ galaxies at $z < 1.5$ and a significant fraction (38.3$\\pm0.3$\\%) of $K > 19$ systems at $z > 1.5$.  We further investigate the stellar masses and morphological properties of K-selected galaxies, particularly extremely red objects, as defined by $(R-K) > 5.3$ and $(I-K) > 4$.  One of our conclusions is that the ERO selection is a good method for picking out galaxies at $z > 1.2$, and within our magnitude limits, the most massive galaxies at these redshifts.  The ERO limit finds 75\\% of all M$_{*} > 10^{11}$ \\solm galaxies at $z \\sim 1.5$ down to K$_{\\rm vega} = 19.7$. We further find that the morphological break-down of $K < 19.7$ EROs is dominated by early-types (57$\\pm3$\\%) and peculiars (34$\\pm3$\\%).  However, about a fourth of the early-types are distorted ellipticals, and within CAS parameter space these bridge the early-type and peculiar population, suggesting a morphological evolutionary sequence.  We also investigate the use of a $(I-K) > 4$ selection to locate EROs, finding that it selects galaxies at slightly higher average redshifts ($<z> = 1.43\\pm0.32$) than the $(R-K) > 5.3$ limit with $<z> = 1.28\\pm0.23$.  Finally, by using the redshift distribution of $K < 20$ selected galaxies, and the properties of our EROs, we are able to rule out all monolithic collapse models for the formation of massive galaxies. ",
        "introduction": "Deep imaging and spectroscopic surveys in the optical have become the standard method for determining the evolution of the galaxy population (e.g., Kron 1980; Steidel \\& Hamilton 1993; Williams et al. 1996; Giavalisco et al. 2004). These surveys have revolutionised galaxy formation studies, and have allowed us to characterise basic properties of galaxies, such as their luminosity functions, stellar masses, and morphologies, and how these properties have evolved (e.g., Lilly et al. 1995; Ellis 1997; Wolf et al. 2003; Conselice et al. 2005a).    However, due to technological limitations with near-infrared arrays, most deep surveys have been conducted in optical light, typically between $\\lambda =$ 4000-8000 \\AA.  This puts limits on the usefulness of optical surveys, as they select galaxies in the rest-frame ultraviolet at higher redshifts where many of the galaxies contributing to the faint end of optical counts are located (e.g., Ellis et al. 1996).  Galaxies selected in the rest-frame ultraviolet limit our ability to trace the evolution of the galaxy population in terms of stellar masses. As the properties of galaxies can be quite different between the rest-frame optical and UV (e.g., Windhorst et al. 2002; Papovich et al. 2003, 2005; Taylor-Manger et al. 2006; Ravindranath et al. 2006), it is desirable to trace galaxy evolution at wavelengths where most of the stellar mass in galaxies is visible. To make advances in our understanding of galaxy evolution and formation at high redshifts, ($z > 1$), therefore requires us to search for, and investigate, galaxy properties in the near-infrared (NIR).  Studying galaxies in the NIR has many advantages, including minimised  K-corrections which are often substantial in the optical, as well as giving us a more direct probe of galaxy stellar mass up to $z \\sim 3$ (e.g., Cowie et al. 1994).  This has been recognised for many years, but most NIR surveys have been either deep pencil beam surveys (e.g., Gardner 1995; Djorgovski et al. 1995; Moustakas et al. 1997; Bershady, Lowenthal \\& Koo 1998;  Dickinson et al. 2000; Saracco et al. 2001; Franx et al. 2003), or large-area, but shallow, surveys (e.g., Mobasher et al. 1986; Saracco et al. 1999; McCracken et al. 2000; Huang et al. 2001; Martini 2001; Drory et al. 2001; Elston et al. 2006).  This is potentially a problem for understanding massive and evolved galaxies at high redshifts, as  red objects are highly clustered (e.g., Daddi et al. 2000; Foucaud et al. 2007), as are the most massive galaxies (e.g., Coil et al. 2004a).  Previous deep NIR surveys can detect these unique galaxy populations, but they typically do not have a large enough area to probe the range of galaxies selected in the near-infrared. Likewise, large area, but shallow surveys may not be deep enough to detect with a high enough signal to noise these unique populations. In this paper we overcome this problem by presenting a  1.5 deg$^{2}$ survey down to 5 $\\sigma$ depths of K$_{\\rm vega}$ = 20.2-21.5 and J$_{\\rm vega} = 22.5$. This  brings together the properties of both deep and wide surveys. In this paper we explore NIR galaxy counts, and study the properties of faint NIR galaxies, which are often red in near-infrared/optical colours.  Unique galaxy populations have long been known to exist in near-infrared selected surveys. These include the extremely red objects (Elston, Reike, Reike 1988) and the distant red galaxies (Saracco et al. 2001; Franx et al. 2003; Conselice et al. 2007a; Foucaud et al. 2007), both of which are difficult to study at optical wavelengths.  The existence of these galaxies reveals a large possible differential in the galaxy population between optical and near-infrared surveys. While these objects can be located in deep optical surveys, they are often very faint with $R > 26$ (e.g., van Dokkum et al. 2006), making it difficult to understand these objects in any detail without NIR imaging or spectroscopy. In this paper we analyse the properties of galaxies selected in moderately deep K-band imaging.  We also investigate the redshift distributions, structures and properties of the near-infrared selected galaxy population down to $K_{\\rm vega} \\sim 20$.  One of our main conclusions are that the faint K-band population spans a range of redshift and properties.  We find that galaxies with magnitudes $K_{\\rm vega} < 17$ are nearly all at $z < 1.4$.  Galaxies with magnitudes $K_{\\rm vega} = 17-21$ are found at high redshift, up to at least $z \\sim 4$.  The colours of these galaxies span a wide range, with in particular redder galaxies seen at higher redshifts.  Finally, we investigate the properties of the extremely red objects in our sample, finding that they include most, but not all of, the highest mass galaxies at $z \\sim 1.5$.  The morphologies of these EROs, and the redshift distribution and dust properties of $K_{\\rm vega} < 20$ sources, show that hierarchical galaxy formation is the dominate method by which massive galaxies form.  This paper is organised as follows: in \\S 2 we discuss the data sources used in this paper, including our Palomar imaging, DEEP2 redshifts, and HST ACS imaging. This section also gives basic details of the Palomar survey.  \\S 3 includes our analysis, which contains information on the K and J-band counts, the redshift and colour distributions of K-selected galaxies. \\S 4 includes an analysis of the extremely red galaxy population and its properties, including stellar mass, dust content, and redshift distributions. \\S 5 includes a detailed discussion of our results in terms of galaxy models, while \\S 6 is a summary of our findings.  This paper uses Vega magnitudes unless otherwise specified, and assumes a cosmology with H$_{0} = 70$~km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm m} = 0.3$ and $\\Omega_{\\lambda} = 0.7$. ",
        "conclusions": "Here we discuss the results from this paper in the context of galaxy formation models and scenarios. We include in this discussion the redshift distribution of K-selected galaxies, as well as the stellar mass and structural/morphological properties of EROs to address how the stellar mass assembly of galaxies is likely taking place.  By examining these galaxies we are not relying on assumptions about stellar masses to find the most massive and evolved galaxies at high redshift. In this sense a K-band selected and colour selected population are an alternative approach from stellar mass selection (Conselice et al. 2007b) for understanding galaxy formation due to the simplicity, and reproducibility, of their selection.  We first examine the redshift  distribution of $K < 20$ galaxies, and compare this to models. We use colour information of our faint K-selected galaxies to rule out all monolithic collapse formation scenarios for galaxies. We then examine the properties of the EROs themselves to further argue that these systems are, due to their stellar masses, likely an intrinsically homogeneous population, with the peculiar EROs evolving into the ellipticals at lower redshifts.  \\subsection{K-band Redshift Distributions}  The number of K-band selected galaxies per redshift at a given K-magnitude limit is an important test of when the stellar masses of galaxies were put into place.  In general, ideas for how massive galaxies form explicitly predict how much stellar mass galaxies would have in the past. In a rapid formation, such as with a monolithic collapse, the stellar mass for the most massive galaxies, which in a K-band limited sample will always probe the massive systems, the number of galaxies within a bright K-band limit, say $K < 20$ at high redshift should be much larger than the number of sources seen in a hierarchical model, which predicts that the stellar masses of galaxies grows with time. Therefore the number of bright galaxies at higher redshifts in a hierarchical model is less than that predicted in a monolithic collapse.  This test was first performed by Kauffmann \\& Charlot (1998) who concluded, based on an early hierarchical model, that the predicted counts exceed the observations, a result which has remained despite improvements in data and models (e.g., Kitzbichler \\& White 2006).   \\begin{figure} \\vbox to 90mm{ \\includegraphics[angle=0, width=84mm]{conselice.fig18.ps} \\caption{Redshift distribution of our sample with a $K < 20$ cut.  Comparisons to previous redshift distributions published in the K20 and GOODS surveys are shown.  The two lines demonstrate predictions of fiducial monolithic collapse (Kitzbichler \\& White 2006), and hierarchical model predictions (Kitzbichler \\& White 2007) for this evolution within the same K-band limit.} } \\label{sample-figure} \\vspace{2cm} \\end{figure}   In Figure~18 we show the number of $K < 20$ galaxies as a function of redshift for systems at $z < 4$. We further plot on Figure~18 model predictions for how $N(z)$ changes in a standard pure luminosity evolution monolithic collapse model from Kitzbichler \\& White (2006), and within a standard hierarchical formation scenario from Kitzbichler \\& White (2007).  We have used in this figure our entire K-band distribution, including galaxies for which we had to measure photometric redshifts without an optical band (\\S 2.2). We also plot on Figure~18 the $K < 20$ magnitude distribution for galaxies seen in the GOODS and K20 surveys.  While we generally agree with these previous results, we find a slightly lower number of systems at higher redshifts compared to GOODS and K20. The reason for this could simply be cosmic variance, as the GOODS and K20 samples at these redshifts also differ by a large amount. Otherwise, this difference is likely created by errors within our photometric redshifts. However, it must be noted that the integrated number of $K < 20$ galaxies is similar for our survey and GOODS and K20, but is still much smaller than that predicted by the monolithic model.  As can be seen in Figure~18 there is clearly a large difference between the observed distribution, and the predicted  monolithic collapse distribution. We can rule out this basic monolithic collapse model, which does not include dust, at $>$ 10 $\\sigma$, based on the comparisons to our $K < 20$ redshift distribution.  It is possible to match with a monolithic collapse model if extreme dust content is included in these galaxies, or if there are `hidden' galaxies (Kitzbichler \\& White 2006).  However, as we argue below, there is no evidence that dusty galaxies dominate our K-band selected sample.  From Figure~18 it appears that a basic hierarchical model agrees better with the data (e.g., Somerville et al. 2004; Stanford et al. 2004; Kitzbichler \\& White 2007).  There is perhaps a slight excess of galaxies in the hierarchical formation model, which has been noted before (e.g., Kauffman \\& Charlot 1998; Somerville et al. 2004;  Kitzbichler \\& White 2007).  The origin of this difference is not clear. It perhaps simply implies that too much mass is produced early in the hierarchical model, yet this would seem to contradict the fact that the most massive galaxies with M$_{*} >$ \\mass are nearly all in place by $z \\sim 1.5 - 2$ (Conselice et al. 2007b).  We address this apparent conflict in more detail in \\S 5.2.  Although the basic hierarchical model agrees better with the observed redshift distribution, there are scenarios where the monolithic model can fit the data as well as the hierarchical scenario (see also Cimatti et al. 2002a).   These scenarios require that the K-band selected galaxies have a significant amount of dust extinction in the observed K-band.  The amount of extinction needed varies from 1.0 to 0.7 mag at rest-frame $z$ and $R$ bands from $z = 1.5 - 2$ (Kitzbichler \\& White 2006).  This extinction needed is even higher at $z = 1 - 1.5$.  By using the colours of galaxies that fit within the $K < 20$ criteria, we can determine what the contribution of dusty galaxies are to these counts.  Stellar population synthesis models show that only old stellar populations, and galaxies with dust extinctions with A$_{\\rm v} > 1$, have a colour of $(R-K) > 5$ at $z = 1 - 4$.  We can use this information, and the model results using various dust extinctions from Nagamine et al. (2005) (Figures~7 and 8), to argue that these galaxies are not dominated by dust extinction, which would need to be the case to reconcile the observed K-band distribution with a monolithic collapse model.  First, the model results compared to data shown in Figures~7 and 8, and discussed in \\S 4.3 clearly show that K-selected galaxies are not dominated by heavy dust extinction. At best, only models with E(B-V) = $0 - 0.4$ are able to reproduce the location of real galaxies. The E(B-V) = 0.4 models are however even too red to account for most galaxies in the $(R-K)$ vs. M$_{*}$ diagram (Figure~8).  A representative value of R$_{\\rm V}$ = A$_{\\rm V}$/E(B-V) = 3.1 reveals values of A$_{\\rm V} = 0 - 1.28$ for this range in E(B-V). For the Calzetti et al. (2000) extinction law, this gives lower values, with A$_{\\rm V} = 0 - 0.25$. Thus, it does not appear likely that our $K < 20$ galaxy sample is dominated by enough dust extinction (A$_{\\rm V} = 1$ needed, on average) to match the monolithic collapse galaxy count model.  We find that only 28.3$\\pm0.6$\\% of $K < 20$ selected galaxies at $1 < z < 2$ have a colour $(R-K) > 5$, required to meet the minimum condition for dust extinction.  We have further argued in \\S 4, that a significant fraction of these EROs are old passively evolving stellar populations, which are unlikely to have a screen of dust with A$_{\\rm V} > 1$ (see e.g., White, Keel \\& Conselice 2000). As discussed in \\S 4.5.2 over half of the EROs at $1 < z < 2$ appear as early-types, thus at most only $\\sim 14\\%$ of the $K < 20$ galaxies at $1 < z < 2$ have a significant amount of dust extinction. The dusty pure-luminosity evolution monolithic collapse models from Kitzbichler \\& White (2006) predict that all $K < 20$ galaxies have this amount of dust extinction, which clearly they do not.  It is thus impossible to reconcile the K-band redshift distribution with any monolithic model.  In summary our K-band redshift distribution, and those from the K20 and GOODS surveys, appear to be much closer to the basic hierarchical model of Kitzbichler \\& White (2007) than to the monolithic collapse model.  We can also rule out pure luminosity evolution models with significant dust extinction. Therefore, down to $K = 20$ at $z < 4$ it appears that the monolithic model cannot be the dominate method for forming galaxies.   \\subsection{Structural Evolution}  As discussed in \\S 4.5, a large fraction of the EROs we have HST imaging appear to be undergoing some type of evolution based on their structural appearance. A logical conclusion is that many of the peculiar systems are in some phase of a merger (Conselice et al. 2003a).  We can quantify this by examining the merger fraction as derived through the CAS definition of $A > 0.35$ and $A > S$ (Conselice 2003). This is a strict definition, and will allow us to measure a lower limit on the merger fraction, even when observing in the rest-frame ultraviolet (Taylor-Mager et al. 2007).  \\begin{figure} \\vbox to 80mm{ \\includegraphics[angle=0, width=84mm]{conselice.fig19.ps} \\caption{The fraction of galaxies identifiable as a merger using the CAS system, and the fraction of early-types which appear to have a distorted structure.} } \\label{sample-figure} \\vspace{2cm} \\end{figure}  Figure~19 shows the evolution of the merger fraction derived through this CAS definition. We also include the fraction of early-types which appear to have a distorted structure. It appears from this that between $5-10\\%$ of the EROs at $1 < z < 2$ would be identified as a major merger through the CAS parameters. This is a factor of $\\sim 3$ less than the peculiar fraction (Figure~13), which is what we would expect for a population which is undergoing major mergers (Conselice 2006; Bridge et al. 2007).  The reason for this is that the CAS approach is sensitive to $\\sim 0.4$ Gyr of the merger process (Conselice 2006), while eye-ball estimates of merging last for $\\sim 1.2$ Gyr (Conselice 2006). This is another indication that the peculiar galaxies we are seeing in the ERO population are in fact undergoing merging.  Figure~20, as well as Figure~15 demonstrate what evolution is likely occurring within the ERO population. As can be seen, the EROs, within our $K < 19.7$ selection, have the same upper range of masses at $z \\sim 0.8 - 2$ (Figure~20). Therefore, little mass growth occurs within this population at this K-band limit.  We also know that a large fraction of the EROs with M$_{*} >$ \\mass are peculiar in some way. This varies from 64$\\pm$24\\% at $z > 1.5$ to 41$\\pm$6\\% at $1 < z < 1.25$ (\\S 4.5).  However, at $z \\sim 0$ the fraction of M$_{*} >$ \\mass red galaxies which are morphologically early-types is $\\sim 85$\\% (Conselice 2006a; Conselice et al. 2007b).  It is clear that a significant fraction of EROs must have undergone morphological evolution as they cannot lose mass.   What likely is occurring is that the distorted elliptical and peculiar galaxies which dominate the population at $z > 1.5$ (Figure~13 and 20) transform into morphologically and spectrally evolved systems at $z \\sim 1$.  The reason this is likely the case is that there is no difference in the masses of the peculiar and the early-type EROs.  This also explains why this population, despite having a mixed morphology clusters so strongly (e.g., Daddi et al. 2000; Roch et al. 2003).  Calculations based on the merger rate in Conselice et al. (2007b) for the most massive galaxies suggest that on average about one or two major mergers are occurring for the M$_{*} >$ \\mass population at $z < 2$, but fewer at lower redshifts.  Finally, these major mergers are what may make the K-band counts in the hierarchical model redshift distribution higher than the observed.  The reason is that in the standard hierarchical model, the number and mass densities of the most massive galaxies with M$_{*} >$ \\mass underpredict the observed number of massive systems by up to two orders of magnitude (Conselice et al. 2007b).  Within these models the stellar masses for these systems are largely already formed, but are in distinct galaxies that have not yet merged (De Lucia et al. 2006).  It is thus easy to see that if a single massive galaxy at $z \\sim 1.5$ was in several pieces, all of which would still meet the criteria of $K < 20$ based on the relation of stellar mass and K-mag (Figure~7), then the number of galaxies with $K < 20$ at higher redshift in the hierarchical model would be higher than the observed number. This is consistent with the number densities of massive galaxies being higher than in the models, as well as for a rapid formation of massive systems through major mergers at $z > 2$ (Conselice 2006b).  \\begin{figure} \\vbox to 80mm{ \\includegraphics[angle=0, width=84mm]{conselice.fig20.ps} \\caption{The distribution of ERO galaxy mass with redshift. Labeled on this figure are the morphological types for each galaxy as determined through our ACS imaging. Note that massive EROs are detected all the way out to $z = 2$.} } \\label{sample-figure} \\vspace{2cm} \\end{figure}"
    },
    "0711/0711.1560_arXiv.txt": {
        "abstract": "The use of Type Ia supernovae (SNe Ia) as cosmological standard candles is a key to solving the mystery of dark energy. Improving the calibration of SNe Ia increases their power as cosmological standard candles. We find tentative evidence for a correlation between the late-time lightcurve slope and the peak luminosity of SNe Ia in the $B$ band; brighter SNe Ia seem to have shallower lightcurve slopes between 100 and 150 days from maximum light. Using a Markov Chain Monte Carlo (MCMC) analysis in calibrating SNe Ia, we are able to simultaneously take into consideration the effect of dust extinction, the luminosity and lightcurve width correlation (parametrized by $\\Delta m_{15}$), and the luminosity and late-time lightcurve slope correlation. For the available sample of 11 SNe Ia with well-measured late-time lightcurves, we find that correcting for the correlation between luminosity and late-time lightcurve slope of the SNe Ia leads to an intrinsic dispersion of 0.12 mag in the Hubble diagram. Our results have significant implications for future supernova surveys aimed to illuminate the nature of dark energy. ",
        "introduction": "One of the most important problems in cosmology today is to solve the mystery of dark energy, the unknown reason for the observed cosmic acceleration \\citep{Riess98,Perl99}. The discovery of cosmic acceleration was made using Type Ia Supernovae (SNe Ia), as SNe Ia can be calibrated to be good cosmological standard candles \\citep{Phillips93,Riess95}. Improving the calibration of SNe Ia increases their power as cosmological standard candles.  A SN~Ia is a thermonuclear explosion that completely destroys a carbon/oxygen white dwarf as it approaches the Chandrasekher limit of 1.4 $M_{\\odot}$.\\footnote{The exception recently found by \\cite{Howell06} is super-luminous and can be easily separated from the normal SNe~Ia used for cosmology.} This is the reason SNe~Ia can be calibrated to be good standard candles. The first challenge to overcome when using SNe~Ia as cosmological standard candles is properly incorporating the intrinsic scatter in SN~Ia peak luminosity. The usual calibration of SNe~Ia reduces the intrinsic scatter in SN~Ia peak luminosity to about 0.17~mag \\citep{Phillips93,Riess95}. The calibration techniques used so far are based on one observable parameter, the lightcurve width, which can be parametrized either as $\\Delta m_{15}$ (decline in magnitudes in the $B$-band for a SN~Ia in the first 15 SN Ia restframe days after $B$-band maximum, see \\cite{Phillips93}), or a stretch factor (which linearly scales the time axis, see \\cite{Goldhaber01}). The lightcurve width is associated with the amount of $^{56}$Ni produced in the SN~Ia explosion, which in turn depends on when the carbon burning makes the transition from turbulent deflagration to a supersonic detonation \\citep{Wheeler03}.  There are theoretical reasons to expect the existence of additional physical parameters that characterize the lightcurve of a SN~Ia. For example, \\cite{Milne01} modeled the late light curves of SNe~Ia, and found that the nonlocal and time-dependent energy deposition due to the transport of Comptonized electrons can produce a 0.10 to 0.18~mag correction to the late-time brightness of SNe~Ia, depending on their position in the $\\Delta m_{15}$ sequence.  In this paper we show that late-time lightcurves of SNe Ia can be used to tighten the calibration of SNe~Ia, and reduce their effective (post-calibration) intrinsic scatter in the Hubble diagram. We describe our method in Section 2, and present our results in Section 3. We summarize and conclude in Section 4. ",
        "conclusions": "\\label{sec:discussion}  Improving the calibration of SNe Ia increases the power of SNe Ia as cosmological standard candles. We have shown the tentative evidence for a correlation between SN Ia peak luminosity and late-time lightcurve slope in the $B$ band; brighter SNe Ia have shallower lightcurve slopes between 100 and 150 days from maximum light (see Fig.8). Using this correlation, we are able to reduce the intrinsic scatter of SNe Ia (measured by the Hubble diagram residue) to 0.12 mag (see Fig.7).  We use the MCMC algorithm in our likelihood analysis; this allows us to explore all the parameters simultaneously and efficiently in calibrating the SN Ia peak luminosity for our sample of SNe Ia. We only used $\\Delta B_{max}$ in calculating $\\chi^2$, since including the errors in color, $\\Delta m_{15}$, and late-time lightcurve slope introduces spurious pdfs in the MCMC analysis. Our results are not sensitive to the uncertainty in $B_{max}$ (setting $\\Delta B_{max}$ to a constant gives similar results to using the $\\Delta B_{max}$ from Table 1). The parameter set is ($M$, ${\\cal R}$, $\\alpha$, $\\alpha\\prime$, $\\beta$), which contains the nuisance parameter and correction coefficients for color, $\\Delta m_{15}$, and late-time lightcurve slope respectively. The pdf's we have found for (${\\cal R}$, $\\alpha$, $\\alpha\\prime$, $\\beta$) are very close to being Gaussian (see Fig.1 and Fig.3), except for the pdf's for $\\alpha\\prime$ (which are cutoff by the prior of $\\alpha\\prime\\geq -2$). Late-time lightcurve slope $b_{lt}$ appears uncorrelated with color (see Fig.2), which supports $b_{lt}$ being an intrinsic observable that correlates with the peak luminosity of SNe Ia.  The apparent correlation between peak luminosity and late-time lightcurve slope in the $B$ band that we have found is consistent with the results by \\cite{Capp97}.  They studied the late-time lightcurves of five SNe Ia in the $V$ band with the focus of modeling lightcurves at very late times (later than 150-200 days from maximum light), and found a similar trend in the peak luminosity and late-time lightcurve slope (which they parametrized with the difference in $V$ magnitude from maximum to 300 days, see Table 1 of \\cite{Capp97}).  The coefficients for correcting for the correlation between decline-rate and peak luminosity that we have found are very different from previous work (see Table 2 and \\cite{Phillips93,Hamuy96a,Phillips99,Lifan06}). This is likely partly a consequence of the particular sample of 11 SNe Ia we have used for our analysis. This is not surprising, since previous work by various groups also found very different $\\Delta m_{15}$ corrections \\citep{Phillips99,Lifan06}, indicating that this correction is highly sensitive to sample selection.  Note that the color correction coefficient we found was not sensitive to whether we correct for late-time lightcurve slope, and is consistent with the color correction coefficient of ${\\cal R}=2.59\\pm 0.24$ found by \\cite{Lifan06}. This is very different from a coefficient of ${\\cal R}=4.1$ that is typical of Milky Way dust. We found that fixing ${\\cal R}=4.1$ gives a Hubble diagram residue almost twice as large as the Hubble diagram residue found when allowing ${\\cal R}$ to vary (see Fig.7 and its caption). This strengthens the evidence that the color correction coefficient ${\\cal R}$ of SN data cannot be explained by extinction by typical Milky-Way dust. The deviation of ${\\cal R}$ from ${\\cal R}=4.1$ could be due to the mixing of intrinsic SN Ia color variation with dust extinction \\citep{Conley07}, or variations in the properties of dust. The definitive resolution of this issue will likely require NIR observations of a large number of SNe Ia \\citep{KK04,Phillips06}.   Our results show the benefit of extending the observations of SN Ia lightcurves beyond 100 days in the SN Ia restframe, with several data points between 50 to 150 days to enable a robust measurement of the late-time lightcurve slope. Our sample only contains 11 SNe Ia, limited by currently available data. The expansion in the scope of SN observations to obtain late-time lightcurves is straightforward for surveys which cover the same areas in the sky in regular intervals \\citep{Wang00,Phillips06}.  Our results have significant implications for the survey strategy of future SN surveys, such as those planned for the Advanced Liquid-mirror Probe for Astrophysics, Cosmology and Asteroids (ALPACA)\\footnote{http://www.astro.ubc.ca/LMT/alpaca/}; Pan-STARRS\\footnote{http://pan-starrs.ifa.hawaii.edu/}; the Dark Energy Survey (DES)\\footnote{http://www.darkenergysurvey.org/}; the Large Synoptic Survey Telescope (LSST)\\footnote{http://www.lsst.org/}, and the Joint Dark Energy Mission (JDEM).   \\bigskip  {\\bf Acknowledgments} We are grateful to Lifan Wang for providing the lightcurve fits to the SN data we used, Alex Conley for sending us his compilation of nearby SN data, and David Branch for helpful discussions and comments on drafts of the manuscript. This work was supported in part by the NSF CAREER grant AST-0094335 and AST 05-06028."
    },
    "0711/0711.2563_arXiv.txt": {
        "abstract": "We solve for the behaviour of cosmological perturbations in the Dvali-Gabadadze-Porrati (DGP) braneworld model using a new numerical method.  Unlike some other approaches in the literature, our method uses no approximations other than linear theory and is valid on large scales.  We examine the behaviour of late-universe density perturbations for both the self-accelerating and normal branches of DGP cosmology.  Our numerical results can form the basis of a detailed comparison between the DGP model and cosmological observations. ",
        "introduction": "In the braneworld paradigm our universe is a lower-dimensional object embedded in a higher-dimensional bulk spacetime. Standard Model fields are assumed to be confined to the brane while gravity is allowed to propagate in the bulk. The Dvali-Gabadadze-Porrati (DGP) \\cite{Dvali:2000rv,Dvali:2000hr} model postulates that we live in a 4-dimensional hypersurface in a 5-dimensional Minkowski bulk. General Relativity (GR) is recovered at small scales (smaller than the crossover scale $r_\\c$) due to the inclusion of an induced gravity term in the action.  It was quickly realized that this model has two distinct classes of cosmological solutions \\cite{Deffayet:2000uy}. One of them exhibits accelerated expansion at late times without the need to include any exotic cosmological fluids, such as dark energy, or any brane tension that acts as an effective 4-dimensional cosmological constant. Hence, this branch of solutions is called ``self-accelerating''. Several attempts to confront the self-accelerating universe with observations have been made \\cite{Fairbairn:2005ue, Maartens:2006yt, Song:2006jk} (also see \\cite[review]{Koyama:2007rx} and references therein). To explain the observed acceleration we require $r_\\c \\sim H_0^{-1}$, where $H_0$ is the current value of the Hubble parameter.  It is expected that structure formation will help to distinguish the self-accelerating DGP universe from dark energy models based on 4-dimensional GR. This is because the growth of cosmological perturbations is very sensitive to the existence of an extra dimension.  A full 5-dimensional treatment is required to model these perturbations, which is why obtaining observational predictions for the behaviour of fluctuations in the DGP model is technically challenging.  Several authors have considered the problem of the dynamics of perturbations in the DGP model, but they have all relied on some sort of approximation or simplifying \\emph{ansatz}.  Two examples of this are the quasi-static (QS) approximation developed by \\citet{Koyama:2005kd} and the dynamical scaling (DS) \\emph{ansatz} proposed by \\citet{Sawicki:2006jj}.  The former approximation scheme solves the perturbative equations of motion by focussing on the extreme subhorizon regime.  In contrast the DS method, which assumes that perturbations evolve as power laws of the scale factor with time-varying power law indices, is supposed to be valid on all scales.  It has been shown that the DS solution approaches the QS solution for subhorizon perturbations.   In this paper we present a complete numerical analysis of the evolution of scalar perturbations in the DGP model.  Mathematically, the problem involves the solution of a partial differential equation in the bulk coupled to an ordinary differential equation on the brane.  A numerical method for dealing with such systems has previously been developed for cosmological perturbations in the Randall-Sundrum (RS) model \\cite{Cardoso:2006nh,Cardoso:2007zh}. However, the DGP problem is more complicated than the RS case due to a non-local boundary condition on the bulk field. Hence, the algorithm used in this paper represents a significant generalization of the one used in Refs.~\\cite{Cardoso:2006nh,Cardoso:2007zh}.   Unfortunately, some theoretical issues cast doubt on the validity of the self-accelerating DGP solutions.  Specifically, the existence of a perturbative ghost perhaps suggests that this solution cannot describe our Universe \\cite{Koyama:2005tx, Gorbunov:2005zk}, though there has been a debate on the physical implication of the ghost mode (see \\cite[review]{Koyama:2007za} and references therein). However, as alluded to above there is another ``normal'' branch of solutions in the DGP model. We cannot explain the late time accelerated expansion of the Universe using the normal branch without including an effective cosmological constant induced by the brane tension $\\sigma$. However, by allowing a non-zero $\\sigma$, normal branch solutions can mimic dark energy models with the equation of state $w$ smaller than $-1$ \\cite{Sahni:2002dx, Lue:2004za}. Cosmological constraints on the background dynamics of these models have been studied \\cite{Lazkoz:2006gp, Lazkoz:2007zk}. Unlike 4-dimensional models that realize $w < -1$ by the introduction of a phantom field, the normal branch of DGP cosmology is ghost-free \\cite{Charmousis:2006pn}. This unique feature is what motivates us to numerically study the perturbations of the normal branch of DGP cosmology in the penultimate section of this paper. During the preparation of this manuscript, we learned that the behaviour of perturbations in the normal branch has been independently obtained by \\citet{Song0} using the DS method.   The structure of the paper is as follows: The background cosmology of the DGP model is discussed in \\S\\ref{sec:background}.  In \\S\\ref{sec:master}, we express the equations of motion for scalar and tensor perturbations of the DGP model in the dimensionless canonical form introduced in \\S\\ref{sec:canonical}. The numerical method used to solve these canonical equations of motion is developed in \\S\\ref{sec:numeric}.  Our algorithm is tested in \\S\\ref{sec:tensor results}, where we numerically recover analytic results for the behaviour of tensor perturbations in matter-free DGP models.  In \\S\\ref{sec:scalar SA} and \\S\\ref{sec:scalar NB}, we solve the scalar perturbations problem in the self-accelerating and normal branches, respectively.  Finally, \\S\\ref{sec:conclusions} is reserved for our conclusions. ",
        "conclusions": "\\label{sec:conclusions}  In this paper, we have presented numeric solutions for cosmological perturbations in the DGP braneworld model both in the self-accelerating and the normal branches.  We extended the algorithm developed for the Randall-Sundrum (RS) model to handle the nonlocal boundary conditions characteristic of the DGP model. The numerical code was tested for tensor perturbations and the agreement with the analytic solutions was found to be excellent.  We confirmed that on small scales $k > 0.01 h$ Mpc${}^{-1}$, the quasi-static (QS) approximation reliably predicts the evolution of perturbations with relative errors less than around $5\\%$ at $z=0$. Our results are quite insensitive to the initial conditions as long as we start our simulations early enough.  On larger scales, the potential $\\Phi_-$, which determines the integrated Sach-Wolfe (ISW) effect, shows more suppression than the QS prediction. We find that our numerical solutions agree well with the dynamical scaling (DS) solution both in the self-accelerating and normal branches, except in the asymptotic de Sitter phase of the normal branch where the dynamical scaling solution fails to exist \\cite{Song}.  Our numeric solutions provide the basis for studying observational signatures of the model, especially in the normal branch where the influence of the extra dimension on the evolution of large scale structure has not yet been explored."
    },
    "0711/0711.0799_arXiv.txt": {
        "abstract": "{}{ To study the variation of the angular momentum and the rotational kinetic energy of the Sun, and associated variations in the gravitational multipole moments, on a timescale of the solar cycle.} {Inverting helioseismic rotational splitting data obtained by the Global Oscillation Network Group and by the Michelson Doppler Imager on the Solar and Heliospheric Observatory.} {The temporal variation in angular momentum and kinetic energy at high latitudes ($>\\pi/4   $) through the convection zone is positively correlated with solar activity, whereas at low latitudes it is anticorrelated, except for the top 10\\% by radius where both are correlated positively.} {The helioseismic data imply significant temporal variation in the angular momentum and the rotational kinetic energy, and in the gravitational multipole moments.  The properties of that variation will help constrain dynamical theories of the solar cycle.} ",
        "introduction": "Helioseismic data from the Global Oscillation Network Group (GONG) and the Michelson Doppler Imager (MDI) have been used to infer the rotation rate in the solar interior (Thompson et al.~1996; Schou et al.~1998). With the accumulation of these data for about 11 years it has now become possible to study its temporal variations. Previous studies have revealed a distinct pattern of bands of faster and slower than average rotation (Howe et al.~2000a; Antia \\& Basu 2000), which at low latitudes move towards the equator, and at high latitudes move towards the poles (Antia \\& Basu 2001). This pattern, which is similar to the torsional oscillations observed at the solar surface (Howard \\& LaBonte 1980), has been well studied, and it has been found to penetrate through much of the body of the convection zone (Vorontsov et al.~2002; Basu \\& Antia 2003; Howe et al.~2005, 2006). Below the convection zone, however, there is no generally accepted unambiguous detection of significant temporal variation in the rotation, although Howe et al.~(2000b, 2007) have reported a periodic variation near the equator between $r=0.65R_{\\odot}$ and $r=0.75R_{\\odot}$  in the early years of the last sunspot cycle with a period of 1.3 years, and Gough (2007) has reported that there is evidence that the oscillation penetrates even more deeply. Other studies (Antia \\& Basu 2000; Basu \\& Antia 2003) have not confirmed such a variation.  A direct consequence of temporal variation in the angular velocity is that global quantities like angular momentum and kinetic energy of rotation must also exhibit variation with time (cf., Komm et al.~2003). Because of the steep increase in density with depth, the major contribution to these global quantities is naturally likely to arise from the deeper layers, where a temporal variation in rotation has not been determined reliably; and as a result, it is difficult to draw any firm conclusion about the global variation of angular momentum and kinetic energy. However, following Komm et al.~(2003), we may consider the contributions from different layers within the convection zone separately.  We hope that such a study will shed some light on the rotational dynamics and its role in driving the activity cycle. Indeed, the study we present here is very similar to that of Komm et al.~(2003), although we are now able to extend it over an entire solar cycle. Apart from angular momentum and kinetic energy, we study also the temporal variation of the gravitational quadrupole and higher-order multipole moments due to the distortion in the solar figure induced by the known rotation.  In this work we concentrate principally on temporal variations in angular momentum and rotational kinetic energy in different regions in the solar convection zone, and compare their variation with the solar activity.  We describe in Section 2 the basic technique and the data used in the study; the main results are described in Section 3. In Section 4 we present the variations in the multipole moments of the external gravitational potential.  We draw our conclusions in Section 5. ",
        "conclusions": "We have studied the temporal variation in the rotational kinetic energy and angular momentum of the Sun in different regions of the convection zone. We find that in the outer layers ($r\\ga 0.9R_\\odot$) the variation is well correlated with solar activity. At high latitudes this correlation continues throughout the convection zone, while at low latitudes (and also when integrated over all latitudes), there is a transition to anticorrelation around $r=0.9R_\\odot$. Through the bulk of the convection zone the high latitudes seem to be rotating faster during the solar activity maximum, and the equatorial regions rotate more slowly.  It is difficult to draw any definite conclusions about the temporal variation below the convection zone, however, because the errors are large.   \\begin{table} \\begin{center} \\caption{Rotational kinetic energy, thermal energy and the estimated rms magnetic field using Eq.~(5) in different layers} \\begin{tabular}[h]{@{}ccrc} \\hline \\noalign{\\smallskip} Radial extent&$\\!\\!\\!$Mean rot.~KE&rms~$ B$\\hfill&Thermal energy\\\\ &(erg)&(G)&(erg)\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} $0.99<r/R_\\odot<1.00$&$5.3\\times10^{37}$&250&$6.2\\times10^{40}$\\\\ $0.95<r/R_\\odot<0.99$&$6.5\\times10^{39}$&1400&$2.4\\times10^{43}$\\\\ $0.85<r/R_\\odot<0.95$&$9.3\\times10^{40}$&3700&$1.4\\times10^{45}$\\\\ $0.75<r/R_\\odot<0.85$&$2.0\\times10^{41}$&6100&$7.6\\times10^{45}$\\\\ $0.65<r/R_\\odot<0.75$&$2.6\\times10^{41}$&7900&$2.2\\times10^{46}$\\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{center} \\end{table}  The amplitude of the temporal variation in the rotational kinetic energy integrated over the entire convection zone is of order of $3\\times10^{38}$ erg.   This implies a rate of variation whose amplitude is about $5\\times10^{30}$ erg s$^{-1}$. It is not clear if this variation is compensated for by opposite variation in deeper layers or whether it represents exchange of energy between rotation and magnetic field or some other form. It is interesting to note that this value is roughly comparable to (actually about 3 times) the amplitude of the observed variation in the solar irradiance, which is about 0.04\\%, although the correspondence might well be fortuitous.  It is, however, tempting to speculate that the temporal variation in rotational kinetic energy, and the magnetic energy, has some role to play in the irradiance variation. A part of the variation in the rotational kinetic energy would be expected to be compensated by the magnetic energy stored in the solar convection zone. If that were the case at all radii, we can estimate the spatial variation of the rms magnetic field using the relation \\begin{equation} {B^2\\over 8\\pi}4\\pi r^2\\delta r\\approx \\delta T. \\end{equation} The total variation $\\delta T$ in $T$ over the solar cycle is about 0.002 $T$, where $T(r_1,r_2)$ is the mean rotational kinetic energy in the region between radii $r_1$ and $r_2$.  Using this, together with the mean kinetic energy integrated over all latitudes within several spherical shells, we list in Table~2 these admittedly very crude estimates of the magnetic field.  The gravitational quadrupole moment $J_2$ does not exhibit any significant temporal variation, whereas some other, higher-order, multipole moments do. The variation in $J_4$ is correlated with solar activity, but the amplitudes of variation seen in GONG and MDI data are somewhat different. Clear variation in $J_8$ -- $J_{12}$ is inferred from both GONG and MDI data, the variation in $J_{10}$ being correlated with solar activity and that in $J_8$ and $J_{12}$ anticorrelated. The absence of temporal variation in $J_2$ is consistent with observed absence of variation in the $P_2(\\cos\\theta)$ component of distortion at the solar surface (Kuhn et al.~1998).  This is to be expected, since $J_2$ is determined essentially by the spherically averaged component of the angular velocity, which suffers little relative variation. The higher-order multipole moments are determined also by the latitudinally varying components of the angular velocity, which are known to suffer temporal variation correlated with solar activity. Note that the predominant contribution to $J_2$ arises from regions near the base of the convection zone (Gough, 1981; Pijpers 1998), whereas the higher-order multipole moments arise largely from contributions from outer layers of the Sun, where the temporal variation in the rotation rate is not insubstantial.  Of course, the magnitude of $J_2$ is much larger than that of the higher multipole moments, and it is likely that $J_2$ also has temporal variations with comparable absolute magnitude."
    },
    "0711/0711.0750_arXiv.txt": {
        "abstract": "We examine the possibility of observing gravitational lensing in the weak deflection regime by the supermassive black hole in the center of the galaxy M31. This black hole is significantly more massive than the black hole in the center of our Galaxy qualifying itself as a more effective lens. However, it is also more distant and the candidate stellar sources appear consequently fainter. As potential sources we separately consider stars belonging to the bulge, to the disk, to the triple nucleus formed by P1+P2 and by the recently discovered inner cluster P3. We calculate the number of simultaneously lensed stars at a given time as a function of the threshold magnitude required for the secondary image. For observations in the K-band we find 1.4 expected stars having secondary images brighter than $K=24$ and 182 brighter than $K=30$. For observations in the $V$-band we expect 1.3 secondary images brighter than $V=27$ and 271 brighter than $V=33$. The bulge stars have the highest chance to be lensed by the supermassive black hole, whereas the disk and the composite nucleus stars contribute by 10$\\%$ each. The typical angular separation of the secondary images from the black hole range from 1 mas to $0.1''$. For each population we also show the distribution of the lensed sources as a function of their distance and absolute magnitude, the expected angular positions and velocities of the generated secondary images, the rate and the typical duration of the lensing events. ",
        "introduction": "As soon as the evidence of a supermassive black hole in the center of the Milky Way (identified with the radio source Sgr A*) became overwhelming, the scientific community started to investigate its ability of acting as a gravitational lens for background stellar sources. The first work by \\citet{WarYus} considered lensing of sources in the immediate environment of Sgr A* and showed that the background sky should appear slightly depleted by a lensing effect. They suggested the possibility that gravitational lensing could be strong enough to generate an observable secondary image for several stars at a given time. More precise estimates by \\citet{Jar}, \\citet{AleSte}, and \\citet{CGM}, point out that about 10 stars belonging to the Galactic bulge should give rise to secondary images with $K<23$ at a given time. Deep and high resolution images in the near infrared of the Galactic center are being obtained by several advanced observatories such as Keck, VLT, and Gemini North. The progress registered by these observations encourages to search for signatures of lensing effects by Sgr A*. For example, identifying the two images of a background source \\citep{Ale}, or measuring the astrometric shift of the background stars due to gravitational lensing \\citep{NusBro}, could help to determine the position of Sgr A* very accurately. Enhanced microlensing caused by black holes surrounding Sgr A* has also been considered \\citep{AleLoe,CGM}. A fascinating possibility is offered by stars orbiting very close to Sgr A*, which have now been followed very accurately across several years \\citep{EckGen,Genzel,Ghez1,Eckart,Sch1,Sch2,Ghez2,Ghez3,Weinberg,Eis,Paumard,Reid}. In fact, the precise knowledge of the position of the source allows to predict the time, the position and the brightness of the secondary image \\citep{DeP,BozMan1,BozMan2}.  Meanwhile, our knowledge about the central regions of the Andromeda galaxy (M31) has grown considerably. \\citet{Kent} has provided a detailed brightness and density profile for the bulge and the disk of M31, which is still the reference for present works. \\citet{Lauer93} discovered that the nucleus of M31 is actually constituted by two components, which they called P1 and P2. \\citet{King95} noticed that P2 is much brighter than P1 in the ultraviolet. Then \\citet{Tremaine} hypothesized that P1 and P2 are actually parts of the same eccentric disk orbiting a supermassive black hole, with P1 being formed by stars at the apocenter and P2 being formed by stars at the pericenter. Later on, \\citet{Bender05} discovered a bright cluster of young stars (named P3) embedded within P2. They also managed to give the most precise estimate for the mass of the central black hole of M31, $M=1.4^{+0.9}_{-0.3}\\times 10^8$ M$_\\sun$. Recently, \\citet{DemVir} have cast doubt on the hypothesis that the P3 cluster is made of young stars of spectral class A5 -- B5, proposing that it is actually made of old stars. On the other hand, \\citet{CMCQ} have supported the young-stars hypothesis suggesting a mechanism for the refuelling of gas in the neighborhood of the central black hole.  It is very interesting to note that the supermassive black hole in the center of M31 is nearly two orders of magnitudes heavier than the black hole in our Galaxy, estimated to $3.6\\times 10^6$ M$_\\sun$ by \\citet{Eis}. From the lensing point of view, the greater distance of M31 with respect to Sgr A* is therefore partly compensated by the larger mass of its central black hole. This motivates a deep investigation of the probability of having stellar sources lensed by the supermassive black hole in M31.  Gravitational lensing effects have been successfully employed in several contexts as a tool to study the structure of galaxies, large-scale structures and cosmological parameters \\citep{SaasFee}. As we will show in this paper, it is not unlikely that gravitational lensing by the supermassive black hole in M31 will follow the wake of its ancestors. With a conspicuous number of events it will be possible to undertake a precise reconstruction of the mass distribution in the inner core of M31. Gravitational lensing would then provide an independent and unbiased method to be crossed with other investigation methods such as object counting, spectroscopic or proper motion measurements and so on. Indeed the study of the physics of the core of M31 would greatly benefit from these new data, which will help to understand the physics of the stellar environment of supermassive black holes and possiblly shed light on the true origin of these enigmatic objects.  In this paper, we shall exploit the present knowledge of the bulge and the disk populations in M31 to calculate the expected number of simultaneously lensed sources for a given threshold magnitude in either the $K$-band or the $V$-band. We will also consider sources belonging to the central clusters, modelling P1 and P2 as components of a single eccentric system and P3 as a separate inner stellar cluster surrounding the central black hole. For each of the four populations (disk, bulge, P1+P2, P3), we shall present several probability distributions characterizing the properties of the lensing events. The paper is structured as follows: in \\S~2 we indicate our reference models for the source populations and describe their features. In \\S~3 we review some basics on gravitational lensing, with particular reference to the black hole in M31. In \\S~4 we present our estimates for the number of lensed sources at any given time for a given threshold magnitude of the secondary image. In \\S~5 we examine the contribution of sources of different magnitude to the total number of events. In \\S~6 we give the distribution of events as a function of the source distance. In \\S~7 we show the distribution of the angular positions of the images. In \\S~8 we present the distribution of their apparent angular velocities. In \\S~9 we estimate the rate and the average duration of the lensing events. In \\S~10 we discuss some issues concerning the identification of the lensing events. \\S~11 contains the conclusions. ",
        "conclusions": "In this paper we have undertaken a deep investigation of the possible observation of gravitational lensing effects due to the black hole in the center of M31. We have considered stars belonging to four different populations (bulge, disk and the central clusters P1+P2 and P3) as candidate sources. Through detailed modelling of the spatial distributions and the luminosity functions of these four populations we have calculated several quantities that can be used to quantify the power of the central black hole as a gravitational lens using present and future observational facilities. We have carried out our analysis both in the $K$-band and in the $V$-band.  The main outcome is the number of expected lensed sources at a given time whose secondary image is brighter than a specified threshold magnitude. The results are summarized in Figure \\ref{Fig NEvents} and Tables \\ref{Tab NEventsK} and \\ref{Tab NEventsV}, showing that indeed we expect 1.4 lensed sources at a threshold $K_\\mathrm{thr}=24$, 16 events at $K_\\mathrm{thr}=27$ and 180 events at $K_\\mathrm{thr}=30$. In the $V$-band we would have 1.3 events at $V_\\mathrm{thr}=27$, 25 events at $V_\\mathrm{thr}=30$ and 270 events at $V_\\mathrm{thr}=33$.  We have also presented the distribution of the lensed sources as a function of their absolute magnitude, showing the contribution of different stellar branches as sources of gravitational lensing events.  The distribution of the lensed sources as a function of the distance shows that the disk stars are mostly lensed at about 1 kpc distance from the black hole, whereas the bulge stars can be lensed at distances ranging from a few pc to 1 kpc. The inner cluster stars are lensed at fractions of a pc.  This difference is reflected in the angular separation of the images from the black hole, which ranges from 1 mas to $0.1''$ for bulge and disk stars, whereas it stays of the order 1 mas or below for the inner clusters. The resolution needed to get a significant number of events is thus of the order of a few mas, requiring the employment of long baseline interferometers or extremely large telescopes. This justifies our choice to present our analysis in the $K$-band, for which the present infrared interferometers such as Keck, LBT and the future space telescope JWST are optimized. The VLTI, though representing the most advanced interferometer operating in IR bands, cannot observe M31 efficiently. The parallel analysis in the $V$-band provides a deeper comprehension of all the effects coming into play. In this band, the expectations for the number of lensing events are comparable to those in the $K$-band, provided that $V$-band interferometers are constructed in the future with sizes of the same order as in the $K$-band.  The gravitational lensing events discussed in this paper have a very slow evolution, with typical angular velocities of the images between $0.01$ and $0.1$ mas$/$y. We have outlined a possible observational strategy to select gravitational lensing events within the expected background stars, by checking the alignment of pairs of images, the achromaticity of the flux ratio and setting up follow-up observations to detect any secular development.  Our analysis can be repeated with different models of the source populations. Given the present uncertainties on the morphology of the components of M31, we believe that our estimates can be corrected by a factor of a few at most.  The chances for concrete observations of gravitational lensing events by the supermassive black hole in M31 are rather low (though not null) even with the best facilities available today. The perspectives will be definitely increased with the realization of future projects such as the JWST, the extremely large ground telescopes or the realization of new long baseline optical interferometers in the northern hemisphere. With such facilities it should be possible to reach the higher thresholds indicated in this paper while keeping a very high angular resolution. This would open the way to an intensive research of gravitational lensing events and to their use in the investigation of the environment of the supermassive black hole."
    },
    "0711/0711.1340_arXiv.txt": {
        "abstract": "We present new infrared imaging of the NGC\\,2264\\,G protostellar outflow region, obtained with the InfraRed Array Camera (IRAC) on-board the \\emph{Spitzer} Space Telescope. A jet in the red outflow lobe (eastern lobe) is clearly detected in all four IRAC bands and, for the first time, is shown to continuously extend over the entire length of the red outflow lobe traced by CO observations. The redshifted jet also extends to a deeply embedded Class\\,0 source, VLA\\,2, confirming previous suggestions that it is the driving source of the outflow \\citep{gomez94}. The images show that the easternmost part of the redshifted jet exhibits what appear to be multiple changes of direction. To understand the redshifted jet morphology we explore several mechanisms that could generate such apparent changes of direction. From this analysis, we conclude that the redshifted jet structure and morphology visible in the IRAC images can be largely, although not entirely, explained by a slowly precessing jet (period $\\approx$ 8000\\,yr) that lies mostly on the plane of the sky. It appears that the observed changes in the redshifted jet direction may be sufficient to account for a significant fraction of the broadening of the outflow lobe observed in the CO emission. ",
        "introduction": "\\label{sec:introd} It is well established that the early stages of star formation are accompanied by an energetic molecular outflow that is driven by a fast underlying wind. In contrast, the nature of this wind remains disputed and is typically modelled as either a wide-angle wind \\citep[e.g.][]{shu91,shu95,lee01} or a collimated jet \\citep[e.g.][]{raga93,masson93,stahler94,konigl00}. The latter class of model is made particularly appealing by the observation that protostellar jets and molecular outflows are often associated with the same young stellar objects \\citep[e.g.][]{bachiller96}. However, models of jet-driven outflows are unable to produce the wide cavities often displayed in CO maps. It has frequently been proposed that variation in the direction of the jet propagation may act to broaden the outflow \\citep{raga93}. Indeed, there is mounting observational evidence that jets do change direction within their outflows \\citep[e.g.][]{bence96,schwartz99,davis00,arce01,arce02}. Along with the possibility that jets can act to broaden outflows, there is recent evidence of a precessing jet excavating a protostellar envelope \\citep{ybarra06}. Nevertheless, clear evidence that a jet can broaden a molecular outflow to an extent comparable to that seen in the CO observations is still lacking.  \\citet{fich97} found indications of jet induced broadening of an outflow in the case of NGC\\,2264\\,G. This is a powerful, early-stage, bipolar molecular outflow discovered in a survey of high-velocity gas in NGC\\,2264 \\citep{margulis86}, situated at a distance of 800\\,pc \\citep{walker56}. \\citet{gomez94} suggested that the driving engine for this outflow is VLA\\,2, a protostar that is located between the two CO lobes. These researchers found that this radio continuum source, observed at 6\\,cm and 3.6\\,cm, is spatially coincident with an ammonia clump (T $\\approx$ 15\\,K). They additionally found that the source observed at 3.6\\,cm presented an elongated axis that was approximately parallel to that of the NGC\\,2264\\,G outflow and proposed that the continuum emission correspond\\textcolor{blue}{s} to a collimated jet. \\citet{ward-thompson95} observed VLA\\,2 with the JCMT and classified it as a Class\\,0 protostar based on its ratio of submillimeter (L$_\\mathrm{submm} \\approx$0.3\\,L$_\\odot$) and bolometric luminosities (L$_* \\approx$12\\,L$_\\odot$). They estimate a mass for the dust continuum source (VLA\\,2+envelope) of \\hbox{2-4\\msun}.   \\begin{figure*} \\centering \\includegraphics[width=\\textwidth]{f1.eps} \\caption{Colour composite image of the NGC\\,2264\\,G outflow region, where red corresponds to 8\\um, green to 4.5\\um, and blue to 3.6\\um\\ IRAC data. We see the jet associated with the red lobe of the outflow in greenish hues, extending over 1.1\\,pc \\citep[assuming a distance of 800\\,pc;][]{walker56}. The numbered dashed vertical lines mark the locations where the jet appears to change direction in the red CO lobe (coordinates $(\\alpha_1,\\delta_1),\\ (\\alpha_2,\\delta_2)$, and $(\\alpha_3,\\delta_3)$, see text for details), while the bottom horizontal arrows indicate the location of the red and blue CO emission lobes \\citep{margulis86}. The driving source of the jet, VLA\\,2, is identified as well as VLA\\,1 and bow shock structures of the jet.} \\label{fig:IRACcolour} \\end{figure*}   NGC\\,2264\\,G extends $\\approx$ 1.6\\,pc along an east-west direction \\citep{CoG06}, is highly collimated, has a mechanical luminosity of $\\approx$1\\,L$_\\odot$ and a mass of $\\approx$1\\msun\\ \\citep{lada96}, and has been extensively mapped in lines of CO \\citep[e.g.][]{margulis86,fich98}. The outflow displays a ``Hubble law'' velocity distribution, from which \\citet{lada96} were able to estimate an upper limit to the dynamical time of 3-4$\\times10^4$\\,yr.  In addition, there is an `\\emph{S}-shaped' symmetry to the outflow, with low velocity ($\\approx$ 6\\kms) CO emission tracing a cavity near to the source while the more collimated high velocity ($\\approx$ 40\\kms) CO emission lies along a different axis toward the end of the outflow \\citep{lada96}.  Finally, the high degree of collimation and highly bipolar nature of the CO outflow lead \\citet{fich97} to conclude that the outflow must lie within 20 degrees of the sky.  NGC\\,2264\\,G has also been imaged in \\h2\\ v=1-0 S(1) by \\citet{davis95}, who interpreted bright knots at the east edge of the outflow as a signature of a turbulent mixing layer at the boundary between outflow and ambient medium, and at higher resolution by \\citet{CoG06}, who found chains of knots they attribute to a jet along the outflow. \\citet{fich97} compared the CO maps of \\citet{lada96} with the \\h2\\ images of \\citet{davis95} and noted that the bright \\h2\\ knots coincided with points at which the axis of CO changed direction. On the basis of this, they suggested that the high velocity CO emission in both lobes of the outflow traced the molecular component of a deflected jet. Furthermore, they speculated that such deflections could serve to broaden the CO outflow lobes to an extent comparable to that observed. The aim of this paper is to investigate their hypothesis.  The \\emph{Spitzer} Space Telescope has provided the astronomical community with striking images of protostellar jets \\citep[e.g][]{noriega04}. The plethora of molecular lines that fall within the InfraRed Array Camera (IRAC) broad band filters, combined with the high sensitivity of the detector, show detailed features that are observationally unaccessible to near-infrared or optical imaging, or even to ground-based mid-infrared imaging. Here, we present IRAC images of the NGC\\,2264\\,G outflow that clearly show a jet, which appears to exhibit three changes of direction in its easternmost part.  In conjunction with CO maps from \\citet{fich98}, we reconsider the evidence for deflection of the outflow presented by \\citet{fich97} and explore other jet-bending mechanisms, in particular, precession.  The structure of the paper is the following: in \\S\\,\\ref{sec:obs_red} we give a brief description of the observations and data reduction; in \\S\\,\\ref{sec:res} we present the results obtained from our analysis; and, in \\S\\,\\ref{sec:dis} we discuss several possible jet bending mechanisms with a view to explaining the observed jet morphology. Finally, in \\S\\,\\ref{sec:concl} we summarize our main conclusions. ",
        "conclusions": "\\label{sec:concl}   We analysed \\emph{Spitzer} observations of the NGC\\,2264 region with the aim of investigating the suggestion of \\citet{fich97} that the NGC\\,2264\\,G outflow is shaped by a deflected redshifted jet. Here, we summarize our findings. \\begin{enumerate} \\item  A jet is clearly seen in all four \\emph{Spitzer} IRAC images, in the red lobe of the NGC\\,2264\\,G outflow, and is brightest in the 4.5\\,$\\mu$m band. We find the redshifted jet to extend continuously across the entire length of the red outflow lobe. It appears to emanate directly from VLA\\,2, confirming this Class\\,0 source as the driving source of the outflow. The most striking features of the redshifted jet image are three distinct points where the redshifted jet appears to change direction. \\item We discuss various mechanisms that can result in changes of jet direction and explore further the two most likely, deflection and precession. We are able to explain the majority of the redshifted jet structure, as seen by IRAC, with a simple precessing jet model. This model predicts a precession time of $\\approx$8000\\,yrs for a jet velocity of 200\\kms. The precession of this jet may contribute to the broadening of the red lobe and account for its overall morphology, however, this model cannot satisfactorily explain the entire morphology of the blue CO lobe without invoking an additional change in direction or deflection of the underlying jet similar to that proposed by \\citet{fich97}. \\end{enumerate}"
    },
    "0711/0711.4565_arXiv.txt": {
        "abstract": "% We present numerical simulations of the very late thermal pulse (VLTP) scenario for a wide range of remnant masses. We show that, by taking into account the different possible remnant masses, the fast outburst evolution of V4334 Sgr (a.k.a. Sakurai's Object) and V605 Aql can be reproduced within standard 1D stellar evolution models. A dichotomy in the born again timescales is found, with lower mass remnants evolving in a few years and higher mass remnants (M $\\hbox{\\rlap{\\hbox{\\lower4pt\\hbox{$\\sim$}}}\\hbox{$>$}}$ 0.6 M$_{\\odot}$) failing to expand due to the H-flash and, as a consequence, evolving in timescales typical of He-shell flash driven born agains ($\\sim 100$ yr). ",
        "introduction": "Since the discovery of post-AGB thermal pulses in the early simulations of Paczynski (1970) these events have been object of several studies. Fujimoto (1977) showed that during thermal pulses penetration of the pulse driven convective zone (PDCZ) into the H-rich envelope is possible if the mass of this envelope was small enough. These works were confirmed by Sch\\\"onberner (1979) who performed calculations of low mass stars with steady mass loss. Sch\\\"onberner (1979) found that contact between the PDCZ and the H-rich envelope was possible if the thermal pulse happens when the star enters the white dwarf stage (event now termed VLTP).  Since the work of Renzini (1979) post-AGB thermal pulses are associated with the formation of H-deficient stars, an idea pushed forward by the simulations of Iben et al. (1983) (who introduced the term ``born again AGB stars'' for this objects). Due to numerical difficulties in the treatment of the simultaneous mixing and burning of H at high temperatures during a VLTP, early simulations did not follow the evolution during this event. This was first attempted by Iben \\& MacDonald (1995) who presented the first VLTP simulations that included a time dependent treatment of the convective mixing and burning of H. After the works of Herwig et al. (1997, 1999) it became clear that the abundances of some post-AGB H-deficent stars could be well understood within the born again scenario (see, however, De Marco this proceedings). \\begin{table*}[ht!] {\\footnotesize \\begin{center} \\begin{tabular}{c|c|c|c|c|c} & Iben \\&   &  Herwig   &  Herwig  & Lawlor \\&  & Miller Bertolami \\\\ & MacDonald &  et al.   &          & MacDonald  &  et al.    \\\\ &  (1995)   &  (1999)   &  (2001)  &  (2002)    & (2006)     \\\\ \\hline Mass  [M$_{\\odot}$] & 0.60  & 0.604  & 0.535 & 0.56 to 0.61 &  0.589   \\\\ Born Again Time [yr] & 17 & 350 & 21 & 4.5 to 8.5 & 5 to 10 \\\\ \\end{tabular} \\label{tabla} \\caption{Born Again timescales of previous simulations of the VLTP with standard MLT approach that include the violent H-burning event.} \\end{center} } \\end{table*}  The identification of V4334 Sgr as a star undergoing a VLTP event (Duerbeck \\& Benetti 1996) has renewed the interest in this particular kind of late helium shell flash. V4344 Sgr has showed a very fast evolution of a few years (Duerbeck et al. 1997, Asplund et al. 1999 and van Hoof et al. 2007), not far from the only theoretical value available at the time of its discovery ($\\sim 17$ yr, Iben \\& MacDonald 1995). Since then, the fast evolution of V4334 Sgr has been the object of several works with results far from showing a consistent picture (see Table 1). The works of Herwig (2001) and Hajduk et al. (2005) introduced a new free parameter in the treatment of convection (the mixing efficiency) and showed that, by fine tuning the velocities of convective mixing during the VLTP (to values different to those predicted by the mixing length theory, MLT), the fast outburst evolution and reheating of V4334 Sgr could be reproduced by a model of 0.604 M$_{\\odot}$. In this context, VLTP evolution was proposed as a tool to test convection theory of reactive convective fluids (Herwig 2001). However, born again timescales of previous simulations of the VLTP episode with standard MLT approach (i.e. no reduction in convective velocities) show a wide range of born again timescales, with differences rising up to a factor 70 (see Table 1). If V4334 Sgr observed evolution and VLTP simulations are to be used to test convection theory, these differences have to be understood first.  As can be seen from third and fourth columns of Table 1 the remnant's mass seems to play an important role in determining the born again timescale. However, most of the existing simulations of the VLTP have been performed in a narrow range in mass. The aim of the present article is the exploration of the (neglected) importance of the remnant mass for the born again timescale.    For the present work we have performed numerical simulations of the VLTP for 10 different remnant masses. All of the remnant models are the result of full evolutionary calculations from the ZAMS, through the AGB phase and to the post-AGB phase for a metallicity of Z=0.02.  Overshooting at every convective boundary was considered as in Herwig et al. (1997). Convective mixing was considered as a difussive process and solved simultaneously with nuclear burning and convective velocities were adopted as given by standard MLT approach. For a more extensive discussion about numerical and physical aspects of the present simulations we refer the reader to Miller Bertolami \\& Althaus (2007). ",
        "conclusions": "We have shown that post-VLTP born again times can be divided into two groups: low mass remnants that evolve back to the AGB in a few years and high mass remnants in which the energy released by H-burning fails to drive the expansion back to the AGB and, consequently, show born again timescales of the order of a century. We also showed that V4334 Sgr evolution can be understood within the standard MLT as the VLTP evolution of a $\\sim 0.56$ M$_{\\odot}$ remnant. However our models fail to reproduce the rapid reheating reported by van Hoof et al. (2007) something that may be pointing to the need of a better treatment of the outer layers of the envelope in our models.       \\begin{figure}[ht!] \\begin{center} \\includegraphics[ height=120 pt] {M3B_talk_01.eps}  \\caption{\\footnotesize Left: Comparison between the energy available from the burning of the H-content of the star (shaded zone) with the energy needed to expand the envelope (solid line). Right: Born again timescales of the simulations.} \\end{center} \\end{figure}   \\begin{figure}[h!] \\begin{center} \\includegraphics[ height=152 pt] {M3B_talk02.eps} \\caption{\\footnotesize \\emph{Left:} Outburst evolution (top: luminosity, bottom: effective temperature) of our lower mass sequences as compared with observations of V4334 Sgr. Observations are taken from Duerbeck et al. (1997; $\\diamondsuit$ and +), Asplund et al. (1999; $\\times$), Pavlenko \\& Geballe (2002; $\\circ$) and van Hoof et al. (2007; $\\bullet$). For the 0.561M$_{\\odot}$ model two different mass loss rates were considered at log($T_{\\rm eff}$)$<3.8$ leading to different reheating timescales. Inset shows a zoom of the evolution of the effective temperature during the discovery of V4334 Sgr. \\emph{Right:} Preoutburst location of our sequences ($\\bullet$). Grey dashed line correspond to possible detection of V4334 Sgr in a ESO/SERC survey (Herwig 2001). V4334 Sgr preoutburst locations are from Kerber et al. (1999; black square) and Hajduk et al. (this proceedings; grey diamond). V605 Aql value is taken from Lechner \\& Kimeswenger (2004).} \\end{center} \\end{figure}  {\\footnotesize"
    },
    "0711/0711.4423_arXiv.txt": {
        "abstract": "The quantum behavior of noncommutative eternal inflation is quite different from the usual knowledge. Unlike the usual eternal inflation, the quantum fluctuation of noncommutative eternal inflation is suppressed by the Hubble parameter. Due to this, we need to reconsider many conceptions of eternal inflation. In this paper we study the Hawking-Moss tunneling in noncommutative eternal inflation using the stochastic approach. We obtain a brand-new form of the tunneling probability for this process and find that the Hawking-Moss tunneling is more unlikely to take place in the noncommutative case than in the usual one. We also conclude that the lifetime of a metastable de-Sitter (dS) vacuum in the noncommutative spacetime is longer than that in the commutative case. ",
        "introduction": "Inflation has been widely considered as a remarkably successful theory in explaining many problems in the very early universe, such as the flatness, horizon and monopole problems \\cite{Guth81,Sato:1980yn,Linde82,Steinhardt82,Starobinsky-inf}. During inflation, quantum effects play a crucial role and may bring the universe into a self-reproducing process which is dubbed ``eternal inflation\" \\cite{steinhardt-nuffield,vilenkin-eternal,linde-eternal}. The string theory landscape indicates that there are a huge number of meta-stable vacua surrounded by various kinds of effective potentials \\cite{bousso-landscape,kachru-landscape,susskind-landscape,douglas-landscape}. The realization of string landscape provides an important arena for eternal inflation.   In eternal inflation driven by the false vacuum, the false vacuum is a meta-stable state and would decay through a mix of semiclassical tunneling and stochastic evolution. The probability of finding the inflaton at the top of the plateau in its potential decreases exponentially with time \\cite{Guth:1985ya}. However, the false vacuum is also expanding exponentially while decaying. When the rate of exponential expansion is larger than the decay rate during this process, the total volume of the false vacuum will grow eternally although the false vacuum is decaying. In this case, bubbles form by random nucleation and then start to expand. Every growing bubble can be viewed as an open FRW universe \\cite{Coleman:1980aw}, and we are living in one of such ``pocket universes\" \\cite{Guth:2000ka}. Another approach to eternal inflation is achieved in chaotic inflation when the quantum fluctuation of the inflaton dominates over its classical motion. As the inflaton is rolling down the potential classically, its change during one Hubble time ($\\delta t=\\frac{1}{H}$) can be divided into $\\delta\\varphi=\\Delta\\varphi+\\delta_q\\varphi$, where $\\Delta\\varphi$ denotes the classical value and $\\delta_q\\varphi$ represents the quantum one. For a Gaussian probability distribution, when $\\delta_q\\varphi>0.61 \\Delta\\varphi$, the quantum behavior overwhelms the classical evolution and inflation becomes eternal.  When dealing with the decaying process of false vacua, one have to make the inflaton tunnel from one false vacuum to another. One method was provided by Coleman and De Luccia (CDL) \\cite{Coleman:1980aw}; another method was investigated by Hawking and Moss \\cite{Hawking:1981fz}. During the Hawking-Moss tunneling, the potential between two vacua is so flat that CDL instanton can not exist. It was shown that the probability of tunneling from one false vacuum $\\varphi_0$ to another is given by \\begin{eqnarray}\\label{Pclass} P_C &\\sim& \\exp\\left(-\\frac{24\\pi^2}{V(\\varphi_0)}+\\frac{24\\pi^2}{V(\\varphi_{\\rm top})}\\right)\\nonumber\\\\ &\\sim& \\exp\\left(-8\\pi^2\\cdot\\frac{H(\\varphi_{\\rm top})^2-H(\\varphi_0)^2}{H(\\varphi_0)^2H(\\varphi_{\\rm top})^2}\\right)~, \\end{eqnarray} which is related to the value of the top barrier. A proper scenario of this tunneling can be realized in the stochastic approach to inflation \\cite{Starobinsky:1986fx,Rey:1986zk,Nakao:1988yi,Linde:1991sk,Gratton:2005bi,Linde:2006nw}. Here the potential of the scalar field is flat enough for slow rolling which requires $V''<V$. In the stochastic description, the quantum fluctuations can be simulated by the stochastic noise and the scalar field walks in random.  As is known, eternal inflation happens when the energy scale of the universe is extremely high. Thus we would like to take into consideration more fundamental theories in logic, namely, the string theory. Recently, we have considered the effects of spacetime noncommutativity on slow-roll eternal inflation in Ref. \\cite{Cai:2007et}. In noncommutative inflation, it is generally assumed that the background evolution of inflaton is not modified but the fluctuations are affected by noncommutativity (see Ref. \\cite{Bran2002,Koh:2007rx,qghuang-nc,cai2007}, for a review in \\cite{Brandenberger:2007rg} and references therein). In order to introduce the noncommutativity \\cite{Li:1996rp,Seiberg99} into the 4-dimensional flat Friedmann-Robertson-Walker universe, we would like to define another time coordinate $\\tau$, \\begin{equation} ds^2=dt^2-a^2(t)d{\\vec x}^2=a^{-2}(\\tau)d\\tau^2-a^2(\\tau)d{\\vec x}^2~, \\end{equation} where $a$ is the scale factor. Then the spacetime uncertainty relation can be realized by the commutation relation: \\begin{equation} [\\tau,x]_*=i M_{\\rm N}^{-2}~, \\end{equation} where $M_{\\rm N}$ is the energy scale of noncommutativity and the $*$-product is defined as \\begin{eqnarray} (f*g)(x,\\tau) &=& \\exp \\left(-\\frac{i}{2}M_{\\rm N}^{-2}\\left(\\partial_x\\partial_{\\tau'} -\\partial_\\tau\\partial_y\\right)\\right) \\nonumber\\\\ &\\times& f(x,\\tau)g(y,\\tau') \\left. \\right|_{y=x,\\tau'=\\tau}~. \\end{eqnarray} From the result of the paper \\cite{Cai:2007et}, we can see that the quantum fluctuation still satisfies the Gaussian distribution, but the form of its amplitude in IR region is changed to $\\delta_q\\varphi\\simeq\\frac{1}{2\\pi}\\frac{M_{\\rm N}^2}{H}$. Therefore, when the Hubble parameter is lifted highly enough, eternal inflation would cease. This is strongly different from the normal scenario of eternal inflation. In this paper, we shall use the analysis of noncommutativity mentioned above (especially the IR region) and the stochastic approach to study Hawking-Moss tunneling of eternal inflation.  This paper is organized as follows. In Section 2, we study the Hawking-Moss tunneling in noncommutative eternal inflation using the stochastic approach. In Section 3, we make conclusions that Hawking-Moss tunneling is more unlikely to happen in the noncommutative case than in the usual one and the lifetime of a metastable de-Sitter vacuum in the noncommutative spacetime is longer than that in the commutative case. ",
        "conclusions": "To take a further discussion, we would like to compare the difference of dS decaying processes whether or not the spacetime noncommutativity is present. According to the work of Coleman and De Luccia\\cite{Coleman:1980aw}, the decay time of a metastable dS vacuum has an approximate expression $T\\sim P^{-1}$. By neglecting all the sub-exponential factors, we have \\begin{eqnarray} T_C&=&\\exp\\left\\{24\\pi^2M_p^4(\\frac{1}{V_0}-\\frac{1}{V_{top}})\\right\\};\\\\ \\label{timenc}T_{\\rm N}&=&\\exp\\left\\{\\frac{8\\pi^2}{3M_{\\rm N}^4} (V_{top}-V_0)\\right\\}~, \\end{eqnarray} which represent the dS decay times without the spacetime noncommutativity, and the IR region with noncommutativity respectively. To be clear, we have written $M_p$ explicitly here. It is clear that $T_N>T_C$, which is a very general result indicating that the lifetime of a metastable dS vacuum with noncommutativity is longer than that without noncommutativity.  To summarize, from the results obtained in this note we learn that the Hawking-Moss tunneling effect of noncommutative eternal inflation in the IR region is greatly different from the usual one. Its probability distribution is exponentially suppressed by the top barrier value of the potential and make Hawking-Moss tunneling more difficult to happen than in the usual case. This is because the quantum fluctuation is suppressed by spacetime noncommutativity. Consequently, we may expect the application of noncommutativity would bring a closer sight into high energy physics of eternal inflation. Based on the new form of the probability distribution, we find that the lifetime of a metastable dS vacuum in the noncommutative case is longer than in the usual one. This may leave more clues for investigating the new physics of noncommutativity which is worthy for further studies."
    },
    "0711/0711.1801_arXiv.txt": {
        "abstract": "{With the discovery during the past decade of a large number of extrasolar planets orbiting their parent stars at distances lower than 0.1 astronomical unit (and the launch and the preparation of dedicated space missions such as CoRoT -Bord\\'e {\\it et al.}, 2003- and KEPLER -Borucki {\\it et al.}, 2007-), with the position of inner natural satellites around giant planets in our Solar System and with the existence of very close but separated binary stars, tidal interaction has to be carefully studied.} {This interaction is usually studied with a ponctual approximation for the tidal perturber. The purpose of this paper is to examine the step beyond this traditional approach by considering the tidal perturber as an extended body. To achieve this, the gravitational interaction between two extended bodies and more precisely the interaction between mass multipole moments of their gravitational fields and the associated tidal phenomena are studied.} {Use of Cartesian Symmetric Trace Free tensors, of their relation with spherical harmonics and of the Kaula's transform enables us to analytically derive the tidal and mutual interaction potentials as well as the associated disturbing functions in extended bodies systems.} {The tidal and mutual interaction potentials of two extended bodies are derived. In addition, the external gravitational potential of such tidally disturbed extended body is obtained, using the Love number theory, as well as the associated disturbing function. Finally, the dynamical evolution equations for such systems are given in their more general form without any linearization. We also give, under simplified assumption, a comparison of this formalism to the ponctual case. {\\bf The non-ponctual terms have to be taken into account for strongly deformed perturbers ($J_{2}\\ge10^{-2}$) in very close systems ($a_{\\rm B}/R_{\\rm B}\\le5$).}} {We show how to derive explicitely the dynamical equations for the gravitational and tidal interactions between extended bodies and associated dynamics. Conditions for application of such formalism are given.} ",
        "introduction": "\\par In celestial mechanics,  one of the main approximation done in the modelling of tidal effects (star-star, star-innermost planet or planet-natural satellites interactions) is to consider the tidal perturber as a point mass body. However a large number of extrasolar Jupiter-like planets orbiting their parent stars at a distance lower than 0.1 AU have been discovered during the past decade (Mayor {\\it et al.}, 2005). Moreover, in Solar System, Phobos around Mars and the inner natural satellites of Jupiter, Saturn, Uranus and Neptune are very close to their parent planets. In such cases, the ratio of the perturber mean radius to the distance between the center of mass of the bodies can be not any more negligeable compared to 1 (cf. Fig.~\\ref{Perturber}). Furthermore, it can be also the case for very close but separated binary stars. In that situation, neglecting the extended character of the perturber have to be relaxed, so the tidal interaction between two extended bodies must be solved in a self-consistent way with taking into account the full gravitational potential of the extended perturber, generally expressed with some mass multipole moments, and then to consider their interaction with the tidally perturbed body. In the litterature, not so many studies have been done (Borderies, 1978-1980; Ilk, 1983; Borderies \\& Yoder, 1990; Hartmann, Soffel \\& Kioustelidis, 1994; Maciejewski, 1995). The purpose of this work is then to provide a theoretical procedure to obtain this tidal gravitational interaction as well as its associated tidal dynamical evolution.\\\\  \\par Several years ago, Hartmann, Soffel \\& Kioustelidis (1994) introduced in Celestial Mechanics an interesting tool, based on Cartesian Symmetric Trace Free (STF) tensors, to treat straighforwardly the couplings between the gravitational fields of extended bodies. These tensors are fully equivalent to usual spherical harmonics but in addition a set of STF tensors represents an irreductible basis of the rotation group SO3 (Courant \\& Hilbert 1953, Gelfand {\\it et al.} 1963). It means that using algebraic properties of STF tensors with the index notation of Blanchet \\& Damour (1986), these objects become a powerful tool to determine the coupling between spherical harmonics in an elegant and compact way. However, as these tensors are not widely used in Celestial mechanics, we first recall their definition and fundamental properties and we stress their relation with usual spherical harmonics. Then, we treat the multipole expansion of gravitational-type fields for which each type is related to a given extended body. First, the well-known external field of such body is derived using STF tensors; classical identities are provided. Next, the mutual gravitational interaction between two extended bodies and the associated tidal interaction are derived. We show how the use of STF tensors leads to an analytical and compact treatment of the coupling of their gravitational fields. We deduce the general expressions of tidal and mutual interaction potentials expanded in spherical harmonics. Using the classical Kaula's transform (Kaula, 1962), we express them as a function of the Keplerian orbital elements of the body considered as the tidal perturber. These results are used to derive the external gravitational potential of such tidally perturbed extended body. Introducing a third body, its mutual interaction potential with the previous tidally perturbed extended body is defined that allows us to derive the disturbing function using the results obtained with STF tensors and the Kaula's transform. At this stage, the different type of mutual gravitational interaction are defined. The dynamical equations ruling the evolution of this system are obtained. Finally, we use a reduced form of that equations to qualitatively quantify the influence of non-ponctual terms of the disturbing function in comparison with the ponctual case.  \\begin{figure}[h!] \\centering \\resizebox{12cm}{!}{\\includegraphics{Perturber.pdf}} \\caption{System of two extended bodies. ${\\overline R}_{\\rm A}$ and ${\\overline R}_{\\rm B}$ are the respective mean radius of A and B while $D_{\\rm AB}$ is the distance between their respective center of mass. $\\varepsilon_{\\rm A}$ and $\\varepsilon_{\\rm B}$ are the conical angles with which each perturber is seen from the center of mass of the perturbed body. When $\\frac{{\\overline R}_{i}}{D_{\\rm AB}}\\!<\\!\\!<\\!1$ with $i=$A or B the i$^{\\rm th}$ body could be considered as a ponctual mass perturber.} \\label{Perturber} \\end{figure} ",
        "conclusions": "\\par This work represents another step towards the modelling of the tidal dynamical evolution of planetary systems and of multiple stars by taking into account the extended character of the bodies. We have used the STF tensors which allow us to treat in an analytical and compact way the complex couplings between the multipole behaviour of gravitational fields of two extended bodies. First, the main properties of STF tensors have been defined and derived. Next, after deriving the classical multipole expansion of the self gravitational field of an extended body A, we have derived  the tidal potential related to the interaction between two extended bodies A and B, which is a generalization of the previous results where the perturber is treated through the ponctual mass approximation. Using those results, the external potential of a tidally-perturbed body has been provided and used to obtain the associated mutual interaction potential with a body C for which the dynamics is studied and which could be different from the previous tidal perturber B, as well as the corresponding disturbing function. The disturbing function being expanded into Fourier's series, the dynamical equations are derived. Those allows us to study the dynamical evolution of the angular momentum and the obliquity of the body A simultaneously with the Keplerian oribtal elements of C relative to the center of mass of A, namely the semi major axis, $a_{\\rm C}$, the eccentricity, $e_{\\rm C}$, and the orbital inclination, $I_{\\rm C}$. {The equations have been derived in a general way without any linearization of function of any orbital elements as it has been done in Hut (1981-1982) and in Correia \\& Laskar (2003c).} This allows us to study the strongly non-linear problem of inclined, eccentric orbits. Therefore, this formalism can be useful for elliptic systems.\\\\  \\par The major interest of this modelling concerns the dissipation of the tidal mechanical energy. Here, the response of each extended body is parametrized through the Love number $k_{l}$ for the adiabatic one, and an impedance $Z_{\\rm T}$ and its associated delay $\\delta_{\\rm T}$ which describes the damping and the phase lag due to the (viscous and thermal) dissipation acting on the bulk induced by the perturber. Self-consistent modellings have to be developed for elasto-viscous bodies (Rogister and Rochester, 2004) as well as for fluid bodies (cf. Zahn, 1977) as well as equilibrium and stability conditions. This will be done for extended fluid bodies in a forthcoming sery of papers.\\\\  \\par Moreover, due to the general character of this work, various applications are possible. The first one deals with the dynamics of planets which are very close to their star as observed in several extrasolar planetary systems (cf. Laskar \\& Correia, 2004; Levrard et al., 2007; Fabrycky et al., 2007; {\\bf Correia et al. 2008}). Next, in our own Solar System, dynamical systems, such as giant planets and their internal satellites, can be studied by taking into account their spatial extension and their multipolar behaviour due for example to their rapid rotation. Finally, this work can be applied to the dynamical evolution of close binary stars where the tidal interaction and its dissipation dominate the behaviour of the system until it has reached its lower energy state, where the spins are all aligned, the orbits are circular, and the components are synchronized with the orbital motion. For these coming applications, the potential impact of the spatial extension of bodies on dynamics has to be carefully evaluated and understood.  \\begin{acknowledgement} This is a real pleasure to acknowledge Pr. J.-P. Zahn, Pr. T. Damour and Dr. J. Vaubaillon for helpful comments. We are also grateful to the anonymous referee for his suggestions that allow us to improve the original manuscript. \\end{acknowledgement}"
    },
    "0711/0711.4109_arXiv.txt": {
        "abstract": "Recent measurements of high-redshift QSO clustering from the Sloan Digital Sky Survey indicate that QSOs at $z\\sim 4$ have a bias $\\langle b\\rangle\\simeq 14$. We find that this extremely high clustering amplitude, combined with the corresponding space density, constrains the dispersion in the $L-\\mh$ relation to be less than 50\\% at 99\\% confidence for the most conservative case of a 100\\% duty cycle. This upper limit to the intrinsic dispersion provides as strong a constraint as current upper limits to the intrinsic dispersion in the local $\\mbh-\\sigma$ relation and the ratio of bolometric to Eddington luminosity of luminous QSOs. ",
        "introduction": "\\label{sec:introduction}  It has recently become accepted that quasar activity and black hole growth are an integral part of galaxy evolution, however a detailed understanding of what triggers quasar activity and how they are fueled still eludes us. The leading contender for the identity of luminous, high redshift QSOs is that they are black holes fed by by major mergers of gas-rich galaxies \\citep{Car90}.  Recent incarnations of such models \\citep{HaiLoe98,CavVit00,KauHae00,WyiLoe02,hopkins06} provide a good description of many observed properties of the QSO population.  The situation is particularly interesting at high redshift, where the population of supermassive black holes that powers the QSOs is growing rapidly \\citep[see e.g.][for a recent review]{Fan06}.  To further understand this important phase of black hole and galaxy evolution we would like to build a model in which QSO activity is tied to the evolving cosmic web of dark matter halos.  The relationship between QSOs and dark matter halos, their environments and duty cycles, can be constrained via observations of their space density and large-scale clustering \\citep{ColKai89,HaiHui01,MarWei01}.  These constraints become particularly sensitive if the QSOs inhabit the rarest, most massive halos for which the spatial clustering depends strongly on halo mass \\citep{Kai84,BBKS86,EFWD88,ColKai89}.  At redshifts $z<3$ the advent of large optical surveys for QSOs has led to firm constraints on the clustering as a function of luminosity and redshift \\citep{Cro05,Hen06,PorMagNor04,PorNor06,Mye07a,Mye07b,Mye07c,Ang08,Pad08}. With the Sloan Digital Sky Survey (SDSS) we are now able to measure the clustering of QSOs well even at $z>3$ \\citep[][see also earlier work by \\citealt{Kun97,Ste97}]{She07}. Interestingly, the correlation length of the QSO population increases rapidly with redshift, from $r_0=16.90\\pm 1.73$ at $z\\simeq 3$ to $r_0=24.30\\pm 2.36$ at $z\\simeq 4$. \\citet{She07} demonstrate that they can fit the observed clustering and space density of $z\\sim 4$ QSOs with the model of \\citet{MarWei01} provided the $z\\sim 4$ QSOs are relatively long lived ($\\tq\\sim 160\\,$Myr) and inhabit halos more massive than about $5\\times 10^{12}\\,h^{-1}M_\\odot$. For this calculation \\citet{She07} assume that there is a monotonic relationship between instantaneous QSO luminosity and halo mass, with no scatter. The actual relation between instantaneous QSO luminosity and host halo mass is expected to include some scatter.  Scatter is expected in several of the relationships linking the QSO luminosity and the host halo mass: the relationship between the host halo mass and galaxy bulge, in the relationship between galaxy bulge and black hole mass, in the relationship between black hole mass and peak luminosity and in the relationship between peak and instantaneous QSO luminosity.  Here we demonstrate that the very high correlation length measured by \\citet{She07}, when combined with the rapid increase in bias for the most massive halos, strongly constrains the scatter between instantaneous QSO luminosity and halo mass. The essential idea is that any scatter increases the contribution from lower mass (and less highly biased) halos, so that a measurement of large clustering amplitude limits the contribution from lower mass objects. This constraint is as strong as direct, observational constraints at lower redshift on the amount of scatter in the various relationships mentioned above.  We shall work throughout in the $\\Lambda$CDM framework and adopt the following cosmological parameters: $\\Omega_{\\rm mat}=0.25$, $\\Omega_\\Lambda=0.75$, $h=0.72$ and $\\sigma_8=0.8$. Where appropriate we shall comment on how our results depend on these particular choices.  The next section outlines our formalism and applies it to the data of \\citet{She07}, while \\S\\ref{sec:discussion} interprets our results and discusses future observations. ",
        "conclusions": "\\label{sec:discussion}  It appears that the most luminous, highest redshift QSOs have instantaneous luminosities which are well correlated with their host halo masses. The already strong constraint shown here should be improved by measurements of high-redshift QSO clustering with yet larger samples. New observations will broadly improve the constraint in two ways: through wider area observations to identify a larger number of QSOs similar in redshift to the SDSS sample, and through deeper observations over the same area. Deeper observations will identify slightly fainter QSOs at all redshifts and may provide a sufficient sample to measure the clustering amplitude at even higher redshift than $z \\sim 4$, where we expect the bias to be even larger. Going wider in area will decrease the fractional error on $\\beff$ for high-redshift QSOs. As these measurements of high-redshift QSO bias will remain Poisson limited, the fractional error on $\\beff$ will scale as $1/N_{\\rm qso}$ for a survey of $N_{\\rm qso}$ QSOs. Going deeper to measure luminous QSOs at yet higher redshifts is also expected to be useful, as the data indicate strong evolution in the clustering amplitude with redshift.  Figure~\\ref{fig:b_vs_z} illustrates how $\\beff$ increases with redshift for three values of the scatter between QSO luminosity and host halo mass assuming $\\fon = 1$. At each redshift $\\beff$ was calculated to match the QSO space density evolution parametrized by \\citet{Ric06}, although converted to our cosmological parameters\\footnote{For simplicity we did not correct the \\citet{Ric06} LF parameterization to match their binned LF here.  The space density has therefore been slightly underestimated from $z=3\\rightarrow 4$ and consequently the predicted bias is slightly higher.}. The steepest relation between $\\beff$ and redshift is for the case of the smallest scatter. If $\\fon = 1$ and the scatter is minimal, the relative value of $\\beff$ for different values of the dispersion is a good measure of the fractional uncertainty in $\\beff$ required to improve on our constraint. These ratios demonstrate that the most effective constraint from clustering arises from the lowest redshift at which $\\fon = 1$ and minimal scatter is a reasonable approximation. For example, the ratio of $\\beff(\\sigma=0.1)$ to $\\beff(\\sigma=1.0)$ increases by less than 10\\% from $z=4$ to $z=5$, while the $N_{\\rm qso}$ above the SDSS flux limit drops by approximately an order of magnitude.  Therefore while the clustering amplitude increases, the decline in $N_{\\rm qso}$ for a flux-limited, high-redshift survey more than offsets this gain.  By contrast, color selection of just very high-redshift (e.g.~$z>5$) QSOs could produce a substantial improvement for a given total number of spectra.  The immediate prospect for improvement in this constraint is completion of the SDSS observations beyond the 4000 square degrees employed by \\citet{She07}. These observations will simply increase $N_{\\rm qso}$ and the error bars should scale as $N_{\\rm qso}^{-1}$. On a somewhat longer timescale, the proposed Baryon Oscillation Spectroscopic Survey (BOSS\\footnote{http://cosmology.lbl.gov/BOSS}) should approximately double the number of $z>3.5$ QSOs, although it will mostly achieve this by going deeper. Provided $\\fon = 1$ and small scatter are still reasonable approximations to this slightly fainter population (that will still be well above the break in the QSO LF), BOSS' factor of two improvement in sample size should decrease the error bars by approximately a factor of two and produce a powerful improvement in the constraint.  For example, if the measured bias stayed fixed and the errorbars decreased by a factor of two, the upper limit on the observed scatter would improve to less than 0.5 (i.e.~0.2 dex) for $\\fon=1$. There are also upcoming photometric surveys that go even wider in area and deeper: Pan-STARRS\\footnote{http://pan-starrs.ifa.hawaii.edu}, the Dark Energy Survey (DES\\footnote{http://www.darkenergysurvey.org}), and the Large Synoptic Survey Telescope (LSST\\footnote{http://www.lsst.org}). While these surveys do not include a dedicated plan for spectroscopic observations of QSOs, they will provide the necessary candidate database of these extremely rare objects for spectroscopy.  \\begin{figure} \\begin{center} \\resizebox{3.5in}{!}{\\includegraphics{b_vs_z_phi.ps}} \\end{center} \\vspace{-0.1in} \\caption{({\\it Top}) The large-scale bias, $\\langle b\\rangle$, vs.~redshift of halos that match the integrated QSO space density from \\protect\\citet{Ric06} for $\\fon = 1$ and dispersion $\\sigma=0.1$ (long dashed), 0.5 (dashed), 1 (solid) as in Figure~\\ref{fig:b_vs_n}. Our bias estimates derived from the \\protect\\citep{She07} estimates are also shown. ({\\it Bottom}) Normalized number of QSOs per redshift bin from \\protect\\citet{She07}. The rapid decrease in SDSS QSOs with redshift indicates that purely flux-limited surveys at $z > 3$ can not place as strong a constraint on $\\sigma$ at yet higher redshifts, even though those QSOs are more strongly biased.} \\label{fig:b_vs_z} \\end{figure}  On the theoretical side, the calculations presented here could also be improved.  Better modeling could allow us to use a wider range of scales. As large-scale, high-resolution simulations become increasingly feasible, it will become possible to precisely measure both the bias of all halos and the bias as a function of halo history. The scatter between these various relationships can also be measured in self-consistent numerical simulations of merger-driven black hole growth. \\citet{Hop07} presented a recent analysis of the observed and simulated scatter between black hole mass and various host (bulge) galaxy properties ($\\sigma$, host galaxy mass, effective radius) at lower redshifts than we are considering here.  They find that the intrinsic scatter is generally lower than the observed (upper limit) and as small as $\\sim 0.2\\,$dex for combinations that produce a fundamental plane for black holes. However, they also note that the strong correlation of black hole mass with bulge properties implies that the connection to halo mass is only indirect and dominated by the evolution in the typical gas fractions. Measurement of the evolution of these quantities in full cosmological simulations \\citep[e.g.][]{DiM08} could provide measurements of the expected scatter at the very high redshifts where QSOs are observed to be so strongly biased.  Regardless of these future prospects, the strength of the present constraint on scatter in a monotonic relation between QSO luminosity and halo mass is already surprising. As noted in the introduction, we expect scatter to be present because we expect scatter in at least four relationships that combine to determine the instantaneous luminosity of a QSO in a halo of a given mass: the relationship between instantaneous luminosity and peak luminosity; the relationship between peak luminosity and black hole mass; the relationship between black hole mass and bulge properties; and the relationship between bulge and halo mass. These expectations arise from the observed scatter in these individual relations at lower redshifts--intriguingly, in all cases the observed scatter only sets an upper limit on the intrinsic scatter.  For the relation between black hole mass and bulge velocity dispersion known as the  $\\mbh - \\sigma$ relation \\citep{ferrarese00,gebhardt00}, \\citet{tremaine02} showed that the intrinsic dispersion is no larger than 0.25--0.3 dex (or $0.58 - 0.70$ in our natural log notation). \\citet{ferrarese02} explored the relation between $\\mbh - M_h$ by using rotation curve data and argued that the dispersion between black hole and halo mass could be even less than the $\\mbh-\\sigma$ dispersion, although she notes that the uncertainty in the transformation between circular velocity and halo mass is not well characterized. Finally, \\citet{Kol06} measured the distribution in Eddington ratio, $L_{\\rm bol}/L_{\\rm Edd}$, for a sample of $z=0.3-4$ QSOs and find the distribution is well described as log-normal with a peak at 0.25 and a dispersion of 0.3 dex. As this observed dispersion must account for the uncertainty in the mass estimator and bolometric correction in addition to $L_{\\rm bol}/L_{\\rm Edd}$, this measurement also sets an upper limit to the intrinsic dispersion in $L_{\\rm bol}/L_{\\rm Edd}$.  In addition to the small scatter in $L - M_h$, the data also suggest that the duty cycle is approximately unity. This may be less surprising, given their extremely high luminosities and masses of their supermassive black holes. If QSOs are limited to accreting at no more than the Eddington rate, the luminosities of $z>6$ QSOs imply that their central, supermassive black holes already exceeded $\\mbh \\sim 10^9\\msun$ when the Universe was less than $1\\,$Gyr old \\citep{Fan01}. Subsequent near-infrared spectroscopy provides further evidence that their central black holes are indeed this massive \\citep{barth03,jiang07}. To create such massive black holes at early times requires $\\fon \\sim 1$, modulo the uncertainty in the initial seed mass. Together, these strong constraints on the duty cycle and scatter in $L-M_h$ for high-redshift QSOs provide important new information on how supermassive black holes and their host halos grew at early times.  \\medskip  We thank Yue Shen and Cris Porciani for helpful comments on an earlier draft of this paper. The simulations used in this paper were analyzed at the National Energy Research Scientific Computing Center. We thank CCAPP at The Ohio State University and the participants in the quasar mini-workshop for a stimulating meeting. JDC and MW also are grateful to CCAPP for hosting them for an extended visit. MW is supported by NASA."
    },
    "0711/0711.2944_arXiv.txt": {
        "abstract": "Observations of high-redshift Ly-$\\alpha$ sources are a major tool for studying the high-redshift Universe and are one of the most promising ways to constrain the later stages of reionization. The understanding and interpretation of the data is far from straightforward, however. We discuss the effect of the reionizing intergalactic medium on the observability of Ly-$\\alpha$ sources based on large simulations of early structure formation with radiative transfer. This takes into account self-consistently the reionization history, density, velocity and ionization structures and nonlinear source clustering. We find that all fields are highly anisotropic and as a consequence there are very large variations in opacity among the different lines-of-sight. The velocity effects, from both infall and source peculiar velocity are most important for the luminous sources, affecting the line profile and depressing the bright end of the luminosity function. The line profiles are generally asymmetric and the line centers of the luminous sources are always absorbed due to the high density of the local IGM. For both luminous and average sources the damping wing effects are of similar magnitude and remain significant until fairly late, when the IGM is ionized between 30\\% and 70\\% by mass.  The ionizing flux in the ionized patch  surrounding a high density peak is generally strongly dominated, particularly at late times, by the cluster of faint sources, rather than the central massive galaxy. Our results reproduce well the observed mean opacity of the IGM at $z\\sim6$. The IGM absorption does not change appreciably the correlation function of sources at high redshift. Our derived luminosity function assuming constant mass-to-light ratio provides an excellent match to the shape of the observed luminosity function at $z=6.6$ with faint-end slope of $\\alpha=-1.5$. The resulting mass-to-light ratio implies that the majority of sources responsible for reionization are too faint to be observed by the current surveys. ",
        "introduction": "The reionization of the universe was the last global transition of the Intergalactic Medium (IGM), from fully-neutral after cosmic recombination at $z\\sim1100$ to fully-ionized as we see it today, caused by the radiation from the first stars. Currently there are still only very few direct observational constraints on this epoch. The lack of Gunn-Peterson trough in the spectra of high-redshift sources indicates a low mean neutral fraction $x_{\\rm HI}\\lesssim10^{-4}$ out to redshift $z\\sim6$, which implies overlap was achieved sometime before that, at $z_{\\rm ov}>6$ .  On the other hand, the WMAP 3-year data \\citep{2007ApJS..170..377S} yielded a fairly high value for the integrated Thomson electron scattering optical depth to the surface of last scattering, at $\\tau_{\\rm es}=0.09\\pm0.03$. This requires a significant ionized fraction out to high redshifts, $z>12$, and thus implies an extended reionization. The optical depth by itself does not put very stringent constraints on the possible reionization histories, however. The reason for this is the self-regulated nature of the reionization process \\citep{2007MNRAS.376..534I}, whereby the Jeans-mass filtering of low-mass sources in the ionized regions results in the $\\tau_{\\rm es}$ and $z_{\\rm ov}$, the overlap redshift, being only loosely related. The overlap redshift $z_{\\rm ov}$ is determined by the abundances and efficiencies of the high-mass sources, whose formation is not suppressed by reionization, while $\\tau_{\\rm es}$ depends on both high- and low-mass sources. Thus, varying the ionizing efficiencies of the small sources yields a wide range of $\\tau_{\\rm es}$ values for the same value of $z_{\\rm ov}$.  This relative lack of observational data is set to change dramatically in the near future due to a number of large observational projects which are currently under way. The 21-cm data from high redshifts contains potentially the richest set of information since the signal is inherently three-dimensional, on the sky and in redshift/frequency \\citep[e.g.][see \\citet{2006PhR...433..181F} for a detailed recent review]{1997ApJ...475..429M,2000ApJ...528..597T,2002ApJ...572L.123I, 2003MNRAS.341...81I,2004ApJ...608..622Z,2006MNRAS.372..679M, 2006ApJ...646..681S,wmap3}. The features that could be derived include the full reionization history, geometry, statistics and individual bright features \\citep{2006MNRAS.372..679M, 2006PhR...433..181F,wmap3}. There are significant challenges to be overcome, however, particularly related to precise subtraction of the very strong foregrounds present at low frequencies \\citep[e.g.][]{2006PhR...433..181F,2006ApJ...648..767M,wmap3}.  The patchiness of reionization creates secondary temperature anisotropies in the CMB through the kinetic Sunyaev-Zel'dovich effect \\citep[][]{1998ApJ...508..435G,2000ApJ...529...12H,2001ApJ...551....3G, 2003ApJ...598..756S,2005ApJ...630..643M,kSZ}, as well as polarization anisotropies \\citep{2000ApJ...529...12H,2003ApJ...595....1H, 2003ApJ...598..756S,mortonson06,cmbpol}. Unlike the 21-cm signal, the reionization signatures in the CMB are integrated over the reionization history and contain no frequency information. However, the typical scales of reionization are reflected in a characteristic peak of the kSZ anisotropy signal \\citep{kSZ}, and the shape of the power spectrum is dependent on the reionization parameters (source efficiencies and small-scale gas clumping). CMB anisotropy observations can therefore provide us with key information about the process of reionization and since its systematics are different it would be an important complement to the 21-cm studies. One could also combine these observations more directly, by using 21-cm observations to derive the Thomson optical depth fluctuations \\citep{pol21}. Small-scale CMB anisotropy and polarization measurements would be quite difficult due to the weakness of these signals, but are within the abilities of modern detectors \\citep{kSZ,cmbpol}.  Narrow-band searches for high-redshift Ly-$\\alpha$ emitters have been very successful at finding sources at ever higher redshifts, currently up to $z\\sim7$ \\citep[e.g.][]{2002ApJ...568L..75H,2005pgqa.conf..363H,2003PASJ...55L..17K, 2005PASJ...57..165T,2003AJ....125.1006R,2004ApJ...604L..13S,2004ApJ...617L...5M, 2006NewAR..50...94B,2006PASJ...58..313S,2006ApJ...648....7K,2008ApJ...677...12O}. Together with studies of the Ly-$\\alpha$ resonant absorption in the IGM \\citep[e.g.][]{2002AJ....123.1247F,2003AJ.126..1W,2004AJ....128..515F, 2006AJ....132..117F} they provide important independent approaches for studying reionization \\citep[see][for a recent review]{2006ARA&A..44..415F}. The optical depth for Ly-$\\alpha$ resonant absorption in neutral IGM at high redshifts is quite large \\citep{1964AZh....41..801S,1965ApJ...142.1633G}, thus absorption studies are mostly sensitive to very low hydrogen neutral fractions, typically below $x_{\\rm HI}\\sim10^{-4}$, otherwise the absorption saturates. This technique is thus best suited for studying highly-ionized regions and the very end of reionization and has been very successful for setting low limits for the redshift at which reionization was completed. On the other hand, Ly-$\\alpha$ emitter surveys do not require a very low average hydrogen neutral fraction and thus in principle can probe further into the reionization epoch \\citep{2004ApJ...617L...5M}. Other related probes include damping wing measurements \\citep[e.g.][]{1998ApJ...501...15M,2004ApJ...611L..69M, 2006PASJ...58..485T}, quasar HII region sizes and their evolution \\citep[e.g.][]{2007ApJ...660..923M,2004ApJ...610..117W,2006AJ....132..117F, 2007MNRAS.374..493B,2007MNRAS.376L..34M}, and sizes of dark gaps in high-z spectra \\citep[e.g.][]{2006AJ....132..117F,2008MNRAS.386..359G}.   The correct interpretation of these data for reionization is far from straightforward, however. The strong dependence of Ly-$\\alpha$ absorption on the neutral fraction and gas density means that both should be modelled with certain precision. High-redshift sources are rare and strongly clustered \\citep[see e.g.][]{2004ApJ...609..474B,2004ApJ...613....1F,2007MNRAS.376..534I}. As a result, H~II regions generally contain multiple ionizing sources and grow much larger than the ones created by individual sources. This minimizes the effect of the damping wings and increases the transmission, allowing the detection of more and fainter sources than would be naiively expected \\citep{2005ApJ...625....1W}. However, while this is the generic expectation, the actual effect would be dependent on the exact geometry of reionization, e.g. sources close behind a neutral region will be damped even if they are inside a very large H~II region. Simplified models typically assume spherical ionized regions and either ignore source clustering or assume linear bias \\citep[e.g.][]{2000ApJ...542L..75C,2004MNRAS.349.1137S,2004MNRAS.354..695F, 2005ApJ...623..627H,2005ApJ...628..575W,2006ApJ...649..570K,2007MNRAS.374..960W}. In practice the large ionized regions form by local percolation of multiple smaller ones and as a consequence are highly nonspherical \\citep{2006MNRAS.369.1625I,mellema06,2007MNRAS.376..534I}. The inhomogeneous cosmological density fields and non-equilibrium chemistry effects (particularly in recently-ionized gas) further complicate the picture and point to the need of following the cosmological structure formation and the reionization history of a given region. A proper account of all these effects can only be done through detailed cosmological radiative transfer simulations.     A number of radiative transfer methods have been developed in recent years and now they are reaching a certain level of maturity and are producing fairly reliable results \\citep{comparison1}. However, performing large-scale reionization simulations, as required for Ly-$\\alpha$ studies is still technically very challenging. Recently \\citet{2007MNRAS.381...75M} used large scale structure formation numerical simulations postprocessed with radiation transfer to study the observability of Lyman-$\\alpha$ emitters at high redshifts and what these can tell us about reionization. In order to achieve high dynamic range, these authors employed a subgrid model for the collapse of the smallest halos. Another, semi-numerical approach was used by \\cite{2008MNRAS.385.1348M}, who took the linear density and velocity fields at early time (essentially the initial conditions for an N-body simulation) and used an excursion-set approach combined with a first-order Lagrangian theory to ``paint'' the H~II regions on the density field. This procedure provides large dynamic range at a low cost, but at the expense of making significant approximations and thus cannot fully replace full numerical simulations.  In this paper we use the results of large scale numerical simulations to study the transfer of Ly-$\\alpha$ through the IGM. In \\S~2 we briefly describe our simulation method. In \\S~3 we describe the evolution and environment of a rare, massive source, similar to the ones that are currently observed. In \\S~4 we address the observability of Lyman-$\\alpha$ emitters, considering the reduction of the transferred line flux due to the absorption in the IGM and luminosity functions. In \\S~5 we describe the effects of the patchiness of reionization on the angular correlation function of Lyman-$\\alpha$ emitters, and, finally, in \\S~6 we sum up our conclusions. ",
        "conclusions": "We considered the effects which the reionizing IGM has on the observations of high-redshift Ly-$\\alpha$ sources. For this we utilized detailed structure formation and radiative transfer simulations, which allowed us to evaluate many features which can only be studied by detailed simulations, as well as to quantify better a number of previously-proposed effects. We followed the full reionization history self-consistently and accounted for the actual source distribution, neutral fraction, density and velocity fields.  We find that the density, neutral fraction and velocity fields are all highly anisotropic, which results in large variations in the IGM transmission and source visibility among different LOS. The velocity effects, both gas infall and source peculiar velocity are most important for massive, luminous sources. The most luminous sources are found in highest peaks of the density field, which at late times are significantly overdense out to $\\sim10$~comoving Mpc (cMpc) and are surrounded by infall extending to $\\sim20$~cMpc. The infall of gas blueshifts it in frequency space and results in significant absorption on the red side of the line center, while the peculiar velocity of the source itself can either alleviate or exacerbate this effect, depending on the halo- and infall velocity alignment.  The spherically-averaged local density enhancement and gas infall have been modelled analytically in approximate ways \\citep{2004MNRAS.347...59B}, and thus can be incorporated in semi-analytical models \\citep[e.g.][]{2007MNRAS.377.1175D}. However, such models are unable to account for the strong intrinsic anisotropies of the neutral fraction, density and velocity fields. The analytical and semianalytical models typically assume spherical symmetry and full ionization inside the H~II regions, both of which assumptions are quite unrealistic.  The Ly-$\\alpha$ lines we derive are generally asymmetric and vary hugely from LOS to LOS. The luminous sources form at the highest density peaks and as a consequence their line centers are always highly-absorbed even though their proximity regions are very highly ionized, with typical neutral fractions $x_{\\rm HI}\\sim10^{-5}-10^{-6}$. The luminous sources also more affected by infall and exhibit more pronounced proximity region with higher transmission of the blue wing of the line.  High-redshift sources are strongly clustered around the high peaks of the density field. The central source contributes the majority of the ionizing flux only in its immediate vicinity, within 1-2 comoving Mpc. Beyond that distance the ionizing flux is dominated by the fainter sources clustered around it. This dominance is particularly strong at late times, when both many more sources form and the ionized regions become larger, resulting in the fainter sources contributing up to 2 orders of magnitude more photons than the central source.  Compared to single-source ionized bubbles, the larger H~II regions from clustered sources diminish the effects from the damping wing of the line. Nevertheless, these remain significant until fairly late (ionized mass fraction $x_m=0.3-0.7$, which for the simulation considered here corresponds to redshifts $z\\sim9-8$). Interestingly, the average damping wing effect is similar for luminous and typical sources, even though naiively one might expect that damping could be weaker for the former, since they are typically in the middle of large bubbles, away from the neutral patches, unlike the fainter sources, which are more evenly distributed.  Both the mean IGM transmission and the typical photoionization rates we find are high compared to observations at $z\\sim6$, indicating that our adopted source efficiencies are also high. The mean IGM transmissivity decreases only slowly towards the higher redshifts and the GP transparency occurs significantly after the actual overlap epoch. For the simulation considered here overlap (defined as 1\\% neutral fraction) occurs at $z=6.6$, while average neutral fraction of $10^{-4}$ is reached only by $z=6$ and even then relatively small fraction (few to 10\\%) of the flux is transmitted. By overlap the spectra start showing significant transmission gaps in the mean IGM (i.e. away from the proximity region of a luminous source).  We find that for a given number density of sources (e.g. as determined by observations) the clustering of these sources depends only weakly on the IGM absorption during reionization. As a consequence, the reionization patchiness has little effect on the observed Ly-$\\alpha$ source clustering, which implies that source clustering is not a good indicator for reionization patchiness.  Our derived luminosity function assuming constant mass-to-light ratio provides an excellent match to the shape of the observed luminosity function at $z=6.6$ with faint-end slope of $\\alpha=-1.5$. The resulting mass-to-light ratio implies that the majority of sources responsible for reionization are too faint to be observed by the current surveys."
    },
    "0711/0711.0988_arXiv.txt": {
        "abstract": "We present constraints on the number of Galactic magnetars, which we have established by searching for sources with periodic variability in 506 archival \\chandra\\ observations and 441 archival \\xmm\\ observations of the Galactic plane ($|b|$$<$5\\degree). Our search revealed four sources with periodic variability on time scales of 200--5000 s, all of which are probably accreting white dwarfs.  We identify 7 of 12 known Galactic magnetars, but find no new examples with periods between 5 and 20 s. We convert this non-detection into limits on the total number of Galactic magnetars by computing the fraction of the young Galactic stellar population that was included in our survey. We find that easily-detectable magnetars, modeled after persistent anomalous X-ray pulsars (e.g., either with X-ray luminosities $L_{\\rm X}$=$3\\times10^{33}$ \\ergs\\ [0.5--10.0 keV] and fractional rms amplitudes $A_{\\rm rms}$=70\\%, or $L_{\\rm X}$=$10^{35}$ \\ergs\\ and $A_{\\rm rms}$=12\\%), could have been identified in $\\approx$5\\% of the Galactic spiral arms by mass. If we assume that there are 3 previously-known examples within our random survey, then there are 59$^{+92}_{-32}$ in the Galaxy. Barely-detectable magnetars ($L_{\\rm X} = 3\\times10^{33}$ \\ergs\\ [0.5--10.0 keV] and $A_{\\rm rms}$$=$15\\%) could have been identified throughout $\\approx$0.4\\% of the spiral arms. The lack of new examples implies that $<$540 exist in the Galaxy (90\\% confidence). Similar constraints are found by considering the detectability of transient magnetars in outburst by current and past X-ray missions. For assumed lifetimes of $10^{4}$ yr, we find that the birth rate of magnetars could range between 0.003 and 0.06~yr$^{-1}$. Therefore, the birth rate of magnetars is at least 10\\% of that for normal radio pulsars. The magnetar birth rate could exceed that of radio pulsars, unless the lifetimes of transient magnetars are $\\ga$$10^5$~yr. Obtaining better constraints will require wide-field X-ray or radio searches for transient X-ray pulsars similar to XTE~J1810--197, AX~J1845.0--0250, CXOU~J164710.2--455216, and 1E~1547.0-5408. ",
        "introduction": "Astronomers have only started to appreciate the diversity of the properties of neutron stars that are produced when a massive star collapses and explodes \\citep{ptp06}. The list of different manifestations of neutron stars now includes: radio pulsars that are powered by the rotation of their $10^{8}-10^{13}$ G magnetic fields \\citep[e.g.,][]{lor06}; accreting X-ray pulsars \\citep[e.g.,][]{bil97,wvdk98} and thermonuclear bursters \\citep[e.g.,][]{sb06} that are accreting matter from binary companions; magnetars that are powered by the decay of their $\\ga$$10^{14}$ G fields \\citep[which historically have been categorized as either anomalous X-ray pulsars or soft gamma repeaters; e.g.,][]{wt06}; intermittently-detectable ``Rotating RAdio Transients'' \\citep[RRATs;][]{mcl06}; isolated, cooling neutron stars that shine primarily in soft X-rays \\citep{wwn96,hab06}, and central compact objects that are seen as point sources near the centers of supernova remnants \\citep[][]{cha01,sew03,pst04}. Understanding the properties of these compact objects and their birth rates provides important constraints on the late-time evolution of massive stars, and on the processes that occur during stellar collapse. The relationships among the different classes of compact object could reveal how their magnetic fields decay and their interiors cool.   In this paper, we present a search for magnetars. This search is timely for two reasons. First, recent evidence suggests that magnetars are the products of unusually massive progenitors. Three magnetars have been found to be in clusters of massive, young stars \\citep{fuc99,vrb00,eik04}, and the turn-off masses of two of these clusters imply that the progenitors to the neutron stars were very massive, 30--40 \\msun\\ \\citep{fig05,mun06}. A fourth magnetar has been associated with a bubble of neutral hydrogen that was probably blown by the wind of a $>$30 \\msun\\ progenitor \\citep{gae05}.  This suggests that massive stars may be more likely to produce magnetars, whereas ordinary radio pulsars are generally presumed to be left by lower-mass, 8--20 \\msun\\ progenitors \\citep[e.g.,][]{heg03}. Given that less massive stars are much more common \\citep[e.g.,][]{krou02}, if massive stars produce magnetars, one would expect that their birth rates should be much lower than those of radio pulsars \\citep{gae05}.  Second, there is significant debate about how the strong magnetic fields that characterize magnetars are produced. The original hypothesis is that magnetars are born with millisecond periods, and that the strong fields are produced by a dynamo in the rapidly-rotating proto-neutron star \\citep{td93,heg05}. However, observations of supernova remnants associated with magnetars rule out the expected input of energy from neutron stars with initial spin periods $\\la$3 ms \\citep{vk06}. At the same time, the discovery of a few OB stars with $10^3$ G surface fields \\citep{don02,don06a,don06b} has motivated the alternative hypothesis that magnetar-strength fields are primordial, having been amplified only by the collapse of the core \\citep{fw06}. This second hypothesis is attractive because it makes a straightforward prediction, that the birth rate of magnetars should be equal to that of highly-magnetized OB stars. However, it cannot yet explain why some highly-magnetized neutron stars are radio pulsars instead of magnetars \\citep[e.g.,][]{pkg00}. This discrepancy is one of the main points in favor of the $\\alpha$--$\\Omega$ dynamo process that would act in a rapidly-rotating proto-neutron star, because it could produce $\\ga$$10^{15}$~G internal fields that would power the magnetars \\citep{tlk02}.  Here, we report the results of our search for new Galactic magnetars in archival observations taken with the {\\it Chandra X-ray Observatory} and the {\\it XMM-Newton} Observatory. Previous searches for periodic sources with these observatories have been limited to small fields, such as the Small Magellanic Cloud \\citep{mac03,edge04} and the central 20 pc of the Galaxy \\citep{mun03}. Our search incorporates observations through the entire Galactic plane. Although we find four new sources with significant periodic signals, none are likely to be magnetars. Therefore, we use the survey to place limits on the number of active magnetars in the Galaxy, and discuss the implications for their birth rates.  \\begin{figure*}[htb] \\centerline{\\psfig{file=survey_map.ps,width=\\linewidth}} \\caption{The locations of archival observations used for this survey. Black squares are \\chandra\\ observations, schematically represented with boxes that have an area equal to that of the ACIS-I. Grey rectangles are \\xmm\\ observations, schematically represented with boxes that have an area equal to that of one EPIC MOS.} \\label{fig:coverage} \\end{figure*} ",
        "conclusions": "We have serached archival \\chandra\\ and \\xmm\\ observations for X-ray pulsars, in order to constrain the Galactic population of magnetars. Although we found four objects with periodic variability on time scales of 200--5000~s (Figs.~\\ref{fig:cxcfft}--\\ref{fig:xmmprof} and Table~\\ref{tab:signals}), we found no sources that were obviously magnetars. The archival observations that we used covered a moderate fraction of the young stellar populations from which magnetars descend. Our search was sensitive to the standard, easily-detectable AXP-like magnetar ($L_{\\rm X}$=$10^{35}$ \\ergs\\ [0.5--10 keV] and a sinusoidal pulse profile with $A_{\\rm rms}$=0.12, or equivalently $L_{\\rm X}$=$3\\times10^{33}$ \\ergs\\ and $A_{\\rm rms}$=0.7) with 90\\% completeness for $\\approx$5\\% of the mass in the Galactic spiral arms. Searching for transient AXPs in quiescence ($L_{\\rm X}$=$3\\times10^{33}$ \\ergs\\ and $A_{\\rm rms}$=0.15), however, we were only sensitive over 0.4\\% of the Galactic spiral arms.  Based on the number of known Galactic magnetars, we then placed constraints on the total population. Our archival search placed strict constraints on the number of standard, persistent AXPs, which number $59^{+92}_{-32}$ with 90\\% confidence. We also placed constraints on the number of SGR-like magnetars as a function of the rate at which they produce outbursts (Fig.~\\ref{fig:numsgrs}), and find that for a likely recurrence rate of 0.01 $yr^{-1}$, there would be $13^{+18}_{-7}$ in the Galaxy. However, these populations may be dwarfed by the number of transient AXPs similar to XTE~J1810--197 \\citep[see also][]{ibr04}. No new examples were detected in archival \\chandra\\ and \\xmm\\ observations, so up to 540 quiescent AXPs could be present in the Galaxy. We also considered the fact that two transient AXPs have been discovered in outburst by \\asca\\ and \\rxte. The fact that \\asca, \\rxte, \\chandra, and \\xmm\\ together have surveyed 2\\% of the Galaxy each year for the last 11 years allows us to constrain the number of transients as a function of their recurrence time (Fig.~\\ref{fig:numtaxps}). For a recurrence time of once every 10 years, this calculation suggests there are between 80 and 580 transient AXPs, with 90\\% confidence. Assuming a lifetime of $10^4$ yr for persistent magnetars \\citep{kou99,ggv99}, their birth rate is at least 10\\% of that of radio pulsars, and could equal the radio pulsar birth rate \\citep[e.g.,][]{lor06,fgk06}. The birth rate of transient magnetars could exceed those of radio pulsars, unless their lifetimes are significantly longer than $10^4$ years.  The main uncertainty in the number of Galactic magnetars is introduced by the unknown activity level of transient examples. Further monitoring of known transient AXPs would constrain their duty cycles and reduce this uncertainty, so long as the known examples can be taken as representative of all magnetars. Ultimately, though, a wide-field monitoring campaign is needed to identify more transient AXPs in outburst. This can be accomplished in principle in the X-ray band, although the sensitivity of current and planned wide-field monitors \\citep[$\\approx$$2\\times10^{-11}$ \\ergcms, 0.5--10 keV;][]{lev96,rem00,bm06} is only sufficient to barely detect a $10^{35}$ \\ergs\\ magnetar to a distance of 4 kpc from Earth. On the other hand, the detections of XTE J1810--197 \\citep{cam06} and 1E~1547.0--5408 \\citep{cam07} in the radio with characteristic intensities of $\\ga$100 mJy kpc$^{2}$ suggests that future radio surveys with the Low-Frequency Array, the Mileura Wide-field Array, or the Square Kilometer Array \\citep[e.g.,][]{clm04} could provide the best constraints on the numbers of magnetars in the Galaxy."
    },
    "0711/0711.2513_arXiv.txt": {
        "abstract": "In this paper, we study the large scale structure formation of the normal branch in DGP model (Dvail, Gabadadze and Porrati brane world model) by applying the scaling method developed by Sawicki, Song and Hu for solving the coupled perturbed equations of motion of on-brane and off-brane. There is detectable departure of perturbed gravitational potential from LCDM even at the minimal deviation of the effective equation of state $w_{\\rm eff}$ below $-1$. The modified perturbed gravitational potential weakens the integrated Sachs-Wolfe effect which is strengthened in the self-accelerating branch DGP model. Additionally, we discuss the validity of the scaling solution in the de Sitter limit at late times. ",
        "introduction": "Dvali, Gabadadze and Porrati proposed a brane world model embedded in the Minkowski 5D dimension~\\cite{dvali00}, which has been taken as an alternative theory to explain cosmic acceleration. There are two distinct branches of this model: one is the self-accelerating branch (hereafter sDGP) and the other is the normal branch (hereafter nDGP).  Deffayet discovered the accelerated phase in sDGP without cosmological constant~\\cite{deffayet00}, which has been tested geometrically by many authors~\\cite{dvali03a,Avelino01,fairbairn05,Maartens06,Song06}. The detectability will be enhanced by the structure formation test ~\\cite{song04,Schimd04,gabadadze04,ishak05,Knox:2006fh,linder05,gabadadze05,Song06sa,tang06,huterer06,Carroll06,Stabenau06,Pal06hg,Pal07,Hu07}. The distinct evolution of perturbed gravitational potential of sDGP has been studied by many authors at scales in quasi-static limit~\\cite{deffayet02a,lue04,koyama05,Sawicki06} and scales at nearly horizon size~\\cite{Sawicki06}. sDGP strengthens the integrated Sach-Wolfe (hereafter ISW) effect, and leaves a detectable signature on CMB power spectra and ISW-galaxy cross-correlation~\\cite{Song06}.  There is no cosmic acceleration induced spontaneously in nDGP without introducing cosmological constant. But it's attracting features are that nDGP has a spectrum leading to the stability in quantum level~\\cite{Charmousis06}, and that the effective equation of state $w_{\\rm eff}$ crosses $w_{\\rm eff}=-1$ without phantom dark energy. nDGP has been tested by many authors geometrically as a natural theory crossing $w_{\\rm eff}=-1$ line below~\\cite{Sahni02,Chimento06,Lazkoz06}. But no detailed study has been done to show how to formulate the large scale structure formation of nDGP in order to improve the detectability.  We study the large scale structure formation of nDGP by using the scaling method developed by Sawicki, Song and Hu~\\cite{Sawicki06}. On-brane equations of motion are not in the closed form without the information of the gradient crossing the bulk direction. Off-brane equation of motion should be coupled in order to supply this gradient of perturbations at the location of brane. We assume the scaling ansatz for the perturbation on the brane, and solve it iteratively until it converges. We find that the quasi-static limit is valid at scales relevant to ISW-galaxy cross-correlations. The modified perturbations weaken the ISW effect, which is the opposite behavior of perturbations found in sDGP. While we are preparing this paper, Cardoso, Koyama, Seahra and Silva independently derives the identical solution by using different approach~\\cite{Cardoso07}. They verify that the scaling method is valid for both branches of DGP. ",
        "conclusions": "We have solved the dynamic equations of nDGP by using the scaling method. A detectable departure from LCDM is found with $w_{\\rm eff}$ below $-1$. The matter fluctuations are enhanced due to the deeper potential well in nDGP than in LCDM, which weaken the ISW effect. We are able to probe nDGP by detecting this weakened ISW effect by ISW-galaxy or ISW-weak lensing cross-correlation. It is highly recommended to select high redshift bins where the modified perturbed potential increases to provide the negative cross-correlation. If found, this would be strong evidence for the presence of the nDGP model~\\cite{Giannantonio07}."
    },
    "0711/0711.2039_arXiv.txt": {
        "abstract": "We now have a good measurement of the cooling rate of G117-B15A. In the near future, we will have equally well determined cooling rates for other pulsating white dwarfs, including R548. The ability to measure their cooling rates offers us a unique way to study weakly interacting particles that would contribute to their cooling. Working toward that goal, we perform a careful asteroseismological analysis of G117-B15A and R548. We study them side by side because they have similar observed properties. We carry out a systematic, fine grid search for best fit models to the observed period spectra of those stars. We freely vary 4 parameters: the effective temperature, the stellar mass, the helium layer mass, and the hydrogen layer mass. We identify and quantify a number of uncertainties associated with our models. Based on the results of that analysis and fits to the periods observed in R548 and G117-B15A, we clearly define the regions of the 4 dimensional parameter space ocuppied by the best fit models. ",
        "introduction": "G117-B15A and R548 are pulsating white dwarfs with atmospheres dominated by hydrogen. These stars are called DAVs or ZZ Ceti stars. Their effective temperatures are close to 12,000 K. They are non-radial, g-mode pulsators, where the restoring force is buoyancy. Because of its potential use to study axions and other exotic physics \\citep{isern, Kepler2}, G117-B15A was the object of several asteroseismological studies. R548 is similar to G117-B15A both spectroscopically and pulsationally so studying them side by side may provide additional clues. \\citet{paul98} performed such a parallel study. More recently \\citet{benvenuto} published an asteroseismological fit to G117-B15A's period spectrum. \\citeauthor{benvenuto} evolved models from the Main Sequence and included detailed physics, such as time-dependent diffusion. Their best fit model has an effective temperature and a mass consistent with spectroscopic results, but the periods of the model do not match the observed period spectrum as well as \\citeauthor{paul98}'s fits do.  We take a new approach to white dwarf asteroseismology, enabled by the computing resources available today. Using models similar to \\citeauthor{paul98}'s, we perform a systematic, fine grid search for the best fits to the period spectra of G117-B15A and R548. We study both stars side by side. While we do not ignore the spectroscopy entirely, we take a naive approach and base much of our results on the period fitting alone. At the end we compare our results with the spectroscopy and previous studies. In the same naive spirit, we treat the helium layer mass as another free parameter. \\citet{paul98} and \\citet{benvenuto} both fixed its value at $10^{-2} M_*$, where $M_*$ is the total mass of the model.  We organize our paper as follows. In section 2, we present the observables we have for G117-B15A and R548. In section 3, we summarize concisely the asteroseismological results of \\citet{benvenuto} and \\citet{paul98}. In section 4, we explain our method. We include a discussion of our models and a quantitative study of what input parameters matter most. We present our results in section 5, along with a discussion of how they compare with the fits done by \\citet{paul98} and \\citet{benvenuto}. We summarize and discuss our results in section 6. ",
        "conclusions": "We performed a systematic fine grid search for best fit models to the two DAVs G117-B15A and R548. We find best fit models for both stars between 11600 and 12700 K and between 0.59 and 0.66 \\msunm. In both cases, the region occupied by the best fit models in the \\massteff plane is related to the thickness of the hydrogen layer. Treating the helium layer mass as a free parameter, we discovered that the lowest period mode for each star (215s for G117-B15A and 213s for R548) singlehandedly sets the helium layer mass to $4 \\times 10^{-3}$. Both stars are well fit with thin hydrogen layer models ($\\rm 10^{-7} < M_H < 10^{-8}$). For G117-B15A, we find a second family of solutions between 11300 and 12300 K and between 0.65 and 0.68 \\msunm. Those fits have slightly thicker hydrogen layers ($\\rm M_H = 6.3 \\times 10^{-7}$).  For R548 we find a unique, robust mode identification. The dominant modes (213s and 274.5s) are $\\ell=1$, $k=1,2$ modes. The 318s mode is also an $\\ell=1$ mode, with $k=4$. The last two modes (187s and 334s) are $\\ell=2$, $k=4$ and 8 respectively. For G117-B15A, we find two distinct families of best fit models. Models with $\\rm M_H = 6.3 \\times 10^{-7}$ all have the same mode identification, namely $k=2,3$, and 4 for the three observed periods (215s, 271s, and 304s). Models with $\\rm M_H \\simeq 1-4 \\times 10^{-8}$ are consistent with G117-B15A's periods being $k=1,2$, and 3. R548 and G117-B15A have similar observed properties and based on that fact, it is likely that they have similar structures. In this case, the second class of models appears better, but we cannot discard the first class of models based on that fact alone.  By sampling parameter space systematically and homogeneously, we found that our models were offset in mass and temperature compared to the spectroscopy (the models are hotter and/or more massive). \\citeauthor{reid96} (1996) measured the gravitational redshift of spectral lines for G117-B15A. The resulting mass is in tune with the spectroscopic mass ($0.536 \\pm 0.010 \\msun$). Both different equation of state tables and opacities could alter the structure of the models and possibly modify our results. We do not expect either to have a large effect, however. Just recently, \\citet{cassisi} published new electron-conduction opacities that treat the partially degenerate regime relevant in white dwarf envelopes better. According to \\citeauthor{cassisi}, the new opacities differ only by at most a factor of 2 in white dwarf envelopes and have very little effect on the mass-radius relation. We also tried to revert back to the \\citet{oldeos} envelope equations of state and found that this had a negligible effect on the gravities in our models.  Results from section 4.2 suggest a more promising avenue. They show that structure in the core and the shape of the helium composition profile at the carbon-helium interface can have a large effect on the pulsations periods of a model. We adopted core composition profiles from stellar evolution calculations and did not vary them. Our asteroseismological results suggest that perhaps we should try different core composition profiles. With precise enough determinations of G117-B15A and R548's mass and effective temperature from spectroscopy or other independent methods, one can turn the problem around and attempt to determine what composition profiles are needed in order to reconcile the asteroseismology with the spectroscopy. While preliminary results appear promising, a full investigation requires the clever analysis of a more extensive grid of models and we leave that for a future publication.  Once we are satisfied that we have models that match both asteroseismologically and spectroscopically, we can calculate \\pdots \\ for our models and use them together with the observed \\pdots \\ to constrain the energy loss rate due to any weakly interacting particles, such as axions. The fine grid approach allows us to formaly assess the uncertainties in our models' parameters and to obtain tight constraints on the emission rate of weakly interacting particles in G117-B15A."
    },
    "0711/0711.3515_arXiv.txt": {
        "abstract": "We  show that collisions between the outer Galactic H~{\\sc i}  disk and the leading arms (LAs) of the Magellanic stream (MS) can create giant H~{\\sc i}  holes and chimney-like structures in the disk. Based on the results of our  N-body simulations on the last 2.5 Gyr evolution of the Large and Small Magellanic Clouds (LMC and SMC, respectively) interacting with the Galaxy, we  investigate  when and where the LAs can pass through the Galactic plane after the MS formation. We then investigate hydrodynamical interaction between  LAs and the Galactic H~{\\sc i}  disk (``the Magellanic impact'') by using our new hydrodynamical simulations with somewhat idealized models of the LAs. We find that about 1-3\\% of the initial gas mass of the SMC, which consists of the LAs, can pass through the outer part ($R=20-35$ kpc) of the Galactic H~{\\sc i}  disk  about 0.2 Gyr ago. We also find that the Magellanic impact can push out some fraction ($\\sim 1$\\%)  of the outer Galactic H~{\\sc i}   disk to form 1-10 kpc-scale H~{\\sc i}  holes and chimney-like bridges between the LAs and the disk. ",
        "introduction": "H~{\\sc i}  gas in the Magellanic stream, its leading arms, and the Magellanic bridge have long been discussed in the context of dynamical and hydrodynamical interaction between the Magellanic Clouds (MCs) and the Galaxy (e.g., Murai \\& Fujimoto 1980; Putman et al. 1998; Mastropietro et al. 2005; Muller \\& Bekki 2007). Such intense interactions between the three galaxies have been also  suggested to be closely associated not only with the long-term star formation histories of the LMC and the SMC (e.g.,  Harris \\& Zaritsky 2004; Cioni et al. 2006) but also with structural and kinematical properties of the MCs (e.g., Bekki \\& Chiba 2005, BC05). It is however not fully understood how the interactions can change dynamical  and chemical properties  and star formation histories of  the MCs (e.g., Westerland 1997).     Although many previous numerical/theoretical works discussed influences of the Galaxy on the evolution of the MC (e.g., Gardiner \\& Noguchi 1995; BC05; R\\. u\\v zi\\v cka et al. 2007), only a few works investigated the possible gravitational influences of the MCs on the dynamical evolution of  the Galactic {\\it the stellar disk} (e.g., Weinberg 1998; Tsuchiya 2002) and the gaseous one (Weinberg \\& Blitz 2006). For example, these works discussed whether the observed Galactic warp can be due to the dynamical interaction between the LMC and the Galaxy (e.g., Tsuchiya 2002). Although previous numerical and observational studies on the MS formation suggested possible interaction between the LAs and the Galaxy (e.g., Yoshizawa \\& Noguchi 2003; Br\\\"uns et al. 2005; Connors et al. 2006), the details of the interaction processes have not been investigated yet. In particular, it is totally unclear {\\it how the gaseous components of the Galaxy are influenced by the gaseous components of the MCs}.   Recent observational studies on the physical properties of the outer H~{\\sc i}  disk of the Galaxy have reported a  significant ($\\sim 10$ km s$^{-1}$) velocity perturbation in the outer H~{\\sc i}  spiral arm of the Galaxy at $l=275^{\\circ}-295^{\\circ}$ (McClure-Griffiths et al. 2004, M04). Furthermore, it remains unclear why the outer H~{\\sc i}  disk ($19<R<28$ kpc) appears to show  significantly lower densities between the Galactic longitudes $l=285^{\\circ}$ and $300^{\\circ}$ (See Figure 2 in M04). No theoretical works have however discussed so far the origin of these unique properties in M04.    The purpose of this Letter is to show, for the first time, how collision between the LAs and the Galactic H~{\\sc i}  disk (hereafter referred to as ``the Magellanic impact'') influences the evolution of the Galactic H~{\\sc i}  disk. We first show when and how the LAs can pass through the Galactic disk based on our collisionless  N-body simulations on the evolution of the MCs for the last 2.5 Gyr (e.g., Bekki \\& Chiba 2007, BC07). We then show that the Magellanic impact can create unique, kpc-scale H~{\\sc i}  holes and chimney-like bridges connecting the LAs and the outer part of the Galactic H~{\\sc i}  disk based on new yet somewhat idealized hydrodynamical simulations on the Magellanic impact. ",
        "conclusions": "We first have shown that the Magellanic impact can create 10kpc-scale H~{\\sc i}  holes and bridges connecting the LA and the Galaxy. The physical properties of the simulated unique structures in the outer part ($R\\sim 30$ kpc) of the Galactic gas disk can be tested against previous and ongoing observational studies on H~{\\sc i}  properties of the Galaxy (e.g., M04; Levine et al. 2006; M06). Although detailed comparison of ongoing and future observations will be given in our future papers, here we briefly discuss previous observations shown in Figures 1 and 2 of M04 compared to the  Model B in a more quantitative manner.   Figure 4 clearly shows the longitude-velocity  ($l-v$) diagram of particles of the Galactic gas disk in the Model B, where $v$ denotes the line-of-sight velocities of the particles  with respect to the Sun. Clearly the  $l-v$ diagram shows a bifurcation around at $l=280^{\\circ}-300^{\\circ}$, the most clearly  at $l \\sim 290^{\\circ}$. This bifurcation is due to the dynamically disturbed gas disk of the Galaxy by the Magellanic impact and  may correspond to a part of the observed  significant ($\\sim 10$ km s$^{-1}$) velocity perturbation in the outer H~{\\sc i}  spiral arm of the Galaxy at $l=275^{\\circ}-295^{\\circ}$ (see Figure 1 of M04). Figure 4 also shows a significantly less populated region  at  $l=280^{\\circ}-300^{\\circ}$ in  the particle distribution on the $(l-b$) plane owing to the presence of the giant H~{\\sc i}  hole there created by the Magellanic impact. This is broadly consistent with observations shown in Figure 2 of M04, which shows that the outer H~{\\sc i}  disk of the Galaxy shows  significantly lower densities between the Galactic longitudes $l=285^{\\circ}$ and $300^{\\circ}$.   Velocities of particles pushed out from the Galactic plane can be more than 50 km s$^{-1}$ higher than those within the Galactic disk in Figure 4, which can not be seen in Figure 1 of M04. This apparent inconsistency possibly means that such high-velocity gas with low-density and low-mass ($\\sim 10^3 {\\rm M}_{\\odot}$) can be soon changed into ionized gas owing to possible interaction with warm/hot  halo gas of the Galaxy so that it can not be observed as H~{\\sc i}. Although the observed $l-v$ diagram  in M04 does not clearly show bifurcation, the H~{\\sc i} seems to be distributed on two belts for $260^{\\circ} <  l < 295^{\\circ}$ and $v>$  60 km s$^{-1}$. It remains observationally unclear whether this intriguing distribution can correspond to a bifurcated distribution of H~{\\sc i} on the $l-v$ diagram.    The direction of the orbit of the LA (with respect to the Galaxy) can significantly change at the Galactic plane (i.e., $b=0^{\\circ}$) as a result of hydrodynamical interaction between the LA and the Galactic H~{\\sc i}  disk.  The orbital change at  $b=0^{\\circ}$ due to the back reaction of the Magellanic impact can explain why the observed distribution of the LA in the Galactic coordinate shows a  ``kink'' at  $b=0^{\\circ}$ (e.g., Br\\\"uns et al. 2005). Furthermore  the total H~{\\sc i}  mass of the LA can be significantly reduced by photoionization of the H~{\\sc i}  gas of the LA by the warm gas of the Galactic disk during the Magellanic impact. This significant reduction of the H~{\\sc i}  mass in the LA can solve the well known problem as to the observed small mass ratio ($\\sim 0.2$) of the LA to the MS (e.g., Br\\\"uns et al. 2005).    It is possible  that  the Magellanic impact can trigger some amount  of  star formation in the outer part of the  Galactic gas disk owing to collisions between gas clouds from the Galaxy and those from the LA. We  thus suggest that the presence (or the absence) of B-type stars with ages of $0.1-0.2$ Gyr at $270^{\\circ} <  l < 300^{\\circ}$ and $b \\sim 0^{\\circ}$  would prove (or disprove) the past star formation activities of the Galaxy induced by the Magellanic impact. We also suggest that hydrodynamical interactions between gaseous  tidal tails and debris  from disrupting low-mass galaxies (like the SMC) and outer parts of galactic H~{\\sc i}  disks can be proved by observations on detailed spatial distributions  of  H~{\\sc i} in galaxies. In particular,  future structural and kinematical studies on the observed giant H~{\\sc i} holes in the very outer parts of galaxies (e.g., NGC 6822;  de Blok \\& Walater 2003) will allow us to discuss whether or not hydrodynamical interaction between gaseous tidal tails and outer gas disks of galaxies can be important for the evolution of the outer disks."
    },
    "0711/0711.3209_arXiv.txt": {
        "abstract": "We study the influence of a strong AGN outburst on the surrounding galaxies. The AGN is assumed to reside in a group of galaxies, and an outburst excites a shock wave in the hot gas in the group. We calculate the impact of the shock wave on the galaxies. We find that if the energy of the outburst is extremely large ($E_{\\rm AGN}\\sim 6\\times 10^{61}$~erg) as the one recently observed in clusters, the impact is strong enough to strip the cold interstellar medium in the disc of the galaxies in the inner region of the group. Moreover, even in the outer region of the group, the warm gas in the halo of the galaxies would be stripped, even if the energy of the outburst is $\\sim 6\\times 10^{60}$~erg. These would decrease star formation activity of the galaxies. If these galaxies fall into the group centre through dynamical friction and their interstellar medium is the fuel of the supermassive black hole in the AGN, the outburst would serve as feedback. While this mechanism works only when $E_{\\rm AGN}$ is extremely large, such outbursts have not been observed in groups at low redshift; it would work at high redshift rather than at low redshift. ",
        "introduction": "\\label{sec:intro}  X-ray observations have shown that hot gas in groups and clusters of galaxies has been heated by some sources in addition to gravity. This was shown by the fact that the luminosity and temperature of a group or cluster follow a scaling relationship (the $L_\\rmn{X}$--$T$ relation), $L_\\rmn{X}\\propto T^3$, \\citep{edg91,all98,mar98,arn99}, which is at odds with that expected for groups and clusters formed by gravitational structure formation, with $L_\\rmn{X}\\propto T^2$ \\citep{kai86}. More recently, it was shown that the entropies of the hot gas in groups and clusters, especially groups, are higher than those predicted by models of gravitational structure formation \\citep*{pon99}.  Supernova-driven galactic winds have been considered the heating source \\citep*{wu98,men00,loe00}. However, the energy from supernovae alone seems to be insufficient to heat the hot gas to the observed level \\citep*{val99,wu00,bow01}. Therefore, active galactic nuclei (AGNs) are now recognised as another promising candidate of the heating source \\citep{ino01}.  AGNs are often found at the centres of groups and clusters \\citep{mcn00,fab00,bla01}. The energy ejected by an AGN can cancel radiative cooing of the hot gas and may prevent development of a cooling flow in the central region of a group or a cluster. For most of the AGNs at $z\\sim 0$, however, the power is not enough to heat the hot gas on a group or cluster-scale, and would be insufficient to account for the entropy excess found in groups and clusters. Recently, however, shock waves associated with extremely powerful AGN outbursts have been found in some clusters \\citep*{mcn05,nul05a,nul05b,wis07,git07}. They seem to be powerful enough to heat the gas on a group or cluster-scale.  Such strong outbursts would also affect the surrounding environment in a form other than heating. For example, \\citet{fuj07} showed that the shock wave excited by an outburst accelerates particles and that the emission from the accelerated particles could be responsible for radio mini-halos observed in clusters. \\citet{raw04} discussed the regulation of galaxy formation by outbursts.  In this letter, we consider another effect of strong AGN outbursts on the environment. We focus on the interaction between the shock wave produced by an outburst and the galaxies surrounding the AGN. We consider an AGN and galaxies in a small group of galaxies with the mass of $10^{13}\\rm\\: M_\\odot$ rather than a cluster of galaxies. In clusters, another interaction between the hot gas and the galaxies, called ram-pressure stripping, is effective (see Section~\\ref{sec:discuss}). Thus, the influence of AGN outbursts would be obscured.  It should be noted that such strong outbursts are rare phenomena at low redshift. \\citet{git07} estimated that strong outbursts ($E_{\\rm AGN}\\sim 10^{61}$~erg) are likely to occur only $\\sim 10$\\% of clusters at low redshift. For groups, the fraction may even be smaller because such outbursts have not been observed. However, at high redshift, they would be more common \\citep[][see Section~\\ref{sec:discuss}]{ued03}.  In this letter, the cosmological parameters are $\\Omega_0=0.3$, $\\lambda_0=0.7$, and $h=0.7$, where $H_0=100\\: h\\rm\\: km\\: s^{-1}\\: Mpc^{-1}$.  \\section[]{Models} \\subsection{Dark Matter and Gas Profile}  Using one-dimensional hydrodynamic simulations, we estimate the impact of the shock wave created by an AGN outburst on galaxies. We assume that an AGN is located at the centre of a group of galaxies with the virial mass of $M_\\rmn{vir}$ and that the group is spherically symmetric.  For the mass distribution of the group, we adopt the so-called NFW profile \\citep*{nav97}, although later studies indicated that the central cusp would be steeper \\citep[e.g.][]{fuk07}. The mass profile is written as \\begin{equation} \\label{eq:NFW} M(R) \\propto \\left[\\ln \\left(1+\\frac{R}{R_s}\\right) -\\frac{R}{R_s (1+R/R_s)} \\right]\\:, \\end{equation} where $R_s$ is the characteristic radius of the group.  The normalisation can be given by $M(R_{\\rm vir})=M_{\\rm vir}$, where $R_{\\rm vir}$ is the virial radius of the group.  The virial radius is given by \\begin{equation} R_\\rmn{vir}=\\left[\\frac{3M_\\rmn{vir}} {4\\pi\\Delta_c(z)\\rho_c(z)}\\right]^{1/3}\\:, \\end{equation} where $\\Delta_c(z)$ is a spherical over-density of the virialized dark halo within $R_\\rmn{vir}$ at redshift $z$, in units of the critical density of the Universe at $z$, or $\\rho_c(z)$. For $\\Delta_c(z)$, we use the fitting formula of \\citet{bry98} for a flat Universe with a non-zero cosmological constant, $\\Delta_c(z)=18\\pi^2 +82x -39x^2$, where $x=\\Omega(z)-1$ and $\\Omega(z)$ is the cosmological density parameter at redshift $z$. The concentration parameter of the group, $c_\\rmn{vir}=R_\\rmn{vir}/R_\\rmn{s}$, is given by \\begin{equation} c_\\rmn{vir}=\\frac{9}{1+z} \\left(\\frac{M_\\rmn{vir}} {1.5\\times 10^{13} h^{-1}\\rmn{M_\\odot}}\\right)^{-0.13} \\end{equation} \\citep{bul01}.  Initially, hot gas or intragroup medium (IGM) is in pressure equilibrium with the gravitational potential formed by the group. We assume that the initial IGM density and temperature profiles follow the `universal profile' derived by \\citet{kom01}. They can respectively be written as \\begin{equation} \\rho_\\rmn{IGM}(R)=\\rho_\\rmn{IGM}(0)y_\\rmn{IGM}(R/R_\\rmn{s})\\:, \\end{equation} \\begin{equation} T_\\rmn{IGM}(r)=T_\\rmn{IGM}(0)y_\\rmn{IGM}^{\\gamma'-1}(R/R_\\rmn{s})\\:, \\end{equation} where \\begin{equation} y_\\rmn{IGM}^{\\gamma'-1}(x)=1-\\frac{3}{\\eta_0}\\frac{\\gamma'-1}{\\gamma'} \\frac{c_\\rmn{vir}}{m(c_\\rmn{vir})} \\left[1-\\frac{\\ln(1+x)}{x}\\right]\\:, \\end{equation} \\begin{equation} m(x)=\\ln(1+x)-x/(1+x)\\:. \\end{equation} The parameters $\\eta_0$ and $\\gamma'$ can be derived from the condition that the IGM and dark matter profiles are the same in the outermost region of the group \\citep{kom01}.  We solve the following equations: \\begin{equation} \\frac{\\partial \\rho_\\rmn{IGM}}{\\partial t} + \\frac{1}{R^2}\\frac{\\partial}{\\partial R} (R^2 \\rho_\\rmn{IGM} V_\\rmn{IGM})=0 \\;, \\end{equation} \\begin{eqnarray} \\frac{\\partial (\\rho_\\rmn{IGM} V_\\rmn{IGM})}{\\partial t} +\\frac{1}{R^2}\\frac{\\partial}{\\partial R} (R^2 \\rho_\\rmn{IGM} V_\\rmn{IGM}^2)\\nonumber\\\\ =-\\rho_\\rmn{IGM} \\frac{G M(R)}{R^2}-\\frac{\\partial p}{\\partial R} \\;, \\end{eqnarray} \\begin{eqnarray} \\label{eq:energy2} \\frac{\\partial e}{\\partial t} +\\frac{1}{R^2}\\frac{\\partial}{\\partial R}[R^2 V_\\rmn{IGM}(p+e)] =-n_\\rmn{e}^2 \\Lambda(T_\\rmn{IGM}) \\nonumber\\\\ -\\rho_\\rmn{IGM} V_\\rmn{IGM} G M(R)/R^2 \\;, \\end{eqnarray} where $G$ is the gravitational constant, and $p$ and $V_\\rmn{IGM}$ are the pressure and velocity of the IGM, respectively. The total energy is defined as $e=p/(\\gamma-1)+\\rho_\\rmn{IGM} V_\\rmn{IGM}^2/2$, where $\\gamma=5/3$. Although we include the cooling function $\\Lambda(T_\\rmn{IGM})$, the evolution of the shock is faster than the radiative cooling. The electron density is defined as $n_\\rmn{e}=0.86\\:\\rho_\\rmn{IGM}/m_\\rmn{p}$, where $m_\\rmn{p}$ is the proton mass. We ignore the self-gravity of IGM.    \\subsection{The Criterion for Stripping}  The cold interstellar medium (ISM) in the disc of a galaxy would be stripped when a shock wave excited in the IGM passes the galaxy. Since the shock wave passes the galaxy in a short time (less than the dynamical time or the rotation time of the galaxy), momentum transfer causes stripping to occur. The integrated momentum from the IGM per unit area at radius $R$ is \\begin{equation} \\label{eq:mom} s(t,R)= \\int_0^t \\rho_\\rmn{IGM}(t,R)V_\\rmn{IGM}^2(t,R)dt\\;. \\end{equation} Thus, the criterion of stripping for a galaxy at radius $R$ is \\begin{equation} \\label{eq:criterion} s(t,R)>\\Sigma_\\rmn{ISM}v_\\rmn{esc}\\;, \\end{equation} where $\\Sigma_\\rmn{ISM}$ is the column density of the cold ISM in the galactic disc, and $v_\\rmn{esc}$ is the escape velocity of the galaxy. Strictly speaking, this relation is valid only when the galaxy is face-on; if not, the stripping would be less efficient. In this relation, we assume that the galaxy is not moving relative to the group. The AGN explodes at $t=0$.  The escape velocity is given by $v_\\rmn{esc}\\sim \\sqrt{2}v_\\rmn{rot}$, where $v_\\rmn{rot}$ is the rotation velocity of the galaxy. (The ISM may stay in the halo of the galaxy with this velocity.) \\citet*{mo98} indicated that for a given $v_\\rmn{rot}$ the total disc surface density of a galaxy, $\\Sigma_\\star$, is proportional to the Hubble constant at redshift $z$: \\begin{equation} H(z)=H_0[\\lambda_0+(1-\\lambda_0-\\Omega_0)(1+z)^2 + \\Omega_0(1+z)^3]^{1/2} \\end{equation} Following \\citet{fuj04}, we assume that the ISM column density is proportional to the disc surface density ($\\Sigma_\\rmn{ISM}\\propto \\Sigma_\\star$). Thus, the former is given by \\begin{equation} \\label{eq:Sig_ISM} \\Sigma_\\rmn{ISM}(z)=\\Sigma_\\rmn{ISM}(0)H(z)/H_0 \\:. \\end{equation} Moreover, we assume that the disc radius of a galaxy has a relation of \\begin{equation} \\label{eq:r_gal} r_\\rmn{gal}(z)=r_\\rmn{gal}(0)[H(z)/H_0]^{-1}\\: \\end{equation} \\citep{mo98}. ",
        "conclusions": "We have investigated the influence of an AGN outburst at the centre of a galaxy group on the galaxies surrounding the AGN. If the energy of the AGN is extremely large ($E_{\\rm AGN}\\sim 6\\times 10^{61}\\rm\\; erg$) as the one recently observed in clusters, the shock wave excited by the outburst strips the cold ISM in the disc of the surrounding galaxies locating at $R\\sim 100$--200~kpc from the AGN. The effect of the stripping can be much stronger than that of usual ram-pressure stripping owing to the motion of the galaxies in a group. However, the stripping of the cold ISM in the disc is not effective if $E_{\\rm AGN}\\sim 6\\times 10^{60}\\rm\\; erg$.  We also showed that even if the cold ISM in the disc of the galaxies is not stripped, the warm gas in the halo would be stripped by the shock wave. This could be effective for the whole galaxies in the group up to the virial radius even when $E_{\\rm AGN}\\sim 6\\times 10^{60}\\rm\\; erg$. The star formation activities of those galaxies would decrease gradually because the ISM consumed by the star formation is not compensated from the gas in the halo (strangulation).  Our results indicate that even one strong AGN outburst significantly affects the evolution of galaxies in groups. After the stripping, the galaxies should become red because of the decrease of the star formation rate.  The success of this model depends on the assumption of an extremely strong outburst. Since such outbursts have not been observed in galaxy groups at low-redshift, the stripping by them would not have strong influence on the evolution of galaxies at low-redshift. However, we suppose that it would be at high redshift where extremely strong AGN activities have often been observed."
    },
    "0711/0711.1984_arXiv.txt": {
        "abstract": "Four eclipsing binaries which show apparent changes of period have been studied in case of possible presence of the light-time effect. With the least squares method we calculated new light elements of these systems and mass function of predicted third body and its minimum mass. We discussed probability of a presence of such bodies in case of mass function, changes in radial velocity and third light in solution of light curves, respectively. ",
        "introduction": "Light-time effect (hereafter LITE) in eclipsing binaries, caused by the orbital motion of the eclipsing pair around the barycenter of the triple system, produces period variation in minima timings. It was explained by Irwin in 1959 and its necessary criteria have been mentioned by Friebos-Conde \\& Hertzeg in 1973. They are following: agreement with a theoretical LITE curve, secondary minima behave identically with primary ones, reasonable value of the mass function and corresponding variations in radial velocity, respectively.  We found, that most of the necessary criteria listed above are satisfied in all systems presented below. Their $O\\!-\\!C$ diagrams are also presented. The curves represent computed LITE versus epoch. The individual primary and secondary minima are denoted by dots and circles, respectively. Larger symbols correspond to the photoelectric or CCD measurements which were given higher weights in our computations. All data were found in published papers.  \\begin{figure}[t!] \\begin{center} \\scalebox{0.68}{\\includegraphics{ARAur1.eps}} \\caption[]{The O-C diagram of AR Aur.} \\label{ARAur} \\end{center} \\end{figure} ",
        "conclusions": "We have derived new LITE parameters for four eclipsing binaries by their $O\\!-\\!C$ diagram analysis. But for the final decision of the presence of LITE we need further accurate timings of minima or appropriate confirmation by another method (radial velocity variations, third light in a light curve analysis or astrometric observations)."
    },
    "0711/0711.4045_arXiv.txt": {
        "abstract": "We describe an extensive observational project that has obtained high-quality and homogeneous photometry for a number of different Galactic star clusters (including M$\\,$92, M$\\,$13, M$\\,$3, M$\\,$71, and NGC$\\,$6791) spanning a wide range in metallicity ($-2.3\\lesssim\\FeH\\lesssim+0.4$), as observed in the $\\ugrizprime$ passbands with the MegaCam wide-field imager on the Canada-France-Hawaii Telescope.  By employing these purest of stellar populations, fiducial sequences have been defined from color-magnitude diagrams that extend from the tip of the red-giant branch down to approximately 4 magnitudes below the turnoff: these sequences have been accurately calibrated to the standard $\\ugrizprime$ system via a set of secondary photometric standards located within these same clusters.  Consequently, they can serve as a valuable set of empirical fiducials for the interpretation of stellar populations data in the $\\ugrizprime$ system. ",
        "introduction": "Recently, the Sloan Digital Sky Survey (SDSS) officially ended its planned five years of sky scanning operations to obtain an unprecedented amount of imaging and spectroscopic data for approximately one-quarter of the sky.  The SDSS was carried out on a dedicated 2.5$\\,$m telescope equipped with a large-format mosaic CCD to image the entire northern Galactic cap (i.e., $b>30^{\\circ}$) in five photometric bands and two digital spectrographs to provide spectra for $\\sim1$ million stars, galaxies, and quasars scattered throughout the imaging area.  Although designed to primarily investigate the large-scale structure of the universe, the imaging component of the SDSS has obtained high-quality multicolor photometry for about $10^8$ stellar objects in the Milky Way, which represents the largest and most homogeneous database on Galactic stellar populations ever obtained.  A notable feature of this database is that it was compiled in a new photometric system consisting of five unique passbands ($\\uprime$, $\\gprime$, $\\rprime$, $\\iprime$, and $\\zprime$) that were specifically designed for the SDSS to provide continuous coverage over the $entire$ optical wavelength range \\citep{Fukugita1996}.  While the SDSS was the first to implement the standard $\\ugrizprime$ photometric system, analogous versions of these same filters are also currently in use with CCD imagers installed on the Gemini Telescopes, the Canada-France-Hawaii Telescope, and the Hubble Space Telescope.  In addition, the very fact that the SDSS has already provided such a large database of photometry of stars and galaxies implies that this photometric system is also becoming widely accepted as the filter set of choice for many planned ground-based observational projects and large-scale sky surveys (e.g., LSST, OmegaCam, Pan-STARRS, VST).  Despite this, much of our current observational and theoretical knowledge of resolved stellar populations is based largely on the conventional Johnson-Kron-Cousins $UBV(RI)_C$ photometric system, with a few other studies relying on niche systems like the Str\\\"{o}mgren, DDO, Vilnius, and Geneva systems.  Consequently, the empirical and theoretical tools that would tie the standard $\\ugrizprime$ system to the fundamental properties of observed stellar populations have yet to be defined.  Specifically, neither well-calibrated fiducial stellar population sequences nor reliable color--$\\Teff$ relations are currently available, and yet, without them, it is impossible to fully exploit the capabilities of the SDSS data set as well as complementary studies employing these same filters.  In order to remedy these deficiencies, there is good reason to rely on $\\ugrizprime$ observations of star clusters within our own Galaxy.  Clusters are the ideal stellar population templates because, despite a handful of exceptions (e.g., $\\omega\\,$Cen and M$\\,$22), their constituents are believed to be effectively coeval, equi-distant, and nearly identical in terms of their heavy elemental abundances.  As a result, their color-magnitude diagrams (CMDs) generally exhibit extremely tight and well-populated sequences of stars that span several orders of magnitude in brightness.  Their wide distribution in metallicity is also suitable for characterizing how the photometric properties of stellar populations vary as a function of $\\FeH$.  Consequently, cluster observations offer the perfect data sets to define fiducial stellar population sequences that cover a broad range of stellar parameter space.  These sequences serve as a set of empirical ``isochrones\" that not only facilitate the analysis of other stellar populations data, but also provide calibrators for stellar evolutionary models that are transformed to the observed CMDs via theoretically-derived color-$\\Teff$ relations (e.g., see \\citealt{Brown2005}).  Given the fact that the standard $\\ugrizprime$ system was introduced only a short time ago, however, the photometric database for star clusters remains too small to accomplish the tasks mentioned above.  Unfortunately, the SDSS alone cannot provide a sufficient database since (1) the main Sloan survey with the 2.5$\\,$m telescope is based on instrumental $ugriz$ passbands that are very similar to, but not quite identical to, the $\\ugrizprime$ passbands with which the $standard$ Sloan photometric system was defined on the 1.0$\\,$m telescope of the USNO Flagstaff Station \\citep{S02}.\\footnote{\\citet{Tucker2006} have published simple linear transformation equations that relate the instrumental $ugriz$ magnitudes to standard $\\ugrizprime$ with reasonable precision (2--3\\%).  At some time in the future more sophisticated transformations may be derived that will enable users to predict standard-system indices from the main-survey observations with a high level of accuracy, at least for normal stellar spectral-energy distributions.  Our own instrumental observations from the CFHT have been similarly transformed to the standard $\\ugrizprime$ system rather than to the $ugriz$ main-survey system; see below, and \\citet{Paper1}.}  In addition, (2) the brightest stars lying on the red-giant branches (RGBs) of the nearest clusters saturate during $\\sim60\\,$s drift-scan exposure times, (3) the SDSS imaging data do not extend deep enough to provide good photometric precision for some of the fainter stars lying on the main sequences in the more distant metal-poor globular clusters, (4) the survey footprint does not reach down to the Galactic plane where the majority of metal-rich open clusters reside.  These reasons imply that the SDSS may not be the most ideal source of star cluster photometry for the derivation of fiducial stellar population sequences.  Hence, we recently began an extensive observational project aimed at thoroughly exploring the nature of stellar populations in the $\\ugrizprime$ photometric system via observations of Galactic star clusters.  Our first paper \\citep[][hereafter Paper I]{Paper1} presented a network of fainter secondary standard stars for the $\\ugrizprime$ system in selected star cluster fields to better aid observers on large-aperture telescopes in calibrating their photometry.  In fact, the current investigation relies on these standards to calibrate a sample of high-quality and homogeneous $\\ugrizprime$ photometry obtained on the 3.6$\\,$m Canada-France-Hawaii Telescope (CFHT).  This photometry is subsequently employed here to derive a set of accurate stellar population fiducial sequences that span a broad range in both magnitude and metallicity.  In our third and final paper (Clem et al. 2008, in preparation) we will utilize these fiducials to test a new grid of theoretical color--$\\Teff$ relations and bolometric corrections for the $\\ugrizprime$ system that have been calculated from synthetic spectra.  The following sections present the details related to the observation, reduction, and compilation of the photometry collected at the CFHT that will be used to derive a set of fiducial stellar population sequences for the $\\ugrizprime$ system.  In Section 2 we describe the observational setup employed at the CFHT to collect the cluster photometry as well as the data reduction procedure, including the important step of calibrating the observed cluster photometry to the standard $\\ugrizprime$ system.  Section 3 presents the details of the fiducial sequence derivation process.  Finally, a short summary of our results, and a discussion of the usefulness of these fiducials for stellar populations research is given in Section 4. ",
        "conclusions": "Using high-quality, homogeneous observations obtained with the wide-field MegaCam imager on the 3.6m Canada-France-Hawaii Telescope, we have derived fiducial stellar population sequences for the Galactic star clusters M$\\,$92, M$\\,$13, M$\\,$3, M$\\,$71, and NGC$\\,$6791 in the $\\ugrizprime$ photometric system.  These sequences, which span a wide range in both metallicity and magnitude, have been accurately calibrated to within 1\\% of the standard $\\ugrizprime$ system using a set of secondary standard stars derived in Paper I.  As a result of our efforts, we anticipate that these fiducial sequences will serve as valuable tools for the interpretation of other stellar population investigations involving the $\\ugrizprime$ bandpasses by virtue of the fact that they represent a set of empirical isochrones for both metal-poor and metal-rich stars having wide-ranging physical parameters.  Indeed, a preliminary set of the fiducials presented in this work has already been employed in the interpretation of CMDs for a number of newly discovered Milky Way satellites (see, for example, \\citealt{Belokurov2006, Belokurov2007a, Belokurov2007b}).  In addition to the usefulness of these fiducial sequences for the interpretation of observed data, they also provide an excellent test of the accuracy of color-temperature relations and bolometric corrections that have been derived from model atmospheres and synthetic spectra.  In a future paper (J.~L. Clem et al. 2008, in preparation) we intend to perform tests of such synthetic color and magnitude transformations by comparing isochrones models to the fiducials derived here.  Our aim is to assess the quality of the color-temperature relations and bolometric corrections in providing isochrones that can reproduce the observed cluster photometry when reasonable estimates of the cluster reddening, metallicity, and distance are assumed, as well as test their consistency when isochrone fits to the $\\ugrizprime$ photometry are compared to those in other photometric systems (e.g., BV(RI)$_C$ and $uvby$).  It is important for the reader to note, however, that the fiducials derived in this investigation are presented on the $\\ugrizprime$ system and $not$ on the natural photometric system of the 2.5$\\,$m SDSS survey telescope (i.e., the $ugriz$ system).  Subtle differences exist between the two systems such that photometry reported on both systems for an identical stellar sample can differ systematically by as much as a few hundredths of a magnitude in some filters.  Therefore, we caution against using these fiducials to interpret and/or analyze $ugriz$ photometry from the SDSS without first applying appropriate transformation relations (see \\citealt{Tucker2006}).  Although these transformations may not be appropriate for stars with strong emission features in their spectra or stars with extreme colors (i.e., later than M0 spectral class), we expect they are good enough to transform the fiducials presented here on the $\\ugrizprime$ system to the SDSS 2.5$\\,$m $ugriz$ system, while keeping the uncertainties in the photometric zero points on the AB system \\citep[see][]{OkeGunn1983} to within a few percent."
    },
    "0711/0711.3179_arXiv.txt": {
        "abstract": "{} {We analyse $V$-filter ASAS-3 photometry of 41 known $\\beta$~Cephei-type stars. The ASAS-3 photometry is combined with the archival data, if available, to determine long-term stability of periods and amplitudes of excited modes.} {Frequencies of modes are derived by means of Fourier periodograms with consecutive prewhitening.  The results are examined in the context of detection threshold.} {We detected amplitude changes in three $\\beta$~Cephei stars, BW Cru, V836 Cen, and V348 Nor.  Period changes were found in KK~Vel and V836 Cen. Our analysis shows that intrinsic period changes are more common among multiperiodic stars, apparently because they are caused by some kind of mode interaction. In addition, we found new modes for seven stars, and for ten others we provide new solutions or remove ambiguities in the detected frequencies. One candidate hybrid $\\beta$~Cephei/SPB star, HD\\,133823, is discovered.} {} ",
        "introduction": "A group of main-sequence early B-type pulsating stars, known as $\\beta$~Cephei stars, has been studied for over a hundred years. The discovery of the mechanism driving pulsations in these stars \\citep{modz92, dzpa93, pamy99} advanced our understanding of their nature and opened the possibility of addressing more complex questions concerning their pulsations.  In particular, studies by means of asteroseismology became possible.  The attractiveness of $\\beta$~Cephei stars for asteroseismology results from the fact that -- as pulsators -- they are neither too simple nor too complex.  It is also one of the first groups of pulsating stars in which nonradial pulsations were studied in detail, as nonradial pulsations are for them the rule than the exception.  Moreover, ty\\-pically less than a dozen modes are observed in a single star which seems to be a sufficient number in seismic modeling.  The properties of $\\beta$~Cephei stars have recently been summarised in a review paper of \\citet{stha05}. Among other topics, these authors discussed the boundaries of the instability strip of $\\beta$~Cephei stars and pulsations in O-type stars. There are many more interesting, yet still not well understood problems related to $\\beta$~Cephei stars, for example, their relation to other variable stars in the upper main-sequence, influence of fast rotation on pulsations, dependence on metallicity. One such problem, the long-term stability of periods and amplitudes in these stars, is addressed in the present paper.  A reliable study of long-term changes of periods and/or amplitudes requires long-time data with a good coverage. There are many $\\beta$~Cephei stars that show stable pulsation(s) on a time scale of decades. Only a handful are known for their variation of period and/or amplitude \\citep{jepi98,jerz99}. Sometimes, the variation is fast enough to be detected using only observations from two consecutive seasons. As possible causes of the period variation include evolutionary effects and mode interaction, these changes merit being studied.  An all-sky survey covering several years of observations offers an excellent opportunity to detect period and amplitude changes, especially if the data can be combined with the archival ones. In the present paper we take advantage of such an opportunity by analysing the photometry obtained within the ASAS-3 survey for all known $\\beta$~Cephei stars in the magnitude range it covered. ",
        "conclusions": "It is not well known how common amplitude and period changes are among $\\beta$~Cephei stars.  The main reason for this is the lack of homogeneous observations covering long time intervals. From the analysis of observations for stars with the longest observational records it seems, however, that, if detectable, changes of periods and amplitudes of modes excited in $\\beta$~Cephei stars have a typical time scale longer than a few months. In many cases it is much longer than a decade as in the well-documented case of 16 (EN) Lac \\citep{jepi96, jepi99}.  \\subsection{Changes of periods} The evidence for long-term period variations among $\\beta$~Cephei stars was summarised a few years ago by \\citet{jepi98} and \\citet{jerz99}. One of the main reasons for the long-term monitoring of periods in $\\beta$~Cephei stars is the hope that evolutionary changes in these stars will be detected and confronted with the predictions of the theory.  According to the theory \\citep{egpe73, leai74}, in the core-hydrogen burning stage of evolution, where most $\\beta$~Cephei stars are believed to occur, the rate of period change, $\\dot{P}$, should be smaller than 0.3~seconds per century and positive. Although this value is small, it is detectable provided that observations cover at least a decade. $\\dot{P} >$ 0 arises as a consequence of increase of stellar radius during the evolution on the main sequence.  Stars with small positive $\\dot{P}$, consistent with theoretical predictions for core-hydrogen burning phase, are indeed observed \\citep{jerz99}, although there are cases (BW Vul, $\\sigma$ Sco) where $\\dot{P}$ is much larger than 0.3~s\\,cen$^{-1}$ which was used as an argument in favour of the hypothesis that they had already evolved off the main sequence \\citep{pigu92,pigu93}. In this group, there was also $\\delta$~Cet. For this star, $\\dot{P}$ was equal to 0.47 $\\pm$ 0.09~s\\,cen$^{-1}$ \\citep{jerz88}, slightly too large for the main-sequence evolution \\citep[see Fig.~1 of ][]{jerz99}. In view of the new observations, however, it is clear that for this star the change of period can no longer be interpreted in terms of a constant $\\dot{P}$ \\citep{jerz07}.  There are, however, $\\beta$~Cephei stars in which the period behaviour was much more complex. \\citet{jerz99} listed three such stars: $\\beta$~CMa, 12 (DD) Lac, and 16 (EN) Lac. Indeed, the evolutionary changes of period may contribute to the observed period changes in these stars only partly because the changes are different for different modes. Moreover, for some modes the period was found subsequently to increase and decrease. The pattern of period changes for modes excited in V836~Cen (Fig.~\\ref{v836cen-aoc}) resembles that observed in the three above-mentioned stars, especially 12 Lac \\citep{pigu94}. Thus, V836\\,Cen may be regarded as the fourth member of this group of $\\beta$~Cephei stars with complex period behaviour. It will be shown in Paper II that one of the new $\\beta$~Cephei stars, HD\\,168050, also exhibits a complex period behaviour. All five stars from this group are multiperiodic. It seems therefore that when many modes are excited, some kind of mode interaction causes period and, as we shall comment below, amplitude changes.  In this context it is worth noting that in the study of period variations we assume that pulsations are coherent. In a general case, this might not be true. If random phase shifts are generated in a star, e.g., by the presence of other mode(s), period variations similar to those seen in 12 Lac and V836\\,Cen can be observed. For such stars we can expect to detect evolutionary period changes only on a time scale which would average random changes. For $\\beta$~Cephei stars this probably means centuries.  Period variations in $\\beta$~Cephei stars might be additionally complicated by the presence of a star in a wide binary system. In such a case, additional contribution to the apparent period changes comes from the light-time effect. This is indeed observed in $\\beta$~Cep \\citep{pibo92}, $\\sigma$~Sco \\citep{pigu92} and possibly BW Vul \\citep{pigu93}. As KK Vel is known to have a close visual companion at separation of about 0$\\farcs$3, the large changes of period seen in Fig.~\\ref{kkvel-pc}, apparently in both directions, can be mostly due to the light-time effect.  \\subsection{Changes of amplitudes} Changes of amplitudes were reported for only a handful of $\\beta$~Cephei stars. \\citet{jerz99} listed five stars with long-term amplitude variations: $\\beta$~CMa, 27 (EW)~CMa, V381~Car, $\\alpha$~Vir A, and 16~Lac. Evidence for amplitude change was also found for $\\nu$~Cen \\citep{aspa92,scte02} and for one mode in 12 Lac \\citep{pigu94}. As modes in V836\\,Cen also show amplitude changes (Fig.~\\ref{v836cen-aoc}), we conclude that all stars with complex period behaviour also show amplitude changes. The outstanding examples from this group are the primary of Spica ($\\alpha$~Vir) in which pulsation amplitudes dropped to undetectable level or ceased altogether \\citep{lomb78,ster86} and 16 Lac which was monitored photometrically many times since the time of discovery. The changes of amplitudes and periods of its four dominant modes are therefore known very well \\citep{jepi96,jepi99}. Apart from V836~Cen, we report amplitude changes for BW~Cru and V348~Nor. Note that both BW~Cru and V348~Nor are not known as multiperiodic.  In this paper, combining ASAS-3 photometry with the ar\\-chi\\-val data, we found evidence for period and amplitude changes in four known $\\beta$~Cephei stars. In addition, in Paper II, complex period changes in HD\\,168050 were found using only ASAS-3 data. This shows that period and amplitude changes might be quite common in $\\beta$~Cephei stars but, obviously, can be detected only if the data cover sufficient time intervals, typically more than a decade. Fortunately, ASAS is continuing and will be soon extended into the northern hemisphere. In addition, other photometric surveys are expected to provide abundant time-series data. In view of the fact that over 200 new $\\beta$~Cephei stars were detected in the ASAS-3 data, the sample of stars in which we will be able to monitor long-term changes of amplitudes and periods will increase considerably over the next decades."
    },
    "0711/0711.4790_arXiv.txt": {
        "abstract": "Millihertz quasi-periodic oscillations reported in three neutron-star low mass X-ray binaries have been suggested to be a mode of marginally stable nuclear burning on the neutron star surface. In this Letter, we show that close to the transition between the island and the banana state, 4U~1636--53 shows mHz QPOs whose frequency systematically decreases with time until the oscillations disappear and a Type I X-ray burst occurs. There is a strong correlation between the QPO frequency $\\nu$ and the occurrence of X-ray bursts: when $\\nu\\gtrsim9$~mHz no bursts occur, while $\\nu\\lesssim9$~mHz does allow the occurrence of bursts. The mHz QPO frequency constitutes the first identified observable that can be used to predict the occurrence of X-ray bursts.  If a systematic frequency drift occurs, then a burst happens within a few kilo-seconds after $\\nu$ drops below 9~mHz. This observational result confirms that the mHz QPO phenomenon is intimately related with the processes that lead to a thermonuclear burst. ",
        "introduction": "\\label{sec:intro}    \\citet{Revnivtsev01} discovered a new class of low frequency quasi-periodic oscillation (QPO) in three neutron star X-ray binary sources (Aql X-1, 4U~1608--52 and 4U~1636--53). This new QPO has frequencies between 7 and 9 $\\times 10^{-3}$ Hz, i.e. it is in the milli-Hertz range, and its other properties also differ from those of the other QPOs found in the neutron star systems \\citep[e.g. energy dependence, see][]{Vanderklis06}. Although \\citet{Revnivtsev01} could not discard an interpretation related with  disk instabilities, they  concluded that the mHz QPO is likely due to a special mode of nuclear burning on the neutron star surface. This interpretation was strengthened by the results of \\citet{Yu02}, who showed that the kHz QPO frequency is anti-correlated with the luminosity variations during the mHz oscillation, suggesting that the inner edge of the disk slightly moves outward as the luminosity increases during each mHz cycle due to stresses generated by radiation coming from the the neutron star surface. This is contrary to the correlation observed between X-ray luminosity ($L_x$) and kHz QPO frequency, where the inner disk edge is thought to move in, as the accretion rate and hence $L_x$, increases \\citep[][and references within]{Vanderklis06}.  The properties of the mHz QPOs as observed up to now can be summarized as follows: \\textit{(1)} the fractional rms amplitude strongly decreases with energy, from $\\approx2$\\% at 2.5 keV, down to an almost undetectable $ < 0.2$\\% at $\\approx5$ keV; \\textit{(2)} mHz QPOs occur only in a particular range of  X-ray luminosity: $L_{2-20~keV} \\approx 5-11 \\times 10^{36}$ erg/s; \\textit{(3)} the frequency of the mHz QPOs is  between 7 and 9 mHz; \\textit{(4)} the mHz QPOs disappear with the occurrence of a type I X-ray burst; \\textit{(5)} as noted above, the kHz QPO frequency is approximately anti-correlated with the $2-5$ keV count rate variations that constitute the mHz oscillation.    \\citet{Heger07} suggested that the mHz QPOs could be explained as the consequence of marginally stable nuclear burning on the neutron star surface.  They found an oscillatory mode of burning, with a period $P_{osc}$ close to the geometric mean of the thermal\\footnote{The thermal timescale is defined as $t_{thermal}=c_p T / \\epsilon$ where $c_p$, $T$ and $\\epsilon$ are the heat capacity at constant pressure , the temperature and the nuclear energy generation rate, respectively} and accretion\\footnote{The accretion timescale is defined as $t_{accr}=y /\\dot{m}$ where $y$ and $\\dot{m}$ are the column depth of the burning layer and local accretion rate, respectively.} timescales of the burning layer. For typical parameters, $P_{osc}\\equiv\\sqrt{t_{thermal} \\cdot t_{accr}}\\approx2$ minutes, in accordance with the characteristic frequency of the mHz QPOs.  The burning is oscillatory only close to the boundary between stable burning and unstable burning (in Type I X-ray bursts), explaining the observation that the mHz QPOs were seen within a narrow range of luminosities.   Two of the three sources in which \\citet{Revnivtsev01} found the mHz QPOs are transient atoll sources (Aql X-1 and 4U~1608--52) while the third one, 4U~1636--53, is a persistent atoll source \\citep{Hasinger89}. The object of our current study, 4U~1636--53, has an orbital period of $\\approx3.8$ hours \\citep{Vanparadijs90} and a companion star with a mass of $\\approx0.4 \\ M_{\\odot}$ \\citep[assuming a NS of 1.4~$M_{\\odot}$; ][]{Giles02}.  4U~1636--53 is an X-ray burst source \\citep{Hoffman77} showing asymptotic burst oscillation frequencies of $\\approx581$~Hz \\citep[][]{Zhang97,Strohmayer02}.  The aperiodic timing behavior of 4U~1636--53 has been studied with the EXOSAT Medium Energy instrument \\citep[][]{Prins97} and with the Rossi X-ray Timing Explorer  \\citep[RXTE, e.g.][]{Wijnands97,Disalvo03,Altamirano07}.    \\begin{figure}[t] \\center \\resizebox{1\\columnwidth}{!}{\\rotatebox{0}{\\includegraphics{./f1.eps}}} \\caption{Color-color diagram as \\citet{Altamirano07}. Each data point represents the average of an observation ($\\approx2$ to $\\approx30$~ksec). The ellipse marks the region in which mHz QPOs with decreasing frequency were found. The labels A, B and C correspond to those in Figure~\\ref{fig:lc}. } \\label{fig:ccd} \\end{figure}  \\begin{figure}[t] \\center \\resizebox{1\\columnwidth}{!}{\\rotatebox{0}{\\includegraphics{./f2.eps}}} \\caption{Dynamical power spectrum smoothed with a 750 seconds sliding window with steps of 200 seconds, showing the mHz QPOs during the last 12~ksec before the X-ray burst occurs. This sequence corresponds to case B in Figures~\\ref{fig:lc}~\\&~\\ref{fig:ccd} -- ObsId: 60032-05-02-00. The three black vertical lines correspond to the times of occurrence of the X-ray bursts. For clarity, we plot only powers above 10 which correspond to $\\gtrsim2\\sigma$ (single trial per 750 seconds window but normalized to the number of possible frequencies in the range 0--0.5~Hz).} \\label{fig:dynspec} \\end{figure}   \\begin{figure*}[t] \\center \\resizebox{2\\columnwidth}{!}{\\rotatebox{0}{\\includegraphics{./f3.eps}}} \\caption{2--60~keV PCA light curve of observations 60032-05-01,02,03 and 04. $A$, $B$ and $C$ indicate intervals in which the mHz QPOs are detected. In all three cases, the frequency of the mHz QPO decreases with time until the appearance of an X-ray burst. } \\label{fig:lc} \\end{figure*}    4U~1636-53 is a reference source for studying nuclear burning on the surface of a neutron star since it shows the full range of burst behavior: single and multi-peaked Type I X-ray bursts, superbursts, burst oscillations, photospheric radius expansion, regular and irregular burst sequences \\citep[e.g.][]{Galloway06} and mHz QPOs. As such, it is an ideal source to understand the relation between these different observational manifestations of nuclear burning.  Recently, \\citet{Shih05} reported that 4U~1636--53 has shown a significant decrease in its persistent $L_x$ during the years 2000 and 2001. \\citet{Altamirano07} show that during the low $L_x$ period, 4U~1636--53 is observed in its (hard) island states. This provides an opportunity to study the mHz QPOs in harder and lower luminosity states than was possible up to now. ",
        "conclusions": "\\label{sec:discussion}    We have shown that close to the transition between the island and the banana state 4U~1636--53 exhibits mHz QPOs whose frequency systematically decreases with time until the oscillations disappear with the occurrence of a Type I X-ray burst. The mHz QPO frequency $\\nu$ constitutes the first identified observable that can be used to predict the occurrence of X-ray bursts: when $\\nu\\gtrsim9$~mHz no bursts occur, while $\\nu\\lesssim9$~mHz does allow the occurrence of bursts. If a systematic frequency drift occurs, then a burst happens within a few kilo-seconds after $\\nu$ drops below 9~mHz. This observational result confirms that the mHz QPO phenomenon is intimately related with the processes that lead to a thermonuclear burst.  The fact that the observation of a systematic frequency decrease with time implies the occurrence of a future X-ray burst, strongly suggests that the frequency of the mHz QPOs is related to the burning processes on the neutron star surface. One possibility is that the frequency of the QPO is somehow a measurement of the accumulation of fresh fuel on the neutron star surface which will be available for a future thermonuclear burst. To our knowledge, there has been only one attempt to theoretically explain the mHz QPOs phenomena \\citep{Heger07}. In this model the frequency of the QPO depends, among others, on the amount of available fresh fuel, on the local accretion rate and the composition of the material. It is beyond the scope of this Letter to perform numerical simulations as those reported by \\citet{Heger07}. In the rest of this discussion we briefly compare these authors' model predictions with our observations and propose some more complex scenarios.  Analytical and numerical results based on the simplified one-zone model of \\citet{Paczynski83} in the \\citet{Heger07} marginally stable burning model (see Section~\\ref{sec:intro}) predict that (i) close to the boundary between stable and unstable burning, the NS surface will show temperature fluctuations with constant frequency $\\nu$ if the local accretion rate $\\dot{m}$ remains constant; (ii) this marginally stable burning regime will occur at $\\dot{m}$ near Eddington, hence accretion must be confined to a surface area $S_A$ that is much smaller that the total area of the NS; (iii)  $\\nu$ correlates with $\\dot{m}$ \\citep[see figure 4 in ][]{Heger07} and (iv) thermonuclear bursts and mHz QPOs should not be observed at the same luminosity and therefore presumably at the same $\\dot{m}$.  In this paper, we show that for constant luminosity the QPO frequency can systematically decrease in time and that instantaneously measured frequencies can be the same for different luminosities. We also show that mHz QPOs and thermonuclear bursts do in fact occur at the same luminosity and that both phenomena are clearly related. This means that we are dealing with a more complex scenario than that introduced by \\citet{Heger07}.  The amount of time between the preceding X-ray burst and the onset of mHz QPOs is variable ($>6$ksec) and apparently independent of source state. If the system is locally accreting at $\\dot{m}\\simeq\\dot{m}_{Edd}$ and if none of the accreting fuel is burnt, only $\\approx1000$~seconds are required to accrete a fuel layer of column depth $y_f$ capable of undergoing marginally stable burning \\citep[$y_f\\approx10^8$~g~cm$^{-2}$ and $\\dot{m}\\approx8\\cdot10^4$ g~cm$^{-2}$~s$^{-1}$ -- see e.g.][]{Heger07}. One possible explanation for the observed longer intervals between burst and onset of oscillations, is that a large fraction of the accreted fuel is burnt as it is accreted on the neutron star surface. Of course the burning fraction could vary in time, and this estimate is assuming that all the fuel was burnt during the last X-ray burst, which is not always true \\citep{Bildsten98}. Interestingly, if this interpretation is correct and low partial burning fractions can occur, under certain conditions the mHz QPOs could appear in much less than a 1000 seconds after an X-ray burst.   The fact that the amount of time between the preceding X-ray burst and the onset of mHz QPOs is variable may be also an indication that not all the accreted fuel is burnt nor available to participate in the marginally stable burning. For example, accretion could occur onto an equatorial region occupying less than 10\\% of the surface area of the star \\citep{Heger07}. A possibility is that part of the fresh fuel burns stably at a rate $B(t)$ per unit area while the other part leaks away from this region at a rate $R(t)$. While the material accumulated at a rate $R(t)$ would serve as fuel for a thermonuclear burst, marginally stably burning of the matter on the equatorial belt is (in principle) still possible. Although such scenario cannot explain the frequency drifts we observe, it can explain why mHz QPO and X-ray bursts do occur at the same $\\dot m$. If mHz QPOs can only occur at a certain local accretion rate $\\dot{m}$ ($\\simeq\\dot{m}_{Edd}$), a small change in effective local accretion rate will lead to an absence of mHz QPOs. This might explain why the mHz QPOs are not always present between X-ray bursts.  Another possibility (which is not taken into account in \\citet{Heger07}'s model) is that there is a significant heat flux from deeper in the star that heats the region undergoing marginally stable burning. For example, changes in heat flux due to energy that is first conducted into deeper layers during an X-ray burst and then slowly outwards towards the surface might be possible. Such a change in the heat flux could affect the conditions of the burning layer (e.g. temperature or burning rate $B(t)$) and therefore could affect the characteristics of the burning processes on the neutron star surface.     Other aspects of the observations offer further challenges for theoretical models that explain burning processes on the neutron star surface as well those which explain atoll sources states. In particular, (i) why the systematic frequency drifts are observed close to the transition between the island and the banana state while the frequencies are approximately constant in the banana state. This may be another indication that the disk geometry of the system is changing during the state transition \\citep[see e.g. ][]{Gierlinski02}; (ii) why in the transition between island and banana state the oscillations disappear \\textit{only} when an X-ray burst occurs, while in the banana state they can also disappear without an X-ray burst (see Section~\\ref{sec:dataanalysis}). Clearly, further theoretical work is needed. More observational work on the interactions between mHz QPOs and X-ray bursts is in progress and will provide further clues for theoretical models.        { \\footnotesize { \\hspace{0.03cm} \\textbf{Acknowledgments:} DA wants to thank A. Patruno, P. Casella, P. Uttley, M. Linares for very helpful discussions. This work was supported by the ``Nederlandse Onderzoekschool Voor Astronomie'' (NOVA), i.e., the ``Netherlands Research School for Astronomy'', and it has made use of data obtained through the High Energy Astrophysics Science Archive Research Center Online Service, provided by the NASA/Goddard Space Flight Center. AC is grateful for support from NSERC, Le Fonds Qu\\'eb\\'ecois de la Recherche sur la Nature et les Technologies, the Canadian Institute for Advanced Research, and as an Alfred P.~Sloan Research Fellow.}  }"
    },
    "0711/0711.0270_arXiv.txt": {
        "abstract": "We study how the two-point density correlation properties of a point particle distribution are modified when each particle is divided, by a stochastic process, into an equal number of identical ``daughter'' particles. We consider generically that there may be non-trivial correlations in the displacement fields describing the positions of the different daughters of the same ``mother'' particle, and then treat separately the cases in which there are, or are not, correlations also between the displacements of daughters belonging to different mothers. For both cases exact formulae are derived relating the structure factor (power spectrum) of the daughter distribution to that of the mother. These results can be considered as a generalization of the analogous equations obtained in \\cite{displa} for the case of stochastic displacement fields applied to particle distributions. An application of the present results is that they give explicit algorithms for generating, starting from regular lattice arrays, stochastic particle distributions with an arbitrarily high degree of large-scale uniformity. ",
        "introduction": "\\label{sec1} Point processes, i.e., stochastic spatial distributions of identical point particles, provide a very useful mathematical scheme for many different $N-$body and granular physical systems such as crystals (both regular, and perturbed) \\cite{ziman,ashcroft,groma}, quasi-crystals \\cite{radin}, structural glasses, fluids \\cite{hansen}, self-gravitating systems in astrophysics and cosmology (see, e.g., \\cite{millionbody, pee80, libro}). They find also many applications in a wide range of scientific fields: computer image  problems \\cite{cpu}, and bio-metrical studies \\cite{renshaw} are only some examples of systems which are usually represented as specific point processes with appropriate spatial correlation properties. The extension of knowledge about this class of stochastic processes can therefore be of fundamental importance for the description and analysis of many scientific topics.  Indeed this is the reason why considerable mathematical effort has already been invested in the study of this class of systems, and many useful results have been derived (e.g., see \\cite{daley,kertscher,torquato98}).  In this paper we present the equations for the structure factor (or power spectrum) of a point process obtained as follows. We start with an arbitrary uniform spatial point particle distribution with a known structure factor (SF).  We now suppose that each of these ``mother'' particle splits into a cloud of $m$ identical ``daughter'' particles, where $m$ is a cloud-independent constant.  Each daughter particle is then assumed to be displaced from its mother position by a stochastic displacement which may, or may not, be correlated with the displacements of other particles.  In other words each set of $m$ particles initially lying at the same spatial point ``explodes'' forming a ``cloud'' of particles around it.  For this reason we call the point process so generated a {\\em cloud process}.  We suppose that the displacements applied to different particles belonging to the {\\it same} mother are symmetrically distributed with arbitrary pair correlations.  One can choose, for instance, these correlations in order to fix certain moments of the mass dispersion of each cloud.  We will distinguish in the following between the two cases in which the displacements applied to particles originating from {\\it different} mothers are, or are not, correlated. We note that the results obtained for these cases are generalizations of those obtained in \\cite{displa}, where the case of a single daughter for each mother was analyzed and solved exactly.  The main immediate application we have in mind of the present study, and the one we discuss at some length, is the determination of the constraints on the stochastic displacement field and the number of daughters $m$ required to produce a target behavior of the SF at large scale (i.e. small wave number $k$) in the daughter particle distribution.  More specifically we are interested in the case of a large scale SF inherent to {\\em superhomogeneous} (or {\\em hyper-uniform}) stochastic point particle distributions \\cite{glass,lebo,torquato}. This class of distribution is characterized by the convergence of the SF to zero as $k \\rightarrow 0$. We will show explicitly that, for the case that the mother distribution is a regular lattice array, the associated (positive) exponent in the $k \\rightarrow 0$ limit of the SF is related to the conservation of the local mass moments in the passage from the single point particle to the cloud of daughter particles\\footnote{We note that approximate heuristic derivations of some of these results can be found in the cosmological literature (see, e.g., \\cite{pee80, wandelt}) in which constructions of this kind are considered in discussions of ``causality bounds'', i.e., bounds imposed on the large scale behavior of the SF by the existence of a causal horizon in cosmological models.}.  The paper is structured as follows. In the next section we introduce our notation and the essential definitions, and then derive a general expression for the SF of the daughter distribution, before averaging over the realizations of the cloud process. In Sect.~\\ref{sec3} we derive our general result for the SF for the case that there are only correlations between displacements of particles derived from the same mother. In the following section we apply this result to the specific case that the mother distribution is a regular lattice, and derive the small $k$ behavior of the SF of the daughter distribution. In Sect.~\\ref{sec5} we then derive our result for the more general case that there is also arbitrary correlation between the displacements of daughters of different mothers, and then also consider the specific case of an initial regular lattice.  As an example we derive the small $k$ behavior of the SF of a lattice of correlated dipoles. In the final section we summarize our results and discuss briefly both possible developments and further applications of the results reported here. ",
        "conclusions": "\\label{sec6}  In this paper we have introduced and analyzed a wide class of non-trivial stochastic point processes for which it is possible to write exactly the two-point correlation function and/or SF.  They are obtained from a given ``mother'' particle distribution, the SF of which is assumed known, by substituting each particle with a cloud of a fixed number $m$ of other particles.  The position of the new particles composing the clouds is determined by that of the related mother particle plus a stochastic displacement vector. An important assumption in all our calculations of the SF is that the stochastic process determining the displacements of the particles is independent of the mother distribution, i.e., the displacements of the cloud particles do not depend on the properties, statistical or otherwise, of the mother distribution to which they are applied. In practice this means that our SF is defined with respect to an ensemble average over two independent ensembles: one describing the realizations of the mother distribution, the other those of the displacement process.  We have distinguished two cases in which: (i) only the displacements of different particles belonging to the same cloud may be correlated, and (ii) the displacements of particle belonging to different clouds may also be statistically dependent.  In both cases we obtain a direct generalization of the relations found in \\cite{displa} for $m=1$.  In the first case, once the average over realizations is taken, the SF of the final particle distribution is related to the SF of the mother distribution by a local relation in the wave vector $\\bk$, while the second case leads to a more complex relation.  A detailed analysis of case (i) led us to find and to discuss, in the case of an initial lattice mother distribution, the relations linking the exponent of the final SF to the number of conserved mass moments in each cloud: we have seen that such an exponent, if larger or equal to $4$, tells us directly which are the locally conserved mass moments in the distribution. When we move to case (ii), the presence of cloud-cloud displacement correlations ``interact'' with the local mass moments conservation in determining the small $k$ behavior of the SF of the particle distribution.  One application we have in mind of the results derived here for this class  of ``stochastically ordered'' point processes is in the systematic study of the dynamics of particle systems driven by long range pair interactions. More specifically, in the case of gravity, it is expected (see, e.g., \\cite{pee80}) that the large scale fluctuations in an infinite particle system dominate the dynamics of the gravitational clustering for an initial SF with a small $k$ behaviour proportional to $k^\\gamma$ and $\\gamma <4$. This ``hierarchical'' behaviour has been observed numerically for a range of such $\\gamma$, up to a maximal value of $\\gamma=2$ (see \\cite{sl1} for a recent discussion, and further references). No study of the regime of initial conditions $\\gamma >4$ has been performed up to now, as no algorithm has been given in the literature, to our knowledge, which can generate such an initial condition\\footnote{Explicit algorithms generating point distribution with $\\gamma=4$ have, on the other hand, been discussed. See, notably, \\cite{fratzl_etal1999, wandelt, gabrielli_etal}.}. It is expected that the gravitational clustering will be qualitatively different in this case, with structures being built up from smaller to larger scales. Indeed the reason why we expect such a difference can be understood easily in the context of our constructions here of such point processes. For $\\gamma > 4$ the fluctuations in mass are so suppressed that gravity is effectively ``screened'' (at least in the initial conditions): in a multipole expansion of the mass far away from a given point in the ``uncorrelated cloud lattice'' only the leading moment varying from cloud to cloud will contribute (as the contributions from the moments which are fixed will cancel out). If this leading contributing moment is the second moment, this gives an effectively short-range force (decaying as the inverse of the fourth power of the distance).  Our analysis here has  been limited to the case of ``analytic''  exponents for the SF derived from short tailed displacements probability density functions (PDF).  In principle one can consider also the case in which such PDFs have a long power law tail. In this case a singular part of the small $k$ expansion of the related displacement characteristic function arises. This leads (see \\cite{displa, libro}) to a final SF characterized by non-analytic (e.g. fractional) small $k$ exponents. We have seen, however, that in the uncorrelated cloud lattice, to attain powers larger than $\\gamma=4$ we must in fact take limited displacements, and thus we necessarily obtain  ``analytic'' exponents (and, in fact, a analytic behaviour of the SF at small $k$). In the case of the correlated cloud lattice, nevertheless, we have given an example (random correlated dipoles) which shows how such non-analytic powers should be attainable by including appropriate correlations between the clouds.  We return finally to an important feature of Eq.~(\\ref{eq24}) which we have discussed at some length in Sect.~\\ref{sec3}.  This is that the exponent in the small $k$ scaling behavior of $S_\\rho(\\bk)$ cannot be larger than that in $S_n(\\bk)$.  As we explained, this can be understood physically as it means that the replacements of particles by clouds cannot make the initial particle distribution more ordered. We underline, however, that this is true given the assumptions we have made, and specifically assuming that the stochastic process generating the clouds is independent of the mother distribution. In a forthcoming article with another collaborator \\cite{tiling} we will report results on a related kind of construction of superhomogeneous point processes, starting from tilings of space with equal volume tiles. In this case it turns out that one can, in certain circumstances, ascribe a cloud of particles to each particle of a given point distribution and as a result increase the exponent $\\gamma$. The reason why this becomes possible is that the displacements of the particles are not applied independently of the correlation properties of the underlying point process, as we have assumed here. Indeed in order to increase the exponent requires that the moments of the clouds are ``tuned'' appropriately to the tile in which the mother point lies. We note also that, in the case that the initial tiling is taken to be a lattice, the algorithm described coincides with that given here and similar results to those given here are recovered.  Details will be reported in \\cite{tiling}.  We thank B. Jancovici, J. Lebowitz and S. Torquato for useful discussions.  \\appendix"
    },
    "0711/0711.2333_arXiv.txt": {
        "abstract": "This paper begins a series in which we examine the structures of distant galaxies to directly determine the history of their formation modes. We start this series by examining the structures of $z_{F850LP} < 27$ galaxies in the Hubble Ultra-Deep field, the deepest high-resolution optical image taken to date.  We investigate a few basic features of galaxy structure using this image.  These include: (1) The agreement of visual eye-ball classifications and non-parametric quantitative (CAS, Gini/\\m20) methods;  (2) How distant galaxy quantitative structures can vary as a function of rest-frame wavelength; and (3) The evolution of distant galaxy structures up to $z \\sim 3$.  One of our major conclusions is that the majority of galaxies with $z_{850} < 27$ are peculiar in appearance, and that galaxy assembly is rapidly occurring at these magnitudes, even up to the present time.   We find a general agreement between galaxy classification by eye and through quantitative methods, as well as a general agreement between the CAS and the Gini/\\m20 parameters.  We find that the Gini/\\m20 method appears to find a larger number of galaxy mergers than the CAS system, but contains a larger contamination from non-mergers. We furthermore calculate the merger rate of galaxies in the UDF up to $z \\sim 3$, finding an increase with redshift as well as stellar mass, confirming previous work in the Hubble Deep Field. We find that massive galaxies with M$_{*} >$ \\lmass undergo 4.3$_{+0.8}^{-0.8}$ major galaxy mergers at $z < 3$, with all of this merging occurring at $z > 1$. ",
        "introduction": "Understanding the formation of galaxies is one of the most intriguing and unanswered problems in astronomy today. Studying galaxies in the distant universe is one of the major approaches for deciphering their formation and evolution.  Traditionally, distant galaxies are examined in terms of their gross properties, such as luminosities (e.g., Faber et al. 2007), colours (e.g., Moustakas et al. 2004), and sizes (Trujillo et al. 2007), and stellar masses (e.g., Bundy et al. 2005; Conselice et al. 2007a,b).  The least developed approached, but potentially one of the most exciting for revealing unique information, is how the structures and morphologies of galaxies change through cosmic time.  The reason a structural and morphological approach is potentially the most fruitful for understanding galaxy formation and evolution is that the structures of galaxies correlate with formation modes, such as star formation and galaxy mergers (Conselice 2003).  It is thus possible to determine the physical processes behind galaxy formation directly without having to compare observables with physical models. This direct approach has been used in the past to calculate the role of galaxy mergers in the formation of massive galaxies (e.g., Conselice et al. 2003a; Lotz et al. 2006; Conselice 2006).  This previous work showed that the most massive galaxies with M$_{*} >$ \\lmass appear to form nearly all their stellar mass in mergers, which occur mostly, and in large numbers, at $z > 1.5$.  There are however still many outstanding problems with using galaxy structure to derive evolution, including the reliability of the approaches used, and the uniqueness and repeatability of results thus far obtained. For the most part, it is necessary to use deep Hubble Space Telescope imaging to measure the structures and morphologies of a large number of distant galaxies, although adaptive optics, particularly in the near-infrared, is becoming another powerful approach (e.g., Wright et al. 2007).  Since the imaging cameras on the Hubble Space Telescope are limited in field of view, only a small fraction of distant galaxies have been studied thus far (e.g., Williams et al. 1996; Giavalisco et al. 2004). Particularly, most of the merger history of galaxies at $z > 1$ has been measured within the Hubble Deep Field-North, which currently has the only high resolution deep near-infrared imaging of distant galaxies (e.g., Dickinson et al. 2000; Conselice et al. 2003a; Papovich et al. 2005).  Thus, it is important to study the structures and morphologies of galaxies in as many fields as possible, using different techniques.  The Hubble Ultra Deep Field (UDF) represents perhaps our best opportunity to study the most distant galaxies at observed optical wavelengths for the foreseeable future. The UDF is a major Hubble imaging program utilising the Advanced Camera for Surveys (ACS) to image a single pointing in the sky to the deepest depth ever probed in the optical.  As the ACS UDF images are very deep, and have a higher resolution than the WFPC2 HDFs, we can address some fundamental questions concerning galaxy structure and morphology.  These include: understanding the nature of galaxy morphologies seen in deep optical surveys; the agreement, or lack thereof, between apparent morphology and structural parameters; and the evolution of galaxy structure through time.  We are particularly interested in determine how major galaxy mergers are driving the formation of the galaxy population at early times.  Our major finding in this paper is that galaxies in the past are significantly more irregular, peculiar, asymmetric, and clumpy than galaxies today, and that this likely results from the past merger and assembly activity that established these systems.   We show that a large fraction of this irregularity cannot be produced simply by star formation, and derive the merger fraction and rate for galaxies of various stellar masses. We confirm the findings of Conselice et al. (2003a) and Conselice (2006) that massive galaxies with M$_{*} >$ \\lmass have a steeply increasing merger fraction, which evolves as $\\sim$ $(1+z)^{6}$ between $z = 0.8$ and 3.   This suggests that these massive galaxies undergo on average 4.3$^{+0.8}_{-0.8}$ major mergers between $z \\sim 3$ to 0.  This paper is organised as follows: \\S 2 includes a discussion of the data sources we use in this paper, \\S 3 is a description of our morphological and structural analyses, and the stellar masses we utilise, \\S 4 is a discussion of our results, including how different structural analysis techniques compare with each other, especially how merger finding methods differ, \\S 5 gives our description of the merger history up to $z \\sim 3$, and \\S 6 is our summary and conclusions.   We use a standard cosmology of H$_{0} = 70$ km s$^{-1}$ Mpc$^{-1}$, and $\\Omega_{\\rm m} = 1 - \\Omega_{\\lambda}$ = 0.3 throughout. ",
        "conclusions": "This paper begins a series in which we examine the structures of distant galaxies, and determine the likely role of galaxy merging in the formation of galaxies. In this paper we examine structural parameters using the CAS and Gini/\\m20 system on galaxies found within the Hubble Ultra Deep Field (UDF).  We also examine the stellar masses and eye-ball estimates of morphological types for these systems. Our major conclusions and findings include:  \\noindent I. Down to $z_{850} = 27$ the majority of galaxies found in the UDF are peculiar in appearance.  This suggests that galaxy formation is actively ongoing for these systems.  \\noindent II. We compare how visual estimates of morphology agree with positions in CAS and Gini/\\m20 space. We find a generally good agreement between these methods with the merger region in CAS space nearly totally occupied by galaxies which  are visually classified as peculiars/mergers.  \\noindent III. We examine the merger fraction of galaxies in the UDF using the CAS system and compare with our previous results of the merger fraction and rate derived using the Hubble Deep Field. We confirm our earlier measurements of the merger history, including the fact that the highest mass galaxies with M$_{*} >$ \\lmass have a steeply increasing merger fraction up to $z \\sim 3$. This increase can be fit as a power-law $\\alpha (1+z)^{m}$, with $m \\sim 6$ for the combined UDF and HDF-N sample.  We find that the merger fraction for lower mass galaxies are lower at high redshift, but reach their peak fractions at lower redshifts. In total we find that M$_{*} >$ \\lmass galaxies undergo on average 4.3$^{+0.8}_{-0.8}$ major mergers at $0 < z < 3$, with most of this merging occurring at $z > 1$.  We thank Dan Coe for making his analysis results of the UDF freely available, and Kevin Bundy for assistance with the stellar masses. We also acknowledge support from the University of Nottingham."
    },
    "0711/0711.2105_arXiv.txt": {
        "abstract": "Thick, fully depleted p-channel charge-coupled devices (CCDs) have been developed at the Lawrence Berkeley National Laboratory (LBNL). These CCDs have several advantages over conventional thin, n-channel CCDs, including enhanced quantum efficiency and reduced fringing at near-infrared wavelengths and improved radiation tolerance. Here we report results from the irradiation of CCDs with 12.5 and 55~MeV protons at the LBNL 88-Inch Cyclotron and with 0.1 - 1~MeV electrons at the LBNL $^{60}$Co source. These studies indicate that the LBNL CCDs perform well after irradiation, even in the parameters in which significant degradation is observed in other CCDs: charge transfer efficiency, dark current, and isolated hot pixels. Modeling the radiation exposure over a six-year mission lifetime with no annealing, we expect an increase in dark current of 20~e$^-$/pixel/hr, and a degradation of charge transfer efficiency in the parallel direction of $3 \\times 10^{-6}$  and $1 \\times 10^{-6}$ in the serial direction. The dark current is observed to improve with an annealing cycle, while the parallel CTE is relatively unaffected and the serial CTE is somewhat degraded. As expected, the radiation tolerance of the p-channel LBNL CCDs is significantly improved over the conventional n-channel CCDs that are currently employed in space-based telescopes such as the Hubble Space Telescope. ",
        "introduction": "The SuperNova/Acceleration Probe (SNAP) is a proposed space-based telescope dedicated to the study of dark energy through the observations of Type Ia supernovae (Ia SNe) and a deep, wide area weak lensing survey\\cite{greg}. From its orbit at the second Earth-Sun Lagrange point (L2), SNAP will carry out two surveys: a deep survey of $7.5$ square degree field with 4-day cadence repeat visits over a period of 22 months to discover and obtain light curves and spectra of over 2000 Ia SNe in the redshift range $0.3 < z < 1.7$; and a wide area weak lensing map to study the growth of large scale structure that will cover 1000 square degrees per year to a depth of AB magnitude $28.0$ in the optical filters. In an extended 6 year SNAP mission, the weak lensing survey covers 4000 square degrees and the mission lifetime.  The telescope is designed with a $0.7$ square degree instrumented field of view divided evenly between 36 CCDs and 36 HgCdTe detectors. The SNAP observing strategy implements a four-point dither pattern with an exposure time of 300 seconds to recover spatial information from the undersampled optics and to reject cosmic rays. The focal plane will be passively cooled to 140~K.  Nine fixed filters cover the wavelength range 400~nm to 1700~nm. With a diffraction limited point spread function (PSF) of $0.1$ arcseconds at~800 nm and zodiacal-dominated background, SNAP will have significantly improved resolution and decreased contamination from sky background compared to ground based telescopes.  The SNAP focal plane design uses thick, fully depleted CCDs developed at LBNL~\\cite{steve1,steve} for visible to near IR observations in six bandpass filters. In space, these detectors will be exposed to significant radiation, primarily solar protons. In this paper we investigate the effects of six years of radiation at L2 on SNAP CCDs in order to qualify them for use in a space mission. In \\S\\ref{section:CCD} we describe the SNAP CCDs and the specifications for performance. The space environment and expected radiation exposure are discussed in \\S\\ref{section:space}. Irradiation using the 88-Inch Cyclotron and the $^{60}$Co source at LBNL is described in \\S\\ref{section:irrad} and \\S\\ref{section:Co60} respectively. CCD performance after proton irradiation is reported in \\S\\ref{section:results} and after $^{60}$Co irradiation in \\S\\ref{section:Co60 results}. Finally, we present an interpretation of the results in the context of the SNAP mission in \\S\\ref{section:discussion} and the conclusions in \\S\\ref{section:conclusion}. ",
        "conclusions": "\\label{section:conclusion}  The behavior of thick, fully depleted, p-channel LBNL CCDs designed for the SNAP satellite has been investigated using irradiation at the LBNL 88-Inch Cyclotron and LBNL $^{60}$Co source. We have performed extensive tests of charge transfer efficiency, generation of dark current, and hot pixel formation from proton exposure. A summary of the results scaled to the expected exposure at L2 can be found in Table \\ref{tab:rad}. CTE performance after irradiation is calculated assuming pre-radiation parallel and serial CTEs of $0.999\\, 999$. The radiation studies show that the LBNL CCDs designed for use in the SNAP satellite will develop negligible contamination from dark current and hot pixels during the course of a six year mission.  \\begin{table}[h] \\renewcommand{\\arraystretch}{1.0} \\caption{Expected CCD Performance after Six Years at L2} \\label{tab:rad} \\centering \\begin{tabular}{|c||c||c|} Quantity & Pre-irrad  & Nominal Dose \\\\ \\hline Defect Pixels & $<0.001$  &  $4.1 \\times 10^{-6}$ \\\\ Dark Current & $3-4$ e$^-$/hr &  $20$ e$^-$/hr \\\\ Ser CTE-no anneal & 0.999 999 & 0.999 998  \\\\ Ser CTE-w/anneal & 0.999 999 & 0.999 997  \\\\ Parallel CTE & 0.999 999 & 0.999 996  \\\\ \\hline \\end{tabular} \\end{table}  Monte Carlo simulations of radiation exposure by propagation of solar protons through the complex shielding of the SNAP satellite will finalize estimates of radiation dose over the mission lifetime. Additional analysis is required to quantify the impact of the degraded performance on science observations: CTE degradation impact on galaxy shapes for weak lensing science goals and CTE dependence on sources signal strength. Future studies of the effects of $^{60}$Co irradiation over a longer time span would be useful in better understanding the mechanism for the generation of dark current.  Nevertheless, the results reported here show that the LBNL CCDs are significantly more radiation tolerant than n-channel CCDs currently in use in space-based observatories. This makes the LBNL CCDs an excellent choice for use in future space-based missions such as SNAP."
    },
    "0711/0711.1089_arXiv.txt": {
        "abstract": "An intensive monitoring program of 54 6.7-GHz methanol maser sources was carried out at the Hartebeesthoek Radio Astronomy Observatory from January 1999 to April 2003.  The monitoring program was subsequently continued on 19 sources of interest.  Analysis of the resulting time-series stretching over eight years shows that six of the sources are periodic, with periods ranging from 133 days to 504 days.  The waveforms in individual sources range from sinusoidal fluctuations to sharp flares and there can be other long term trends in the time-series. The amplitudes of the variations can also change from cycle to cycle.  The time-series of the periodic masers will be presented, and possible causes of the variability discussed. ",
        "introduction": "The existence of variability in 6.7-GHz masers was established by \\cite{Cas95c} and \\cite{Mac96}. However, the nature of the variability had not been fully characterised until the study of \\cite{Goe04} who monitored 54 sources for four years. The results of this study showed a range of variability behaviour, including monotonic, quasi-periodic and  what appeared to be periodic variations. The monitoring programme has been continued on the most variable sources, enabling the confirmation of periodicity in six sources, despite considerable upgrades to the telescope during this time.  The time-series of the periodic candidates are presented here. ",
        "conclusions": "Periodicity was found in $\\sim$ 10\\% of the original sample of 54 sources monitored. No clear cause of the variations can be found at this stage.  High-resolution, multi-wavelength observations of the methanol maser sources are necessary to understand conditions in these regions.  Detailed modelling will also be necessary to understand the effects of radiative transfer in different morphologies and maser models with time-dependent components would have to be developed. This may be a way to gain greater insight into processes occurring in regions that can not be directly observed."
    },
    "0711/0711.3273_arXiv.txt": {
        "abstract": "A viable class of magnetogenesis models can be constructed by coupling the kinetic term of the hypercharge to a spectator field whose dynamics does not affect the inflationary evolution.  The magnetic power spectrum is explicitly related to the power spectrum of (adiabatic) curvature inhomogeneities when the quasi-de Sitter stage of expansion is driven by a single scalar degree of freedom. Depending upon the value of the slow-roll parameters, the amplitude of smoothed magnetic fields over a (comoving) Mpc scale can be as large as $0.01$--$0.1$ nG at the epoch of the gravitational collapse of the protogalaxy. The contributions of the magnetic fields to the Sachs-Wolfe plateau and to the temperature autocorrelations in the Doppler region compare favourably with the constraints imposed by galactic magnetogenesis. Stimulating lessons are drawn on the interplay between magnetogenesis models and their possible CMB signatures. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.1276_arXiv.txt": {
        "abstract": "{Hydrodynamical model of the $\\kappa$-mechanism in a purely radiative case.} {First, to determine the physical conditions propitious to $\\kappa$-mechanism in a layer with a configurable conductivity hollow and second, to perform the (nonlinear) direct numerical simulations (DNS) from the most favourable setups.} {A linear stability analysis applied to radial modes using a spectral solver and DNS thanks to a high-order finite difference code are compared.} {Changing the hollow properties (location and shape) lead to well-defined instability strips. For a given position in the layer, the amplitude and width of the hollow appear to be the key parameters to get unstable modes driven by $\\kappa$-mechanism. The DNS achieved from these more auspicious configurations confirm the growth rates as well as structures of linearly unstable modes. The nonlinear saturation follows through intricate couplings between the excited fundamental mode and higher damped overtones.} {} ",
        "introduction": "Since the beginning of the studies concerning Cepheids, it is well known that convection occurs in Cepheids' envelopes, and thus changes pulsation properties \\citep[e.g. the reviews of][]{Gautschy-Saio,buchler-cepheid-review}. The coldest ones, which are located next to the red edge of the instability strip, have the more extended surface convective zones. However, during many years Cepheids' oscillations models have used the so-called ``frozen-in convection'' approximation which claims that convective flux perturbations are negligible \\citep{Baker1962}. Such kind of models well predict the blue edge of the instability strip but fail to explain the red edge as in this case, the strong existing couplings between the convection and the oscillations are not taken into account. This discrepancy becomes obvious with the accumulation of accurate observations which show a narrower instability strip than the theoretical one, i.e. modes which are linearly unstable in the models are located outside the observational strip \\citep[][hereafter YKB98]{YKB}.  The main theoretical difficulty comes from the fact that convection plays a crucial role on the pulsations while we know that convection itself remains roughly described by mixing-length theories \\citep{BV2,BV1}. However, several attempts in the direction of time dependent convection models (TDC) have been developed \\citep[e.g.][]{Unno,Gough,St}. Recent studies \\citep[YKB98;][]{Bono} rely on Stellingwerf's convection model \\citep{St} or very similar newer approaches to compute linear and nonlinear time evolution of amplitudes of modes \\citep{K,GW,WF}. The major problem of TDC is the choice of the many free parameters introduced by the convection model\\footnote{e.g. the seven coefficients $[\\alpha_c,\\ \\alpha_t,\\ \\alpha_{\\nu},\\ \\alpha_{\\Lambda},\\ \\alpha_s,\\ \\alpha_d$,\\ $\\alpha_p]$ in YKB98 or the eight ones in \\cite{Koll02}.}. As these parameters are not theoretically well determined, one should adjust them by fitting the observations.  Another way to study the convection-pulsation interaction is to achieve (nonlinear) direct numerical simulations (DNS) of the whole hydrodynamical problem. The final aim of our work is to realise such kind of simulations in 2-D and 3-D where a convective zone will be coupled with a radiative one and unstable radial acoustic modes will be self-consistently excited by $\\kappa$-mechanism. However, as DNS are highly time consuming, it is necessary to get in a first step the appropriate initial conditions. That is why we have tried to determine in this paper the physical conditions propitious to an excitation based on $\\kappa$-mechanism.  \\citet{Eddington}, and then \\citet{Zhevakin53} and \\citet{1958Cox} have introduced a mechanism linked to the opacity in ionisation regions, the $\\kappa$-mechanism, where $\\kappa$ denotes opacity \\citep[see also the review of][]{Zhevakin}. They have shown that Cepheids' radial oscillations are driven by a thermal heat engine as radial pulsations have to be maintained thanks to a sustained physical process. That is why, they imagined a Carnot-like thermodynamic cycle which stores heat during compression phases while releasing it during the decompression ones. This mechanism, now called the \\textit{Eddington's valve}, can only occur in specific regions of a star where the opacity varies so as to block the radiative flux during compression phases. Now, opacity tables such as the OPAL ones show strong increases in ionisation regions of main (i.e. Hydrogen and Helium) or heavy elements that are commonly called ``ionisation bumps'' \\citep[e.g.][]{Seaton-Badnell}. As a consequence these ionisation zones are locally responsible for modes amplification. Moreover, beyond this criterion on opacity, these ionisation regions have to be located in a very precise region of a star.  Indeed, if they are too deep or too close to the surface, the driving they cause can be balanced by the damping occurring in other regions. Therefore, an efficient ionisation region has to be located in a specific place, called the \\textit{transition region}, which marks the shift from a quasi-adiabatic interior to a strongly non-adiabatic surface. In classical Cepheids oscillating on the radial fundamental mode, the overlap of this transition region with the ionisation one is around $40\\ 000$ K, which corresponds to the temperature of the helium second ionisation \\citep{Baker-Kippenhahn}. For first overtone Cepheids, things are more intricate as one must take into account first, the respective position between this HeII ionisation region and the nodal line and second, the HeI/H region which also contributes to the driving \\citep{Bono}.  However, the location of these opacity bumps are not solely responsible of the acoustic mode destabilisation. A careful treatment of the $\\kappa$-mechanism in Cepheid stars would also involve the possible dynamical couplings with convective zones or, say, metallicity effects through a realistic equation of state and opacity tables. The corresponding physics is actually not fully understood from a theoretical point of view and the purpose of our work is to sufficiently simplify the hydrodynamical approach while keeping at the same time the \\textit{leading order} phenomenon responsible of the instability, that is, the opacity bumps location.  That's for why we adopt a fully radiative layer of a perfect gas in which a ionisation region is represented by a ``hollow'' in radiative conductivity, corresponding to a ``bump'' in opacity. Strictly speaking, the layer stability will depend on both temperature and density variations, as the radiative conductivity is a function of these two physical quantities. However, as the opacity strongly depends on temperature, this state variable mainly controls the instability. As our final aim is to realise an hydrodynamical study of $\\kappa$-mechanism, we will therefore neglect any dependence of radiative conductivity on density and thus, our conductivity profile is merely a function of temperature. The inferred advantage is to allow easy changes in the different parameters of the ionisation region by setting both position and shape of the conductivity hollow (i.e. its slope, width and amplitude). As a consequence, it is possible to investigate a complete parametric study of $\\kappa$-mechanism in order to determine precisely the physical conditions required by the instability.  We first consider the radiative and hydrostatic equilibria of our layer with an appropriate conductivity profile. By adjusting both the density value at the top of the layer and the flux at its bottom, we then obtain a transition region in the middle of the computational domain. Secondly, we investigate the linear stability analysis by solving perturbation equations for the oscillations. We thus obtain the whole spectrum, and can sort out unstable modes from stable ones. Therefore, we are able to check the relevance of the transition region concept by making every parameters of the conductivity hollow vary.  The first main result of our linear stability study is the confirmation of the underlying conditions defining the transition region. With different parametric studies, we obtain instability strips corresponding to the fundamental mode, as both a minimum hollow width and amplitude are needed to obtain unstable modes. These results are interpreted thanks to the work integral to exactly determine the location of the driving zone. As a consequence, this parametric study of $\\kappa$-mechanism provides us the physical quantities responsible for the instability.  Secondly, these auspicious conditions needed to drive the fundamental mode constitute the starting point of 2-D DNS. Indeed, we check the growth rates as well as the structure of the linearly unstable modes by performing direct simulations until the nonlinear saturation of modes.  In \\S\\ref{hydro_eqs}, we first introduce the general oscillation equations while the different conditions leading to the instability are determined in \\S\\ref{ref:inst-cond}. In \\S\\ref{sec:hydro_model}, we present our hydrodynamical model. Linear stability analysis and DNS results are thus given and compared in \\S\\ref{sec:results} and \\S\\ref{sec:DNS}, respectively, before concluding in \\S\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  We have modeled $\\kappa$-mechanism in Cepheids by a similar physical problem: the propagation of radial acoustic waves in a partially ionised shell. Our model consists in a perfect gas embedded in an entirely radiative layer and the ionisation region has been depicted by a configurable hollow in radiative conductivity. This approach allows to quickly change the shape and location of this ionised region in the layer through the hollow parameters $\\Tbump,\\ \\Amp, \\ \\sigma$ and $e$ (Eqs.~\\ref{conductivity-profile1}-\\ref{conductivity-profile2}). We first have checked that a hollow with a sufficient amplitude and width may lead to  \\begin{equation} \\dfrac{d\\KT}{dz} < 0, \\label{crit1} \\end{equation} which corresponds to the classical instability criterion derived from the work integral and quasi-adiabatic considerations. Nevertheless, another condition is needed to obtain a thermal heat engine in Eddington's sense: the ionisation region has to be located at a certain place in the layer, neither to deep nor to close to the surface. We have shown that this propitious area is located in the so-called \\textit{transition zone} separating the quasi-adiabatic interior from the strongly non-adiabatic surface. This thermodynamical criterion can be summarised by  \\begin{equation} \\Tion \\sim \\Ttr  \\ltex{and} \\dfrac{\\langle c_v \\Ttr \\rangle \\Delta m}{\\Pi L} \\sim 1. \\label{crit2} \\end{equation}  With this second criterion, appropriate boundary conditions have been chosen in order to have a transition region in the middle of our box for the fundamental acoustic mode. Both radiative and hydrostatic equilibria have then been discretized on the (spectral) Gauss-Lobatto grid and solutions have been computed thanks to a linear solver. Once done, we have performed a linear stability analysis by computing the whole spectrum as well as associated eigenvectors for the radial oscillations equations.  The main advantage of our approach lies in its ability to make the parameters of the conductivity hollow vary. The major result of parametric surveys is the checking of the previous conditions by way of the appearing of instability strips: only configurations satisfying both conditions (\\ref{crit1}-\\ref{crit2}) have led to unstable fundamental modes (e.g.~Figs.~\\ref{fig:ampli-Tbump} and \\ref{fig:wi2}).  Then, these most favourable setups found in the linear stability analysis have been the starting points of our first 2-D (nonlinear) DNS of $\\kappa$-mechanism. However, because of a too restrictive Courant-Friedrichs-Levy (CFL) condition at the surface, these setups were not well-suited for a fully explicit code such as the Pencil Code one. We therefore have developed a new \\textit{implicit} module to deal with the radiative diffusion term and thus soften the CFL constraint. These DNS have confirmed with a great agreement both growth rates and structures of linearly unstable modes (see Figs.~\\ref{fig:growth-rate} and \\ref{fig:vert-prof}). Moreover, we have been able to reach the nonlinear saturation which involves an intricate coupling between the fundamental mode and (at least) the damped second overtone of which the period is a multiple of the fundamental one.  This work constitutes a first step in our Cepheids' project devoted to the modelling of the convection-pulsation interaction in the coldest Cepheids close to the red edge of the instability strip. Indeed, it is well known that convection occurs in a non-negligible part of these stars and modify the pulsation properties. For now, models based on time-dependent convection theories reproduce quite well the position of this red edge or, say, observational periods and light curves \\citep[YKB98;][]{Bono}. However, they involve many free parameters which are either fitted to the observations or hardly constrained by theoretical values, leading to almost similar results for different parameter sets \\citep{Koll02,Szabo07}.  Despite their own limitations (their weak contrasts in pressure through the computational domain or their thermal timescale problem, see e.g. \\cite{Bran00}), DNS are without a doubt a good way to address this convection-pulsation interaction as they fully take into account the crucial nonlinearities. With a conductivity profile solely based on temperature (Eq.~\\ref{conductivity-profile1}), it is easy to \\textit{locally} shape one (or several) convective zone by increasing the temperature gradient above the adiabatic one. Indeed, as an ionisation region corresponds to a local increase in opacity, it is well known that convection can develop there. This occurs in cold Cepheids where two separate convectively unstable regions superimpose with the HeI/H and HeII ionisation regions. By adjusting in DNS the convective zone width or the strength of gravity, it will be possible to match the (local) turnover timescale of convection with the mean mode period, and then to study the coupling between convection and pulsation."
    },
    "0711/0711.1781_arXiv.txt": {
        "abstract": "%  On Earth, the Dry Valleys of Antarctica provide the closest martian-like environment for the study of extremophiles. Colonies of bacteries are protected from the freezing temperatures, the drought and UV light. They represent almost half of the biomass of those regions. Due to there resilience, endolithes are one possible model of martian biota.  We propose to use infrared spectroscopy to remotely detect those colonies even if there is no obvious sign of their presence. This remote sensing approach reduces the risk of contamination or damage to the samples. ",
        "introduction": "Space exploration is a difficult task and the search for life is no different. The equipement size a probe can bring severely limit the scope of the search. Every sample cannot be analyzed and selecting those that can be is not trivial. This project investigates the possibilities of remote detection using infrared spectroscopy in order to select those few samples to pick and analyse further. ",
        "conclusions": "Our results show that it is possible to detect biological signatures using a spectrometer operating in the middle infrared range. However, it is not possible to highlight a particular region, or regions, in a spectrum to be used to identify biomarkers. The interdependance of the absorption bands related to the living are too complex to simply isolate a few.  The proposed technique calls for a more subtle approach by comparing witness spectra and an unknown spectrum by plotting them in the PCA's space. If the test spectrum, once projected into the PCA's space, is close to the reference group then the probability of it is containing biomarkers is high. Adding a non-organic spectra to the PCA space as references may improve the idenfitication scheme.  This technique could be used to pinpoint potential life harboring rock for more detailed analysis. It could be possible to extend the method to better idenfity the unknown spectrum using a more precise reference database."
    },
    "0711/0711.4995_arXiv.txt": {
        "abstract": "We present the first detection of molecular gas cooling CO emission lines from ordinary massive galaxies at $z=1.5$. Two sources were observed with the IRAM Plateau de Bure Interferometer, selected to lie in the mass-star formation rate correlation at their redshift, thus being representative of massive high-$z$ galaxies. Both sources were detected with high confidence, yielding $L'_{\\rm CO}\\sim2\\times10^{10}$~K~km~s$^{-1}$~pc$^2$. For one of the sources we find evidence for velocity shear, implying CO sizes of $\\sim10$~kpc. With an infrared luminosity of $L_{\\rm FIR}\\sim10^{12}L_\\odot$, these disk-like galaxies are borderline ULIRGs but with star formation efficiency similar to that of local spirals, and an order of magnitude lower than that in submm galaxies. This suggests a CO to total gas conversion factor similar to local spirals, gas consumption timescales approaching 1~Gyr or longer and molecular gas masses reaching $\\sim10^{11}M_\\odot$, comparable to or larger than the estimated stellar masses. These results support a major role of {\\em in situ} gas consumption over cosmological timescales and with relatively low star formation efficiency, analogous to that of local spiral disks, for the formation of today's most massive galaxies and their central black holes. Given the high space density of similar galaxies, $\\sim10^{-4}$~Mpc$^{-3}$, this implies a widespread presence of gas rich galaxies in the early Universe, many of which might be within reach of detailed investigations of current and planned facilities. ",
        "introduction": "Galaxies in the distant Universe were forming stars with much higher rates than today (Madau et al 1996; Elbaz et al 2002; Le Floc'h et al. 2005). For galaxies of a fixed stellar mass, e.g. of the order of $M^*\\sim10^{10-11}M_\\odot$, the typical star formation rate (SFR) has been steadily increasing with respect to local values, up to $\\sim30$ times higher at $z=2$ (Daddi et al. 2007a; D07a hereafter). This increase can be accounted for either by postulating a higher gas content or a higher star formation efficiency, or both. Galaxy mergers and interactions, expected to be much more common in the distant Universe than today, can increase the density of gas in galaxies, and thus help transform gas into stars more rapidly and with a higher star formation efficiency. The relative importance for massive galaxy formation and assembly of the hierarchical merging of small sub galactic units, versus the in situ formation of stars through gas consumption is largely unknown observationally.  The molecular gas content of galaxies is best probed through the detection of redshifted cooling emission lines of CO. Combining observations for local and distant galaxies, all the way from spirals to (ultra) luminous infrared galaxies  (LIRGs, $L_{\\rm FIR}\\simgt10^{11}L_\\odot$, and ULIRGs, $L_{\\rm FIR}\\simgt10^{12}L_\\odot$), a nonlinear correlation is observed (e.g., Solomon \\& van den Bout 2005) between CO luminosity (proportional to molecular gas content) and FIR luminosity (proportional to the SFR), with more luminous sources having higher FIR to CO luminosities, i.e. higher efficiencies than spiral disks  in converting gas into stars, likely due to their higher gas densities (Tacconi et al. 2006; Downes \\& Solomon 1998). Mapping the physics of molecular gas in representative galaxies at high redshifts  would be a major step forward into understanding the processes regulating galaxy formation, and is a key target of future astrophysical facilities such as ALMA.  However, if the local correlations hold at high redshifts, current techniques and sensitivities allow us to detect CO lines, and thus study the molecular gas content, only for extreme objects such as bright ($S_{850\\mu m}>5$~mJy) submm-selected galaxies (Genzel et al. 2004; Neri et al. 2003; Greve et al. 2005). (SMGs), quasars (Walter et al. 2003; Maiolino et al. 2005), or in the case of strong gravitational lensing. Indirect evidence has been mounting very recently that this might not be the general case. Vigorous starbursts at high redshift appear to be the ordinary star formation mode in massive galaxies. The existence of tight correlations between SFRs and galaxy masses (Elbaz et al. 2007; Noeske et al. 2007; D07a) and the predominance of ULIRGs inside mass limited samples of $M\\simgt5\\times10^{10}M_\\odot$ galaxies (Daddi et al 2005; D07a; Caputi et al 2006) suggests that vigorous starbursts were typically long lasting among ordinary massive galaxies in the distant Universe, which therefore  might be hosting large amounts of gas.  \\begin{figure} \\centering \\includegraphics[width=8.0cm,angle=0]{f1.eps} \\caption{Top panels: CO [2-1] emission line detections in two ordinary, massive high-$z$ galaxies. We show (uncleaned) velocity averaged spatial maps of CO emission, from -350~km~s$^{-1}$ to 425~km~s$^{-1}$ ($BzK-4171$, left) and between  -575~km~s$^{-1}$ and -25~km~s$^{-1}$ and from -950~km~s$^{-1}$ and -800~km~s$^{-1}$ ($BzK-21000$, right). The contours are from -3 to $+6\\sigma$, in steps of $1\\sigma$ (0.1~mJy/beam). The field of view is 1$'$. The size of the beam is shown in the lower-left. Bottom panels: position-velocity maps, extracted along the major axis direction measured on the HST+ACS UV rest frame imaging. The spatial field of view is $\\sim9''$. Red/white colors are for positive pixels. Contours are at the level of 1~mJy/beam. } \\label{Pic} \\end{figure} ",
        "conclusions": "In order to evaluate the implications of these results for massive galaxy formation,  recall that the two observed galaxies are representative of massive galaxies at high redshifts. They were selected  for CO observations simply due to the availability of known spectroscopic redshifts. The two galaxies lie on the SFR to stellar mass correlation inferred at $1.4<z<2.5$, and their colors and morphologies are fully typical of massive $z\\sim2$ galaxies. Star forming galaxies in that redshift range with similar stellar masses are common in the distant Universe, with space densities of order of $10^{-4}$~Mpc$^{-3}$, quite a bit larger than the $2-6\\times10^{-6}$~Mpc$^{-3}$ observed for SMGs by Chapman et al. (2003). Clearly, these results point to the existence of a much larger population of gas rich galaxies in the distant Universe than previously thought.  We notice that the far-IR luminosity estimates for our objects are based on the UV, radio and mid-IR luminosities, and are solid given that self-consistent measurements of SFR activity have been shown to be obtainable for these kind of objects at high redshift (D07a). It is very unlikely that we are strongly underestimating their star formation activity (which would lead us to overestimate the duration timescales). Instead, the $L_{\\rm FIR}$ for SMGs from Greve et al. (2005) is based only on the $850\\mu$m flux and is likely more uncertain. Using radio and Spitzer, Pope et al. (2006) suggest lower $L_{\\rm FIR}$ for SMGs. This could imply that SMGs are closer in their SFR modes to massive $BzK$ galaxies, but would change little about our conclusion because it is the latter which represent the dominating star formation mode at high redshift.  The large molecular gas fraction in these galaxies, possibly up to 50\\% or more, would likely give rise to large scale instabilities, with the formation of clumps (e.g., Immeli et al. 2004; Bournaud, Elmegreen \\& Elmegreen 2007). This could explain the clumpy structure of these two galaxies, and in general of a large fraction of the high-$z$ massive galaxies (e.g., Daddi et al. 2004a), and the high velocity dispersions observed in many of the high-$z$ massive disk galaxies (e.g., Genzel et al. 2006; Forster-Schreiber et al. 2006).   The long duration timescale for star formation activity implies, as a corollary, long duration activity for the galactic nuclei. In fact, it has been recently unveiled by D07b that up to 20-50\\% of massive galaxies in $1.4<z<2.5$ likely contain strongly obscured but relatively luminous AGNs. Some of the large  gas reservoir is likely slowly funneled, in some way, into the central black hole with similar timescales to star formation, at a rate roughly consistent with what is required to assemble the correlation between stellar and black hole masses as we observe today (D07b). Our results underline the major role of gas consumption over long timescales and with low efficiencies, as opposed to rapid and strong merger driven bursts, as a major growth mode for both stellar mass and black holes in the distant Universe.    \\smallskip We thank Roberto Neri and Arancha Castro-Carrizo for assistence with taking and reducing IRAM Plateau de Bure Interferometer data. We  acknowledge very useful discussions, at various stages of this work, with Kartik Sheth, Frederic Bournaud, Pierre-Alain Duc and Fabian Walter. Comments by an anonymous referee improved the paper. We acknowledge the use of GILDAS software (http://www.iram.fr/IRAMFR/GILDAS). Support for this work, part of the Spitzer Space Telescope Legacy Science Program, was provided by NASA, Contract Number 1224666 issued by the JPL, Caltech, under NASA contract 1407."
    },
    "0711/0711.1865_arXiv.txt": {
        "abstract": "Young stars are expected to gain most of their mass by accretion from a disk that forms around them as a result of angular momentum conservation in the collapsing protostellar cloud. Accretion initially proceeds at high rates of $10^{-6}-10^{-5}$ M$_\\odot$ yr$^{-1}$ resulting in strong irradiation of the stellar surface by the hot inner portion of the disk and leading to the suppression of the intrinsic stellar luminosity. Here we investigate how this luminosity suppression affects evolution of the protostellar properties. Using simple model based on the energy balance of accreting star we demonstrate that disk irradiation causes only a slight increase of the protostellar radius, at the level of several per cent. Such a weak effect is explained by a minor role played by the intrinsic stellar luminosity (at the time when it is significantly altered by irradiation) in the protostellar energy budget compared to the stellar deuterium burning luminosity and the inflow of the gravitational potential energy brought in by the freshly accreted material. Our results justify the neglect of irradiation effects in previous studies of the protostellar growth via disk accretion. Evolution of some other actively accreting objects such as young brown dwarfs and planets should also be only weakly sensitive to the effects of disk irradiation. ",
        "introduction": "\\label{sect:intro}   It is currently established that circumstellar disks are quite ubiquitous around young stellar objects of all masses (Muzerolle \\etal 2003; Cesaroni \\etal 2007). They represent an important ingredient of the star formation since initially protostars must be growing predominantly by accretion through the disk: angular momentum conservation forces the infalling protostellar cloud material to form a centrifugally supported disk which then accretes onto a star. It is quite likely that only a small fraction of the final stellar mass gets acquired by the direct infall onto the protostellar surface, so that almost all of the stellar mass gets processed through the disk.  This picture of star formation implies very high initial mass accretion rates in the disk, at the level of $10^{-6}-10^{-5}$ M$_\\odot$ yr$^{-1}$ as the Solar-type stars are thought to gain most of their mass during the first several $10^5$ yr. The presence of such a high-$\\dot M$ accretion flow just outside the protostar immediately raises an issue of its possible effect on the protostellar properties.  There are several ways in which disk accretion affects the protostar. First, star gains mass from the disk which increases its binding energy and tends to make the star more compact. Second, accreting gas brings in some amount of thermal energy with it which contributes to the pressure support in the star. The exact amount of heat advected into the star with the accreted material is unknown but it seems likely that because of the disk geometry the accreted gas would have enough time to radiate away most of its thermal energy and would join the convective interior of the star with temperature much smaller than the stellar virial temperature.  Third, intense energy dissipation taking place in the innermost parts of the accretion disk leads to strong irradiation of the stellar surface by the disk (Frank \\& Shu 1986; Popham 1997). It has been recently realized (Rafikov 2007) that irradiation can act to suppress the internal luminosity of the protostar similar to the suppression of the cooling of hot Jupiters by the radiation of their parent stars (Guillot \\etal 1996; Burrows \\etal 2000; Baraffe \\etal 2003; Chabrier \\etal 2004). Disk accretion at rates $\\dot M \\sim 10^{-6}-10^{-5}$ M$_\\odot$ yr$^{-1}$ can easily reduce internal stellar luminosity by a factor of several which may have important implications for the early stellar evolution.  The goal of the present paper is to assess these implications by following the evolution of accreting protostar properly taking into account effects of disk irradiation. ",
        "conclusions": "\\label{sect:disc}   Weak sensitivity of $R$ to disk irradiation given the strong effect that irradiation has on the internal luminosity $L$ of a protostar may seem surprising. However, one should bear in mind that besides $L$ there are other contributors to the stellar energy budget, namely the D burning luminosity $L_D$ and the gravitational potential energy gained with the accreted material. It turns out that these contributions dominate the energy budget over $L$.  To see this we rewrite equation (\\ref{eq:master}) using definition (\\ref{eq:Lambda}) in the following form: \\ba \\frac{\\dot R}{R}=\\frac{28\\pi}{3}\\Lambda^{-1} \\frac{\\dot M}{M}\\left[\\frac{L_D}{L_0}-\\chi(\\Lambda)- \\frac{\\Lambda}{28\\pi}\\right]. \\label{eq:rewr_master} \\ea In the right-hand side of this equation the first term in brackets describes the relative role of $D$ burning in the total energy budget, second term represents intrinsic stellar luminosity, while the third term is due to the inflow of the gravitational potential energy $GM\\dot M/R$ -- the ratio of this energy inflow to $L_0$ differs from $\\Lambda$ only by a constant factor.   \\begin{figure} \\plotone{f4.eps} \\caption{ Same as Fig \\ref{fig:f2} but for $\\dot M\\propto M$ and the growth time $10^5$ yr to reach $M=1$ M$_\\odot$. \\label{fig:f4}} \\end{figure}  Suppression of $L$ is largest when $\\Lambda$ is highest which is easy to see by inspecting Figures \\ref{fig:f1}-\\ref{fig:f4}. But this automatically means that the maximum deviation of the suppression factor $\\chi$ from unity occurs precisely when the inflow of the gravitational potential energy far exceeds the stellar luminosity. Apparently, under these circumstances $L$ is a subdominant contribution to the stellar energy budget and thus even a significant reduction of $L$ compared to $L_0$ is going to be negligible compared to the gravitational energy influx.  Moreover, Figure \\ref{fig:f1} demonstrates that $\\Lambda$ reaches its maximum when $R$ is at its minimum while $f_D$ is still very close to unity. At this point vigorous D burning commences inside the star giving rise to very high $L_D/L_0$. As a result, at the evolutionary stage when $\\chi$ is minimal $L$ is subdominant in comparison to not only $GM\\dot M/R$ but also $L_D$. This additionally downplays the role of the luminosity suppression by irradiation in the early protostellar evolution.  This line of reasoning also explains why at $M\\sim 1$ M$_\\odot$ we have found $\\Delta R/R$ to be larger for lower $\\dot M$ (see \\S \\ref{sect:res}). First, smaller $\\dot M$ means lower $\\Lambda$ so that the ratio of $L$ to the gravitational energy inflow rate $GM\\dot M/R$ in the low $\\dot M$ case is larger than in the high $\\dot M$ case. Also, at $M\\sim$ M$_\\odot$ one generally finds $f_D\\ll 1$ (see Figure \\ref{fig:f1}e) so that $L_D$ is mainly due to the burning of the freshly accreted D (rather than the D that remained in the protostar from previous accretion). Since in the low $\\dot M$ case less fresh $D$ is supplied to the protostar $L_D$ must also be lower than in the high $\\dot M$ case. As a result, in the lower $\\dot M$ case $L$ plays a more significant role compared to $L_D$ (in which case $\\Delta R/R$ should be more sensitive to changes in $L$ caused by irradiation) than in the high $\\dot M$ case.  This conclusion immediately raises the following question: since $\\Delta R/R$ increases as $\\dot M$ decreases  would one find $\\Delta R/R\\sim 1$ at low enough $\\dot M$? The answer is no, and it has to do with the fact that $\\chi$ appreciably differs from unity (obviously, a necessary condition for getting $\\Delta R/R\\sim 1$) only at rather high $\\Lambda$. This is a generic feature of disk irradiation which is illustrated in Figure \\ref{fig:f5} where we display $\\Lambda_{50}$ -- the value of $\\Lambda$ at which $\\chi(\\Lambda_{50})=0.5$ -- as a function of assumed opacity law represented by the parameter $\\xi$, see equation (\\ref{eq:xi}). One can see that $\\Lambda_{50}\\gtrsim 10^2$ for all $\\xi>4$, meaning that significant luminosity suppression requires rather high $\\dot M$. This inefficiency of irradiation in suppressing $L$ is caused by the specific geometry of disk irradiation in which the irradiation flux is a very sensitive function ($\\propto \\theta^5$) of the latitude at the stellar surface $\\theta$, see Rafikov (2007). Because of that stellar polar caps can stay cool even at rather high $\\dot M$ allowing unsuppressed flux to be emitted over a significant portion of the stellar surface.  In the case of $\\xi=6.5$ as appropriate for cool, low-mass protostars one finds that $\\Lambda_{50}=2.2\\times 10^3$ which according to equation (\\ref{eq:Lambda}) immediately implies that $GM\\dot M/R\\approx 175 L_0$ when $\\chi=0.5$. Clearly, in this case stellar luminosity should have small effect on the protostellar evolution. If $\\dot M$ is reduced\\footnote{At small enough $\\dot M$ ($\\dot M\\lesssim 10^{-6}$ M$_\\odot$ yr$^{-1}$) protostellar magnetic field is likely to disrupt accretion flow (K\\\"onigl 1991) thus invalidating our assumption of the direct disk accretion onto the star.} so that $GM\\dot M/R\\sim L_0$ (and $\\Lambda\\sim 1$) stellar luminosity would be playing a significant role in the stellar energy budget, however $\\chi$ would be very close to unity (see Rafikov 2007) and the $L$ suppression by irradiation would be negligible. Thus, under no circumstances should one expect $\\Delta R/R$ larger than several per cent, meaning that quite generally the irradiation by accretion disk is unlikely to play a significant role in the evolution of the protostellar properties.   \\begin{figure} \\plotone{f5.eps} \\caption{ Plot of $\\Lambda_{50}$ (the value of $\\Lambda$ at which the suppression factor $\\chi=0.5$) as a function of parameter $\\xi$ characterizing the opacity law in the outer layers of irradiated protostar. Note that $\\Lambda_{50}$ never decreases below several hundred. \\label{fig:f5}} \\end{figure}  Looking at this conclusion under slightly different angle, our results also imply that when considering evolution of the protostars accreting at $\\dot M\\sim 10^{-6}-10^{-5}$ M$_\\odot$ yr$^{-1}$ one may completely neglect $L$ in the stellar energy budget and still get rather decent description of the protostellar evolution.  Our results are obtained assuming a specific value of $\\xi$ typical for the low-mass stars. However, one expects our major conclusions to remain valid also in the case of other accreting objects such as young brown dwarfs and giant planets. Although these objects are likely to be characterized by values of $\\xi$ different from $6.5$, our Figure \\ref{fig:f5} clearly demonstrates that the efficient suppression of stellar luminosity requires very large $\\Lambda$ for {\\it any} value of $\\xi$. As a result, we generically expect $L$ to be a subdominant energy contributor for any actively accreting object meaning that its radius (and other properties) should be different from the radius of its non-irradiated counterpart by at most a few per cent.  Previous evolutionary studies of the protostellar assembly by accretion (Mercer-Smith \\etal 1984; Palla \\& Stahler 1992; Siess \\& Forestini 1996; Hartmann \\etal 1997; Siess \\etal 1997, 1999) have always neglected the effect of disk irradiation on stellar properties. Results of this work demonstrate that this simplification should not have led to major deviations from reality since disk irradiation affects stellar properties only at the few per cent level. Protostellar evolution may more sensitively depend\\footnote{As long as $\\alpha$ is not very small. Whether this is true is debatable (Popham 1997), and we assume $\\alpha=0$ in this work.} (Prialnik \\& Livio 1985; Hartmann \\etal 1997) on the exact amount of thermal energy that gets accreted with the disk material which is parametrized by the parameter $\\alpha$, see equation (\\ref{eq:alpha}). Disk irradiation is likely to play a very important role is setting the value of $\\alpha$ since the conditions in the external radiative zone near the stellar surface regulate the thermodynamic properties of the gas at the convective-radiative boundary of the protostar, which is likely to be crucial in determining $\\alpha$."
    },
    "0711/0711.0925_arXiv.txt": {
        "abstract": "Motivated by the recent discovery of 30 new millisecond pulsars in Terzan~5, made using the Green Bank Telescope's S-band receiver and the Pulsar Spigot spectrometer, we have set out to use the same observing system in a systematic search for pulsars in other globular clusters. Here we report on the discovery of five new pulsars in NGC~6440 and three in NGC~6441; each cluster previously had one known pulsar. Using the most recent distance estimates to these clusters, we conclude that there are as many potentially observable pulsars in NGC~6440 and NGC~6441 as in Terzan~5. We present timing solutions for all of the pulsars in these globular clusters. Four of the new discoveries are in binary systems; one of them, PSR~J1748$-$2021B (NGC~6440B), has a wide ($P_b\\,=\\,20.5\\,$d) and eccentric ($e\\,=\\,0.57$) orbit. This allowed a measurement of its rate of advance of periastron, $\\dot{\\omega}\\,=\\,0.00391(18)^\\circ \\rm yr^{-1}$. If due to the effects of general relativity, the total mass of this binary system is $2.92\\,\\pm\\,0.20\\,M_{\\sun}$ (1 $\\sigma$), implying a median pulsar mass of $2.74\\,\\pm\\,0.21\\,M_{\\sun}$. There is a 1\\% probability that the inclination is low enough that pulsar mass is below $2\\,M_{\\sun}$, and 0.10\\% probability that it is between 1.20 and 1.44~$M_{\\sun}$. If confirmed, this anomalously large mass would strongly constrain the equation of state for dense matter. The other highly eccentric binary, PSR~J1750$-$37A, has $e\\,=\\,0.71$, and $\\dot{\\omega}\\,=\\,0.0055(3)^\\circ \\rm yr^{-1}$, implying a total system mass of $1.97\\,\\pm\\,0.15\\,M_{\\sun}$ and, along with the mass function, maximum and median pulsar masses of 1.65 and 1.26\\,$M_{\\sun}$ respectively. ",
        "introduction": "\\label{sec:intro}  A recent survey of the globular cluster (GC) Terzan~5 with the S-band (1.7 to 2.6 GHz) receiver of the Green Bank Telescope (GBT) and the Pulsar Spigot spectrometer \\cite{kel+05} has met with extraordinary success, finding 30 new millisecond pulsars (MSPs). The first 21 discoveries were announced in Ransom et al. (2005)\\nocite{rhs+05}. More recently, nine more pulsars were discovered\\footnote{See \\url{http://www2.naic.edu/$\\sim$pfreire/GCpsr.html}}, including PSR~J1748$-$2446ad, the fastest spinning pulsar known \\cite{hrs+06}. These discoveries show that this observing system (henceforth GBT/S/PS) has an unprecedented sensitivity to MSPs outside the Arecibo sky, due in great part to the factor of 3 higher gain of the GBT as compared to the Parkes telescope.  Perhaps even more importantly, the higher observing frequency ($\\sim$2\\,GHz), large observing bandwidth ($600$\\,MHz), and relatively fine frequency resolution ($\\sim$0.78\\,MHz) retain high sensitivity and time resolution for MSPs with large dispersion measures (DM).  The discovery of the 1.39\\,ms PSR~J1748$-$2446ad at a DM of 235.6 cm$^{-3}$pc dramatically demonstrates the sensitivity of the GBT/S/PS system to fast-spinning pulsars at high DMs. This motivated us to search for MSPs in other promising GCs with the GBT/S/PS system.   \\begin{deluxetable*}{ l r c r c c r r c} \\scriptsize \\tablecolumns{8} \\tablewidth{0pc} \\tablecaption{The six globular clusters with highest $\\Gamma_c$} \\tablehead{ \\colhead{Globular} & \\colhead{$D$\\tablenotemark{a}} & \\colhead{$\\rho_0$\\tablenotemark{b}} & \\colhead{$\\Gamma_c$} & \\colhead{DM\\tablenotemark{c}} & \\colhead{$\\tau_{\\rm scatt, 2 GHz}$\\tablenotemark{d}} & \\colhead{$N_{\\rm psr}$\\tablenotemark{e}} & \\colhead{$\\bigtriangleup T$\\tablenotemark{f}} & \\colhead{DM Ref.} \\\\ \\colhead{Cluster} & \\colhead{(kpc)} & \\colhead{(log [L$_{\\sun}$\\,pc$^{-3}$])} & \\colhead{(\\%)} & \\colhead{(cm$^{-3}$pc)} & \\colhead{($\\mu$s)} & \\colhead{} & \\colhead{(hr)} & \\colhead{\\tablenotemark{g}} } \\startdata NGC 6388 &    10.0   &  5.34  &  9.5 & 318 & 1.4   & 0  & 2.2 & \\\\ Terzan 5 &     10.3  &  5.06  &  8.6 & (234.3--245.6) & 6 & 3 $+${\\bf 30} & 8.0 & (1)\\\\ NGC 6441 &    11.7   &  5.25  &  8.0 & (230.1--234.4) & 2  & 1 $+${\\bf 3}  & 5.3 & (2)\\\\ NGC 6440 &     8.4   &  5.28  &  6.4 & (219.4--227.0) & 5 & 1 $+${\\bf 5}  & 8.4 & (3)\\\\ NGC 6266 &      6.9  &  5.14  &  5.1 & (113.4--115.0) & 4  & 6  &   7.0 & (4)\\\\ Liller 1 &      9.6  &  5.53  &  4.2 & 771 & $\\sim 3 \\times 10^2$ & 0  & 6.2  & \\\\ \\enddata \\tablenotetext{a}{Distance from the Solar System from Harris (1996). These values have significant systematic uncertainties.} \\label{tab:clusters} \\tablenotetext{b}{Central density of GC; see last column of third table in \\url{http://physwww.mcmaster.ca/$\\sim$harris/mwgc.dat}.} \\tablenotetext{c}{For GCs with known pulsars, we present the DM range for the known pulsars, with references in the last column, otherwise we predict the DM for the GC's estimated distance using the Cordes \\& Lazio (2001) model. These can be off by a factor of up to about 2.} \\tablenotetext{d}{Extrapolated from other frequencies using $\\tau_{\\rm scatt, 2 GHz} \\, = \\, \\tau_{\\rm scatt}(F) (F / {\\rm 2 GHz})^{4.4}$. In the case of Terzan 5, the scattering was measured by Nice \\& Thorsett (1992)\\nocite{nt92} as 700\\,$\\mu$s at 685\\,MHz. Otherwise, we scale the scattering time at 1GHz predicted by the Cordes \\& Lazio (2001) model for the DM in the previous column. These can be off by up to an order of magnitude.} \\tablenotetext{e}{Number of previously known pulsars $+$ {\\bf number of new pulsars}.} \\tablenotetext{f}{Amount of time the GC is visible at the GBT per day} \\tablenotetext{g}{(1) Ransom et al. 2005\\nocite{rhs+05}, (2) Possenti et al. 2006\\nocite{pcm+06}, this work, (3) Lyne, Manchester \\& D'Amico 1996, this work, (4) Possenti et al. 2003\\nocite{pdm+03}, Chandler 2003.} \\end{deluxetable*}  \\begin{deluxetable*}{ l c c c} \\scriptsize \\tablecolumns{4} \\tablewidth{0pc} \\tablecaption{Parameters for three globular clusters} \\tablehead{ \\colhead{ } & \\colhead{NGC 6388} & \\colhead{NGC 6440} & \\colhead{NGC 6441}} \\startdata Right Ascension of center, $\\alpha$ \\dotfill & $17^{\\rm h}36^{\\rm m}17\\fs 0$ & $17^{\\rm h}48^{\\rm m}52\\fs 7$ & $17^{\\rm h}50^{\\rm m}12\\fs 9$ \\\\ Declination of center, $\\delta$ \\dotfill & $-44^\\circ 44\\arcmin 06\\arcsec$ & $-20^\\circ 21\\arcmin 37\\arcsec$ & $-37^\\circ 03\\arcmin 05\\arcsec$ \\\\ Galactic longitude, $l$ ($^\\circ$) \\dotfill & 345.56 & 7.73 & 353.53 \\\\ Galactic latitude, $b$ ($^\\circ$) \\dotfill & $-$6.74 & 3.80 & $-$5.01 \\\\ Cluster distance, $D$ (kpc) \\dotfill & 10.0 & 8.2\\tablenotemark{a} & 13.5\\tablenotemark{a} \\\\ Angular core radius, $\\theta_c$ (arcmin) \\dotfill & 0.12 & 0.13 & 0.11 \\\\ Half-mass radius, $\\theta_h$ (arcmin) \\dotfill & 0.67 & 0.58 & 0.64 \\\\ Metallicity ($\\log$[Fe/H]) \\dotfill & $-0.60$ & $-0.34$ & $-0.53$ \\\\ $v_z(0)$ (km s$^{-1}$) \\dotfill & 19.13\\tablenotemark{b} & 13.01\\tablenotemark{b} & $19.5^{+9.4}_{-5.1}$\\tablenotemark{c} \\\\ $a_G$\\tablenotemark{d} ($10^{-9}$m\\,s$^{-2}$) \\dotfill & $-$0.46 & $-$0.029 & $-$0.46 \\\\ \\enddata \\tablecomments{Parameters as in Harris (1996), except where indicated.} \\tablenotetext{a}{Valenti et al. (2007).\\nocite{vfo07}} \\tablenotetext{b}{Webbink (1985).\\nocite{web85}} \\tablenotetext{c}{Dubath et al. (1997).\\nocite{dmm97}} \\tablenotetext{d}{Calculated using the data above and the Kuijken \\& Gilmore (1989) model of the Galaxy.} \\label{tab:GC} \\end{deluxetable*}  \\subsection{Cluster selection}  In the high stellar density environments in the cores of GCs, binary systems with at least one main sequence (MS) star are occasionally disrupted by the intrusion of a neutron star (NS).  The most likely outcome is an ``exchange encounter'', where the lightest component of the previous binary is ejected and a NS-MS binary remains.  The resulting NS-MS systems can become much more numerous than those formed from evolution of primordial binaries.  With the evolution of the MS star, it can happen that matter will accrete onto the NS during a Low-Mass X-ray Binary (LMXB) phase, particularly among the more compact systems. When the LMXB accretion stops, we are left with an MSP orbiting a white dwarf (WD) star.  Pooley et al.~(2003)\\nocite{pla+03} showed a strong correlation between the number of X-ray point sources in a GC (which are generally compact binaries) and the rate of stellar encounters within the half-mass radius, a parameter that they designate $\\Gamma$. This correlation is stronger than for any other individual GC parameter, and applies to active LMXBs as well (although with only a dozen systems in the Galactic GC system, the statistics are poorer). $\\Gamma$ might therefore be a good indicator of the present rate of MSP formation. Pooley et al.~(2003) provide a simple way of calculating the stellar interaction rate at the core of any GC: \\begin{equation} \\label{eq:gamma} \\Gamma_c \\propto \\rho_0^{1.5} r_c^2, \\end{equation} where $\\rho_0$ is the central mass density and $r_c$ is the core radius. This parameter is perhaps more relevant than $\\Gamma$ for the formation of LMXBs because mass segregation constrains neutron stars to within a few core radii of the center of the GC.  We calculated $\\Gamma_c$ for all Galactic GCs with the central luminosity densities and core radii given in Harris~(1996)\\nocite{har96}\\footnote{See http://physwww.mcmaster.ca/$\\sim$harris/mwgc.dat for an updated list of GC parameters.}, and assuming a constant conversion factor between central luminosity density and central mass density. The six best GCs are listed in Table \\ref{tab:clusters}, with their $\\Gamma_c$ presented as a percentage of the total $\\Gamma_c$ summed over all Galactic GCs. This table represents our initial survey expectations; we hypothesized that the percentages should be closely related to the fraction of the total GC MSP population contained in each GC.  For the first phase of the survey, we selected the GCs in Table \\ref{tab:clusters} that had not been previously searched with the GBT/S/PS system: NGC~6388, NGC~6440 and NGC~6441.  Liller~1 was searched at the same time as Terzan~5 by Ransom et al.~(2005) and NGC~6626 has been searched by Jacoby et al.~(2002)\\nocite{jcb+02}. Despite the large predicted stellar interaction rates, only two pulsars were known in these GCs: PSR~B1745$-$20A (NGC~6440A; Lyne et al.~1996)\\nocite{lmd96} and PSR~J1750$-$37A (NGC~6441A; Possenti et al.~2006)\\nocite{pcm+06}. The spin periods and DMs of NGC~6440A and NGC~6441A are 288.6\\,ms and 111.6\\,ms and 220 and 234~cm$^{-3}$\\,pc, respectively.  We assumed that the lack of other known pulsars in these high-DM GCs was largely due to a bias of the earlier surveys against the detection of faster-spinning MSPs. The high central frequency and improved time resolution of the GBT/S/PS system greatly diminish pulse broadening due to multi-path propagation and dispersive smearing, and would therefore alleviate the bias against the detection of other MSPs in these GCs.  \\subsection{Data taking} \\label{sec:datataking}  From 2005 April 18 to August 21, we observed NGC~6388, NGC~6440 and NGC~6441 on three different occasions each, using the GBT's S-band receiver with the Pulsar Spigot. The longest integration times used for each GC are very close to their maximum track times indicated in Table~\\ref{tab:GC}.  These observations take advantage of the relatively radio-frequency-interference (RFI) clean band between 1650 and 2250\\,MHz. The Pulsar Spigot is built around the GBT correlation spectrometer and is available in a number of modes \\cite{kel+05}. For most observations, we used the Spigot in mode 2 where two orthogonal polarizations with a total bandwidth of 800\\,MHz are digitized with three levels, autocorrelated with 1024 lags, and integrated for 81.92\\,$\\mu$s. The resulting autocorrelation functions are summed and then written to disk as 16-bit integers. After Fourier transforming the lags, we have 768 usable channels across the 600-MHz bandpass (0.78125-MHz channels). At a frequency of 2\\,GHz and a DM of 220~cm$^{-3}$\\,pc, the dispersive smearing across one channel is 178\\,$\\mu$s. Recently, we have been using mode 14, this mode writes the 81.92-$\\mu$s integrations to disk as 8-bit integers, but with twice the frequency resolution, resulting in a dispersive smearing across one channel for a DM of 220~cm$^{-3}$\\,pc of 89\\,$\\mu$s.  In the case of NGC~6440 and NGC~6441, the discovery of new pulsars lead us to observe these two GCs on a total of 37 and 24 extra occasions respectively, from 2005 November 5 to 2007 March 28 (Spigot mode 2), and then from 2007 October 10 to November 1 (Spigot mode 14), the latter data has only been added to the ephemeris of NGC~6440B (see \\S \\ref{sec:timing}).  The typical observing session spends the first 2-3 hours on NGC~6440, then observes NGC~6441 for 2.5 to 4 hours, and then spends 2 to 3 hours again on NGC~6440. The total usable observing times are about 90--100\\,hr for each pulsar in NGC~6440 and NGC~6441.  On two occasions, we observed NGC~6440 and NGC~6441 at 820\\,MHz with a bandwidth of 50\\,MHz in order to determine the DMs of the pulsars more accurately and to search for pulsars with spectral indices too steep to be detected at 1950\\,MHz, or that lie outside the area covered by the S-band beam. Because the telescope beam widths at these frequencies (6\\farcm5 at 1.95 GHz and 15\\arcmin \\ at 820~MHz) are significantly larger than the half-mass radii of these GCs (see Table~\\ref{tab:GC}), we observed only single positions centered on the GCs.  \\begin{figure*}[htp] \\begin{center} \\includegraphics[width=7in,angle=0]{fig1.eps} \\figcaption{Average 1950 MHz pulse profiles for the 10 pulsars known in the GCs NGC~6440 and NGC~6441, obtained by adding all observations made with mode 2 of the Pulsar Spigot. These profiles cover one full rotation. The pulsars in binary systems have an asterisk after their names, and the newly discovered pulsars have their names in boldface. The horizontal width of the rectangles indicates the system's total time resolution in mode 2, including the effects of dispersive smearing, relative to each pulsar's spin period. The dip in power after the main pulse of NGC~6440A is likely a Spigot artifact. \\label{fig:profiles}} \\end{center} \\end{figure*} ",
        "conclusions": "We have discovered eight new pulsars in the GCs NGC~6440 and NGC~6441. Their pseudo-luminosities indicate that there may be as many pulsars in these GCs as in Terzan~5, but less pulsars are observed in the former because of their larger distances. The pulsar population in these GCs seems to have distinctly lower spin frequencies. To some extent, this is due to the presence of apparently young objects like NGC~6440A. At least in the case of NGC~6440A and D, and possibly in the case of the other slow MSPs, the larger spin periods could be related with many of them having higher magnetic fields than the MSPs in other GCs.  Four of the new pulsars are in binary systems. NGC~6440D is an eclipsing binary, with a companion star with a mass of about 0.12~$M_{\\sun}$ or larger. NGC~6440B is a 16.7-ms pulsar in a 20.5-day orbit with an orbital eccentricity of 0.57. A measurement of the rate of advance of periastron for this system suggests that this binary hosts the most massive neutron star to date: $2.74\\,\\pm\\,0.21\\,M_{\\sun}$, with only a 1\\% probability that the inclination is low enough that pulsar mass is below $2\\,M_{\\sun}$. This result depends on the (likely) possibility that the observed periastron advance is fully relativistic. Such a large mass would introduce the strongest constraints to date on the equation of state; it also suggests that the products of the coalescence of double neutron stars might be stable neutron stars. The rate of advance of periastron was also measured for NGC~6441A, indicating a significantly less massive pulsar ($M_p\\,<\\,1.65\\,M_{\\sun}$).  The pulsars in NGC~6440 will be monitored carefully over the next few years. In the case of NGC~6440C, the new data will allow a detailed investigation of the origin of its timing irregularities. In the case of NGC~6440B, the continued monitoring will lead to an improvement in the precision of $\\dot{\\omega}$. Simulations indicate that the Einstein delay ($\\gamma$) may become measurable with about 15 years of data; this would provide unambiguous mass measurements and determine the nature of the companion. Deep optical studies of NGC~6440 might be extremely valuable in determining the nature of the companion of NGC~6440B."
    },
    "0711/0711.2327_arXiv.txt": {
        "abstract": "Using the HiVIS spectropolarimeter built for the Haleakala 3.7m AEOS telescope in Hawaii, we are collecting a large number of high precision spectropolarimetrc observations of stars.  In order to precisely measure very small polarization changes, we have performed a number of polarization calibration techniques on the AEOS telescope and HiVIS spectrograph.  We have extended our dedicated IDL reduction package and have performed some hardware upgrades to the instrument.  We have also used the ESPaDOnS spectropolarimeter on CFHT to verify the HiVIS results with back-to-back observations of MWC 361 and HD163296.  Comparision of this and other HiVIS data with stellar observations from the ISIS and WW spectropolarimeters in the literature further shows the usefulness of this instrument. ",
        "introduction": "High-resolution linear spectropolarimetry measures the change in linear polarization across a spectral line.  Many models show this measurement is a useful probe of circumstellar environments at small spatial scales.  Circumstellar disks, rotationally distorted winds, magnetic fields, asymmetric radiation fields (optical pumping), and in general, any scattering asymmetry can produce a change in linear polarization across a spectral line such as $H\\alpha$ (cf. McLean 1979, Wood Brown \\& Fox 1993, Harries 2000, Ignace et al. 2004, Vink et al. 2005a, Kuhn et al. 2007).  These signatures can directly constrain the density and geometry of the circumstellar material and probe the near-star environment.  Typical spectropolarimetric signals are small, often a few tenths of a percent change in polarization across a spectral line.  Measuring these signals requires very high signal to noise observations and careful control of systematics to measure signals at the 0.1\\% level.  This technique probes small spatial scales, typically being sensitive to the region within 10 stellar radii.  Even for the closest young stars (150pc), these spatial scales are smaller than 0.1 milliarcseconds across and will not be imaged directly, even by 100m telescopes.  Since the circumstellar material is involved in accretion, outflows, winds and disks, with many of these phenomena happening simultaneously, spectropolarimetry can put unique constraints on the types of densities and geometries of the material involved in these processes.  In order to address these issues, we built a spectropolarimeter for the HiVIS spectrograph on the 3.7m AEOS telescope and performed a preliminary telescope polarization calibration (Thornton 2002, Thornton et al. 2003, Harrington et al. 2006).  Since then, we have made many improvements.  We have upgraded much of the hardware, calibrations, and reduction package.  A large number of stellar observations have been partially reported in Harrington \\& Kuhn 2007 to illustrate the capability of the spectropolarimeter.  We present the reduction package, polarization calibrations, and implications of these calibrations for stellar observations with this instrument. ",
        "conclusions": "We have presented our dedicated IDL-based reduction package and the polarization properties of the HiVIS spectropolarimeter.  This package compares very favorably with other dedicated reduction packages.  The Esprit package on ESPaDOnS and our scripts give very similar results when applied to ESPaDOnS data.  We have found that the telescope induces 0.5\\% to 3.0\\% continuum polarization in an unpolarized source with a strong dependence on secondary focus.  The telescope can also strongly depolarize a source or rotate the plane of polarization as was shown by highly-polarized twilight observations.  All of these are functions of altitude, azimuth, and wavelength.  Optical ray trace models with Zeemax were presented to illustrate the pointing dependence of the telescope polarization properties, but were ultimately shown to be quantitatively inaccurate due to the very common problem of oxide coatings.  While these telescope effects complicate the analysis of data taken with this instrument, useful information about the amplitude and wavelength dependence of the polarization across a spectral line can still be obtained.  The performance of the instrument was compared to other spectropolarimeters, particularly ESPaDOnS on CFHT, ISIS on WHT, and WWS on OMM.  HiVIS found to reproduce the spectropolarimetric signatures, sometimes strikingly well given the possible complications.  This instrument will be able to make sensitive spectropolarimetric observations, despite some complications in the interpretation of the polarization after so many oblique reflections.  \\subsection{Acknowledgements} This program was supported by the NSF AST-0123390 grant, the University of Hawaii and the AirForce Research Labs (AFRL).  This program made use of observations obtained at the Canada-France-Hawaii Telescope (CFHT) which is operated by the National Research Council of Canada, the Institut National des Sciences de l'Univers of the Centre National de la Recherche Scientifique of France, and the University of Hawaii.  This program also made use of the Simbad data base operated by CDS, Strasbourg, France.  The authors also wish to thank Don Mickey for many useful discussions about polarimetry and optical design."
    },
    "0711/0711.0322_arXiv.txt": {
        "abstract": "%  We present preliminary results from Smooth Particle Hydrodynamics (SPH) simulations of common envelope evolution. We qualitatively compare the interaction between a $0.9M_{\\sun}$ red giant with two different companion masses: a $0.05M_{\\sun}$ brown dwarf and a $0.25M_{\\sun}$ white dwarf companion. ",
        "introduction": "Most stars (60$\\%$) are part of binary or multiple systems. If the binary separation is small enough, the more massive star will eventually engulf the companion as it expands during its giant phase, and form a so-called common envelope (CE) system. Friction then transfers orbital energy and angular momentum into the envelope, which may eventually lead to ejection of the entire envelope. Most tight binary systems are believed to have been in a CE phase at one point in their life: high- and low-mass X-ray binaries, double degenerate white dwarf and neutron star binaries, cataclysmic variables and supersoft X-ray sources. Many stellar types -- such as hydrogen deficient stars, early R-stars, or Wolf-Rayet type central stars of planetary nebulae for example -- have also been proposed to be the result of a CE phase, as isolated stellar evolution models are not able to fully explain their properties \\citep[e.g.][]{DeMarcoWolfRayet, IzzardRstars}. New studies have even proposed that the majority of planetary nebulae may be the result of a CE interaction \\citep{DeMarcoPNbinaries}.  The concept of CE evolution has been around since the pioneering work of \\citet{PaczynskiCE}. Earlier simulation efforts \\citep{RasioCE, SandquistCE} were restricted to high mass companions due to the low numerical resolution achievable at that time. Only a handful low-resolution simulations are available as of now, covering only a tiny fraction of the vast parameter space available. Only very recently, \\citet{RickerCE_AMR2} have restarted simulating these dynamic events at high resolution with the adaptive mesh refinement code {\\it FLASH}, so far concentrating again on interactions with relatively high-mass companions. Theoretical considerations expect lower-mass companions (such as brown dwarfs, or even Jupiter-mass planets) to penetrate much deeper into the giant's envelope or even merge into the core, requiring a much higher numerical resolution at the center, \\citet[see e.g.][]{DeMarcoWolfRayet2}. Yet, these binaries should be much more common and may easily prove powerful enough to shed the rather losely bound envelope of a red giant (RG) or asymptotic giant branch (AGB) star. Here we will present a qualitative comparison between two CE evolution simulations involving a $0.9M_{\\sun}$ red giant interacting with a low ($0.05M_{\\sun}$ brown dwarf, BD) and medium mass ($0.25M_{\\sun}$ white dwarf, WD) companion. The BD-RG system was chosen to be the progenitor system of a recently discovered tight binary system between a $0.05M_{\\sun}$ BD and a $0.39M_{\\sun}$ WD \\citep{MaxtedRGBD}. Our RG was chosen such that its degenerate core matches this WD mass. We will use this system as a validation case for our code, and thus expect the BD to spiral in as far as the observed $0.65R_{\\sun}$ separation.     \\begin{figure}[t] \\hfil \\includegraphics[width=0.25\\textwidth]{WD025_000} \\includegraphics[width=0.25\\textwidth]{WD025_040} \\includegraphics[width=0.25\\textwidth]{WD025_080} \\hfil \\\\ \\null\\hfil \\includegraphics[width=0.25\\textwidth]{WD025_120} \\includegraphics[width=0.25\\textwidth]{WD025_160} \\includegraphics[width=0.25\\textwidth]{WD025_200} \\hfil \\caption{Time sequence of SPH particle plots for a common envelope between a $0.9M_{\\sun}$ red giant and a $0.25M_{\\sun}$ white dwarf. The small white dots show individual SPH prarticles, while the grey filled circles show the degerate core of the red giant (center) and the companion (right, on the surface). The individual plots are spaced 40 days apart from 0 to 200 days. Note how the companion transfers energy and angular momentum into the envelope, which is transported outward in a spiral density wave. At the end of this simulation, the binary has almost stalled its evolution, as there is only little material left at its radius to interact with, and most of the envelope corotates with the companion. \\label{f.WD025}} \\end{figure} ",
        "conclusions": ""
    },
    "0711/0711.2866_arXiv.txt": {
        "abstract": "{ We report on the diffuse X-ray emissions from the Galactic center (GCDX) observed with the X-ray Imaging Spectrometer (XIS) on board the Suzaku satellite. The highly accurate energy calibrations and extremely low background of the XIS provide many new facts on the GCDX. These are (1) the origin of the 6.7/7.0~keV lines is collisional excitation in hot plasma, (2) new SNR and super-bubble candidates are found, (3) most of the 6.4~keV line is fluorescence by X-rays, and (4) time variability of the 6.4~keV line is found from the Sgr~B2 complex. } ",
        "introduction": "The iron K-shell complex in the Galactic center diffuse X-rays (GCDX) is composed of four lines at 6.4, 6.7, 7.0 and 7.1~keV (Koyama et al. 1986). These are K$\\alpha$ lines of neutral(Fe\\textsc{i}), He-like (Fe\\textsc{xxv}) and H-like (Fe\\textsc{xxvi}) iron, and K$\\beta$ line of neutral iron (Fe\\textsc{i}). The origin of these iron lines or the origin of the GCDX has been a big problem. In order to solve the problem, high quality data of the iron K-shell lines are essentially important. Using the X-ray Imaging Spectrometer (XIS) (Koyama et al. 2007a) on board Suzaku (Mitsuda et al. 2007), we have obtained accurate X-ray spectra in the hard X-ray band including these key lines. This paper reports on the line analysis and discussion on these observational data. ",
        "conclusions": ""
    },
    "0711/0711.2261_arXiv.txt": {
        "abstract": "The spectral energy distributions (SEDs), spanning the mid-infrared to millimeter wavelengths, of a sample of 13 high-mass protostellar objects (HMPOs) were studied using a large archive of 2-D axisymmetric radiative transfer models. Measurements from the Spitzer GLIMPSE and MIPSGAL surveys and the MSX survey were used in addition to our own surveys at millimeter and submillimeter wavelengths to construct the SEDs, which were then fit to the archive of models. These models assumed that stars of all masses form via accretion and allowed us to make estimates for the masses, luminosities and envelope accretion rates for the HMPOs. The models fit the observed SEDs well. The implied  envelope accretion rates are high, $\\approx 10^{-2.5} \\msun/yr$, consistent with the accretion-based scenario of massive star formation. With the fitted accretion rates and with  mass estimates of up to $\\sim 20\\ \\msun$ for these objects, it appears plausible that stars with stellar masses $M_{\\ast} > 20 \\msun$ can form via accretion. ",
        "introduction": "The process of massive star-formation is a subject of current debate (McKee \\& Ostriker, 2007; Zinnecker, 2007). While scaled up versions of the standard accretion model and competitive accretion models require high accretions rates ($\\sim 10^{-3} \\msun/yr$), competing coalescence models need high stellar densities ($\\sim 10^8 /pc^3$). Presumably, the models are applicable in different regimes, but it is unclear if there is a limit to the stellar mass in the accretion models.  The ubiquitous detection of outflows with derived high accretion rates supports an accretion picture (Churchwell 1999, Henning et al 2000, Beuther et al 2002, Zhang et al 2005). Observations of spectral line infall signatures also indicate high infall rates (Zhang \\& Ho, 1997, Keto, 2002, Fuller, Williams \\& Sridharan, 2005, Beltran et al 2006, Keto \\& Wood, 2006, Zapata et al 2008).  However, these results are subject to assumptions and questions. Therefore, other independent methods of deducing accretion and determining the accretion rates are important. In this paper, we model the SEDs of a subset of 13 HMPOs and find results that are consistent with the accretion scenario for massive star-formation.  The study takes advantage of (1) the Spitzer legacy surveys GLIMPSE and MIPSGAL to enhance wavelength coverage for the SEDs at high spatial resolutions compared to previous data from IRAS and MSX surveys and (2) new modeling capabilities. Thus, the current work is a significant improvement over previous 1-D modeling of the same objects (Williams, Fuller \\& Sridharan, 2005). ",
        "conclusions": "Combining a large archive of 2-D radiative transfer models (Robitaille et al 2006, 2007) and SEDs with wide wavelength coverage using the Spitzer GLIMPSE and MIPSGAL surveys, and the MSX and our published millimeter and sub-millimeter surveys, a carefully chosen subset of 13 HMPO candidates was modeled. The main findings are: (1) the models fit the data well with high implied envelope accretion rates ${\\dot{M}_{env}} \\approx 10^{-2.5} \\msun/yr$, required for massive star-formation, lending credence to accretion based massive star-formation. (2) stars of masses  $> 20 \\msun$ may form by accretion. We also find a possible mass dependence of the accretion rates as a power law and determine the fraction of the envelope accretion lost to the the outflow jets, both of which are subject to large uncertainties. These results can be improved by extension to larger samples using archival data and better PSF fitting to overcome saturation."
    },
    "0711/0711.4267_arXiv.txt": {
        "abstract": "{Observations of quasars shining through foreground galaxies, offer a way to probe the dust extinction curves of distant galaxies. Interesting objects for this study are found in strong gravitational lensing systems, where the foreground galaxies generate multiple images.} {The reddening law of lensing galaxies is investigated by studying the colours of gravitationally-lensed quasars, and a handful of other quasars where a foreground galaxy is detected.} {We compare the observed colours of quasars reported in the literature, with spectral templates reddened by different extinction laws and dust properties. The data consists of {\\nosys} quasar-galaxy systems, with a total of {\\noim} images. The galaxies, which are both early- and late-type, have redshifts in the interval $z=\\ ${\\gzinterv}.} {We measure a difference in rest-frame $B-V$ between the quasar images we study, and quasars without resolved foreground galaxies. This difference in colour is indicative of significant dust extinction in the intervening galaxy. Good fits to standard extinction laws were found for {\\noimwell} of the images, corresponding to {\\nosyswell} different galaxies. Our fits imply a wide range of possible values for the total-to-selective extinction ratio, $R_V$. The distribution was found to be broad with a weighted mode of $\\bar R_V$={\\ideorv} and a FWHM of $\\Delta{R{_V}}$={\\ideorvfwhm} ($\\sigma_{R_V} \\sim$ {\\ideorvsigma}). Thus the bulk of the galaxies for which good reddening fits could be derived, have dust properties compatible with the Milky Way value ($R_V=3.1$).} {} ",
        "introduction": "\\label{sec:intro}  Measuring the reddening due to dust in galaxies, provides information on dust grains in the interstellar medium. It is in addition important for dimming corrections, when measuring cosmological distances using Type Ia supernovae, the most direct probe of the expansion history of the Universe.  It is often assumed that the wavelength dependence of extinction in distant galaxies is identical to the average in the Milky Way. However, there are indications that this is not always the case \\citep[e.g.][]{prevot84,calzetti00}. The purpose of this work is to examine the reddening relations in galaxies over a wide range of redshifts, to improve our understanding of galaxy dust properties.  For galaxies in which individual stars are resolved, dust extinction can be studied using the pair method. In this method, the galaxy extinction curve is determined by comparing the spectral energy distribution of reddened and non-reddened stars of the same spectral type and luminosity class \\citep{massa83}. For more distant galaxies, other methods must be used, such as comparing the galaxy colours with a dust-free model (although this is subject to large model dependencies) \\citep[see e.g.][]{patil06}, determining the differential extinction of multiply-imaged quasars \\citep[e.g.][]{falco99}, or measuring the absolute extinction by foreground galaxies of quasars \\citep{ostman06}. There are problems using the method of differential extinction for multiply-imaged quasars. To obtain a unique $R_V$ value, one of the images must have negligible dust extinction when compared to the other, or the extinction curves must be similar along the two separate lines of sight. The potential problem of this limitation is discussed in \\citet{mcgough05} and \\citet{ardis}. If these conditions are not met, the $R_V$ value derived will be a function of the extinctions along the two lines of sight, and their $R_V$ values (see Eq.~6 in \\citet{mcgough05}). In this paper, this restriction is circumvented by studying the individual colours of background quasar images in order to investigate the reddening by foreground galaxies. This is possible because of the fairly homogeneous spectral properties of quasars. The method is very similar to the one used in \\citet{ostman06}, where the focus was on finding systems by matching the positions of quasars with galaxies. In this paper, the focus is on using systems found through gravitational lensing.  Studies of colours of Type Ia supernovae have also been used to constrain the dust properties of their host galaxies \\citep[e.g.][]{joeveer83,branch92,riess96,reindl05,guy05}. When the reddening correction has been optimised to minimise the scatter of the Type Ia supernova Hubble diagram, a quite different wavelength dependence of the extinction curve, when compared to the Milky Way, has been derived \\citep[e.g.][]{2006A&A...447...31A}. It is unclear whether these findings indicate the presence of non-standard dust along the line of sight to supernovae, or if they point to a different source of the brightness-colour relation, perhaps intrinsic to supernovae themselves. We aim to establish if the reddening laws derived using supernovae are applicable to different astrophysical scenarios, such as quasars shining through galaxies.  In Sec.~\\ref{sec:general}, different parameterisations of dust attenuation are discussed. The method for estimating the dust properties is presented in Sec.~\\ref{sec:method}. The results and a discussion of the results are given in Sec.~\\ref{sec:results}. Finally, we summarise and conclude in Sec.~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  The reddening of quasars by foreground galaxies has been investigated. At our disposal, we have had {\\nosys} quasar-galaxy systems with a total of {\\noim} quasar images. By looking at the rest-frame $B-V$, we conclude that the foreground galaxy, close to the line of sight, must affect the quasar colours. Using redshifted reddened quasar templates, we fit for $R_V$ and $E(B-V)$ using the measured colours.  Of the data that we consider, {\\noimwell} images could be fitted by dust extinction. Assuming dust extinction in the intervening galaxy following the {\\fitzpatrick} law, we find a most probable $R_V$ value of {\\ideorv}. This value is lower than the Galactic mean value of 3.1. However, the FWHM of our distribution of $R_V$ was determined to be {\\ideorvfwhm} and thus the Galactic mean value is within the range of the spread. We note that our sample will not be representative of the universal galaxy population because the majority of our quasars were observed by strong lensing surveys, which are particularly efficient in finding massive galaxies. The large spread in values indicates that the dust content can vary significantly between different galaxies. However, we found no strong correlation between the values of $R_V$ and redshift, colour excess or galaxy type. In the future, when a larger dataset with superior quality is available, a similar analysis could be of considerable benefit to, for example, the use of extinction-corrected supernovae as distance indicators. It would not be surprising if a correlation was found, because metallicity and elemental abundance ratios are expected to affect dust properties, and both quantities evolve with redshift and galaxy type. Furthermore, there are at least two major mechanisms that create dust grains, the first is in AGB stars \\citep[e.g.][]{mathis90} and the second in core-collapse supernovae \\citep[e.g.][]{bank03}, which could give rise to different types of dust.  The uncertainties in the determination of the amount of dust are large, when taking into account the different possible dust models. This fact may be of importance for flux anomalies, used as a probe of cold dark matter substructure \\citep{kochanek04}.  Another interesting observation is that when several images of the same quasar could be fitted by dust, the preferred values of $R_V$ were similar. This could indicate that $R_V$ has little spread within individual galaxies, but that it varies between galaxies. A caveat, however, is that for many of our quasars one or several of the images could not be fitted by the standard reddening laws we tried, potentially challenging the conclusion of homogeneous dust properties within individual lensing galaxies, unless the colour outliers can be attributed to something different from dust. In total, there were seven quasars in our sample that had several images, one or several of which could be fitted with dust. Three of those quasars had similar values of $R_V$, for all images, regardless of extinction law, two quasars had at least one image that could not be reproduced by assuming the presence of dust, one quasar had images where the $R_V$ values did not agree with each other, and one quasar was consistent with experiencing no dust extinction.  Several systems were found with low values of $R_V$, comparable to what has been measured from global fits, using Type Ia supernova data. However, the distribution of $R_V$ for our fits suggests that a large fraction of the galaxies considered in this study are compatible with Milky Way dust. Due to the limited number of objects in our study, as well as the possibility of selection effects and analysis bias, no strong conclusion can be drawn as to whether the extinction properties in foreground galaxies, along the line of sight to quasars, differ or not from extinction in host galaxies of Type Ia supernovae."
    },
    "0711/0711.4051_arXiv.txt": {
        "abstract": "{ Emission-lines in the form of filamentary structures is common in bright clusters characterized by short cooling times. In the Perseus cluster, cold molecular gas, tightly linked to the H$\\alpha$ filaments, has been recently revealed by CO observations. In order to understand the origin of these filamentary structures, we have investigated the evolution of the hot ICM gas perturbed by the AGN central activity in a Perseus like cluster. Using very-high resolution TreeSPH simulations combined with a multiphase model and a model of plasma bubbles, we have been able to follow the density and temperature evolution of the disturbed ICM gas around the bubbles. Our simulations show that a fraction of the $1-2\\,\\rm{keV}$ gas present at the center of clusters is trapped and entrained by the rising buoyant bubble to higher radius where the AGN heating is less efficient. The radiative cooling makes it cool in a few tens of Myr below $10^4\\,\\rm{K}$, forming cold filamentary structures in the wake and in the rim of the bubbles. The predicted cold gas formed outside the cluster center is in agreement with the total CO mass and density profile of the observed molecular gas as well as with the kinematics of the H$\\alpha$ filaments. This scenario explains the H$\\alpha$ and CO filaments observed in luminous clusters without contradicting the observed lack of $1\\,\\rm{keV}$ gas. It also emphasizes that if the AGN feedback provides some heating (negative feedback) it also perturbs the ICM, increasing its cooling (positive feedback). } ",
        "introduction": "Cooling flows in clusters of galaxies has been a puzzling topic for more than 30 years \\citep{fabian94,peterson06}. While early models of X-ray clusters \\citep{cowie77} applied to observations predict a large cold gas mass accretion rate, the recent data from Chandra and XMM-Newton ruled out such large amounts of cooling \\citep{peterson03}. There are now numerous observations showing that the AGN activity may be a clue to the cooling flow problem \\citep{fabian00,mcnamara00,owen00,mcnamara00,fabian03b,fabian03,blanton03,wise06,fabian06,mcnamara07}. The emerging picture is that cooling flows are considerably reduced by the heating provided by the AGN feedback \\citep{bruggen01,churazov01,quilis01,bohringer02,reynolds02,omma04b,omma04a,sijacki06a,heinz06,cattaneo07}. The bulk of the ICM gas remains hot ($\\ge 2\\,\\rm{keV}$) down to the cluster center, where it is accreted by the central black hole (Bondi accretion).  However, in this picture, it is difficult to understand the origin of the observed cold neutral gas \\citep{edge01,edge03,salome03,salome04,salome06} and the filamentary H$\\alpha$ structures \\citep{conselice01,crawford05}. The connection between radio source and cold CO gas \\citep{salome04} as well as the excess of optical/UV/IR light in clusters with central cooling time less than $3\\,\\rm{Gyr}$ are signs of the link between the filaments, the cooling and the AGN activity. Very deep Chandra observation of the Perseus cluster shows a clear correlation between the soft X-ray emission (between $0.5$ and $1\\,\\rm{keV}$) and the optical filamentary H$\\alpha$ structures \\citep{fabian03,fabian06}. As the filaments cover a large range of cooling time within the hot gas, it is unlikely that they form by direct cooling of the $2\\,\\rm{keV}$ component. Moreover, the kinematics of the filament in H$\\alpha$ \\citep{hatch06} and soft X-rays \\citep{fabian06} matches the gas flow under a buoyantly rising bubble. This led \\citet{fabian03,fabian06} and \\citet{hatch06} to propose that the filaments have been pulled out from a central reservoir of cold and warm gas by the action of radio plasma bubbles inflated by AGN jets.  In this letter, without answering the question of the inner quenching of the cooling flow, we present a series of very high resolution TreeSPH simulations of a Perseus type cluster including AGN feedback. We show that buoyantly rising bubbles entrain a fraction of the $1-2\\,\\rm{keV}$ central ICM gas to much higher radius where it has time to cool and fall back to the cluster center, forming cold filamentary structures ($T < 10^4$). ",
        "conclusions": "New very-high resolution SPH simulations of the Perseus cluster show that the cold gas observed in H$\\alpha$ and CO filaments outside the cluster core is a natural product of the perturbations of the ICM by buoyantly rising bubbles. During the rise of a bubble, a fraction of the low cooling time ambient ICM gas around the bubble is entrained and cools at higher radius, forming cold gas below $10^4\\,\\rm{K}$ in the trunk of the bubble. The streaching of the filaments observed in H$\\alpha$ by \\citep{hatch06} is qualitatively reproduced.  Because no central heating is taken into account, all our models are subject to the cooling catastrophe. Is the proposed scenario still valid in a model where the cooling flow is quenched by heating processes ? The role of the heating source is to suppress the formation of gas with a temperature lower than $2\\,\\rm{keV}$, but $1-2\\,\\rm{keV}$ gas will still be present in abundance in the cluster center \\citep{fabian06}. As seen in section \\ref{cold_trunk}, this is precisely this gas that is trapped by the bubbles and later forms the cold filaments. The heating at the origin of the cooling flow quenching will thus have only a very weak influence on the amount of cold gas formed. In the new emerging picture of cooling flow, AGN feedback provides some heating (negative feedback), but it is also responsible for a positive feedback, where ICM perturbations induce cold gas production. This is a step forward in the `cold feedback` scenario \\citep{soker01,pizzolato05} where the feedback takes place within a large region ($R\\lesssim 5-30\\,\\rm{kpc}$)."
    },
    "0711/0711.1959_arXiv.txt": {
        "abstract": "{We report on the first wide-field, high-precision astrometric analysis of the 13 extragalactic radio sources of the complete S5 polar cap sample at 15.4\\,GHz. We describe new algorithms developed to enable the use of differenced phase delays in wide-field astrometric observations and discuss the impact of using differenced phase delays on the precision of the wide-field astrometric analysis. From this global fit, we obtained estimates of the relative source positions with precisions ranging from 14 to 200\\,$\\mu$as at 15.4\\,GHz, depending on the angular separation of the sources (from $\\sim$1.6 to $\\sim$20.8\\,degrees). These precisions are $\\sim$10 times higher than the achievable precisions using the phase-reference mapping technique.} ",
        "introduction": "In the past few years, we have carried out a series of very-long-baseline-interferometry (VLBI) observations, using the Very Long Baseline Array (VLBA) at 8.4, 15.4, and 43\\,GHz, aimed at studying the absolute kinematics of a complete sample of extragalactic radio sources using astrometric techniques. The target of our programme is the ``complete S5 polar cap sample'', consisting of 13 radio sources from the S5 survey (K\\\"uhr et al. \\cite{Kuhr1981}, Eckart et al. \\cite{Eckart1986}). All sources in this sample have flux densities higher than 0.2\\,Jy at the epoch of our observations and well-defined {\\em International Celestial Reference Frame} (ICRF-Ext.2) positions (see Fey et al. \\cite{Fey2004}). The relative source separations, less than about 15$^{\\circ}$ for neighbouring sources, should allow for typical astrometric precisions in the range of 20 to 100\\,$\\mu$as, depending on the observing frequency. (Lower frequencies translate into lower precisions.) Even though the essence of our global differenced phase delay astrometry is the same as that of phase-reference astrometry (e.~g., Beasley \\& Conway \\cite{Beasley1995}), there are important differences between them.  On the one hand, the coordinates of one source (the target source) are determined in phase-reference astrometry with respect to the coordinates of another source (the reference source). In our global analysis, we use data from all the 13 sources of the S5 polar cap sample simultaneously, in a unique fit, which increases the precision of the astrometry by a factor $\\sim$10, as we will see in Sect.~4. On the other hand, phase-reference astrometry requires small angular separations between the sources (up to a few degrees). In our global analysis, this requirement is dramatically relaxed. Consequently, we should be able to apply a global differenced phase delay astrometric analysis, similar to the analysis reported here, to a set of sources spread across the whole sky.  We have already described the goals of our astrometric programme in Ros et al. (\\cite{Ros2001}, hereafter Paper~I) and presented VLBA maps obtained at two different epochs at 8.4\\,GHz in Paper I and at 15.4\\,GHz in P\\'erez-Torres et al. (\\cite{Perez-Torres2004}, hereafter Paper~II). In this paper, we present the first conclusive phase delay astrometric results of this programme. We describe in Sect. 3 the details of our astrometric analysis and our algorithm developed in-house from which we automatically derive the 2$\\pi$$-$ambiguities inherent in the phase delay observable. We also discuss the contribution of differenced observables in the astrometric precision for sources distributed across a large portion of the sky. ",
        "conclusions": "We report here on the first high-precision, wide-field astrometric results at 15.4\\,GHz of our multi-frequency monitoring of the S5 polar cap sample. To obtain those results we first developed the package {\\sc uvpap}, an extensively improved version of the {\\sc vlbi3} program (Robertson \\cite{Robertson1975}). The ability of {\\sc uvpap} to use differenced phase delays, along with newly developed phase-connection algorithms, enables us the use of differenced phase delays in global astrometric observations. We discuss the impact of the differenced phase delays on our global astrometric analysis and show that their use increases the precision of the relative source positions by a factor of $\\sim$10 compared with the precision achievable using the phase-reference technique with a pair of sources (Pradel, Charlot \\& Lestrade \\cite{Pradel2006}). The astrometric precisions obtained linearly decrease (from 14 to 200\\,$\\mu$as) as the separations between the sources increase (from $\\sim$1.6 to $\\sim$20.8\\,degrees), with the result that the fractional errors in determining the separations of all the studied source pairs are similar ($\\sim$3$\\times10^{-9}$).  We obtain some large corrections for the relative coordinates and separations of the sources. From all the 24 pairs studied, 10 have separation corrections above 500\\,$\\mu$as and, of those, 4 have corrections above 900\\,$\\mu$as. These corrections could be caused by opacity effects (our observations are at 15.4\\,GHz, and the ICRF-Ext.2 positions are based on 8.4 and 2.3\\,GHz observations) and by source structure effects (we are relating our astrometric positions to the phase centres of the source maps, that is, the peaks of brightness; the ICRF-Ext.2 positions are not well-defined on the source structures). Other wide-field, high-precision astrometric analyses of these sources at other frequencies and other epochs are currently under way. This multi-epoch and multi-frequency study will eventually provide spectral information and the absolute kinematics for all sources in the sample. Ultimately, we expect to provide a definitive test of the stationarity of the innermost radio-source cores, associated to the massive black holes, which is a basic tenet of the standard jet interaction model (Blandford \\& K\\\"onigl \\cite{Blandford1979}). In addition, our results will be an excellent complement to future $\\mu$as-precise astrometry at optical wavelengths."
    },
    "0711/0711.0996_arXiv.txt": {
        "abstract": "As in the case of the other gauge field theories, there is so called ``gauge'' also in general relativity. This ``gauge'' is unphysical degree of freedom. There are two kinds of ``gauges'' in general relativity. These are called the first- and the second-kind of gauges, respectively. The gauge of the first kind is just coordinate system on a single manifold. On the other hand, the gauge of the second kind arises in the general relativistic perturbations. Through the precise distinction of these two concepts of ``gauges'', we develop second-order gauge-invariant general relativistic perturbation theory. ",
        "introduction": "General relativity is regarded as a gauge theory. In 1956, Uchiyama pointed out that gravitational field is introduced by the invariance of the action under local Lorentz transformations\\cite{Uchiyama1956}. This is due to general covariance in general relativity. Since this local Lorentz group is a group of coordinate transformations, coordinate system is called ``gauge'' in general relativity from this history.     General covariance intuitively states that there is no preferred coordinate system in nature and it also introduce ``gauge'' in the theory. This ``gauge'' is the unphysical degree of freedom and we have to fix the ``gauge'' or to extract some invariant quantities to obtain physical results. Thus, treatments of ``gauge'' are crucial in general relativity. This situation becomes more delicate in general relativistic perturbation theory.     On the other hand, the general relativistic higher-order perturbation theory is a topical subject in recent physics, for example, cosmological perturbations, perturbations of a black hole and a star. So, it is necessary to formulate the higher-order general relativistic perturbation theory form general point of view.     In this article, we clarify the notion of ``gauges'' in general relativity, which is necessary to develop general relativistic higher-order ``gauge-invariant'' perturbation theory. The details of this perturbation theory can be seen in Ref.\\cite{KNs}. ",
        "conclusions": "We have shown the general procedure to find gauge-invariant variables in the second-order general relativistic perturbation theory through the precise treatments of ``gauges''. We also showed that this general procedure is applicable to cosmological perturbations and developed the second-order cosmological perturbation theory in gauge invariant manner\\cite{KNs}. Due to the general covariance in general relativity, all equations are given in terms of gauge invariant variables. We are planning to apply this second-order perturbation theory to clarify the non-linear physics in Cosmic Microwave Background\\cite{WMAP}.     Besides the application to cosmology, we are also planning to apply the above general framework to black hole perturbations, perturbations of a star, and post-Minkowski description of a binary system. In conclusion, the definitions (\\ref{eq:matter-gauge-inv-def-1.0}) of gauge invariant variables have very wide applications."
    },
    "0711/0711.4466_arXiv.txt": {
        "abstract": "%  We studied the bulge and the disk kinematic of the giant Low Surface Brightness (LSB) galaxy ESO 323-G064 in order to investigate its dynamics and its Dark Matter (DM) content. We observed the galaxy with the integral field spectroscopy (VLT/VIMOS, in IFU configuration). Results for the gaseous kinematics (bulge and disk) and stellar kinematics (bulge) are presented, together with a Jeans model for the stellar bulge kinematics. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.3593_arXiv.txt": {
        "abstract": "We describe the WFCAM Science Archive (WSA), which is the primary point of access for users of data from the wide--field infrared camera WFCAM on the United Kingdom Infrared Telescope (UKIRT), especially science catalogue products from the UKIRT Infrared Deep Sky Survey (UKIDSS). We describe the database design with emphasis on those aspects of the system that enable users to fully exploit the survey datasets in a variety of different ways. We give details of the database--driven curation applications that take data from the standard nightly pipeline--processed and calibrated files for the production of science--ready survey datasets. We describe the fundamentals of querying relational databases with a set of astronomy usage examples, and illustrate the results. ",
        "introduction": "\\label{intro}  The term `science archive' is first seen in the astronomy literature in Barrett~(1993) which describes the High Energy Astrophysics Science Archive Research Centre (HEASARC). This system is much more than a simple repository of data -- HEASARC provides an online resource to enable scientific exploitation of high--energy astronomy missions via provision of science data, software, analysis tools and descriptive information. For example, the data holdings in the HEASARC amount to many terabytes (TB; 1~TB=10$^{12}$~bytes) so wholesale download is impractical; recognising this, a {\\em server--side} analysis facility (i.e.~a facility co-located with the data and hence remote to the typical user) is provided to enable large--scale processing given an arbitrary astronomical usage scenario. In this way, data download is limited to user--defined subsets, sometimes processed in a manner specified by the user at access time.  The advent of the large Schmidt photographic plate digitisation programmes (Hambly et al.~2001a and references therein) and infrared surveys such as DENIS (Epchtein et al.~1994) and 2MASS (Kleinmann et al.~1994) presented similar challenges for ground--based missions. Data distribution for the digitised Schmidt surveys was originally done on removable, permanent storage media (`compact disc' read--only memory), but this became impractical so online services rapidly developed for these also. However, it is probably fair to say that it was with the challenges posed by the Sloan Digital Sky Survey (SDSS; York et al.~2000) that the ground--based astronomical science archive became fully developed (Gray et al.~2002). The first SDSS (the so--called Sloan Legacy Survey) is now complete, and $\\sim2\\times10^{8}$ sources have been measured and characterised, producing a catalogue of several~TB in size with associated imaging data and $\\sim10^6$ spectra amounting to a total volume of $\\sim10$~TB (Adelman--McCarthy et al.~2007 and references therein); the state--of--the--art SDSS science archive is described in Thakar et al.~(2003a).  The challenges and opportunities presented by the current generation of ground--based infrared surveys were noted by Lawrence et al.~(2002). In particular, they cited the advent of a new wide--field camera for the 4-m United Kingdom Infrared Telescope (WFCAM for UKIRT; Casali et al.~2007) and the even greater challenges posed by the new dedicated 4-m telescope for infrared surveys VISTA (Emerson~2001). These ambitious survey missions gave rise to a systems--engineered data management project, the VISTA Data Flow System (VDFS; Emerson et al.~2004) which included provision of pipeline processing and science archiving for WFCAM and VISTA data. Here, we concentrate on the first VDFS science archive known as the WFCAM Science Archive (WSA). From the outset, the design of the WSA has been science--driven with the main science stake--holders being users of the UKIRT Infrared Deep Sky Survey (UKIDSS; e.g.~Warren~2002).  This paper is one of a set of five which provide the reference technical documentation for UKIDSS, although it is of direct relevance to any user of the WFCAM Science Archive. The other four papers in the series describe the infrared survey instrument itself (WFCAM; Casali et al.~2007); the WFCAM photometric system (Hewett et al.~2006); the UKIDSS surveys (Lawrence et al.~2007); and the pipeline processing system (Irwin et al.~2007).   This paper is arranged as follows. In Section~\\ref{design}, we describe the design of the WSA, concentrating on the development of the science requirements into data models (i.e.~the database design) as presented to the end--user at access time. In Section~\\ref{browser} we discuss various detailed implementation issues that in particular inform the user as to how science--ready survey catalogues are generated from the standard flat file products processed by the nightly pipeline. Section~\\ref{queries} then goes on to discuss some illustrative science examples by concentrating on the expression of certain specific science usage modes in Structured Query Language (SQL), the {\\em lingua franca} of relational database users. Following the usual conclusion, acknowledgements and bibliography we present as appendices some supplementary information to aid first--time users of the WSA. ",
        "conclusions": "\\label{concs}  We have described the WFCAM Science archive (WSA), which is the end point in the data flow of UKIRT WFCAM data in the VISTA Data Flow System, and the primary point of access for users of survey science products, especially those of the United Kingdom Infrared Deep Sky Survey (UKIDSS). In particular, we have described: \\begin{itemize} \\item how the top--level requirements and typical usage modes informed the design of the WSA; \\item the arrangement of survey data in terms of a set of related tables; \\item the implementation of the archive within a commercial relational database management system; \\item the curation procedures employed to create science--ready survey catalogues from standard pipeline--processed products; \\item example real--world astronomy usage modes along with typical results. \\end{itemize}  The WSA is the prototype science archive for the VISTA surveys, and the design of the VISTA Science Archive will follow closely that described here."
    },
    "0711/0711.1902_arXiv.txt": {
        "abstract": "We present Spitzer/IRS mid-infrared spectra for 15 gravitationally lensed, 24~\\micron--selected galaxies, and combine the results with 4 additional very faint galaxies with IRS spectra in the literature. The median intrinsic 24~\\micron\\ flux density of the sample is 130~\\microJy, enabling a systematic survey of the spectral properties of the very faint 24~\\micron\\ sources that dominate the number counts of Spitzer cosmological surveys. Six of the 19 galaxy spectra ($32\\%$)  show the strong mid-IR continuua expected of AGN; X-ray detections confirm the presence of AGN in three of these cases, and reveal AGNs in two other galaxies. These results suggest that nuclear accretion may contribute more flux to faint 24~\\micron--selected samples than previously assumed. Almost all the spectra show some aromatic (PAH) emission features; the measured aromatic flux ratios do not show evolution from $z=0$. In particular, the  high S/N mid-IR spectrum of \\imageB\\ agrees remarkably well with low--redshift, lower--luminosity templates. We compare the rest-frame 8~\\micron\\ and total infrared luminosities of star--forming galaxies, and find that the behavior of this ratio with total IR luminosity has evolved modestly from z$=$2 to z$=$0. Since the high aromatic--to--continuum flux ratios in these galaxies rule out a dominant contribution by AGN, this finding implies systematic evolution in the structure and/or metallicity of infrared sources with redshift. It also has implications for the estimates of star forming rates inferred from 24~\\micron\\ measurements, in the sense that at $z \\sim 2$, a given observed frame 24~\\micron\\ luminosity corresponds to a lower bolometric luminosity than would be inferred from low-redshift templates of similar luminosity at the corresponding rest wavelength. ",
        "introduction": "The Spitzer Space Telescope \\citep{werner-spitz} has been tremendously successful at detecting star--forming galaxies and active galactic nuclei (AGN) at $z>1$ by their emission in the observed--frame 24~\\micron\\ band. Down to the confusion limit of $\\sim 50$~$\\mu$Jy at 24~\\micron\\ \\citep{dole04}, the MIPS instrument \\citep{rieke-mips} can detect luminous infrared galaxies (LIRGs)\\footnote{Defined as $11 < \\log L(TIR)$ \\Lsun $< 12$, where L(TIR) is the total infrared luminosity between 8 and 1000~\\micron.} out to $z \\sim 2$, ultra-luminous infrared galaxies (ULIRGs)\\footnote{Defined as $12 < \\log L(TIR)< 13$~\\Lsun.} out to $z\\sim3$, and hyper-luminous infrared galaxies (HLIRGs)\\footnote{Defined as $\\log L(TIR)> 13$~\\Lsun.} out to even higher redshifts.  Detectability at $z>0.7$ is increased by strong aromatic emission features, at rest--frame 6--12~\\micron, that pass through the 24~\\micron\\ band.  Indeed, photometric redshifts place $\\sim 25$--$30\\%$ of the faint 24~\\micron\\ sources at redshifts above 1.4 \\citep{emeric05,pablo05,caputi06,wang06}.  For extremely high--luminosity galaxies, spectra have been obtained with the Infrared Spectrograph (IRS, \\citealt{houck-irs}), e.g.~\\citet{houck-highz,yan07,md07}. However, since spectral properties depend strongly on luminosity, results obtained for hyper-luminous galaxies may have limited applicability for the bulk of the IR--detected population. Spectra of lower--luminosity galaxies have been obtained only in a few cases (e.g. \\citealt{lutz05,teplitz07}).  Numerous authors (e.g. \\citealt{caputi06,choi06,reddy06}) have used the observed 24~\\micron\\ band as an estimator of star formation rates.  Since most of the infrared power is radiated redward of the observed band, the 24~\\micron\\ diagnostic must be calibrated using low--redshift templates (e.g.~\\citealt{charyelbaz, dalehelou,lagache04,brandl06,armus07}). The central assumption in these works is that the spectra and spectral energy distributions of high-z galaxies are matched accurately by the low--redshift templates. For $0.7<z<3$, the 24~\\micron\\ band fluxes  are strongly influenced by aromatic band  emission.  Evolution over time either in the behavior of these aromatic bands, in metallicity, or in the geometry and radiative transfer within these infrared sources, could all undermine this central assumption.  To probe these possibilities, we have been obtaining IRS spectra for intrinsically faint 24~\\micron--selected galaxies at $1\\la z \\la 3$. Our targets are strongly lensed by the gravitational potential of foreground clusters of galaxies,  such that their fluxes are amplified by factors of 3--25. Though these galaxies have observed 24~\\micron\\ flux densities of $\\sim 1$~mJy, sufficient for useful, efficient spectroscopy with IRS, their intrinsic luminosities are down in the LIRG to ULIRG range.  In this paper, we present IRS spectra for the first 15 galaxies in the sample.    The median observed flux density is f$_{\\nu}$(24~\\micron) $=  0.82$~mJy, but the median intrinsic flux density is nearly an order of magnitude lower. We compare their aromatic feature flux ratios, and the ratios of their X-ray, aromatic, and total infrared luminosities to those of low--redshift LIRGs and ULIRG samples. We examine in detail the high--quality spectrum of \\imageB, a LIRG at $z=2.516$. We merge our sample with four additional objects from the literature \\citep{lutz05,teplitz07} to provide a first systematic look at the IRS--measured properties of the faint 24~\\micron\\ galaxy population.  We assume $\\Omega_m = 0.265$, $\\Omega_{\\Lambda} = 0.735$, and H$_0 = 72$~km~s$^{-1}$~Mpc$^{-1}$ \\citep{spergelwmap1,freedman01}. ",
        "conclusions": "\\subsection{Aromatic Feature Flux Ratios in High--Redshift Galaxies}  We now test whether the mid-IR spectra  of star--forming galaxies, with their prominent aromatic emission features, show evidence for evolution from $z=0$ to high redshift.  Given its high quality, we first examine the IRS spectrum for \\imageB\\ in detail. The image is amplified by a factor of $22 \\pm 2$ \\citep{kneib04}, such that without lensing, its 24~\\micron\\ flux density would have been 53~\\microJy, at the MIPS confusion limit. The galaxy redshift, from H$\\alpha$, is $z=2.5165 \\pm 0.0015$ \\citep{kneib04}. Its IRS spectrum, presented in Figure~\\ref{fig:a2218_spec}, shows highly--significant detections of aromatic features at 6.22, 7.7, 8.33, and 8.61~\\micron. (The lower-quality spectrum from program 3453 confirms the detection of these aromatic features, though the two longer--wavelength features are detected at low significance.) Overplotted with arbitrary normalization are two spectral templates:  a) the starburst galaxy NGC~2798 from the SINGS survey \\citep{dale06}; and b) the average spectrum of 13 nearby starburst galaxies with little apparent AGN contribution \\citep{brandl06}, whose mean IR luminosity is $4.9 \\times 10^{10}$~\\Lsun. Coincidentally, this is also the L(TIR) of NGC~2798.  The $z=2.516$ spectrum  in Figure~\\ref{fig:a2218_spec} closely resembles the low--redshift, lower--luminosity (by a factor of 15) templates. Aromatic flux ratios are reported in Table~\\ref{tab:pahrats}, along with those we measure for the Brandl template, and those reported for the SINGS sample. The strength of the 8.6~\\micron\\ feature in \\imageB\\ relative to those at 6.2~\\micron\\ and 7.7~\\micron, is  within the $10\\%$--$90\\%$ variation within the SINGS sample. In addition, the flux ratios of \\imageB\\ are a close match to those of the \\citet{brandl06} average starburst template; while 7.7~\\micron\\ is relatively stronger than the template, the difference is comparable to the standard deviation of line ratios for individual galaxies that were averaged to create the template. {\\bf Thus, we conclude that this $z=2.5$ galaxy has aromatic feature flux ratios that are  consistent with those observed for lower--luminosity, $z=0$ starbursting galaxies.}  While the other spectra in the sample have lower signal-to-noise ratios, they are still sufficient, in the aggregate, to test whether high--redshift galaxies have markedly different aromatic flux ratios.  Figure~\\ref{fig:pahrats} compares, for four aromatic feature flux ratios, the observed ratios in our sample with those measured for the low--redshift \\citet{brandl06} template. Though the errorbars in individual measurements are sizable, Figure~\\ref{fig:pahrats} confirms the result seen for \\imageB: star--forming galaxies at $1<z<3$ have aromatic feature ratios consistent with those of $z=0$ star--forming galaxies.  Variations in aromatic feature ratios have been reported in certain HII regions---this may be attributable to variations in size, composition, and ionization state of the carriers  (\\citealt{draineli07} and references therein). In particular, large PAHs emit more strongly at 11.3~\\micron, and neutral PAHs emit much more strongly at 3.3 and 11.3~\\micron\\ than charged PAHs \\citep{draineli07,galliano06,allam89}. That our measured feature ratios at $z>1$ are consistent with those measured at $z\\sim0$ suggests that the dust grain size and ionization distributions are not strongly evolving.   \\subsection{Evidence for Compact Source Accretion} \\label{sec:nuc}  Active galactic nuclei are detectable in multiple ways---by the dust they heat that radiates in the mid-IR; by X-rays emitted from the accretion disk coronae; and by high--excitation or broad emission lines. We now examine the evidence for AGN activity in our extended sample.  The first AGN diagnostic we examine is X-ray luminosity, plotted in Figure~\\ref{fig:fpahX} against aromatic luminosity. Most of the sample are X-ray non--detections, at limiting luminosities that rule out X-ray--loud QSOs or bright Seyferts, although an X-ray--weak or highly obscured AGN could still be present (c.f. \\citealt{jennradio,almu-plagn}). Two of the four galaxies from the literature, and three of the galaxies in our program are detected in X-rays: two have no apparent aromatic features (sources A2261a and A2390a); and the other, source A2390b, is interesting in that it contains a luminous X-ray--emitting AGN, yet its spectrum still shows aromatic features. That source emits roughly equal power in the 7.7~\\micron\\ aromatic feature (scaled from the 11~\\micron\\ feature) and at 10--30~keV, reminiscent of the more luminous $z=1.15$ source CXO GWS J141741.9+522823 discussed by \\citet{lefloch07}.  We now consider a second AGN diagnostic, the relative contribution of aromatic versus continuum emission to the mid-IR flux. At low redshift, low aromatic feature equivalent widths  have been demonstrated as an effective AGN diagnostic (\\citealt{armus07,brandl06}, and Devost \\etal\\ (in prep.)). However, it is extremely difficult to measure equivalent widths accurately for the high--redshift IRS sources, due to the limited wavelength baseline,  the fact that the aromatic features have broad wings, and most importantly, imperfect sky subtraction which adds a (positive or negative) pedestal to each spectrum. For example, for \\imageB\\ (our highest--quality spectrum), we cannot measure an accurate equivalent width due to the difficulty in determining the true continuum level. We are able to set a lower limit of  $>0.4$~\\micron\\ (rest-frame, conservatively assuming a high continuum level), which is a typical value for low--redshift star--forming galaxies. For the other spectra in the sample, equivalent widths are even harder to measure.  Therefore, we instead create an aromatic-to-mid-IR flux ratio, which we define as the ratio of the flux in the 7.7~\\micron\\ PAH feature (in \\cgsflux, as fit by PAHFIT) to the MIPS photometry at 24~\\micron\\ (in Jy, corrected for bandwidth compression to $z=0$.) This metric does not suffer as strongly from the difficulty in determining the continuum level.  In the low--redshift, star--forming comparison sample we construct in \\S~\\ref{sec:pahtotrat}, this ratio ranges from 7--32 $\\times 10^{-11}$. We thus take $7 \\times 10^{-11}$ as a dividing line between spectra dominated by star formation (above the value), and spectra with substantial AGN contribution (below the value). This discriminant selects the following mid-IR spectra\\footnote{For A2390b, we scale from the flux in the 11~\\micron\\ aromatic feature, using the average 7.7/11~\\micron\\ flux ratio of 3.6 from \\citet{jdsmith07}.} as having substantial AGN contribution: A370a, MS0451a, A1689a, A2218c, A2261a, A2390a. Thus, of 19 galaxies with adequate spectra, 6 show strong indications of AGN contribution to their mid-IR outputs.  A third indication of AGN activity is a mid-IR spectrum that rises steeply with increasing wavelength.  Four spectra (A2261a, A2218c, A1689a, and A370a) show this behavior; none has detected aromatic emission.  A fourth indication of AGN activity is the presence of broad or highly excited lines in an optical spectrum. Of the spectroscopy that has been published, or is in preparation, for our sample (see references in Table~\\ref{tab:sample}), three show evidence for AGN: A1689a has a Keck spectrum showing highly ionized neon (J.~Richard \\etal\\ in prep.); A2390a has a Keck spectrum (see Appendix) which shows Lyman $\\alpha$ with FWHM $\\sim 945$~km~s$^{-1}$, typical of an AGN narrow line region; and the spectrum of A370a is reported to contain AGN lines \\citep{lutz05}.  A fifth indication of AGN activity is a high ratio of [Ne III] 15.5~\\micron to [Ne II] 12.8~\\micron.  Solar--metallicity starbursts are not observed to excite this line ratio above unity \\citep{thornley,rigby-rieke}, whereas the harder continuum of AGN do. Eight spectra cover rest-frame 12.8\\micron; [Ne II] is detected in 6 of these. Six spectra cover rest-frame 15.5~\\micron;  [Ne III] is detected in 4. Of the four galaxies where both lines are detected (A1689a, Ab2218b, A2390b, and A2667a), only in A1689a does the [NeIII]/[NeII] ratio exceed unity:  $2.15 \\pm 0.76$.   This high line ratio offers additional evidence that A1689a hosts a luminous AGN. The other, lower line ratios are more difficult to interpret.  The results for these AGN diagnostics are compiled in Table~\\ref{tab:agn-diag}. Combining these diagnostics, 8 galaxies out of the sample of 19 have at least one indication of AGN activity (X-ray detection, low aromatic contribution to the mid-IR flux, rising mid-IR spectrum, or optical AGN lines).  Five galaxies  have two different diagnostics that indicate AGN. Thus, we find that $42\\%$ of the faint 24~\\micron\\ sources in this sample have some nuclear activity; the $95\\%$ confidence interval is $23\\%$--$63\\%$. Further, in $32\\%$ of the sample, the AGN  strongly affects or dominates the mid-IR spectrum; the $95\\%$ confidence interval is $15\\%$--$53\\%$. (Confidence intervals were calculated using eqn.~26 of \\citet{gehrels86}.)  The high proportion of AGN in our sample of high--redshift LIRGS and ULIRGs mirrors the general behavior of local IR-luminous galaxies, where AGN become substantially more common with increasing luminosity (e.g., \\citealt{lutz98}). Though the errorbars are large given the small sample size, the high proportion ($32\\%$) of AGN sufficiently powerful to dominate the mid--IR spectra may be inconsistent with the frequent assumption that $\\sim 100$~\\microJy, $z>1$ sources are dominated by star formation, with AGN a modest contaminant.  Our result, if confirmed with better statistics, has implications for 24~\\micron--derived star formation rates in deep surveys, where it is often assumed that the mid-IR outputs of high--redshift, faint galaxies are dominated by strong aromatic features from star formation. It also implies that a higher fraction of the integrated 24~\\micron\\ flux on the sky may arise from accretion than previously thought.  \\subsection{Total infrared luminosity}  We now determine the total infrared (TIR) luminosities of the lensed sources, so that we may examine the dependence of aromatic feature emission on total infrared luminosity, and how this behavior may evolve with redshift.  \\subsubsection{\\imageB}  The sub-mm galaxy behind Abell 2218 has exceptionally good photometric coverage from 0.4 to 850~\\micron.  Since it is triply--lensed, we increase the signal-to-noise ratio of its multiband  photometry by using fluxes from all three images, weighed by the amplifications given in \\citet{kneib04}. (Bright neighboring galaxies forced us to use only image A and B for the optical and IRAC photometry.) Figure~\\ref{fig:a2218_phot} plots this photometry.  We fit redshifted \\citet{charyelbaz} and \\citet{dalehelou} templates to the 24, 70, 450, and 850~\\micron\\ photometry, corrected for lensing amplification.  For each template family, we determine the best--fit template and the range of acceptable templates, then determine L(TIR) for those templates in a method consistent with the definition given in Table~1 of \\citet{sanders-mirabel} using 12, 25, 60, and 100~\\micron\\ photometry. We estimate the intrinsic total infrared luminosity of \\imageB\\ as: \\begin{itemize} \\item  $7.6 \\pm1.9 \\times 10^{11}$~\\Lsun\\ for \\citet{dalehelou}  templates. \\item  $1.3^{+1.0}_{-0.2} \\times 10^{12}$~\\Lsun\\ for \\citet{charyelbaz} templates. \\end{itemize}  As Figure~\\ref{fig:a2218_phot} shows, the best--fit \\citet{charyelbaz}  template over-predicts all $>24$~\\micron\\ points, especially at 70~\\micron. The best--fit \\citet{dalehelou} template is a better match to the shape of the overall infrared SED and also has high aromatic feature equivalent widths, consistent with the spectrum of the source. Thus, this galaxy's intrinsic luminosity is that of a LIRG.  \\subsubsection{Other Galaxies}  We adopt a similar approach for the other galaxies with sufficient far--infrared and submillimeter data to provide reasonable constraints. We tried both \\citet{dalehelou} and \\citet{charyelbaz} templates, fitted to fall just below any upper limits so long as they were compatible with the other constraints, i.e., to give a maximal luminosity. Two galaxies are fit poorly by a star--forming template, and much better by an AGN template, Mrk 231; MS0451a may be, as well. The ranges of acceptable luminosities and the preferred templates are listed in Table~\\ref{tab:longwave}. \\citet{dalehelou} templates provided generally better fits to our sample than \\citet{charyelbaz} templates, as expected given the results of \\citet{delphine06} for $0<z<1$ galaxies, fitting ISO, Spitzer, and VLA photometry.   \\subsection{8~\\micron\\ (rest)-to-Total Luminosity Ratio}  % \\label{sec:pahtotrat}  We now consider the aromatic and total infrared (TIR) luminosities of the star--forming galaxies in the lensed sample, and how this ratio behaves as a function L(TIR).  We compare to higher--luminosity star--forming galaxies at $z\\sim 2$, as well as comparable--luminosity star--forming galaxies at $z\\sim 0.8$ and $z \\sim 0$. We interpret the properties of the high redshift galaxies strictly in terms of properties that are well--determined for local analogs, so we can test whether inconsistencies emerge.  \\subsubsection{The $z \\sim 0$ comparison sample}  To probe rest-frame 8~\\micron\\ and total infrared (TIR) luminosities for local galaxies, we use the SINGS sample (photometry from \\citealt{dale05}) supplemented in the ULIRG range by IRAS 09111-1007, 10565+2448, 12112+0305, 14348-1447, 17208-0014, 22491-1808, and Arp220 \\citep{farrah03,armus07}, all selected because they appear to be dominated by young stars rather than AGN.\\footnote{IRAS 12112+0305 and 14348-447 are double nuclei sources.} We also include IRAS 00262+4251, 01388-4618, 02364-4751, 16474+3430, 23128-5919, and 23365+3604 from \\citet{rigo99}. The strength of the aromatic features relative to the continuum indicates that all these galaxies are also dominated by star formation. For the galaxies from \\citet{farrah03} and \\citet{armus07}, we determined 8~\\micron\\ photometry from archival IRAC images. In our reductions, we included a correction for the extended source calibration.  For those galaxies in \\citet{rigo99}, we converted the tabulated continuum and aromatic 7.7~\\micron\\ peak fluxes into 8~\\micron\\ photometric values by comparing with templates for galaxies in common between \\citet{armus07}: IRAS 12112+0305, 14348-1447, 1525+3609, and 22491-1808, plus Arp 220. The peak--to--peak scatter was a factor of two (in agreement with the finding of \\citet{charyelbaz}). We add a variety of LIRGS \\citep{almlirg}, again avoiding any galaxies with indications of significant levels of AGN power. We took IRAS measurements from \\citet{sanders03}, and computed L(TIR) to be consistent with the approach taken in that paper. We define L(8~\\micron) as being proportional to \\nufnu, and use similar definitions for all luminosities at specific wavelengths. Figure~\\ref{fig:TIRL8} shows the relation between rest--frame L(8~\\micron) and L(TIR) for these $z\\sim0$ galaxies.   \\subsubsection{Results from the literature at $z\\sim$ 0.85} We add a sample of infrared galaxies at $z \\sim 0.85$ from \\citet{delphine06}. This paper utilizes 15~\\micron\\ data from ISO as well as 24~\\micron\\ data from Spitzer and, for a subset, radio data from the VLA; the authors found that the galaxies had rest--frame L(8~\\micron)/L(12~\\micron) ratios at the high end of the local galaxy dispersion. We make the assumption that  L(12~\\micron) correlates with L(TIR) in that sample, as it is observed to do well at $z=0$. We computed rest 8~\\micron\\ luminosities for the members of the sample with 0.55 $<$ z $<$ 1.2 and adequate 15~\\micron\\ photometry, and corrected them to the IRAC 8~\\micron\\ band by comparison with the IRS spectrum of IRAS 2249, a local star-formation--dominated ULIRG. We took L(TIR) from the \\citet{dalehelou} template fits \\citep{delphine06}. The results are shown in green in Figure~\\ref{fig:newTIR}.  There is no mid-IR spectroscopy for these galaxies, but if they had strong mid-IR contributions from AGN we would expect substantial dispersion in the L(8\\micron) to L(TIR) ratio.  Such a large dispersion is not seen.   \\subsubsection{Results from the Literature at $z \\sim 2$}  We now compare the behavior of the 8$\\mu$m luminosity vs. L(TIR) at $z \\sim 2$ with that seen in local luminous galaxies.  We first discuss three results from the literature, and then show how our new measurements complement them.  {\\bf First, we took the stacked $z \\approx 1.9$ galaxy template of \\citet{daddi05}} and compared the ratio of observed 850~\\micron\\ to 24~\\micron\\ flux densities, referred to the appropriate rest wavelengths for each of the local star--forming galaxies where there is sufficient data to construct an accurate SED template (IRAS 1211, 1434, 17208, 2249, and Arp 220).  Excluding Arp 220, for which the ratio is 54, the average value for the other four galaxies is 14, whereas the \\citet{daddi05} template indicates a value of 6.5.  Thus, the template indicates stronger emission at rest 8~\\micron\\ relative to the submm than is typical of local star--forming ULIRGs.  We also compare with total infrared luminosity derived from submm measurements.  Since the rest wavelength for the observed 850~\\micron\\ point in the Daddi \\etal\\  template is 293~\\micron, we construct the average ratio of L(TIR)/L(293~\\micron) for local LIRGs and ULIRGs with $4 \\times 10^{11} <  L(TIR) < 2.5 \\times 10^{12}$~\\Lsun\\ and high quality measurements at 350--450~\\micron\\ (IRAS 09111, 1211, 1434, and 17208, Arp 220, and NGC 1614). We correct the flux densities for the differing effective filter bandwidths using the IRAS 2259 template. We correct to 293~\\micron\\ assuming a spectral slope of -3.3, determined from a larger sample of local ULIRG templates.  We find  L(TIR) $=$ ($98 \\pm 40$)  L(293~\\micron)  \\noindent and then plot the Daddi \\etal\\ point as a red square in Figure~\\ref{fig:TIRL8}. We did not fit the observations of lower luminosity galaxies, because the scatter at lower luminosities is very large, probably due to varying amounts of cold and warm dust in the infrared-emitting objects.  {\\bf Second, we make use of the stacking analysis carried out by \\citet{papovich07}.} They report observed--frame 24~\\micron\\ and 70~\\micron\\ stacked photometry for galaxies that are both undetected in X-rays and have non--power--law IRAC SEDS, and are therefore assumed to be powered predominantly by star formation. We computed the rest 8~\\micron\\ and 24~\\micron\\ flux densities for this category (since they quote values for $z = 2$, the rest wavelengths are represented directly by the measurements). We corrected the 8~\\micron\\ flux density (derived from that observed at 24~\\micron) to the equivalent for the IRAC band, again using the SED of IRAS 2249. Finally, we compute L(TIR) from the relation derived from our local galaxy sample:  L(TIR) $=$ $86.5$ $\\times $ L(24~\\micron, rest)$^\\alpha$, where $\\alpha = 0.9758$.  %  The rms scatter around this relation is only $0.14$ dex.  The resulting values are shown in Figure~\\ref{fig:TIRL8} as red diamonds (we omitted the lowest flux density bin from Papovich \\etal\\ because the signal to noise ratio at 70~\\micron\\ is too low).   {\\bf Third, we use the results of \\citet{yan07}.} We select those galaxies  that have $z>1.5$, and rest--frame 6.2~\\micron\\ equivalent widths (measured by \\citealt{sajina07}) greater than 0.3~\\micron. From \\citet{armus07}, this equivalent width threshold should select galaxies dominated by star formation.  The galaxies selected (MIPS 289, 8493, 15928, 16144, 22530) are all categorized as ``type 1'' by \\citet{yan07}, meaning their spectra contain strong aromatic features. The median rest-frame EW(7.7~\\micron) is  2.5~\\micron, corroborating that star formation dominates the mid-IR spectra. We converted the tabulated L(5.8~\\micron) for these galaxies into L(8~\\micron) from an average template for our local ULIRG sample (IRAS 2249).  We estimate L(TIR) for these sources using the radio-infrared relation. We describe our adjustment of this relation for use at high redshift in detail in a forthcoming publication (N. Seymour \\etal, in preparation) but give a short summary here. \\citet{yun01} have quantified the relation for a large number of local galaxies, finding that the average ratio of flux density at 60$\\mu$m to 1.4GHz is 110 independent of luminosity.  The results of \\citet{yun01} have not been K-corrected; since the infrared SEDs drop rapidly toward the blue, while the radio spectrum drops slowly, these corrections are a fairly steep function of redshift.  They enter at a significant level only for the high luminosity galaxies, because they tend to be at higher redshift than the lower luminosity ones. Including them increases the indicated ratio to $\\sim 170$ at a L(TIR) $= 10^{12}$ \\Lsun. A second issue is that the radio outputs of LIRGs and ULIRGs often show signs of free-free absorption, which can decrease the outputs at 1.4GHz. To estimate this effect, we took the 19 star forming galaxies with adequate data from \\citet{condon91} and computed an average slope for them of $-0.58$ between 1.49 and 8.44 GHz. Taking the 8.44 GHz value to represent the un-absorbed radio output and assuming the intrinsic slope is $-0.7$ \\citep{condon91}, we revise the IR/radio ratio downward to 140 to take account of absorption; this value is calculated by extrapolating from 8.44 GHz to 1.4 GHz and comparing results for slopes of $-0.58$ and $-0.7$.  The second concern is that the radio-infrared ratio may be a function of redshift \\citep{kovacs06,vlahakis07}. % We tested this suggestion by determining a K-correction between 1.4 GHz and 850$\\mu$m for an average of local LIRG and ULIRG templates (to be published by Alonso-Herrero et al.) and applying it to a large sample of high redshift sources. We find that the K-correction accurately reproduces the observed behavior with no offset, indicating there is no significant evolution of the relation. We have therefore used a value of 140 for the intrinsic flux density ratio at rest wavelengths to convert the radio measurements to TIR ones. We extrapolated the radio fluxes observed to a rest frequency of 1.4 GHz with a slope of $-0.7$. The redshift usefully makes the observed rest frequency significantly higher than 1.4 GHz, reducing the effects of free-free absorption for the high-z galaxies. To convert to L(TIR), we use the relationship from our local sample,  L(TIR)$ =$ 740 $\\times $ L(60)$^{\\alpha}$, where $\\alpha = 0.9427$  \\noindent with an rms scatter of 0.08 dex.  It is important to point out that both of the corrections we have derived have the effect of {\\it reducing} the change we find in L(8$\\mu$m)/L(TIR). That is, using the relation derived by \\citet{yun01} and applying the proposed change at high redshift \\citep{kovacs06,vlahakis07} would have both resulted in a much larger change than reported. Therefore, we are taking a conservative approach.  We take  radio flux densities for the galaxies observed by \\citet{yan07} from \\citet{condon03} as tabulated on the SSC website\\footnote{http://ssc.spitzer.caltech.edu/fls/extragal/vla.html}. Two of the remaining galaxies are detected at low significance as indicated by inspection of the grey scale images for them (MIPS 429 and 506), so we enter rough values of 0.1~mJy for them. For MIPS 289, which is undetected, we assigned an upper limit of 0.1~mJy. We reject the source MIPS 8493 because of fringing in the VLA image at its position.  The other four (MIPS 289, 15928, 16144, and 22530) are in regions where the radio images are clean. We correct their  observed radio measurements to the rest frame assuming a radio k-correction of $(1+z)^{0.7}$, since a typical slope for a starburst radio spectrum is $-0.7$ \\citep{condon91}. We plot the resulting L(TIR) and L(8~\\micron) values in Figure~\\ref{fig:newTIR} as circles (filled for the detections, open for the upper limit.)       \\subsubsection{8~\\micron-to-TIR ratio for the $z \\sim 2$ lensed galaxies}  The lensed galaxies described in this paper bridge the luminosity divide between the $z\\sim 2$ extremely high--luminosity galaxies in the literature, and the comparison LIRG and ULIRG samples at $z\\sim 0$ and $z = 0.85$. For the lensed sample we can determine which galaxies have mid-IR spectra dominated by star formation, using the aromatic feature strengths. For the subset of lensed galaxies with submm measurements and good upper limits at 70~\\micron, the SEDs are well enough constrained for a direct determination of L(TIR).  Therefore, the lensed galaxies can cleanly test for evolution with redshift in the behavior of the L(8~\\micron)/L(TIR) ratio.  We use the aromatic-to-mid-IR flux ratio introduced in \\S\\ref{sec:nuc} to select galaxies whose mid-IR spectra are dominated by star formation. The ratio selects A2219a, A1689b, A2218a, and A1835a as all having strong aromatic emission; their ratios range from  7--24 $\\times 10^{-11}$, with an average value of $16 \\times 10^{-11}$. By comparison, our sample of local star--forming LIRGs and ULIRGs has a range in ratio of 7--32 $\\times 10^{-11}$, with an average of $21 \\times 10^{-11}$. That is, the relative aromatic feature strength is virtually the same for the two samples, with no evidence that AGN have significantly augmented the 8~\\micron\\ fluxes of our star--forming high--z sub-sample. The EW(6.2~\\micron) criterion we used for the \\citet{yan07} sample would have returned the same sample of star--forming lensed galaxies, except that A2219a lacks 6.2~\\micron\\ coverage and thus would not have been selected.  We plot the four lensed galaxies as red triangles in Figure~\\ref{fig:newTIR}.    \\subsection{Comparison of 8~$\\mu$m-to-TIR ratio at high and low redshift}  The L(8~$\\mu$m)/L(TIR) ratio is important for two reasons: 1) it is one of the most accessible measures of the physical behavior of infrared emission as a function of redshift; and 2) many studies rely on fluxes at observed--frame 15 or 24~\\micron\\ to deduce L(TIR) and related properties of galaxies at $z \\sim 1$ or 2, and are thus relying on the behavior of L(8~$\\mu$m)/L(TIR). Figures~\\ref{fig:TIRL8} and \\ref{fig:newTIR} illustrate the important result that this ratio is well--behaved for star--forming galaxies out to $z \\sim 2.5$. In fact, this good behavior is a reasonably demanding test, since the various samples at $z \\sim 2$ are subject to different selection effects and have L(TIR) derived in different ways, yet they all follow a single trend. To recapitulate the different methods: the lensed galaxies are selected largely on amplification factor, and we measured L(TIR) directly from template fits to rest frame mid- and far-IR and submm photometry. The four galaxies from \\citet{yan07} were selected on the basis of colors and brightness at 24~\\micron, and we estimated L(TIR) from the radio-infrared relation.  The stacked galaxies from \\citet{papovich07} were selected based on faint K-band and 24~\\micron\\ detections; we estimated L(TIR) from rest 24~$\\mu$m measurements.  The stacked galaxies from \\citet{daddi05}  were originally selected from BzK colors indicating they are at $z \\sim 2$ and were stacked based on selections at a variety of wavelengths, and we used rest submillimeter measurements to estimate L(TIR).  Having noted the overall uniform behavior, and the general similarity of this behavior independent of redshift, we now evaluate the evidence from Figures~\\ref{fig:TIRL8} and \\ref{fig:newTIR}   % for modest evolution with redshift. Specifically, there is a tendency for the ratio L(8$\\mu$m)/L(TIR) to decrease with increasing luminosity for the local galaxies for L $> 10^{11}$ L$_\\odot$, but this tendency seems reduced for the galaxies at $z \\sim  2$.  In fact, the high redshift galaxies appear to have SED behavior similar to normal local galaxies of lower luminosity in this regard.  We have made two tests of the significance of this difference. The first is based on the stacking results of \\citet{daddi05} and \\citet{papovich07}. Both studies select the galaxies to stack on the basis of K-band detections, and therefore the mid-- and far--infrared characteristics of the galaxies should be representative of the entire population of IR-active galaxies (i.e., there is no bias toward a specific L(8$\\mu$m)/L(TIR). To have an analogous set of data for the local galaxies, we have averaged their L(8$\\mu$m)/L(TIR) ratios in luminosity intervals of $10^{0.5}$; these averages are shown as the large blue triangles in Figure~\\ref{fig:TIRL8}.  % We have then fitted the averaged data points, shown as a blue line. To test whether the high redshift stacked results were compatible with the same fit, we carried out a second fit in which we included the three stacked points. For the galaxies above $10^{11}$ L$_\\odot$, we then compared $\\chi^2$ {\\it for the low redshift galaxies only} in the two cases. We found that this $\\chi^2$ is 4.5 times larger if the three stacked points at $z\\sim 2$ are included in the fit; that is, there is a strong indication that a single functional fit cannot be used for both the local galaxies and for those stacked at $z \\sim 2$.\\footnote{This analysis makes the assumption that low--luminosity galaxies have similar behavior at low and high redshift.}  To determine whether a few anomalous local galaxies could affect this conclusion, we repeated the procedure but using a more robust average for the luminosity-binned ratios. For each bin with ten or more items, we discarded the two highest and two lowest ratios and then averaged the remainder. Where there were less than ten items, we retained the straight average. The highest luminosity bin, which is responsible for most of the change in slope of the relation between the two luminosities, has ten objects and hence its average was recomputed. The results were virtually identical to those with the straight averages, as shown in Figure 7.  A second test was applied to the high redshift galaxies measured individually, either by \\citet{yan07}, or by us. In the case of the Yan \\etal\\ sample, we expect a strong bias toward objects with a large ratio of L(8$\\mu$m)/L(TIR), because they selected relatively bright objects for spectroscopy, given the sensitivity limitations of the IRS. Our sample was selected on a combination of 24$\\mu$m flux density and lensing amplification, so a bias is likely, although perhaps less severe.  To compare with these two studies, we selected a local sample with relatively high 8~$\\mu$m output by removing all the galaxies that are below the fit made to the averaged values (above). This sample is shown in Figure~\\ref{fig:newTIR}. % It is noteworthy that the scatter in L(8$\\mu$m)/L(TIR) is small, with no outliers toward high L(8$\\mu$m), so this is an appropriate comparison sample for cases where a bias is suspected. As before, we carried out two fits, one just to the local galaxies, and another to them plus the eight individually measured high redshift ones. We found that $\\chi^2$ for the low redshift sample with L $> 10^{11}$ L$_\\odot$ increased by a factor of 2.5 if we included the high redshift galaxies in the fit. Comparing Figure~\\ref{fig:newTIR} with Figure~\\ref{fig:TIRL8}, the discrepancy would be substantially larger if we had included all the local galaxies in our fits.   This result confirms that found for the stacked samples, namely that a single functional fit is not a good fit to the data over all redshifts.  Another comparison can be made with the sample of infrared galaxies at $z \\sim 0.85$ from \\citet{delphine06}. The results fall in a region occupied by both low and high redshift fits in Figure~\\ref{fig:newTIR}, % which confirms the relatively modest change in L(8$\\mu$m)/L(TIR) with redshift, but (not surprisingly given the small difference between $z \\sim 0$ and 2) does not definitively deviate from either fit. In making this comparison, we have assumed that the galaxies studied by \\citet{delphine06} have a bias toward bright 15$\\mu$m observed flux densities. If instead we treat them as an unbiased sample and average them in luminosity intervals, they fall along (or slightly above) the fit for the $z \\sim 2$ sample in Figure~\\ref{fig:TIRL8}.  %  The tendency at high redshift for relatively strong 8$\\mu$m emission could arise either through a contribution of an extra, AGN-powered component at 8~$\\mu$m (e.g. the extra 5.8~\\micron\\ continuum claimed by \\citealt{sajina07} and the mid-IR excess galaxies with hard X-ray emission discussed by \\citealt{daddi07b}), or to a change in the properties of the stellar--powered infrared--emitting regions.  The AGN explanation runs counter to our selection of individual galaxies with strong aromatic features that correspond to equivalent widths typical of purely star forming local galaxies. It is also counter to the precautions taken in the stacking analyses to exclude AGN (e.g., \\citealt{papovich07}). In addition, AGN contributions would come in a broad range, so it would be hard to explain the consistent behavior of the five high redshift samples in Figures~\\ref{fig:TIRL8} and \\ref{fig:newTIR} if AGN were causing the departure from the behavior from local star forming galaxies. Finally, the behavior we see is similar to that found by \\citet{xzheng07} in a stacking analysis at $z \\sim 0.7$; specifically, they report that the far-IR SEDs appear to resemble those of local galaxies of lower luminosity.  Thus, while AGN may boost 8~\\micron\\ emission in other sources, it is unlikely that AGN contamination is responsible for the behavior seen in Figures~\\ref{fig:TIRL8} and \\ref{fig:newTIR}.  We conclude that star-forming galaxies have a mild trend toward increasing L(8~\\micron)$/$L(TIR) ratio with redshift that reaches up to a factor of $\\sim 2$ at $z=2$. In other words, L(8~\\micron) does not turn over with increasing L(TIR) for $z\\sim2$ galaxies as it does at low redshift.  This probably helps explain why some high--redshift galaxies have higher $6.2$\\micron\\ aromatic equivalent widths than are observed for local ULIRGs \\citep{desai07,md07,sajina07}, though direct comparisons cannot be made because there are not local galaxies with such extremely high luminosities.   Because the lensed galaxies and stacked samples directly probe the $10^{11}$--$10^{12.5}$~\\Lsun\\ range at $z\\sim 2$, we can directly compare to $z\\sim 0$ galaxies at similar luminosity rather than extrapolating.  That aromatic luminosities do not turn over with L(TIR) in high redshift galaxies may be related to lower densities and optical depths of the infrared--emitting regions in the high-z galaxies.  Lower optical depths would result in part if the metallicity were lower, resulting in a higher gas-to-dust ratio.  \\citet{draine07} show that the dust content of galaxies has a very steep dependence on their metallicity. The densities and optical depths of the star forming regions would also be lower if the high luminosities are associated with star formation on large spatial scales in the high--redshift galaxies, as opposed to strongly--concentrated, major-merger--associated nuclear star formation in local ULIRGs. Such very dense regions in local ULIRGs appear to account for the scattering of points below the $z = 0$ trend line.  Use of local templates to estimate L(TIR) for high-z galaxies detected at 24~$\\mu$m has indicated the existence of very high luminosities at $z \\sim 2$ and beyond (e.g., \\citealt{yan07}.) The change in L(8~\\micron)/L(TIR) suggested here would reduce L(TIR) for a given L(8~\\micron) by a factor of $\\sim 4$. As shown in the figure, in our sample there is only one galaxy selected to be dominated by star formation that has a L(TIR) significantly exceeding $10^{13}$~\\Lsun.\\footnote{This galaxy was selected as star--forming based on its EW(6.2~\\micron) \\citep{sajina07}, but its EW(7.7~\\micron) is unusually low, which may indicate a substantial AGN contribution.} A luminosity of $10^{13}$~\\Lsun\\ corresponds to a star formation rate of $\\sim 2000$~ \\Msun\\ yr$^{-1}$ if we use the local calibration \\citep{kennicutt98}, and is suggestively close to the theoretical maximum luminosity that a starburst is allowed by causality arguments---that star formation throughout a galaxy should not transpire faster than the dynamical timescale \\citep{lh96,elmegreen99}. (The actual maximum value depends on efficiency and dynamical mass.) Thus, our lower calibration for L(8~\\micron)/L(TIR) may remove the discrepancy between L(TIR) inferred for high redshift galaxies and plausible upper limits on what a starburst can produce.  Spatially--resolved imaging of high--redshift starbursts (e.g. \\citealt{nesvadba07}) can test whether they are in fact forming stars at the dynamical limit."
    },
    "0711/0711.1305_arXiv.txt": {
        "abstract": "\\small The {\\it XMM-Newton Extended Survey of the Taurus Molecular Cloud} (XEST) is a survey of the nearest large star-forming region, the Taurus Molecular Cloud (TMC), making use of all instruments on board the {\\it XMM-Newton} X-ray observatory. The survey, presently still growing, has provided unprecedented spectroscopic results from nearly every observed T Tauri star, and from $\\approx$50\\% of the studied brown dwarfs and protostars. The survey includes the first coherent statistical sample of high-resolution spectra of T Tauri stars, and is accompanied by an U-band/ultraviolet imaging photometric survey of the TMC. XEST led to the discovery of new, systematic  X-ray features not possible before with smaller samples, in particular the {\\it X-ray soft excess} in classical T Tauri stars and the {Two-Absorber X-ray} (TAX) spectra of jet-driving T Tauri stars. This paper summarizes highlights from XEST and  reviews the key role of this large project. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.2130_arXiv.txt": {
        "abstract": "The energy dependence of such characteristics as a ratio of the total number of charged particles to the total flux of EAS Cherenkov radiation, a ratio of $E_{thr} \\ge 1$~GeV muon flux density at the distance of $600$~m from a shower core to charged particle flux density, a ratio of the energy transferred to the electromagnetic component of EAS to the primary particle energy is presented. Their comparison with two-component mass composition of cosmic rays (p-Fe) in the framework of calculations by a QGSJET model is given. ",
        "introduction": "The irregularities in the cosmic ray (CR) energy spectrum of ``knee'' type at $E_{0} \\simeq 3 \\times 10^{15}$~eV and ``ankle'' type at $E_{0} \\simeq 8 \\times 10^{18}$~eV found in~\\cite{ref1, ref2} are yet of special interest from the point of view of interpretation of these phenomena from the position of astrophysics. In recent years a few papers~\\cite{ref3, ref4, ref5} have been published which try to explain such a behavior of CR spectrum with the help of new models of generation and propagation of CRs. There exists also another namely a nuclear-physical point of view for the formation of irregularity at $E_{0} \\simeq3 \\times 10^{15}$~eV~\\cite{ref6, ref7}. In some sense, the answer to the problem on reasons of the formation of breaks is in the detailed study of different EAS characteristics in the region of the first and second irregularities in the spectrum. We have made such a work in~\\cite{ref8, ref9, ref10, ref11, ref12, ref13}. ",
        "conclusions": "At the Yakutsk complex  array for almost the 35-year period of continuous observations the unique experimental data on electron, muon and Cherenkov EAS components in the region of superhigh and ultrahigh energies have been accumulated. Results on longitudinal and radial development of EAS are presented in figures. From the total combination of data one can select two energy regions where, as is seen from Figures, the characteristics of EAS have the complex dependence on the energy. This is the energy ranges $10^{15} - 10^{17}$~eV and $10^{18} - 10^{19}$~eV. As is known, in these energy ranges the irregularities of ``knee'' and ``ankle'' in the energy spectrum of EAS are observed.  From the comparison of all experimental data with calculations by the QGSJET model in~\\cite{ref17} the results of the cosmic ray mass composition in the energy range $10^{15} - 3 \\times 10^{19}$~eV have been obtained. It follows from the analysis that after the ``knee'' the mass composition becomes heavier and in the region of ``ankle'', on the contrary, becomes lighter. Such a conclusion doesn't contradict the hypothesis on a cosmic ray generation up to the energy $\\sim 10^{18}$~eV in our Galaxy and their propagation according to the model of anomalous diffusion in the fractal interstellar medium. Beginning with $E_{0} > 3 \\times 10^{18}$~eV the mass composition becomes lighter and it doesn't contradict the presence of cosmic rays metagalactic origin in the total flux of cosmic rays."
    },
    "0711/0711.2595_arXiv.txt": {
        "abstract": "{ Although known for almost a century, the photophoretic force has only recently been considered in astrophysical context for the first time. In our work, we have examined the effect of photophoresis, acting together with stellar gravity, radiation pressure, and gas drag, on the evolution of solids in transitional circumstellar disks. We have applied our calculations to four different systems: the disks of HR 4796A and HD 141569A, which are several Myr-old AB-type stars, and two hypothetical systems that correspond to the solar nebula after disk dispersal has progressed sufficiently for the disk to become optically thin. Our results suggest that solid objects migrate inward or outward, until they reach a certain size-dependent stability distance from the star. The larger the bodies, the closer to the star they tend to accumulate. Photophoresis increases the stability radii, moving objects to larger distances. What is more, photophoresis may cause formation of a belt of objects, but only in a certain range of sizes and only around low-luminosity stars. The effects of photophoresis are noticeable in the size range from several micrometers to several centimeters (for older transitional disks) or even several meters (for younger, more gaseous, ones). We argue that due to gas damping, rotation does not substantially inhibit photophoresis. ",
        "introduction": "The standard (core accretion) scenario of planet formation implies continuous growth of solids in a protoplanetary disk around a young star \\citep[see, \\eg\\ ][]{safronov,wetherill-1980,shu_etal_87,lissauer_93,strom_edwards_93, weidenschilling_cuzzi,papaloizou_etal_99,blum_wurm_2000,wurm_etal_06,henning_etal_06,meyer_etal_06}. Micron-sized dust grows step by step to larger bodies, until asteroid-sized planetesimals are reached. At sizes well below 1~km,  gravitational interactions between the objects play a minor role and the process is largely determined by interactions of solids with the ambient gas in the disk. Gas causes sedimentation, mixing, radial drift \\citep{weidenschilling_77,kley_etal_93,tak_arty,brauer_etal_07} and other effects that all dictate the spatial, size, and velocity distributions of grains and therefore the conditions for their growth. As the disk evolves, both gas and dust are gradually removed from the systems. In this process, the gas-to-dust ratio reduces from $\\sim 100$ in the initial protoplanetary disk stage, at the ages on the order of $1$~Myr, to vanishing values of $\\ll 1$ at the debris disk stage after $\\sim 10$~Myr \\citep{lawson_etal_04,haisch_etal_05,hollenbach_etal_05, takeuchi_etal_05,jayawardhana_etal_06,balog_etal_07,currie-et-al-2007}.  Transition from gas- and dust-rich, optically-thick protoplanetary disks to nearly gas-free, optically-thin debris disks is currently in the focus of interest of both observational and theoretical effort \\citep{wieneke_clayton_83, morfill_83, morfill_88, weidenschilling_cuzzi,strom_etal_93,zuckerman_becklin_93,simon_prato_95,meyer_beckwith_2000, ardila,calvet_etal_05,hueso_guillot_05,augereau_06,bouwman_etal_06,eisner_etal_06}. Examples of transitional objects are TW Hya, HR 4796A and HD 141569A, while the well-known system $\\beta$ Pic already enters into the realm of almost gas-free debris disks \\citep{thebault-augereau-2005}. Resolved transitional disks exhibit radial structure in the form of rings and gaps, often alternating, and there is a lot of debate whether this structure is caused by gravity of hidden planets \\citep{augereau-et-al-1999a,wyatt-et-al-1999,telesco} or by interaction of solids with the ambient gas component \\citep{klahr-lin-2000,tak_arty,besla_wu_07}. The first possibility is supported by the fact that the amount of gas remaining in these systems is probably no longer sufficient to form Jupiter's gas envelope \\citep{chen-kamp-2004} and so, the planet formation process must already be finished. The second hypothesis is substantiated by simulations based on the observational estimates of the gas contents. These show that gas drag, acting together with other forces~-- stellar gravity and radiation pressure~--- could result in segregation of different-sized solids and radial fractionation of dust.  In this paper, we study the motion of different-sized solids under the combined action of stellar gravity, radiation pressure, and gas drag, to which we add photophoresis, an additional radial force that acts on particles in a gas disk which are exposed to the radiation field of the central star. Photophoresis, one of the lesser known forces of physics, has been first described by \\cite{ehrenhaft}. The force is caused by the following process. If an object is embedded in a thin gas and is exposed to an anisotropic radiation field (\\ie of a star), a gas molecule that becomes accommodated by the object's surface  and rejected again will depart from the illuminated -- and therefore warmer -- side on average with a greater velocity then from the dark, colder one. Net momentum is transferred to the object, accelerating it away from the light source \\citep[note that for very small particles, the force can also be attractive -- see][]{tehranian_etal_2001}. Although known for almost a hundred years and successfully used in technical applications, such as the construction of optical traps \\citep[][]{steinbach_etal_04}, it has only recently been analyzed in astrophysical context. The first attempt to examine its influence on planetary formation was made by \\cite{krausswurm}. In their subsequent paper \\citep[][]{krausswurm_cai}, they investigated how photophoresis might affect the formation of chondrules and the survival of Calcium/Aluminum-rich inclusions. Finally, \\cite{krausswurm_06} analyzed the effects imposed by photophoresis on the inner rim of a dusty protoplanetary disk. They assumed the disk to stay optically thick at all times, thus allowing only its inner edge to be influenced by photophoresis. Using an exponential law to describe particle growth and the $\\alpha$ 1+1D turbulent model \\citep[][]{shakura_sunyaev_73} used in \\cite{alibert_05} and \\cite{papaloizou_terquem_99} for the evolution of the gas disk, Krauss and Wurm computed the motion of the inner edge. This is identical to the motion of the smallest particles, since, as will be shown subsequently, for solar luminosity systems, the velocity of outward motion of an object increases with its size in a certain size range (compare Fig. \\ref{radvel}). In essence, they showed that the dust disk's edge first moves outward, pushed by photophoresis. After a few million years of time, it shrinks towards the star again, because gas pressure declines as disk dispersal proceeds.  This paper focuses on transitional disks. In these systems, the conditions are quite favorable for photophoresis to be efficient. Transitional disks are already optically thin, so that the particles are exposed to the strong radiation field. On the other hand, they are still sufficiently gas-rich. We would like to check whether, and to which extent, photophoresis may affect the radial migration of solids.  We use a single-body dynamics approach to investigate the behavior of a two-dimensional swarm of particles, the sizes of which are distributed in an interval ranging from several micrometers to several meters. The size of an individual particle, though, is being kept constant. We develop a simple theory for the long-term behavior of a distribution of solids experiencing gravity, radiation pressure, gas drag and photophoresis. Methodically, our study is similar to \\cite{tak_arty}, whose computations we generalize to the presence of photophoresis. To solve the equation of motion, we employ a modification of the semianalytic scheme used by \\cite{weidenschilling_77}. The results are applied to circumstellar disks with different gas density around stars of different luminosity.  In Sect. 2, we discuss the astrophysical status and parameters of the systems chosen for this study, as well as the formulas used for gas density and temperature in the disks. Section 3 provides formulas for different perturbing forces, including the photophoretic one. Section 4 deals with the equation of motion. In Sect. 5, we analyze the properties of the particles' radial motion and derive a formula for the radii of stable orbits. We also estimate the size ranges in which photophoresis should be taken into account. In Sect. 6, we examine the problem of particle rotation in order to check whether it might nullify the effect of photophoresis. In Sect. 7, the results are summarized and ideas for future research are presented. ",
        "conclusions": ""
    },
    "0711/0711.2076_arXiv.txt": {
        "abstract": "We calculate the scattering effects associated with the interaction of a surface gravity or {\\it f} mode with a thin magnetic flux tube embedded in a realistically stratified medium. We find that the dominant scattered wave is an $f$ mode with amplitude and phase of 1.17\\% and around $49^\\circ$ relative to the incident wave, compared to the values of 0.13\\% and $40^\\circ$ estimated from observations. The extent of scattering into high-order acoustic $p$ modes is too weak to be accurately characterized. We recover the result that the degree of scattering is enhanced as (a) the frequency of the incident wave increases and (b) the flux tube becomes magnetically dominated. ",
        "introduction": "A problem of considerable interest and more recently controversy relates to the influence of magnetic fields on acoustic waves in the near-surface regions of the Sun. In these sub-photospheric magnetic regions, the ratio of magnetic to gas pressure could be very close to unity \\citep[e.g.][]{mcintosh, jefferies}, leading to the contention that magnetic field effects are systematic and significant. Mode conversion \\citep[e.g.][]{spruit,cally00}, is a phenomenon commonly invoked to describe magneto-hydrodynamic (MHD) wave interactions; from one acoustic mode to another, from acoustic modes to Alfv\\'{e}n waves and so on, it is estimated that acoustic energy is redistributed, contributing perhaps to $p$-mode absorption observed in sunspots \\citep{braun95}.  Thin flux tubes are important building blocks of solar magnetic activity, intimately tied to studies of the solar photospheric dynamo \\citep[e.g.][]{cattaneo} and flux emergence \\citep[e.g.][]{cheung}. One way of deconstructing the structure and dynamics of thin flux tubes is through experimental analyses of changes in the statistics of the wave field in the vicinity of these magnetic features. Applying techniques of time-distance helioseismology \\citep{duvall93}, \\citet{duvall06} estimated the phase and amplitude of the scattered $f$ modes from observations of thousands of isolated, small magnetic features. These parameters contain the scattering properties of an `average' flux tube, giving us insights into the nature of these flux tubes.  However, one must apply caution when invoking causative mechanisms to explain these statistics, taking care to understand the nature of wave scattering in the presence of magnetic fields \\citep[e.g.][]{spruit,bogdan95,bogdan96,cally00}. Once a sound theoretical basis is formed, observations characterizing the extent of scatter can be used to place constraints on theories of the structure and dynamics of flux tubes.  In an attempt to construct an analytical framework to decipher the nature of wave interaction with magnetic regions, \\citet{bogdan96} model a magnetic flux concentration in the thin tube limit \\citep[e.g.][]{spruit91} embedded in a truncated polytrope. The flux tube is assumed to be thin enough that variations across its interior are neglected; moreover, the tube must be a sub-wavelength feature and the radius small in comparison to the pressure scale height. The parameters of the polytrope are selected so as to mimic the solar sub-photospheric layers as closely as possible and is truncated a little below the photosphere because the rapidly diminishing density makes the thin tube approximation inapplicable very close to the surface. These simplifying assumptions reduce the generality of the results presented here and in fact, direct numerical solvers might appear to be a more attractive proposition. However, while numerical simulations can undoubtedly be employed to address a broader range of questions, it is also important to develop an analytical treatment of a comparable situation, as we attempt to do in this paper. Such results serve not only as benchmarks for direct computations of these phenomena but also assist us in developing a theoretical appreciation of MHD-wave interactions.  \\citet{bogdan96} have demonstrated that the dominant flux tube mode is the kink mode and that the $f$ mode couples strongly with the flux tube. Moreover, the measurements of \\citet{duvall06} are for $f$ modes. Consequently, the focus of this paper will be to quantify the interaction between an incoming {\\it f} mode and the resulting kink mode exhibited by the flux tube.  In $\\S$ \\ref{mod}, we describe the model in greater detail, followed by a discussion of the method we apply to extract the scattering coefficients in $\\S$ \\ref{coeffs.num.sec}. The theoretical values of the phase and amplitude of the scattered $f$ mode are compared with those obtained by \\citet{duvall06} (also see appendix~\\ref{howtogetit}). We also demonstrate that the scattering process is predominantly restricted to $f$-$f$ while the $f$-$p$ conversions are very weak. We conclude with a summary of our calculations and a discussion of the relevance and importance of these results in $\\S$ \\ref{conclusions.sec}. ",
        "conclusions": "We have presented a model of wave scattering off a magnetic flux concentration. Earlier efforts to model magneto-acoustic interactions \\citep[e.g.][]{gizon06} have been under less realistic conditions because of the related mathematical difficulties. However, it is noted that this model does not capture the complexities of the solar case where phenomena such as radiative heat transfer, azimuthal asymmetries in the tube, the lack of applicability of the thin tube approximation in near-surface regions, local changes in source properties, downflows, etc. probably play an important role in affecting the wave field. Even under relatively idealized conditions as in this model, the manipulations required to recover the scattering properties are non-trivial. Perhaps this indicates that the formalism is sub optimal when keeping in mind the relatively significant effort required to compute in this basis set of Whittaker functions while coping with the restricted generality of the solutions. With the development of methods of computational forward modeling to address issues of wave interactions with magnetic fields \\citep[e.g][]{cameron07,hanasoge07} comes the ease of studying these problems in greater generality. The estimates for the degree of scattering obtained here can be used to test the validity of numerical forward calculations.  It is shown that the contribution to the scattering process by the so-called {\\it jacket} or evanescent modes is entirely non trivial, and therefore inevitable in studies of scatterers that contain subwavelength variations in their spatial structure. It has been realized by \\citet{maynard} that these near-field evanescent modes may be analyzed to image with subwavelength resolution. Thus, future prospects of studying these flux tubes in greater detail are exciting, especially in the context of upcoming satelite missions.  Lastly, to say something about the nature of the average flux tube in the measurements of \\citet{duvall06} based on this analysis, we first assume that the interactions are well captured by the single scattering approximation. This implies that the scattering magnitude will scale linearly with the number of model flux tubes while the phase of the scattered wave is unchanged from the value listed in Table~\\ref{scatter.coeffs}. While our scattering model over estimates the amplitude of scatter by a factor of 8.8, the phase of the  $f$-mode scattering coefficient agrees well with the observations. Several possibilities arise: either the $\\beta \\gg 1$ for flux tubes in the Sun, or the upper boundary conditions of this model are inaccurate, or there are important pieces of physics, like non-linearities, radiative heat transfer etc., not accounted for here. More conclusive results await further investigations."
    },
    "0711/0711.4888_arXiv.txt": {
        "abstract": "{% At low redshift ($z<2$), almost half of the baryons in the Universe are not found in bound structures like galaxies and clusters and therefore most likely reside in a Warm-Hot Intergalactic Medium (WHIM), as predicted by simulations.  Attempts to detect WHIM filaments at cosmological distances in absorption towards bright background sources have yielded controversial results that I review here.  I argue that a secure detection of absorption features by the WHIM is at the limit of the {\\it XMM-Newton} capabilities, but feasible.  A proper characterisation of the whole WHIM belongs to the realm of future X-ray missions.} ",
        "introduction": "In the current cosmological paradigm, the content of the Universe is dominated by the \"dark sector\" (Dark Matter and Dark Energy), with only $\\sim 4.5\\%$ being in the form of baryonic matter.  This distribution is observationally supported by the angular shape of the Cosmic Microwave Background anisotropies and by counts of type Ia Supernoave as a function of redshift. The small fraction of baryons was inferred long ago by Big Bang nucleosynthesis being responsible for the formation of light elements.  In particular a higher baryon density early Universe would be unable to produce enough Deuterium (and $^3$He) to the level required by its measured abundance and a lower baryon density Universe would not generate enough $^4$He. It is a remarkable success that the value of the baryon density inferred from D/H measurements in the Lyman $\\alpha$ forest towards distant QSOs agrees with that derived from the power spectrum of the anisotropies in the CMB, with a value $\\sim 4.5$\\% (for a Hubble constant $H_0=70\\, {\\rm km}\\, {\\rm s}^{-1}\\, {\\rm Mpc}^{-1}$, a value used throughout this paper).   As noted in a number of works, at redshift $z>2$ the amount of baryons found in the Universe is roughly consistent with the above estimate (see, e.g., Fukugita, Hogan \\& Peebles 1998). A significant fraction of them reside in the Lyman $\\alpha$ absorption systems, and in particular in the Damped Lyman $\\alpha$ Absorbers or DLAs. At lower redshift, it is expected that the amount of baryons contained in these absorbers decreases, as they should be forming stars.  The number of absorbers per unit redshift decreases substantially towards lower redshift in such a way that the amount of baryons contained in the Lyman $\\alpha$ absorbers is negligible at $z=0$. However, the baryon density found in stars, interstellar and the intra-cluster medium in the local Universe is significantly lower than the 4.5\\% of baryons contained in the Universe. This has been summarized by Nicastro et al (2005a) with the conclusion that the grand total of detected baryons at $z<2$ is $2.5\\%$. Therefore this opens a major cosmological issue:  45\\% of the baryons in the local Universe are missing.  On the other hand, it is theoretically inconceivable that all baryons at low $z$ are locked into stars, galaxies and clusters of galaxies. There are at least two reasons for this: first the efficiency of forming galaxies in binding all baryons is surely not 100\\%; second, the evolution of stellar populations in galaxies involves Supernovae, winds and other phenomena that will end up polluting and enriching the surrounding intergalactic medium. Using simulations, it has been confirmed in a number of papers (Ostriker \\& Cen \\break 1996, Cen \\& Ostriker 1999, Dav\\'e\\ et al 1999) that a significant fraction of baryons would reside today in a Warm-Hot Intergalactic Medium (WHIM), with temperatures ranging between $10^5$ and $10^7$K. Using cosmological simulations, Dav\\'e\\ et al. (2001) concluded that this Warm-Hot component of the Intergalactic Medium becomes more abundant towards low redshift, that it becomes hotter with time and that its spatial distribution follows the general Dark Matter filamentary structure. The WHIM is heated to such high temperatures primarily by shocks, as assumed in the above modeling. Extra heating from Supernovae or AGN of the intergalactic baryons has also been invoked by Pen (1999) and Wu, Nulsen \\& Fabian (2001), in such a way that its thermal emission does not exceed the extragalactic component of the soft X-ray background. The required amount of heat deposited in the WHIM by these non-gravitational processes (of the order of 1 keV per nucleon) will increase the temperature of the intergalactic gas to $\\sim 10^7$~K.  Most of the WHIM gas is contained in non-virialized structures. Detailed simulations of the distribution of various ion species have been performed by Fang, Bryan \\& Canizares (2002), showing that higher ionisation species\\\\ tend to appear in higher density environments, but the abundant ions such as OVII or OVIII are to be found in the low-density filaments, and when observed in absorption towards distant X-ray sources might give rise to an ``X-ray forest'' (Hellsten et al 1998). Velocity broadening in these structures is significant ($\\sim 100\\, {\\rm km}\\, {\\rm s}^{-1}$).  In this paper I review some of the attempts conducted during the last few years with the {\\it XMM-Newton} and {\\it Chandra} gratings in attempting to detect WHIM filaments in absorption towards bright background sources. A local WHIM component has been detected and confirmed beyond any doubt.  However, the situation regarding non-local WHIM filament detections is not fully settled. Earlier claims of detections of absorption features at significant redshifts were not confirmed later.  There is now a heated debate around the detection of a couple of intervening WHIM filaments along the line of sight towards Mrk 421, with unclear outcome.  After discussing all this, I will suggest a possible way forward do detect WHIM filaments with contemporary instrumentation. ",
        "conclusions": ""
    },
    "0711/0711.3650_arXiv.txt": {
        "abstract": "Possible origins of the magnetic fields of neutron stars include inheritance from the main sequence progenitor and dynamo action at some stage of evolution of progenitor. Inheritance is not sufficient to explain the fields of magnetars. Energetic considerations point to differential rotation in the final stages of core collapse process as the most likely source of field generation, at least for magnetars. A runaway phase of exponential growth is needed to achieve sufficient field amplification during relevant phase of core collapse; it can probably be provided by a some form of magnetorotational instability. Once formed in core collapse, the field is in danger of decaying again by magnetic instabilities. The evolution of a magnetic field in a newly formed neutron star is discussed, with emphasis on the existence of stable equilibrium configurations as end products of this evolution, and the role of magnetic helicity in their existence. ",
        "introduction": "At all stages of stellar evolution some percentage of stars  appears to be strongly magnetic: the magnetic A- and B -stars, magnetic white dwarfs, and neutron stars. Do these fields have a common origin or are different mechanisms at work? What kinds of mechanisms do we know of, and which would be the most plausible one for neutron stars? Three possible origins have been discussed for pulsar magnetic fields: \\begin{itemize} \\item{} the fossil field hypothesis, also called `flux conservation', \\item{} fields generated by  a dynamo process at some stage in the evolution of the progenitor star, \\item{} the thermomagnetic effect in the neutron star crust. \\end{itemize} The last one will not be discussed here, since the effect is probably not enough to produce pulsar and magnetar fields of the required strength. See, however, Ho et al. (2004).  \\subsection{Fossil fields} The simplest and most popular hypothesis, flux conservation, is that neutron star magnetic fields are simple remnants of the field of their main sequence (MS) progenitors. How does this work out quantitatively? The strongest fields known in MS stars are around $10^4$G. For a MS progenitor of radius $4\\,10^{11}$cm, (10M$_\\odot$) the star would contain a magnetic flux of $\\sim 5\\, 10^{27}$ G\\,cm$^{2}$. A neutron star with the same flux would have surface field strength of $5\\,10^{15}$G, sufficient for a magnetar. But this is an unrealistically optimistic estimate. The neutron star contains only some 15\\% of the progenitor's mass, and the inner $1.4$M$_\\odot$ of the progenitor occupy only some 2\\% of the star's crossection. If the field is roughly uniform, the flux contained in the core would correspond to $10^{14}$G on a neutron star, insufficient for a magnetar.  The second problem is statistics: only a very small fraction of the progenitors has magnetic fields as large as $10^4$G. Magnetars, on the other hand, are born frequently. As shown in P.\\ Wood's talk (this volume), their birth rate is probably comparable to that of normal neutron stars.  The fossil hypothesis is therefore not very attractive, at least not for magnetars. Magnetar fields might have a different origin than normal pulsars, of course, but that raises the obvious question why a process operating in magnetar progenitors could not work, at lower strength, in those of pulsars as well.  In any case, the `flux conservation' of a fossil field would have interesting consequences for the rotation of pulsars, since this scenario implies that a strong magnetic field is present in the star during its entire evolution. This includes the giant/supergiant phases in which the envelope rotates very slowly. Since this envelope also contains almost all the moment of inertia of the star, any coupling between core and envelope will quickly spin the core down to the same low rotation rate. The internal magnetic field needed to explain the magnetic flux of pulsars is more than enough to provide such coupling. Except in the very latest stages of evolution, when the Alfv\\'en travel time through the star becomes longer than the evolution time scale, the core will corotate with the envelope. The angular momentum left in the core is then much too low to explain to rotation rates of pulsars.  This would disqualify the fossil field hypothesis unless the rotation of pulsars is due to something else, for example `birth kicks', as proposed by Spruit and Phinney (1998). Either the angular momentum, or the magnetic field of pulsars can be fossil, but not both.  Birth kicks would generically lead to correlation between spin axis and direction of proper motion (independent of whether the magnetic field is a fossil or not). Interestingly, such correlations are now being found (Rankin 2006). As in the case of the Crab and Vela pulsars, there appears to be a general preference for alignment of proper motion and spin axis on the plane of the sky. In the analysis of Spruit and Phinney, this means that the kicks are of relatively long duration, long compared with the spin period of the (proto-)neutron star at the time of the kicks (Romani 2005, Wang et al. 2006).  \\subsection{Dynamos} A magnetic field might be generated in the core of a SN progenitor at some earlier stage of its evolution, and subsequently amplified during core collapse (for example simply by `flux conservation'). Possible processes are:  \\subsubsection{Convection} Large scale magnetic fields generated by dynamo processes are traditionally associated with convective zones.  Field strengths of such dynamos are often estimated by assuming equipartition of magnetic energy density with the kinetic energy density of convection, $B^2=4\\pi\\rho v^2$ where $v$ is the convective velocity. The energy flux in convection is $F\\approx\\rho v^3$, and if this carries the star's luminosity $L$, the equipartition field strength is found as \\beq B_{\\rm e}\\approx M^{1/6}L^{1/3}r^{-7/6}\\approx 10^9\\,G\\,({M\\over M_\\odot})^{1/6} L_{38}^{1/3} r_8^{7/6},\\label{beq} \\eeq where the mean density $3M/4\\pi r^3$ has been assumed for $\\rho$. Compressing such a field from $r=10^8$ to $10^6$ cm would yield $B=10^{13}$G. This would be sufficient for a normal pulsar.  It is not at all clear that this is a reasonable estimate, however. In the Sun, the {\\em large scale} (dipole) field is some 300 times smaller than equipartition with convection would suggest. If this carries over to a stellar interior, convection-generated fields in the progenitor would be insufficient for producing pulsars fields.  It is also quite unclear if a field generated by a convective zone, even if it were strong enough, would leave a permanent dipole moment in the core when it becomes stably stratified again. Does the magnetic field generated in a convective region simply retreat, together with the convection, or can it leave a net magnetic field behind? Since expansion and retreat of a convective zone is slow compared with the characteristic dynamo reversal time, it is likely to leave behind, if anything a series of magnetic zones of different orientation instead of a net large scale field (much like the famous magnetic stripes along midocean ridges in the earth's crust).  \\subsubsection{Field generation in stable zones} Convection is not an essential ingredient in magnetic field generation in stars. The energy of differential rotation is sufficient. `Closing of the dynamo loop' can be achieved by instabilities in the magnetic field itself, in much the same way as `magnetorotational' fields  are generated in accretion disks (Hawley et al. 1995).  Contrary to conventional wisdom, field generation can thus take place also in stably stratified zones of stars (Spruit 2002, Braithwaite and Spruit 2006). The stable stratification, however, strongly limits the radial length scale on which the process operates (just as it does with hydrodynamic processes, cf. Zahn 1983, 1992). For this reason it is unlikely to produce a large scale field.  As in the case of convection it is unlikely that such a process would produce a net dipole moment of much significance for the neutron star descendant. It would, however, be of critical importance for the evolution of the angular momentum distribution in the progenitor, since even weak magnetic fields with small radial length scales can exert torques that are significant on evolutionary time scales (Heger et al. 2005).  \\subsubsection{Neutrino-driven convection} The convective velocities in the above estimates are driven by the (radiative) luminosity of the star. Much more powerful convection takes place during core collapse, driven by the much higher neutrino flux. The luminosity is then of order $10^{52}$ erg/s, the size of the core some $3\\,10^6$, which yields  (from eq. \\ref{beq}) $B_{\\rm e}\\sim 10^{15}$ G, i.e. in the magnetar range (Thompson \\& Duncan 1993). However, if the actual dipole field generated is as small, compared with this equipartition number, as it is in the Sun, it would be 300 times smaller. Even after contraction from $3\\,10^6$ to $10^6$cm, this would produce a field of only $10^{14}$ G, marginal for a magnetar, though sufficient for ordinary pulsars.  \\subsection{Magnetic fields during collapse: energy estimates}  The arguments above suggest that field generation in the progenitor followed by compression during core collapse is not the most promising scenario for the production of pulsar fields. The field generation processes in the progenitor are important for providing a `seed field', but it still has to be amplified  by a large factor.  Suppose a useful initial magnetic field has somehow been generated in the pre-SN core. If the collapse proceeds with flux conservation, $B\\sim 1/R^2$, the magnetic energy $E_B\\approx B^2R^3$ increases as $1/R$, i.e. as a constant fraction of the gravitational binding energy $E_{\\rm G}$ of the core, $E_{\\rm G}\\sim GM^2/R$. Simple flux conservation therefore cannot increases a dynamically insignificant initial field into a dynamically significant field. A field of of $10^{15}$ G in the newly formed neutron star ($R=10^6$) would require a field of $10^{11}$G in a pre-collapse core of $R=10^8$ cm. As I have argued above, neither a fossil field nor an internally generated (pre-collapse) field is likely to be this high.  It is possible to do better, at least in simple energetic terms, by exploiting differential rotation of the star. Under angular momentum conservation the rotational energy varies as \\beq E_\\Omega={1\\over 2}I\\Omega^2={1\\over 2}J^2/I\\sim 1/R^2.\\eeq This increases more rapidly than the gravitational energy. For the generation of a magnetic field,  only the energy fraction in {\\em differential} rotation is relevant: \\beq E_{\\Delta\\Omega}=({\\Delta\\Omega\\over\\Omega})^2E_\\Omega\\sim 0.1E_\\Omega,\\sim{1/R^2}\\label{eom} \\eeq where I have assumed $\\Delta\\Omega\\approx 0.3\\Omega$ as a plausible rate of differential rotation.  Assuming there is a process that can convert all energy of differential rotation into magnetic energy $E_B=B^2R_{\\rm NS}^3\\sim E_{\\Delta\\Omega}$, this would produce a maximum field strength inside the star \\beq B_{\\rm max}\\approx 10^{17}{\\Delta\\Omega\\over\\Omega}{\\Omega_{\\rm NS}\\over\\Omega_{\\rm K}},\\label{bmax}\\eeq where $\\Omega_{\\rm NS}$ is the spin of the neutron star end product and $\\Omega_{\\rm K}$ the Kepler frequency at the neutron star surface. The observable surface field strength would be lower than this internal field, and the conversion of differential rotation into magnetic energy less than 100\\% efficient. These efficiency factors are hard to guess, but it suggests that field strengths of the order of a magnetar field are plausible, provided the collapsing core contains enough angular momentum.  How large can the angular momentum of the collapsing core be? An upper limit is set by the average energy of observed core collapse supernovae, of the order $10^{51}$ erg. If the rotational energy of the neutron star formed is larger than this, and the the star strongly magnetic, spindown by pulsar emission will be so fast that the rotational energy is dumped already into the supernova itself. This sets an upper limit of $10^{51}$ erg to the rotational energy, corresponding to a rotation period of 4ms for a neutron star of 1.4M$_\\odot$, or $\\Omega_{\\rm NS}/\\Omega_{\\rm K}\\approx 0.1$. The maximum field strength from (\\ref{bmax}) is then $3\\,10^{15}$ G, assuming again $\\Delta\\Omega/\\Omega=0.3$. This shows that magnetar  field strengths, up to $10^{15}$, are possible, but require a fairly efficient conversion of differential rotation into magnetic energy.  A way of estimating the angular momentum available during core collapse is to follow the evolution of the internal rotation in the progenitor up to core collapse. This requires faith in one's quantitative understanding of the angular momentum transport  processes in the star. Internal torques due to magnetic fields are probably more important than purely hydrodynamic processes alone. Heger et al. (2005) have computed the initial spins of neutron stars with the prescription for magnetic torques from Spruit (2002). Rotation periods in the range 8-10ms were obtained, somewhat independent of the initial (main sequence) spin. This is comfortably below the upper limit given above, but still rapid enough for magnetar fields of $10^{15}$G.  \\subsubsection{Consequences and extrapolations} If differential rotation is an essential ingredient for the production of a strong magnetic field, as suggested above, at least magnetars must be formed spinning rapidly. If a magnetar with $B=3\\,10^{14}$G is formed in this way, and becomes visible through the supernova debris after a couple of years, say, it would spin at about $P\\sim 0.1$s, and would be more luminous than the Crab pulsar. It is not clear if this is compatible with observations of supernovae in nearby galaxies.  On the other hand, it is clear from pulsars with independent age estimates (independent from the spindown time scale) that normal pulsars are formed with a range of initial spins, from $\\sim 15$ms (Crab) to 400 ms (Camilo et al. 2007, Gotthelf \\& Halpern 2007, see also Gotthelf, this volume). This shows that there is either a significant spread in angular momentum of the collapsing core, or that an efficient post-collapse spindown process operates, such as friction against a fall-back disk (cf. Li \\& Jiang 2007, Wang et al. 2007).  Some of the slow-born pulsars have rather weak fields (Camilo et al. 2007, Gotthelf \\& Halpern 2007). This  would agree with angular momentum of the pre-SN core playing an important role in determining the field strength of a neutron star, as suggested above.  The rapid initial spin required for magnetar fields can be extrapolated to higher rotation rates. At sufficiently rapid rotation, the spindown of the nascent, highly magnetized neutron star could power a Gamma-ray burst. This magnetar-GRB scenario has recently become popular, cf. Kommissarov \\& Barkov (2007), Bucciantini et al. (2007), Yu \\& Dai (2007) and references therein. ",
        "conclusions": "Three popular scenarios for the origin of  magnetic fields in pulsars and magnetars are i) the fossil field hypothesis, ii) fields generated internally in the progenitor, and iii) field generation during core collapse. Fossil fields are ruled out for magnetars (required field strength of the MS progenitor too large, number of magnetic MS progenitors too small). If a fossil field plays a role for normal pulsars, magnetic coupling between core and envelope of the progenitor will  remove the angular momentum from the core, producing neutron stars with extremely long rotation periods. The spin of pulsars must then be due to the `birth kicks' that also give rise to their proper motions.  Internally generated fields, followed by core collapse under flux conservation, is unattractive because of the rather low field strengths that can plausibly be produced. The most plausible source of energy for field generation is the differential rotation during core collapse; it increases more rapidly than even the gravitational energy. This predicts that the most magnetic stars are also formed rotating intrinsically rapidly (but suffering rapid spindown by a pulsar wind already during the supernova, cf. Metzger et al. 2007).  The most critical factor for the production of  a strong field is the finite time available in the final collapse phase, on the order of seconds. At earlier times in the collapse the energy in differential rotation is not yet high enough, at later times (when the neutrinos have left) there may still be enough energy in differential rotation, but the stratification has then become so stable that even fields of magnetar strength cannot reach the surface any more.  Combined with the lowish initial field strengths in the pre-collapse core, this implies the existence of a very efficient field amplification process during collapse. Linear winding-up in differential rotation is too slow except possibly for normal pulsars. Runaway growth of a magnetic field is possible by the combination of differential rotation with magnetically driven or shear-driven instabilities. Magnetorotational instability works fastest because it is essentially a shear-driven instability operating on the short time scale of differential rotation. The simulations of Akiyama et al. (2003) and Ardeljan et al. (2005) illustrate the feasibility of such a process.  Most simulations of magnetic core collapse (e.g. Shibata et al. 2006, Obergaulinger et al. 2006, Sawai et al. 2007, Burrows et al. 2007, see also the more analytical considerations in Uzdensky and MacFadyen 2007) start with rather optimistic assumptions on the strength and configuration of the pre-collapse field (an aligned dipole of $10^{11}$G at $R=10^8$ cm for example). Such initial states produce interesting results in themselves. But to explain magnetar fields, or  magnetically powered supernovae, let alone the magnetic field strengths needed to produce magnetically powered GRB jets, much more has to happen  than the flux freezing collapse of a dipole field.  It is likely that pre-collapse fields with more realistic configurations and strengths, and processes like those seen in Moiseenko et al. (2006), will produce magnetic fields that behave rather differently. In particular it is still an open question how the ordered, rotation-aligned magnetic field configurations are formed that work so well in producing jets (or if they form at all).  Since the field of the neutron star is formed in a highly dynamic event, it is likely to be far from equilibrium, or at least far from stable. For the same reason it is also likely to be dominated by higher multipoles rather than the dipole component that determines pulsar phenomenology.  For field strengths of $10^{13}$ G and less magnetic instabilities are slow enough for the field to survive until it can be anchored by the newly formed crust (rapid rotation helps here, since the Coriolis force reduces the growth rate of instabilities).  Fields of magnetar strength decay rapidly, and are too strong to be anchored by the crust. Their apparent survival implies that they have found a stable equilibrium configuration within seconds after their formation. The formation of such stable equilibria from complex, strongly unstable initial fields can be demonstrated with numerical simulations. The key ingredients for the existence of such equilibria are i) the stable stratification of a neutron star, and ii) a nonvanishing magnetic helicity.   \\begin{theacknowledgments} I thank Chris Thompson, Mal Ruderman and Joseph Taylor for discussions during the conference, from which the present text has benefited. \\end{theacknowledgments}"
    },
    "0711/0711.4907_arXiv.txt": {
        "abstract": "The atmospheric structure of chemically peculiar stars deviates from that of normal stars with similar fundamental parameters due to unusual chemistry, abundance inhomogeneities and the presence of strong magnetic field. These effects are not considered in the standard model atmospheres, possibly leading to large errors in the stellar parameter determination and abundance analysis. To tackle this problem we used the state-of-the-art opacity sampling model atmosphere code {\\tt LLmodels} to calculate comprehensive grid of new generation model atmospheres for magnetic CP stars. This grid covers the whole parameter space occupied by SrCrEu and Si-peculiar stars, taking into account characteristic temperature dependence of the chemical abundances. Here we present the first results of our model atmosphere calculations. ",
        "introduction": "To calculate realistic models of chemically peculiar stars we have used a 1-D, LTE, hydrostatic model atmosphere code {\\tt LLmodels} (Shulyak. et al. \\cite{llmodels}). This code allows one to treat the line opacity without simplifying approximations and in this way accurately estimate how modification of the line opacity due to anomalous abundances, chemical stratification or the presence of strong magnetic field affects stellar atmospheres.  Our calculations cover the $T_{\\rm eff}$ range between 6500 and 18000~K with a step of 250--500~K and the $\\log g$ range of 3.25--4.5~dex with a step of 0.25~dex. For $\\log g=4.0$ models we have performed additional calculations for 5~kG magnetic field using the method described by Kochukhov et al. (\\cite{llmag}). Individual abundances of 41 elements were adopted based on the results of modern chemical composition analyses of the SrCrEu and Si-type CP stars. For Ca, Si, Cr, Fe and several other elements we include empirical temperature dependence of the element concentrations (Ryabchikova \\cite{r05}). The line lists are extracted from the VALD database (Kupka et al. \\cite{vald}). The Ap-star model grid is complemented with the solar composition models, calculated for an extended range of atmospheric parameters (6000--20000~K in $T_{\\rm eff}$ and 2.5--4.5~dex in $\\log g$). ",
        "conclusions": ""
    },
    "0711/0711.4300_arXiv.txt": {
        "abstract": "We study the procedure of the reconstruction of phantom-scalar field potentials in two-field cosmological models. It is shown that while in the one-field case the chosen cosmological evolution defines uniquely the form of the scalar potential, in the two-field case one has an infinite number of possibilities. The classification of a large class of possible potentials is presented and the dependence of  cosmological dynamics on the choice of initial conditions is investigated qualitatively and numerically for two particular models. ",
        "introduction": "The discovery of cosmic acceleration \\cite{cosmic} has stimulated the construction of a class of dark energy models \\cite{dark,quintessence,darkmodel} describing this effect. Notice that the cosmological models based on  scalar fields were considered long before the observational discovery of cosmic acceleration \\cite{before}. The dark energy should possess a negative pressure such that the relation between pressure and energy density $w$ is less than $-1/3$. The present day value of the parameter $w$ is close to $-1$ (cosmological constant) but some observations indicate that the value of this parameter slightly less than $-1$ provides the best fit. The corresponding dark energy has been named phantom dark energy \\cite{phantom}. According to some authors, the analysis of observations permits the existence of the moment when the universe changes the value of the parameter $w$ from $w > -1$ to $w < -1$ \\cite{phant-obs,phant-obs1}. This transition is called ``the crossing of the phantom divide line''. The most recent investigations  have shown that the phantom divide line crossing is still not excluded by the data \\cite{obs-new}.  It is easy to see, that the standard minimally coupled scalar field cannot give rise to the phantom dark energy, because in this model the absolute value of  energy density is always greater than that of pressure, i.e. $|w| < 1$. A possible  way out of this situation is the consideration of the scalar field models with the negative kinetic term. Thus, the important problem arising in connection with the phantom energy is  the crossing of the phantom divide line. The general belief is that while this crossing is not admissible in simple minimally coupled models its  explanation requires  more complicated models such as ``multifield'' ones or models with non-minimal coupling between scalar field and gravity (see e.g. \\cite{divide}).  In preceding papers some of us \\cite{we,we1} described the phenomenon of the change of sign of the kinetic term of the scalar field implied by the Einstein equations. It was shown that such a change is possible only when the potential of the scalar field possesses some cusps and, moreover, for some very special initial conditions on the time derivatives and values of the considered scalar field approaching to the phantom divide line. At the same time, two-field models  including one standard scalar field and a phantom field can describe the phenomenon of the (de)-phantomization under very general conditions and using rather simple potentials \\cite{two-field,two-field1}. In the present paper we would like to attract  attention to a drastic difference between two- and one-field models.  The reconstruction procedure of the scalar field potential models is well-known \\cite{Star,Barrow,we-tach,Yurov,Ohta,Zhuk,Szydlowski,Vernov}. We recapitulate it briefly. The cosmological evolution in a flat Friedmann model with the metric \\begin{equation} ds^2 = dt^2- a^2(t)dl^2 \\label{Fried} \\end{equation} can be given by the time evolution of the Hubble parameter $h(t) \\equiv \\frac{\\dot{a}}{a}$, which satisfies the Friedmann equation \\begin{equation} h^2 = \\varepsilon, \\label{Fried1} \\end{equation} where $\\varepsilon$ is the energy density and a convenient normalization of the Newton constant is chosen. Differentiating equation (\\ref{Fried1}) and using the energy conservation equation \\begin{equation} \\dot{\\varepsilon} = -3h(\\varepsilon + p), \\label{en-conserv} \\end{equation} where $p$ is the pressure, one comes to \\begin{equation} \\dot{h} = -\\frac32 (\\varepsilon + p). \\label{h-dot} \\end{equation} If the matter is represented by a spatially homogeneous minimally coupled scalar field, then the energy density and the pressure are given by the formulae \\begin{equation} \\varepsilon = \\frac12\\dot{\\phi}^2 + V(\\phi), \\label{energy} \\end{equation} \\begin{equation} p = \\frac12\\dot{\\phi}^2 - V(\\phi), \\label{pressure} \\end{equation} where $V(\\phi)$ is a scalar field potential. Combining equations (\\ref{Fried1}), (\\ref{h-dot}), (\\ref{energy}), (\\ref{pressure}) we have \\begin{equation} V = \\frac{\\dot{h}}{3} + h^2, \\label{poten} \\end{equation} and \\begin{equation} \\dot{\\phi}^2 = -\\frac23 \\dot{h}. \\label{phi-dot} \\end{equation} Equation (\\ref{poten}) represents the potential as a function of time $t$. Integrating equation (\\ref{phi-dot}) one can find the scalar field as a function of time. Inverting this dependence we can obtain the time parameter as a function of $\\phi$ and substituting the corresponding formula into equation (\\ref{poten}) one arrives to the uniquely reconstructed potential $V(\\phi)$. It is necessary to stress that this potential reproduces a given cosmological evolution only for some special choice of initial conditions on the scalar field and its time derivative \\cite{we-tach,we}. Below we shall show that in the case of two scalar fields one has an enormous freedom in the choice of the two-field potential providing the same cosmological evolution. This freedom is connected with the fact that the kinetic term has now two contributions. In order to examine the problem of phantom divide line crossing we shall be interested here in the case of one standard scalar field and one phantom field $\\xi$ , whose kinetic term has a negative sign. In this case the total energy density and pressure will be given by \\begin{equation} \\varepsilon = \\frac12\\dot{\\phi}^2 - \\frac12\\dot{\\xi}^2 + V(\\phi,\\xi), \\label{energy1} \\end{equation} \\begin{equation} p = \\frac12\\dot{\\phi}^2 - \\frac12\\dot{\\xi}^2 - V(\\phi,\\xi). \\label{pressure1} \\end{equation} The relation (\\ref{poten}) expressing the potential as a function of $t$ does not change in form, but instead  of equation (\\ref{phi-dot}) we have \\begin{equation} \\dot{\\phi}^2 - \\dot{\\xi}^2  = -\\frac23 \\dot{h}. \\label{phi-dot1} \\end{equation} Now, one has rather a wide freedom in the choice of the time dependence of one of two fields. After that the time dependence of the second field can be found from equation (\\ref{phi-dot1}). However, the freedom is not yet exhausted. Indeed, having two representations for the time parameter $t$ as a function of $\\phi$ or $\\xi$, one can construct an infinite number of potentials $V(\\phi,\\xi)$ using the formula (\\ref{poten}) and some rather loose consistency conditions. It is rather difficult to characterize all the family of possible two-field potentials, reproducing given cosmological evolution $h(t)$. In the present paper, we describe some general principles of construction of such potentials, then we consider some concrete examples. ",
        "conclusions": "We have considered the problem of reconstruction of the potential in a theory with two scalar fields (one standard and one phantom) starting with a given cosmological evolution. It is known ( see e.g. \\cite{Star,Barrow,we-tach}) that in the case of the only scalar field this potential is determined uniquely as well as the initial conditions for the scalar field, reproducing the given cosmological evolution.  Changing the initial conditions, one can find a variety of cosmological evolutions, sometimes qualitatively different from the ``input'' one (see e.g. \\cite{we-tach}). In the case of two fields the procedure of reconstruction becomes much more involved. As we have shown here, there is a huge variety of different potentials reproducing the given cosmological evolution (a very simple one in the case, which we have explicitly studied here). Every potential entails  different cosmological evolutions, depending on the initial conditions.   It is interesting that the existence of different dynamics of scalar fields corresponding to the same evolution of the Hubble parameter $h(t)$ can imply some observable consequences connected with the possible interactions of the scalar fields with other  fields. Indeed, in this case, the time dependence of the scalar fields considered above can directly affect physically observable quantities. We are going to consider this topic elsewhere \\cite{we-future}."
    },
    "0711/0711.0523_arXiv.txt": {
        "abstract": "We outline the steps needed in to calibrate the Monte Carlo code in order to perform large scale simulations of real globular clusters. We calibrate the results against $N$-body simulations for $N = 2500$, $10000$ and for the old open cluster M67. The calibration is done by choosing appropriate free code parameters. ",
        "introduction": "Most of the modeling of individual star cluster has been focused on static models based on the King model (see Meylan \\& Heggie 1997). There have been also a small number of studies able to follow the dynamical evolution of a system. They were mainly based on variants of a Fokker-Planck technique and small $N$-body simulations (e.g. Grabhorn \\etal\\  1992, Drukier 1993, Murphy \\etal\\ 1998, Giersz \\& Heggie 2003 and Hurley \\etal\\  2005).  In this presentation we show the further developments of the Monte Carlo code needed to properly follow the evolution of real star clusters. The dynamical ingredients of the code are basically the same as those described in Giersz (2006 and references therein). This code is based on an original code by Stod\\'o{\\l}kiewicz (1986), which in turn was based on the code devised by H\\'enon (1971). The extensions to the code were mainly connected with: (i) upgrading the prescription of stellar evolution to the algorithm of Hurley \\etal\\ (2000), (ii) adding the procedures for the internal evolution of binary stars (Hurley \\etal\\ 2002), both using the McScatter interface (Heggie, Portegies Zwart \\& Hurley 2006), (iii) a better treatment of the escape process in the presence of a static tidal field. Then the new version of the code was calibrated against the results of $N$-body simulations so as to allow us to construct dynamical evolutionary models of real star clusters. ",
        "conclusions": "It was shown that the Monte Carlo code can be successfully calibrated against small $N$-body simulations. Calibration was done by choosing the free parameters describing the relaxation process, such as: coefficient in the Coulomb logarithm $\\gamma = 0.02$, minimum deflection angle $\\beta_{min}$, time step $\\tau$, and coefficient in the formula for the critical energy of escaping stars, $\\alpha = 1.5 - 3.0 (ln(\\gamma N)/N)^{1/4}$. The calibrated code successfully reproduced the $N$-body simulations of the old open cluster M67 (Hurley \\etal\\ 2005), which was the main objective of the calibration procedure. The code is able to provide as detailed data as the observations do. However, it showed also some weaknesses, e.g. some important channels of blue stragglers formation are not present (coalescence of binaries due to their dynamical interactions) and too crude treatment of the escape process. The work is in progress to cure these problems (e.g. a few body direct integrations). It was shown also that the Monte Carlo code can be used to model evolution of real star clusters and successfully compare results with observations (see Heggie 2008, in this volume). The very high speed of the code makes it an ideal tool for getting information about the initial parameters of star clusters. It is worth to mention that to complete the model of the M67 cluster only about seven minutes are needed!"
    },
    "0711/0711.0647_arXiv.txt": {
        "abstract": "{ {\\rm Context.} Atomic diffusion is believed to cause the abundance anomalies observed in AmFm stars. However, the detailed process has still not been well-established. For instance, two possible scenarios for the diffusion theory are presently envisaged. They differ mainly by the depth from which the abundance anomalies emanate. The first scenario predicts that the abundances are modified in the superficial regions of the star, just below the hydrogen convection zone. The second scenario predicts that a much deeper extension of the mixing zone exists due to the convection caused by Fe accumulation in regions below the hydrogen convection zone.\\\\ {\\rm Aims.}  We calculate much more accurate radiative accelerations of Sc than previously, to better understand the observed abundance anomalies of this element. We believe that it is a key element to use as a diagnostic tool for understanding AmFm stars. \\\\ {\\rm Methods.} The method employed to obtain these radiative accelerations is based on an interpolation from the parameters of the so-called SVP parametric method. \\\\ {\\rm Results.} The radiative accelerations, shown here in a typical Am stellar model, are discussed in light of the observed anomalies of Ca and Sc. Our results suggest that the deeper mixing scenario is not entirely satisfactory: the mixing zone should be deeper than what is predicted by recent models to account for observed Sc underabundances. Our results seem more  compatible with the scenario where the abundances anomalies are created in the superficial regions. However, only detailed evolutionary modelling with mass loss and diffusion of all important species, including Ca and Sc, with accurate radiative accelerations, will be able to give more insight into where the source of these anomalies occur in AmFm stars. ",
        "introduction": "Among the main observational peculiarities of Am stars are the superficial underabundances of Ca and Sc, as well as an overabundance of iron-peak elements (e.g. Preston 1974). The classical explanation of the AmFm star phenomenon is that these stars, which are slow rotators, are supposed to have smaller large-scale (mixing) motions in their envelope than \\emph{normal} main sequence stars. Since these stars are almost always in binary systems, their slow rotation is believed to be due to tidal forces (Abt 1961). The relative stability of their envelope allows for a large proportion of He atoms to sink towards the deeper regions due to gravitational settling, leading to the disappearance of the He convection zone. The abundance anomalies were then thought to be produced by the diffusion process (Michaud 1970) at the bottom of the hydrogen convection zone (Watson 1971, Smith 1971, Michaud et al. 1983, Alecian 1996). At that depth, Ca is mostly in its Ar-like configuration so its radiative acceleration is weak because of its noble gas configuration, thus leading to its underabundance in the superficial hydrogen convection zone. The radiative acceleration ($g_{rad}$) is due to the momentum transfer from photons to atoms during photoabsorption. The value of $g_{rad}$ of a given element depends on its absorption cross-section, on its abundance and on the local physical conditions of the plasma (e.g. Alecian \\& LeBlanc 2000). The diffusion process that can create over or underabundances at various depths is a non linear, time-dependent process that can affect the structure and the evolution of stars.  Numerical modelling of Ca and Sc time-dependent diffusion in AmFm stars was carried out by Alecian (1996). This modelling corresponds to the scenario first proposed by Watson (1971), but including a bulk velocity due to mass loss (as considered for $\\lambda$ Bootis stars by Michaud \\& Charland 1986). Superficial abundances of these elements appear to depend closely on the age of the star, on the mass loss rate, and they are very sensitive to the position of the bottom of the superficial mixing zone (hydrogen convection zone plus overshooting). In these computations, the stellar model was static and no structure evolution was considered (the star was kept at the same position on the main sequence), only diffusion was time-dependent.  Recent evolutionary models including the effect of the diffusion of the elements show that a Fe convection zone can occur due to its accumulation at certain depths (Richard, Michaud \\& Richer 2001 and Richer, Michaud \\& Turcotte 2000), and suggest that the abundance anomalies emanate from deeper layers (at temperatures of approximately $200\\,000$ K). At that depth, a large proportion of Ca is also in a noble gas configuration (Ne-like) which explains its observed underabundance. However, these computions are done without mass loss. Mass loss has two effects: it smooths the abundance inhomogenietes, and it shifts the abundance stratifications produced by diffusion towards shallower layers in the star (Michaud \\& Charland 1986; Alecian 1996). Therefore, mass loss may affect iron accumulation and thus the Fe convection zone predicted by these evolutionary models.  The study of Ca diffusion is relatively easy because its atomic data are rather well known. Indeed, a good knowledge of the atomic data is essential since radiative accelerations are, with gravity, among the main ingredients in the diffusion equation. The situation for Sc is quite different. For instance, both the OPAL opacities (e.g. Rogers \\& Iglesias 1992) and the Opacity Project (hereafter OP, Seaton et al. 1992) do not include scandium in their databases. The Kurucz (1990) data include only the first nine Sc ions. The detailed study of Sc diffusion has never been done. Alecian (1996) calculated the $g_{rad}$ of Sc based on the parametric method of Alecian \\& Artru (1990a,b). But, as mentioned in his study, the calculated $g_{rad}$  are not accurate enough since extrapolations had to be done to complete the data for highly ionized Sc ions. In this paper, we aim to calculate the $g_{rad}$ of Sc by an interpolation method based on the parametric form of the $g_{rad}$ formulae of Alecian \\& LeBlanc (2002) and LeBlanc \\& Alecian(2004) (hereafter respectively Papers I and II).  Because scandium underabundance, along with the one of calcium, characterizes AmFm stars, it is a key element to better understand the AmFm phenomenon. Scandium is in an ionisation state with noble gas configuration in layers close to the bottom of the H convection zone, and close to the bottom of the Fe convection zone predicted by Richer, Michaud \\& Turcotte (2000), but not at the same positions as for the corresponding calcium configurations. Any model of AmFm stars should be able to explain underabundances for both elements. Therefore, the study of Sc diffusion could possibly shed light on which one of the two diffusion scenarios, diffusion below the H convection zone while including mass loss (Alecian 1996; hereafter \\emph{scenario A}), and diffusion below the Fe convection zone (Richer, Michaud \\& Turcotte 2000; hereafter \\emph{scenario B}), is more likely to explain the abundance anomalies of these stars.  In Sec.~\\ref{sec:grad}, we present $g_{rad}$ calculations for Sc while using an interpolation method related to the new parametric formulae described in Papers I and II. In Sec.~\\ref{sec:abundances}, we discuss our $g_{rad}$ with respect to the observed abundance anomalies of Ca and Sc in AmFm stars. A short conclusion follows. ",
        "conclusions": "\\label{sec:conclusion}  The radiative accelerations of Sc presented here shed new light on the source of Am stars' abundance anomalies and represents a new diagnostic tool to test various theoretical models. Scandium could to be a key element to better understanding these stars.  Our results rekindle the question as to where the abundance anomalies emanate from in Am stars: either from a superficial or a much deeper mixing zone. The results shown here regarding $g_{rad}$ of Sc suggest that for \\emph{scenario B} to be able to reproduce both the observed Ca and Sc abundances, the mixing must be deeper than that found in the models of Talon, Richard \\& Michaud (2006) or those of Richard, Michaud \\& Richer (2001) for instance. In these models, the mixing goes down to approximately $\\log\\,T = 5.3$, while from the results shown here, it has to go down to approximately $\\log\\,T = 5.7$ to be able to reproduce both Ca and Sc observed abundances.  Relating to the question of dilution of the abundances during the subgiant phase of evolution, it is not clear if this argument can eliminate \\emph{scenario A}. Even though the surface underabundances are created at much shallower regions of the star in this scenario compared to \\emph{scenario B}, there should also exist an underabundance of both Ca and Sc in the two regions where the noble gas configurations dominate. The value of the overabundance that should prevail between these two regions as well as the depth of the mixing zone will determine to what extent the surface abundances are affected at the subgiant stage. Only detailed evolutionary calculations while including the deeper mixing at the subgiant stage could show how the surface abundances of \\emph{scenario A} are affected at that stage of evolution. The predicted abundances of other elements should also give insight into the amplitude of the dilution.  Our results show that near $\\log\\,T = 4.6$ (for abundances of 0.1 times solar values), the $g_{rad}$ of Ca are stronger than those of Sc. \\emph{Scenario A} could then produce Ca and Sc abundances consistent with their relative observed abundances in Am stars. The results presented here for Sc thus reinforce the superficial mixing zone scenario but only evolutionary stellar models with mass loss and diffusion, while including Sc, might help to finally solve the Am stars' mysteries. We hope that the $g_{rad}$ of Sc shown here, or others calculated directly from reliable atomic data, will be included in evolutionary models in the near future. For that purpose, the parameters for Sc calculated here are made available at http://www.umoncton.ca/leblanfn/grad."
    },
    "0711/0711.0537_arXiv.txt": {
        "abstract": "{The introduction of infrared arrays for lunar occultations (LO) work and the improvement of predictions based on new deep IR catalogues have resulted in a large increase in sensitivity and in the number of observable occultations.} {We provide the means for an automated reduction of large sets of LO data. This frees the user from the tedious task of estimating first-guess parameters for the fit of each LO lightcurve. At the end of the process, ready-made plots and statistics enable the user to identify sources that appear to be resolved or binary, and to initiate their detailed interactive analysis.} { The pipeline is tailored to array data, including the extraction of the lightcurves from FITS cubes. Because of its robustness and efficiency, the wavelet transform has been chosen to compute the initial guess of the parameters of the lightcurve fit. } {We illustrate and discuss our automatic reduction pipeline by analyzing a large volume of novel occultation data recorded at Calar Alto Observatory. The automated pipeline package is available from the authors.} {} ",
        "introduction": "For decades, lunar occultations (LO) have occupied a special niche as a technique for high-angular resolution with excellent performance, but relatively inefficient yield. The diffraction fringes that are created by the lunar limb as it occults a background source, provide a unique opportunity to achieve milliarcsecond angular resolution with single telescopes also of relatively small diameter. In terms of instrumentation, LO have always been simple, requiring only a fast photometer. Of course, they have the significant drawback that only sources included in the apparent lunar orbit can be observed (about 10\\% of the sky), and then only at arbitrary fixed times and with limited opportunities for repeated observations. If one adds that each observation only provides a one-dimensional scan of the source, it is clear that detailed and repeated observations are better performed with long-baseline interferometry (LBI), when available. One should, however, not forget additional important advantages of LO: even for complicated sources, the full, one-dimensional brightness profile can be recovered according to maximum-likelihood principles without any assumptions on the source's geometry  (Richichi \\cite{1989A&A...226..366R}). Besides, the limiting sensitivity achieved in the near-IR by LO at the 1.5\\,m telescope on Calar Alto is K$\\approx 8$\\,mag (Richichi et al.~\\cite{2006A&A...445.1081R}). When extrapolated to a 4-meter class telescope or larger, LO appear  quite competitive with even the most powerful, LBI facilities (Richichi~\\cite{1997IAUS..158...71R}).  As a result, although the trend is understandably to develop more flexible, powerful and complex interferometric facilities, there is some balance that makes LO still attractive at least for some applications. It should not be forgotten that the majority of the hundreds of directly-measured stellar angular diameters (Richichi~(\\cite{vlti_2007}) listed 688, and the numbers keep increasing) were indeed obtained by LO, and that LO are still the major contributor to the discovery of small separation binary stars.  Two recent developments, however, have provided a significant boost to the performance of the LO technique, and have significantly enlarged its range of applications: a) the introduction of IR array detectors that can be read out at fast rates on a small subarray has made it possible to provide a large gain in limiting sensitivity, and b) IR survey catalogues that have led to an exponential increase of the number of sources for which LO can be computed. Literally, thousands of occultations per night could now be potentially observed with a large telescope. We describe in this paper the details and impact of these two factors for LO work. We also address the new needs imposed on data reduction by the potential availability of a large volume of lunar occultation data per night, by describing new approaches to an automated LO data pipeline. We illustrate both the new quality of LO data and their analysis by means of examples drawn from the observation of two recent passages of the Moon over crowded regions in the vicinity of the Galactic Center, carried out with array-equipped instruments at Calar Alto and Paranal observatories. ",
        "conclusions": "The observation of lunar occultation (LO) events with modern infrared array detectors at large telescopes, combined with the use of infrared survey catalogues for the predictions, has shown that even a few hours of observation can result in many tens if not hundreds of recorded occultation lightcurves. The work to bring these data sets to a stage where an experienced observer can concentrate on accurate interactive data analysis for the most interesting events is long and tedious.  We have designed, implemented, and tested an automated data pipeline that takes care of extracting the lightcurves from the original array data (FITS cubes in our case); of restricting the range from the original tens of seconds to the few seconds of interest near the occultation event; of estimating the initial guesses for a model-dependent fit; of performing the fit; and finally of producing compact plots for easy visual inspection. This effectively reduces the time needed for the initial preprocessing from several days to a few hours, and frees the user from a rather tedious and error-prone task. The pipeline is based on an algorithm for automated extraction of the lightcurves, and on a wavelet-based algorithm for the estimation of the initial parameter guesses.  The pipeline has been tested on a large number of simulated lightcurves spanning a wide range of realistic signal-to-noise ratios. The result has been completely satisfactory: in all cases in which the algorithm converged, the derived lightcurve characterization was correct and consistent with the simulated values. Convergence could not be reached due to poor signal-to-noise ratio in only in ten cases out of 60000. These cases would, of course, be challenging for an interactive data analysis by an experienced observer as well. We also tested the pipeline on a set of real data, with similar conclusions. We identified and discussed the cases that may prove problematic for our scheme of automated preprocessing."
    },
    "0711/0711.1707_arXiv.txt": {
        "abstract": "The idea is considered that a classical non-Abelian gauge field can be considered as a dark matter candidate. It is shown that Yang-Mills equations have solutions with such distribution of the mass density that allows to describe a rotational curve of spiral galaxies. The conditions necessary for such consideration are considered. One parameter is estimated from Yang-Mills equations and from astrophysical observations (stars rotation curve). The agreement is to within 1\\%. ",
        "introduction": "The electromagnetic field behaves at different conditions either as pure quantum or almost as a classical field. But for non-Abelian gauge fields we know only quantum manifestations of these fields. Here we would like to present the idea that classical non-Abelian gauge fields can become visible as DM. \\par The problem of the DM nature is now one of the most  fundamental problems in modern physics. The reason is to explain the invisibility of DM it is necessary to involve such kinds of an exotic matter as, for example, WIMP, supersymmetric particles etc. \\par In this article we offer a new model of DM. In the studied model the DM is an interior of an abortive singularity or black hole. We understand this suggestion as follows. \\par In Ref's~\\cite{Obukhov:1996ry}~\\cite{Dzhunushaliev:1999fy} it is shown that the energy density of spherically symmetric solutions of Yang-Mills equations may have non-standard behavior (behind an exception, of course,  't Hooft - Polyakov monopole). For example, in Ref.~\\cite{Dzhunushaliev:1999fy} it is shown that the energy density is weakly decreasing. It is natural that such distribution of the matter will give diverging solutions by switching on the gravity. For some boundary conditions at the center it will be an interior of a non-Abelian back hole and for other values of the boundary conditions a singularity located at some distance from the center will appear. But one very interesting subtlety exists here. This solutions are classical ones. The careful analysis shows that at the infinity these solutions are strongly  oscillating  in the space (gravity is switched off) and the period of the oscillations increases by moving away from the centre. It is obvious that on some distance from the center the quantum fluctuations of gauge field on the period distance of oscillations become comparable with a field magnitude. It means that at this distance the gauge field becomes quantum one. The problem in calculating of such  distribution of the non-Abelian gauge field is that the gauge field becomes essentially non-perturbative and it is impossible to apply the perturbative Feynman diagram techniques to its description. \\par The idea presented here is that the radius where the gauge field becomes quantum can be less than the radius of a singular or an event horizon corresponding to a non-Abelian black hole. In this case this object does become neither a singularity nor a black hole (abortive singularity/black hole). It looks as follows: in space  there is a sphere filled with a classical gauge field and on border of this sphere the non-Abelian gauge field becomes quantum one and the mass contained in this sphere is not enough for formation of a singularity or an event horizon. The color gauge field in the sphere does not interact with the elementary particles because the particles are colorless and consequently this classical field is invisible and can be applied as a candidate of DM. \\par In fact in this paper we consider the idea that in a non-Abelian gauge theory may exist such space distribution of a gauge field that classical and quantum phases exist simultaneously but spatially separated. In this case the jump condition from classical phase to quantum one is the strong oscillations of classical non-Abelian field. These oscillations leads to the fact that quantum fluctuations become essential at a distance of oscillation period. In other words the space is filled with a non-perturbative gauge vacuum where there exist defects filled with the  classical gauge field. The galaxies are located in these defects and the classical gauge field is the DM. ",
        "conclusions": "In this paper we have considered the idea that the problem of DM probably can be connected with the problem of non-perturbative quantization of strongly interacting fields. It allows us to connect one problem in macroscopical physics (the problem of DM) with another problem in microscopical physics (confinement problem on quantum chromodynamics). In this connection R. Kolb in Ref.~\\cite{Kolb} write: ``Dark matter and dark energy are two of the binding cords I will use to illustrate how collaborations of astronomers and high energy physicists on large astronomical projects can be good for astronomy, and how discoveries in astronomy can guide high-energy physicists in their quest for understanding nature on the smallest scales.''. \\par The features of the model of colored DM presented here are: \\begin{itemize} \\item The estimation of the parameter $\\alpha$ from the Yang-Mills equations (see Appendix \\ref{alpha}) and from the astrophysical (see Appendix \\ref{asymptotic}) point of view is in agreement with remarkable accuracy. \\item On the background of a non-perturbative vacuum of non-Abelian gauge field there exists a bubble of the same  classical gauge field. The classical non-Abelian gauge field filled the bubble is colored DM. \\item Spherically symmetric classical solutions of the Yang-Mills equations have weakly decreasing mass density distribution leading to a good agreement with the Universal Rotation Curve. \\item The distribution of these classical gauge field is that on some distance from the center the transition from classical phase to quantum occurs. \\end{itemize} The problems for the future investigations are: \\begin{itemize} \\item The theoretical estimation of $V_0^2$ parameter and comparing it with the fitted value \\eqref{3-91}. \\item The fitting of the Universal Rotational Curve \\eqref{5b-30} using the function joining \\eqref{5b-50} and \\eqref{3-78}. \\item The most important problem in the model presented here is the calculation of gauge field distribution using a non-perturbative quantization technique. \\item The search for possibility of the classical gauge fields detection  . \\end{itemize}"
    },
    "0711/0711.3534_arXiv.txt": {
        "abstract": "{We performed a new calibration of the \\strom metallicity index $m_1$ based on the $b-y$ color of cluster red giant stars. The current Metallicity-Index-Color (MIC) relation is not linear in the color range $0.40 \\lesssim b-y \\lesssim 1.0$, but provides iron abundances of cluster and field red giants with an accuracy of $\\sim$ 0.25 dex. ",
        "introduction": "Empirical calibrations of the \\strom metallicity index $m_1$=$(v-b)-(b-y)$ were usually based either on field stars (Anthony-Twarog \\& Twarog 1998, hereafter ATT98) or on a mix of cluster and field stars (Schuster \\& Nissen 1989; Hilker 2000). However, empirical spectroscopic evidence suggest that field and cluster stars present different heavy element abundance patterns (Gratton, Sneden \\& Carretta 2004). Moreover, the occurrence of CN and/or CH rich stars in globular clusters (GCs, Grundahl, Stetson, \\& Andersen 2002) along the red giant (RG, Hilker 2000) and the subgiant branches, and the main sequence (Kayser et al. 2006) opens the opportunity for an independent calibration of the \\strom metallicity index only based on cluster stars.  Calamida et al. (2007a, hereafter CA07) provided new empirical calibrations of the $m_1$ metallicity index based only on cluster RG stars. The new Metallicity-Index-Color (MIC) relations have been validated by adopting GC and field RGs with known spectroscopic iron abundances in the Zinn \\& West (1984) scale, and provide metallicity estimates with an accuracy $\\le 0.2$ dex. The main advantage of CA07 MIC relations compared to similar relations available in the literature is that they adopt the $u-y$ and $v-y$ colors instead of the $b-y$. The main advantage of the $u-y$ and $v-y$ colors is the stronger temperature sensitivity, and the MIC relations show a linear and well-defined slope in the $m_1,\\ u-y / v-y$ planes (see CA07). We now perform a new empirical metallicity calibration based on the $b-y$ color, and show that the $m_1,\\ b-y$ relation for cluster RGs is not linear over the color range $0.42<(b-y)_0<1.05$. We then compare the new metallicity calibration with the CA07 $m_1,\\ v-y$ MIC relation.  \\vspace{-0.35truecm} ",
        "conclusions": "We present a new empirical calibration for the \\strom $m_1$ metallicity index based on the $b-y$ color. We included a quadratic color term in the MIC relation in order to account for the non-linearity of the $m_1,\\ b-y$ relation for RG stars in the color range $0.42 < (b-y)_0 < 1.05$. The new MIC relation provides metallicity estimates for both cluster and field RGs with an accuracy of $\\sim$ 0.25 dex.   \\begin{table} \\caption{\\footnotesize{Spectroscopic measurements from Rutledge et al. (1997) and photometric metallicity estimates for the GCs adopted to validate the \\strom metallicity index.}} \\label{table1} \\begin{scriptsize} \\begin{center} \\vspace{-0.55truecm} \\begin{tabular}{lcc} \\hline \\\\ GC  &  $[Fe/H]_{spec}$ & $[Fe/H]_{phot}$ \\\\ \\hline \\\\ NGC~6397 & $-1.91\\pm0.14$ & $-2.04\\pm0.15$ \\\\ NGC~6752 & $-1.54\\pm0.09$ & $-1.67\\pm0.18$ \\\\ NGC~288  & $-1.40\\pm0.12$ & $-1.30\\pm0.11$ \\\\ NGC~362  & $-1.27\\pm0.07$ & $-1.24\\pm0.30$ \\\\ M71\t & $-0.73\\pm0.05$ & $-0.48\\pm0.34$ \\\\ \\\\ \\hline \\end{tabular} \\end{center} \\end{scriptsize} \\end{table}   \\vspace{-0.35truecm}"
    },
    "0711/0711.1641_arXiv.txt": {
        "abstract": "We investigate the observational signatures of the holographic dark energy models in this paper, including both the original model and a model with an interaction term between the dark energy and dark matter. We first delineate the dynamical behavior of such models, especially whether they would have ``Big Rip'' for different parameters, then we use several recent observations, including 182 high-quality type Ia supernovae data observed with the Hubble Space Telescope, the SNLS and ESSENCE surveys, 42 latest Chandra X-ray cluster gas mass fraction, 27 high-redshift gamma-ray burst samples, the baryon acoustic oscillation measurement from the Sloan Digital Sky Survey, and the CMB shift parameter from WMAP three years result to give more reliable and tighter constraints on the holographic dark energy models. The results of our constraints for the holographic dark energy model without interaction is $c=0.748^{+0.108}_{-0.009}$, $\\Omega_{\\mathrm{m0}% }=0.276^{+0.017}_{-0.016}$, and for model with interaction ($% c=0.692^{+0.135}_{-0.107}$, $\\Omega_{\\mathrm{m0}}=0.281^{+0.017}_{-0.017}$ ,$% \\alpha=-0.006 ^{+0.021}_{-0.024}$, where $\\alpha$ is an interacting parameter). As these models have more parameters than the $\\Lambda$CDM model, we use the Bayesian evidence as a model selection criterion to make comparison. We found that the holographic dark energy models are mildly favored by the observations compared with the $\\mathrm{\\Lambda CDM}$ model. ",
        "introduction": "It has been realized that our Universe is experiencing an accelerated expansion, as shown by several astronomical observations \\cite{SN}. The acceleration of the Universe strongly indicates the existence of a mysterious exotic matter, namely the dark energy, which has large enough negative pressure and makes up the largest portion of the total matter in the current Universe. The combined analysis of observational data suggests that the Universe is spatially flat, and consists of approximately 3/4 dark energy, 1/4 dust matter (cold dark matter plus baryons), and negligible amount of radiation \\cite{SN,LSS,CMB}. The simplest candidate of dark energy is the cosmological constant $\\Lambda $ (vacuum energy) which has the equation of state $w=-1$. The cosmological constant-cold dark matter model ($\\Lambda $CDM) works very well, and is in agreement with a large number of resent observations. However, there are two problems in this scenario --- the \\textit{fine-tuning} problem and the \\textit{cosmic coincidence} problem \\cite{coincidence}. The fine-tuning problem asks why the vacuum energy density today is so small ($\\mathrm{10^{-47}GeV^{4}}$) compared with the theoretical value ($\\mathrm{10^{74}GeV^{4}}$) from the quantum gravity. The cosmic coincidence problem is that since the evolution of the energy densities of dark matter and dark energy are so different during the expansion of the Universe, why are they nearly equal to each other today?  The dark energy may be a problem which has to be solved in the context of quantum gravity \\cite{Witten:2000zk}. In the classical gravity theory, the dark energy density (cosmological constant) can be an arbitrary value. However, a complete theory of quantum gravity should be capable of determining the properties of dark energy such as the energy density and the equation of state \\cite{Witten:2000zk}. The holographic dark energy model is an attempt to apply the holographic principle of quantum gravity theory to the dark energy problem \\cite% {Cohen:1998zx,Horava:2000tb,Hsu:2004ri,Li:2004rb}.  It is well known that the holographic principle is an important result from the explorations of the quantum gravity theory and string theory \\cite% {holoprin}, enlightened by investigations on the properties of black holes. For an effective field theory in a box of size $L$ with UV cut-off $\\Lambda _{c}$, the entropy $S$ scales extensively as $S\\sim L^{3}\\Lambda _{c}^{3}$. However, considering the peculiar thermodynamics of black hole \\cite{bh}, the maximum entropy in a box of volume $L^{3}$ may behave nonextensively, growing only as the surface area of the box, i.e.$S\\leq S_{\\text{BH}}\\equiv \\pi M_{\\mathrm{pl}}^{2}L^{2}$ (Bekenstein entropy bound). This nonextensive scaling suggests that quantum field theory breaks down in large volume. To reconcile this breakdown with the success of local quantum field theory, Cohen et al. \\cite{Cohen:1998zx} proposed a more restrictive energy bound. They pointed out that in quantum field theory a short distance (UV) cut-off is related to a long distance (IR) cut-off due to the limit set by forming a black hole. In other words, if the quantum zero-point energy density $\\rho _{\\Lambda }$ is relevant to a UV cut-off, the total energy of the whole system with size $L$ should not exceed the mass of a black hole of the same size, thus we have $L^{3}\\rho _{\\Lambda }\\leq LM_{\\mathrm{pl}}^{2}$. When we take the whole Universe into account, the vacuum energy related to this holographic principle \\cite{holoprin} is viewed as dark energy, usually dubbed holographic dark energy. The largest IR cut-off $L$ is chosen by saturating the inequality so that we get the holographic dark energy density \\begin{equation} \\rho _{\\mathrm{de}}=3c^{2}M_{\\mathrm{pl}}^{2}L^{-2}~,  \\label{Holo} \\end{equation}% where $c$ is a numerical constant, and $M_{\\mathrm{pl}}\\equiv 1/\\sqrt{8\\pi G} $ is the reduced Planck mass. If we take $L$ as the size of the current Universe, for instance the Hubble scale $H^{-1}$, then the dark energy density will be close to the observational result. However, Hsu \\cite% {Hsu:2004ri} pointed out that this yields a wrong equation of state for dark energy. Li \\cite{Li:2004rb} subsequently proposed that the IR cut-off $L$ should be taken as the size of the future event horizon \\begin{equation} R_{\\text{\\textrm{eh}}}(a)=a\\int_{t}^{\\infty }{\\frac{dt^{\\prime }}{% a(t^{\\prime })}}=a\\int_{a}^{\\infty }{\\frac{da^{\\prime }}{Ha^{\\prime 2}}}~. \\label{Horizon} \\end{equation}% Then this problem can be solved, and the holographic model can thus be constructed successfully. For extensive studies of this model (HDE), see Ref. \\cite% {Huang,Enqvist04,Ke,HuangLi,Zhang,Pavon,WangGong,Kim,Nojiri,Hu,Chen,Ma}. The holographic dark energy may also have interaction with matter, this is the so called interacting dark energy model (IHDE) \\cite{Pavon,WangGong}. Both the HDE and IHDE model can be tested by cosmological observations, such as the type Ia supernovae \\cite{HuangGong}, CMB \\cite{Enqvist05,Shen,Kao}, X-ray gas mass fraction in the clusters \\cite{Chang}, differential ages of passively evolving galaxies \\cite{Yi} and combinations of SN Ia, CMB and LSS data \\cite{ZhangWu,WangLin,WuGong,Feng}. Currently, the tightest constraint on the HDE model is given by \\cite{WuGong}, namely $c=0.85^{+0.18}_{-0.02}$ and $\\Omega_{m0}=0.27^{+0.04}_ {-0.03}$, but $c>1$ is still allowed under this constraint. In addition, by using the $\\chi ^{2}$ statistic as a model comparison technique, Ref.\\cite{WuGong} suggests that HDE model is equally favored by the current observational data compared with the $\\Lambda$CDM model.  In this paper we consider the dynamical behavior of the holographic dark energy models, particularly what is the condition for the Universe to end in the so called \\textquotedblleft Big Rip\\textquotedblright\\ \\cite{bigrip}. we then use several new data sets to constraint the holographic dark energy models. These includes a sample of 182 high quality SN Ia data (Despite coincidence in number, this is not exactly the same data set used by Gold06 \\cite{Riess06}), a sample of 42 latest X-rays gas mass fraction ($\\mathrm{% f_{gas}}$) data \\cite{Allen07}, a sample of 27 gamma-ray burst (GRB) data% \\cite{Schaefer07} generated from one of the tightest correlations---the $% E_{peak}-E_{\\gamma }$ correlation \\cite{Ghirlanda04}, the baryon acoustic oscillation measurement from the Sloan Digital Sky Survey\\cite{Eisenstein05}% , and the CMB shift parameter from WMAP 3 years result. With these more diverse set of observations, we can obtain constraints which are not only tighter, but also more reliable.  We compare the observational fit of the HDE and IHDE models with the $% \\Lambda $CDM model. In previous works, the $\\chi ^{2}$ statistic has been used for comparison. However, as the holographic dark energy models (with or without interaction) have more parameters than the $\\Lambda$CDM, the $\\chi^2$ may not be a suitable criterion for making comparison. Here we use the Bayesian Evidence (BE) as the information criterion to assess the strength of the holographic models.  The paper is organized as follows: In \\S 2, we review our models of holographic dark energy, and derive the evolution equations. We also discuss qualitatively the behavior of the HDE and IHDE models, particularly the impact of the interacting term on the fate of the Universe. In \\S 3, we present our method of analysis and the data used (\\S 3A), then briefly discuss the Markov Chain Monte Carlo (MCMC) techniques (\\S 3B) and the model selection criteria (\\S 3C). In \\S 4, we show the results of our constraints, and using the Bayesian evidence we make comparisons between the holographic dark energy models (including the HDE and IHDE) and the $\\Lambda $CDM model. We summarize our result and conclude in \\S 5. ",
        "conclusions": "In this paper we firstly gave a brief review of holographic dark energy model for both the non-interacting case and interacting case. We introduced a new interacting term $Q=3\\alpha H\\rho _{\\mathrm{de}}$ and the non-interacting case could be viewed as the special case with $\\alpha =0$. We derived the equations for the evolution of $\\Omega _{de}$ and $H(z)$, and illustrated the dynamical behavior of these models by chosen some representative values of the parameters $c$ and $\\alpha$. The condition for the model to have ``Big Rip'' is determined.  Secondly, we utilize several data sets from the resent observations to constrain the models. Our data sets consist of the selected 182 high-quality type Ia supernovae, the baryon acoustic oscillation measurement from the Sloan Digital Sky Survey, the latest X-ray gas mass fraction data from $% Chandra$ observations, 27 GRB samples generated with $E_{peak}-E_{\\gamma }$ correlation, and the CMB shift parameter from WMAP three years result. We used the MCMC technique to simulate the posterior probability of the model parameters. The best-fits for the three data sets are given in Table 1 for the HDE model, and Table 2 for the IHDE model. We also give the probability distribution of the parameters and the contour maps for the HDE and IHDE models.  Next, we utilize the Bayesian evidence (BE) as a model selection criterion to compare the holographic models with $\\Lambda $CDM model for the three data sets. The BE is particularly appropriate for comparing models with different number of parameters. Both the HDE and the IHDE model are mildly favored by the current observational data set, although the evidence is weak. For both the HDE and IHDE models, the data favors ``quintom'' behavior slightly, i.e. the dark energy initially has $w_{eff}>-1$, but eventually crossing the phantom dividing line, and the model ends with a ``Big Rip''. However, quintessence-like behavior is also still allowed with the present data.  In brief, we conclude that according to the combined measurements data, the holographic dark energy model, especially the interacting holographic dark energy model is mildly favored by the observations, and for the best fit model the equation of state for both the HDE and IHDE crosses $-1$, for which the Universe ends up in a Big Rip."
    },
    "0711/0711.3158_arXiv.txt": {
        "abstract": "The ANTARES neutrino telescope is presently being built in the Mediterranean Sea at a depth of 2500 m. The primary aim of the experiment is the detection of high energy cosmic muon neutrinos, which are identified by the muons that are produced in charged current interactions. These muons are detected by measuring the Cerenkov light which they emit traversing the detector. Sometimes a high momentum muon produces electromagnetic showers. The subject of this paper is a method to reconstruct these showers which includes several steps: an algorithm for the fit of the muon track parameters, preselection of detected photons belonging to a shower, and a final fit with the preselected detected photons to calculate the electromagnetic shower position. Finally a comparison between data obtained with that part of the detector that is currently in operation and simulations is presented.  \\vspace{1pc} ",
        "introduction": "The ANTARES \\cite{Antares} detection principle relies on the observation of Cerenkov (CK) light emitted by relativistic charged particles in water. The emitted CK photons are detected by a 3-dimensional grid of 900 photomultiplier tubes (PMTs) arranged in 12 detection lines at a depth of 2500 m in the Mediterranean Sea. The detector will be completed in early 2008. Since the end of January 2007, the detector has 5 operating detection lines. Amongst others first data of atmospheric downward going muons have been measured, which are up to $500$ m long tracks with energies larger than $100$ GeV. This gives the unique opportunity to study the average number of showers per track length for high energy muons passing through sea water.   The muon track is reconstructed from the arrival times of the CK light detected by the PMTs, whose positions are known. The muon energy loss mechanism in water are ionization, \\mbox{$e^+e^-$ pair} production, bremsstrahlung, and photonuclear interactions. Above several \\mbox{hundred GeV} the muon energy loss is dominated by the latter three effects \\cite{PDG}, which are discrete and are characterized by large energy fluctuations. Furthermore, the average muon energy loss due to these effects is proportional to the energy of the muon.  In the following a method to reconstruct discrete electromagnetic (EM) showers is presented, which gives additional information about a muon track (e.g. the shower multiplicity) as compared to existing algorithms only measuring the direction and position of the muon track. ",
        "conclusions": ""
    },
    "0711/0711.4587_arXiv.txt": {
        "abstract": "Theories of structure formation in a cold dark matter dominated Universe predict that massive clusters of galaxies assemble from the hierarchical merging of lower mass subhalos.  Exploiting strong and weak gravitational lensing signals inferred from panoramic {\\it Hubble Space Telescope} imaging data, we present a high resolution reconstruction of the mass distribution in the massive, lensing cluster Cl\\,0024+16 at $z = 0.39$. Applying galaxy-galaxy lensing techniques we track the fate of dark matter subhalos as a function of projected cluster-centric radius out to 5 Mpc, well beyond the virial radius. We report the first detection of the statistical lensing signal of dark matter subhalos associated with late-type galaxies in clusters. The mass of a fiducial dark matter halo that hosts an early type L$^*$ galaxy varies from $M = 6.3_{-2.0}^{+2.7}\\,\\times\\,10^{11}\\,M_{\\odot}$ within $r < 0.6$ Mpc, $1.3_{-0.6}^{+0.8}\\,\\times\\,10^{12}\\,M_{\\odot}$ within $r < 2.9$ Mpc and increases further to $M = 3.7_{-1.1}^{+1.4}\\,\\times\\,10^{12}\\,M_{\\odot}$ in the outskirts. The mass of a typical dark matter subhalo that hosts an $L^*$ galaxy increases with projected cluster-centric radius in line with expectations from the tidal stripping hypothesis.  The mass of a dark matter subhalo that hosts a late-type L$^*$ galaxy is $1.06_{-0.41}^{+0.52}\\,\\times\\,10^{12}\\,M_{\\odot}$. Early-type galaxies appear to be hosted on average in more massive dark matter subhalos compared to late-type galaxies. Early-type galaxies also trace the overall mass distribution of the cluster whereas late-type galaxies are biased tracers. We interpret our findings as evidence for the active assembly of mass via tidal stripping in galaxy clusters. The mass function of dark matter subhalos as a function of projected cluster-centric radius, is compared with an equivalent mass function derived from clusters in the Millenium Run simulation populated with galaxies using semi-analytic models. The shape of the observationally determined mass functions based on an I-band selected sample of cluster members and lensing data are in agreement with the shapes of the subhalo mass functions derived from the Millenium Run simulation. However, simulated subhalos appear to be more efficiently stripped than lensing observations suggest. This is likely an artifact of comparison with a dark matter only simulation. Future simulations that simultaneously follow the detailed evolution of the baryonic component during cluster assembly will be needed for a more detailed comparison. ",
        "introduction": "Clusters of galaxies are the most massive and recently assembled structures in the Universe. In the context of the hierarchical growth of structure in a cold dark matter dominated Universe, clusters are the repository of copious amounts of the dark matter. Gravitational lensing, predicted by Einstein's theory of General Relativity, is the deflection of light rays from distant sources by foreground mass structures. In its most dramatic manifestation, strong lensing requires a rare alignment with foreground dense structures and produces highly distorted, magnified and multiple images of a single background source (Schneider, Ehlers \\& Falco 1992). More commonly, the observed shapes of background sources viewed via a foreground cluster lens are systematically elongated, in the so-called weak lensing regime. Strong and weak lensing offer the most reliable probes of the distribution of dark matter on various cosmic scales (Blandford \\& Narayan 1992; Mellier 2002; Schneider, Ehlers \\& Falco 1992).  Strong lensing studies of the core regions of several clusters indicate that the dark matter distribution can be represented by a combination of smoothly distributed, extended cluster mass components and smaller-scale clumps or subhalos associated with luminous galaxies (Kneib et al. 1996; Natarajan \\& Kneib 1997; Natarajan, Kneib, Smail \\& Ellis 1998). The smooth components have been detected using weak lensing techniques out to the turn-around radius (typically of the order of several Mpc) in clusters (Kneib et al. 2003; Gavazzi et al. 2003; Broadhurst et al. 2005; Bradac et al. 2006; Clowe et al. 2006; Wittman et al. 2006; Limousin et al. 2007c; Bardeau et al. 2007). To date, however, attention has largely focused on the lensing derived density profile of the smooth cluster component, and its agreement with profiles computed from high resolution numerical simulations of structure formation in the Universe (Navarro, Frenk \\& White 1997; Navarro et al. 2004; Sand et al. 2004). In fact, the granularity of the dark matter distribution associated with individual galactic subhalos holds important clues to the growth and assembly of clusters. Several earlier studies have explored this issue for the particular case of Cl\\,0024+16 (Tyson, Kochanski \\& dell'Antonio 1998; Broadhurst, Huang, Frye \\& Ellis 2000; Jee et al. 2007; Smail et al. 1996).  The detailed mass distribution of clusters and in particular, the fraction of the total cluster mass associated with individual galaxies has important consequences for the frequency and nature of galaxy interactions in clusters (Merritt 1983; Moore et al.\\ 1996; Ghigna et al.\\ 1998; Okamato \\& Habe 1999). Infalling subhalos suffer a range of violent fates as the strong gravitational potential of the cluster tidally strips dark matter and removes baryons via ram-pressure stripping from them (Cortese et al. 2007). Simulations suggest that subhalos may not be arranged equally around galaxies of different morphologies given their varying histories in the cluster environment (Ghigna et al. 1998; Tormen, Diaferio \\& Syer 1998; Springel, White, Tormen \\& Kauffmann 2001). Moreover, subhalos may become tidally truncated by an amount that will differ substantially over the large dynamic range in cluster density. Observations of tidal stripping offer important clues to key questions regarding the growth and evolution of clusters. How much dark matter is associated with the subhalos in clusters as a function of radius?  To what extent do the luminous cluster galaxy populations trace the detailed mass distribution? And, how significant is tidal stripping for the various morphological galaxy types in the cluster? These are questions we attempt to answer in this work using observational data and by comparing with numerical simulations.  To explore cluster galaxy masses, we exploit the technique of galaxy-galaxy lensing, which was originally proposed as a method to constrain the masses and spatial extents of field galaxies (Brainerd, Blandford \\& Smail 1996), which has been since extended and developed over the years to apply inside clusters (Natarajan \\& Kneib 1996; Geiger \\& Schneider 1998; Natarajan et al. 1998; 2002a; Natarajan, De Lucia \\& Springel 2007; Limousin et al. 2007a).  Previous attempts to measure the granularity of the dark matter distribution as a function of cluster-centric radius from observations have had limited success. Analyzing CFHT (Canada-France-Hawaii Telescope) weak lensing data from the supercluster MS0302+17, Gavazzi et al.\\ (2004) claimed detection of a radial trend in the extents of dark matter subhalos in this supercluster region extending out to a few Mpc from the center. Gavazzi et al. reported that the mass distribution derived from weak lensing was robustly traced by the luminosity of early-type galaxies, although their analysis did not include late-type galaxies or a large-scale smooth component. However, utilizing ground-based CFHT weak lensing data for a sample of massive clusters at $z = 0.2$, Limousin et al.\\ (2005) did not detect any variation of the dark matter subhalo masses with cluster-centric radius out to a significant fraction of the virial radius. The resolution of ground-based data appears to be inadequate to detect this effect.  In this paper, we present the determination of the mass function of substructure in Cl\\,0024+16 (at $z = 0.39$) in 3 radial bins using panoramic Hubble Space Telescope {\\it HST} imaging data. A high resolution mass model tightly constrained by current observations is constructed including individual cluster galaxies and their associated dark matter subhalos. We show that over a limited mass range we can successfully construct the mass function of subhalos inside this cluster as a function of cluster-centric radius. The three bins span from the center to 5 Mpc (well beyond the the virial radius of 1.7 Mpc) providing us insights into the tidal stripping process. We also compare properties of the subhalos that host early-type galaxies with those that host late-type galaxies in Cl\\,0024+16. In addition, we compare the results retrieved from the lensing analysis with results from the largest cosmological simulation carried out so far - the {\\it Millennium Simulation}. N-body simulations in combination with the semi-analytic models that we employ in this work are an invaluable tool for investigating the non-linear growth of structure in detail and to provide insights into the cluster assembly process.  The outline of this paper is as follows: in \\S2, we discuss the theoretical framework of tidal stripping and galaxy-galaxy lensing in clusters; in \\S3 the observations and modeling are described. The analysis for Cl\\,0024+16 is presented in \\S4 including a discussion of the uncertainties; results and the comparison with clusters in the Millennium Simulation are described in \\S5. We conclude with a discussion of the implications of our results for the LCDM model and the future prospects of this work. \\footnote{Throughout this work wherever required we have used the following values for the cosmological parameters: H$_0$ = 72 km/s/Mpc; $\\Omega_{\\rm m}$ = 0.3; $\\Omega_{\\Lambda}$ = 0.7.  At the redshift of Cl\\,0024+16, 1'' = 5.184 kpc.} ",
        "conclusions": "Earlier work on galaxy-galaxy lensing in the field has identified a signal associated with massive halos around typical field galaxies, extending to beyond 100\\,kpc (e.g.\\ Brainerd, Blandford \\& Smail 1996; Ebbels et al.\\ 2000; Hudson et al.\\ 1998; Wilson et al.\\ 2001; Hoekstra et al.\\ 2004). In particular, Hoekstra et al.\\ (2004) report the detection of finite truncation radii via weak lensing by galaxies based on 45.5 deg$^2$ of imaging data from the Red-Sequence Cluster Survey. Using a truncated isothermal sphere to model the mass in galaxy halos, they find a best-fit central velocity dispersion for an $L^*$ galaxy of $\\sigma = 136 \\pm 5$ kms$^{-1}$ (68\\% confidence limits) and a truncation radius of $185 \\pm 30$ kpc. Galaxy-galaxy lensing results from the analysis of the Sloan Digital Sky Survey data (Sheldon et al. 2004; Guzik \\& Seljak 2002; Mandelbaum et al. 2006) have contributed to a deeper understanding of the relation between mass and light. Similar analysis of galaxies in the cores of rich clusters suggests that the average mass-to-light ratio and spatial extents of the dark matter halos associated with morphologically-classified early-type galaxies in these regions may differ from those of comparable luminosity field galaxies (Natarajan et al.\\ 1998, 2002a). We find that at a given luminosity, galaxies in clusters have more compact halo sizes and lower masses (by a factor of 2--5) compared to their field counter-parts. The mass-to-light ratios inferred for cluster galaxies in the V-band are also lower than that of comparable luminosity field galaxies. This is a strong indication of the tidal stripping effect of the dense environment on the properties of dark matter halos. In recent work, using only strong lensing constraints in the inner regions of the Abell cluster A\\,1689 derived from images taken by the Advanced Camera for Surveys (ACS) aboard {\\it HST}, Halkola \\& Seitz (2007) also find independently that the subhalos of cluster galaxies are severely truncated compared to equivalent luminosity galaxies in the field.  The subhalo mass function represents an important prediction of hierarchical CDM structure formation models and has been subject of intense scrutiny since the `satellite crisis' was identified (Moore et al.\\ 1999; Klypin et al.\\ 1999). This crisis refers to the fact that within a radius of $400\\,h^{-1}\\,{\\rm kpc}$, from the Milky Way, cosmological simulations of structure formation predict $\\sim$ 50 dark matter satellites with circular velocities in excess of $50\\,{\\rm kms^{-1}}$ and mass greater than $3\\,\\times\\,10^8\\,M_{\\odot}$.  This number is significantly higher than the dozen or so satellites actually detected around our Galaxy. Several explanations have been proposed to resolve this discrepancy.  The missing satellites could for example be identified with the detected High Velocity Clouds (Kerr \\& Sullivan 1969; Willman et al. 2002; Maller \\& Bullock 2004). {\\it Warm} or {\\it self-interacting} dark matter could also selectively suppress power on the small scales, therefore reducing the predicted number of satellites.  The leading hypothesis however remains that the solution to this problem lies in processes such as heating by a photo-ionizing background that preferentially suppresses star formation in small halos at early times (Bullock et al.\\ 2000; Benson et al.\\ 2002; Kravtsov et al. 2004). On the scale of galaxy clusters, many more dark matter structures are expected to be visible, thus making the comparison with expectation from numerical simulations less affected by uncertainties in the poorly understood physics of the galaxy formation. In earlier work, we showed that there is very good agreement between the lensing observations and the Millenium Run simulations in the inner 1 Mpc or so for a sample of clusters (Natarajan, De Lucia \\& Springel 2007). Now we are able to extend our analysis out to 5 Mpc for Cl\\,0024+16 due to the unique data-set that is available. The diagnostic available for comparison here is the shape of the mass function in each of the three bins. We find that the overall shape of the subhalo mass functions derived from the 2 independent methods is in very good agreement out to beyond the virial radius. We find that the mass of a typical subhalo that hosts an $L^*$ early-type galaxy increases with cluster-centric radius in concordance with theoretical expectations. However, the estimates of the mass of a subhalo that hosts an $L^*$ galaxy derived from simulations is significantly lower than those derived from lensing observations. The origin of this discrepancy lies in the fact that tidal stripping appears to be more efficient in the simulations.  Due to the large area probed by this Cl\\,0024+16 dataset, we are also able to constrain the properties of dark matter subhalos associated with late-type galaxies that preferentially lie in the outer regions of the cluster (Treu et al. 2003). We report the first detection of the presence of a dark matter subhalo associated with late-type galaxies in Cl\\,0024+16. While early-type galaxies appear to trace the overall mass distribution robustly, the subhalos associated with late-type galaxies do not contribute significantly to the total mass budget at any radius. In the cluster Cl\\,0024+16 within 5 Mpc we find the following contributions to the total mass: $\\sim$70\\% of the total mass of the cluster is smoothly distributed, the subhalos associated with early-type galaxies contribute $\\gsim\\,20\\%$, and subhalos hosting late-type galaxies account for the remaining $<\\,10\\%$.  The mass resolution of our technique varies slightly with cluster-centric distance owing to the nature of observational constraints that dominate the likelihood optimization. While the strong lensing constraints in the core are the most stringent and drive the fit in the inner regions, the anisotropy in the shear field is statistically harder to recover. As we progressively step out in radius away from the cluster center, the shear of the large scale smooth component drops, and that of the individual subhalos dominates but the overall summed shear signal is significantly lower than in the inner regions. The current analysis is primarily limited by the quality of the available data. Datasets from the {\\it ACS} will allow mass modeling of lensing clusters at even higher resolutions providing increasing accuracy enabling better mapping of the lower mass end of the mass functions of substructure. However, as described above a vast complement of ground-based observations are also needed for this kind of comprehensive analysis which is extremely time consuming. Ground based data provided many important constraints, for instance, the large number of measured central velocity dispersions for cluster galaxies (Moran et al. 2007) were used as priors in modeling the perturbing subhalos that made the optimization more efficient.  Below we summarize the key results on comparison with simulations, where we mimic-ed the selection process adopted for the observational data of Cl\\,0024+16.  Dividing the simulated clusters drawn from the Millenium Run into 3 equivalent radial bins as the observational data, we were able to estimate (i) the mass function in each bin and (ii) the subhalo masses that host $L^*$ early-type galaxies. The shapes of the lensing derived mass functions are in reasonable agreement with those derived from simulations when we normalize the lensing results to the those of the total number of model Type 1's and Type 2's..  Our results provide strong support for the tidal stripping hypothesis. We also find evidence for the variation in the efficiency of tidal stripping with cluster-centric radius and morphological type. We conclude that dark matter in clusters is assembled by the incorporation of infalling subhalos that are progressively stripped during their journey through the cluster. The finding of kinematically disturbed features in the cluster galaxy population by Moran et al. (2007) corroborates our conclusion. We have significantly improved on previous ground-based studies as space-based data affords greater accuracy in shape measurements. Future space-based surveys coupled with ground-based spectroscopic follow-up will provide an unprecedented opportunity to follow the cluster assembly process."
    },
    "0711/0711.4855_arXiv.txt": {
        "abstract": "Using metallicities from the literature, combined with the Revised Bologna Catalogue of photometric data for M31 clusters and cluster candidates (the latter of which is the most comprehensive catalogue of M31 clusters currently available, including 337 confirmed globular clusters -- GCs -- and 688 GC candidates), we determine 443 reddening values and intrinsic colours, and 209 metallicities for individual clusters without spectroscopic observations. This, the largest sample of M31 GCs presently available, is then used to analyse the metallicity distribution of M31 GCs, which is bimodal with peaks at $\\rm {[Fe/H]}\\approx -1.7$ and $-0.7$ dex. An exploration of metallicities as a function of radius from the M31 centre shows a metallicity gradient for the metal-poor GCs, but no such gradient for the metal-rich GCs. Our results show that the metal-rich clusters appear as a centrally concentrated spatial distribution; however, the metal-poor clusters tend to be less spatially concentrated. There is no correlation between luminosity and metallicity among the M31 sample clusters, which indicates that self-enrichment is indeed unimportant for cluster formation in M31.\\\\ The reddening distribution shows that slightly more than half of the GCs are affected by a reddening of $E(B-V) \\la 0.2$ mag; the mean reddening value is $E(B-V) = 0.28_{-0.14}^{+0.23}$ mag. The spatial distribution of the reddening values indicates that the reddening on the northwestern side of the M31 disc is more significant than that on the southeastern side, which is consistent with the conclusion that the northwestern side in nearer to us. ",
        "introduction": "\\label{Introduction.sec}  The formation and evolution scenarios of the Milky Way Galaxy still remain among the most important unsolved problems in contemporary astrophysics \\citep{per02}. One promising way for us to better understand, and to possibly make progress towards addressing these issues is by studying globular clusters (GCs). GCs are generally considered the fossils of the galactic formation and evolution processes, since they form during the very early stages of their host galaxy's evolution \\citep{bh00}. GCs are very dense, gravitationally bound spherical systems of several thousands to more than a million stars. They can be observed out to much greater distances than individual stars, so that they can be used, and are in fact the ideal tracers, to study the properties of extragalactic systems. The most distant GC systems that have been studied to date are those in the Coma cluster; for instance, \\citet{baum95} presented GC counts in the bright elliptical galaxy NGC 4881 \\citep[at a distance of $\\simeq 108$ Mpc;][]{baum95} based on {\\sl Hubble Space Telescope (HST)}/WFPC2 observations. Using data of similar quality, \\citet{Kavelaars00}, \\citet{harris00} and \\citet{wh00} published a series of papers on the GCs in NGC~4874 and IC 4051, the central cD galaxy and a giant elliptical galaxy in the Coma cluster, respectively.  Located at a distance of about 780 kpc \\citep{sg98, mac01}, M31 is the nearest and largest spiral galaxy in the Local Group of galaxies. It contains 337 confirmed GCs and 688 GC candidates \\citep{gall04}, thus allowing us access to a much larger number of GCs than in our own Galaxy. However, despite the difference in GC numbers, from the observational evidence collected so far \\citep[see, e.g.,] []{rich05}, the M31 GCs and their Galactic counterparts reveal some striking similarities \\citep{ffp94,dj97,bhh02}. More recently, \\citet{kim07} embarked on a new systematic wide-field CCD survey of M31 GCs, and found 1164 GCs and GC candidates -- of which 559 are previously known GCs and 605 newly-found GC candidates; based on survey data from the Canada France Hawaii telescope and Wide Field Camera on the Isaac Newton telescope, \\citet{huxor07} combined his detailed analysis of the M31 GC system with recent results based on the galaxy's stellar halo, and concluded that M31 and the Milky Way are rather more similar than previously thought. Therefore, studying the properties of the GCs in M31 may not only improve our understanding of the formation and structure of our nearest large neighbour galaxy, but also that of our own Galaxy. However, there are also significant differences\\footnote{Huxor (2007) suggests that the primary difference between the Galaxy and M31, and between their GC systems in particular, is likely due to the more vigorous recent merger history of M31.} between the GC systems of the Milky Way and M31: in particular, M31 has a much larger population of GCs than the Milky Way \\citep[see] [and references therein]{kim07}; there are populations of ``faint fuzzies'' and extended GCs in the outer halo of M31 that are not seen in the Milky Way \\citep[] []{huxor05,mackey06,mackey07}, and there is a population of young to intermediate-age GCs in M31 that are (again) not seen in the Milky Way again \\citep[see, eg.,] []{beasley04,puzia05}.  A reliable estimate of the reddening, caused by dust contamination, is important for the study of the stellar population of a given GC, in order to obtain its intrinsic spectral energy distribution. In general, on galaxy-wide scales dust tends to be distributed close the the galactic plane in galactic discs. Therefore, the disc clusters are most affected by extinction due to dust in the galactic disc (in addition to the effects due to Galactic foreground extinction). A reliable estimate of the extinction caused by Galactic material along a given line of sight can be obtained from the reddening maps of \\citet{bh82} and \\citet{Schlegel98}. However, a reliable estimate of the internal extinction in a given galaxy is not easy to obtain. More specifically, although there were some ways in which to deal with the problem of determining the reddening toward the GCs in M31 prior to the publication of \\citet{bh00}, some may not have been fully satisfactory. For example, \\citet{van69} assumed a uniform reddening for the clusters when he pioneered integrated-light spectroscopy of M31 GCs, whereas \\citet{ir85} assumed GCs to have a single intrinsic colour when using them as reddening probes; some authors \\citep[see e.g.][]{km60,ve62,har74,van75,bg77} assumed that the halo GCs in M31 were only affected by Galactic foreground extinction, based on which they then averaged the intrinsic colours of the halo GCs and subtracted this value from the observed colours to obtain the reddening values for all the GCs \\citep[see details in][]{bh00}. There are also two other ways to determine the reddening of M31 clusters, which seem reasonable but they have only been applied to a handful of clusters: \\citet{Frogel80} estimated the individual reddening for 35 clusters using the reddening-free parameter $Q_K$, based on unpublished spectroscopic data by L. Searle; \\citet{Crampton85} used the intrinsic colours of the 35 clusters obtained by \\citet{Frogel80} to calibrate $(B-V)_{\\rm 0}$ as a function of spectroscopic slope parameter $S$ of the continuum between $\\sim$ 4000 and 5000 \\AA, and then determined the intrinsic colours for about 40 GCs and GC candidates.  \\citet{bh00} presented a new catalogue of photometric and spectroscopic data of M31 GCs, and determined the reddening for each individual cluster using correlations between optical and infrared colours and metallicity, and by defining various ``reddening-free'' parameters based on this catalogue. \\citet{bh00} found that the M31 and Galactic GC extinction laws (see their table 6), and the M31 and Galactic GC colour-metallicity relations are similar to each other. They then estimated the reddening to M31 objects with spectroscopic data using the relation between intrinsic optical colours and metallicity as determined for Galactic clusters. For objects without spectroscopic data, they used the relationships between the reddening-free parameters and certain intrinsic colours, based on the Galactic GC data. \\citet{bh00} compared their results with those in the literature and confirmed that their estimated reddening values are reasonable, and quantitatively consistent with previous determinations for GCs across the entire M31 disc. In particular, \\citet{bh00} showed that the distribution of reddening values as a function of position appears reasonable in that the objects with the smallest reddening are spread across the disc and halo, while the objects with the largest reddening are concentrated in the galactic disc. The reddening values for M31 clusters obtained by \\citet{bh00} are widely used \\citep[see, e.g.,][]{jiang03,Ma06b, fan06,rey06}  To study the metal abundance properties of GCs can help us understand the formation and enrichment processes of their host galaxy. For example, if galaxies form as a consequence of a monolithic, dissipative and rapid collapse of a single massive, nearly-spherical spinning gas cloud in which the enrichment time-scale is shorter than the collapse time, the halo stars and GCs should show large-scale metallicity gradients \\citep{eggen62,bh00}. On the other hand, \\citet{sz78} presented a chaotic scheme for the early evolution of a galaxy, in which loosely bound pre-enriched fragments merge with the main body of the proto-galaxy over a significant period, in which case there should be a more homogeneous metallicity distribution. At present, most galaxies are thought to have formed through a combination of both of these scenarios (see also Section 4.5)  {\\sl HST} provides a unique tool for studying GCs in external galaxies. For example, based on data from the {\\sl HST} archive, \\citet{Gebhardt99}, \\citet{Larsen01} and \\citet{kw01} showed that many large galaxies possess two or more subpopulations of GCs that have quite different chemical compositions \\citep[see also][]{west04}. Recently, \\citet{peng06} presented the colour distributions of GC systems for 100 early-type galaxies from the ACS Virgo Cluster Survey, and found that, on average, galaxies at all luminosities appear to have bimodal or asymmetric GC colour/metallicity distributions. The presence of colour bimodality among these old GCs indicates that there have been at least two major star-forming mechanisms in the (early) histories of massive galaxies \\citep{west04,peng06,strader06}.  Based on the newest photometric and spectral data, in this paper we determine reliable reddening values for 443 GCs and GC candidates (the largest GC sample in M31 used to date), and we also determine the metallicities for 209 GCs and GC candidates without spectroscopic observations. We then perform a statistical analysis using this GC sample. We describe the photometric and spectroscopic data for the M31 GCs in Section \\ref{data.sec}. In Section \\ref{red.sec}, we determine and analyse the reddening of our M31 GC sample. Section \\ref{metal.sec} is devoted to our statistical analysis. Finally, the discussion and conclusions of this paper are presented in Section \\ref{Conclu.sec}. ",
        "conclusions": "\\label{Conclu.sec}  In this paper, we have (re-)determined the reddening values for 443 clusters and cluster candidates in M31, as well as metallicities for 209 sample objects without spectroscopic observations. We have followed the methods described by \\citet{bh00}, who found that the M31 and Galactic extinction laws are the same within the observational errors, and that the M31 and Galactic GC C-M relations are also consistent with each other. The sample of spectroscopic and photometric data used in this paper is the newest and largest to date. The spectroscopic data were obtained from the most recent references currently available and the photometric data are from the most comprehensive catalogue of M31 clusters available at present, which includes 337 confirmed GCs and 688 GC candidates. Using the metallicities of the largest sample of clusters and cluster candidates at hand, we studied the properties of the M31 clusters. Our main conclusions are summarised below:  \\begin{enumerate} \\item The reddening distribution shows that slightly more than 50 per cent of the GCs suffer from a reddening of less than $E(B-V)=0.2$ mag, and the mean value is $E(B-V) = 0.28_{-0.14}^{+0.23}$ mag. The spatial distribution of $E(B-V)$ indicates that the reddening on the northwestern side of the M31 disc is greater than that on the southeastern side, which is consistent with the conclusion that the northwestern side in nearer to us.  \\item The metallicity distribution of the M31 GCs is bimodal with peaks at $\\rm {[Fe/H]}\\approx -1.7$ and $-0.7$ dex.  \\item The diagram of metallicities as a function of radius from the M31 centre shows a metallicity gradient for the metal-poor GCs, but no such gradient for the metal-rich GCs.  \\item The metal-rich clusters appear to constitute a centrally concentrated spatial distribution; however, the metal-poor clusters tend to be less spatially concentrated.  \\item There is no correlation between luminosity and metallicity among our M31 sample clusters, which indicates that self-enrichment is indeed unimportant for cluster formation in M31. \\end{enumerate}  We reiterate that in using the method of \\citet{bh00}, there are two major unavoidable assumptions (acknowledged by these authors), i.e. that in the Milky Way and in M31 both the extinction law and the intrinsic colours of the GCs are the same. The latter assumption seems reasonable, since there is no evidence that GCs in different galaxies have different intrinsic colours. Regarding the former assumption, there is inconsistent evidence as to whether or not this is a valid assumption. For example, \\citet{wk88} found that the extinction law in M31 is very similar to that in the Milky Way, by analysing the two major dust lanes on the near side of M31; however, several studies have suggested that the reddening in M31 appears to be peculiar: with $E(U-B)/E(B-V)=1.01\\pm0.11$ \\citep{ir85} and $E(U-B)/E(B-V)\\sim 0.5$ \\citep{Massey95}, compared to 0.72 for the same ratio in the Milky Way. Based on a large sample of GCs with optical and near-infrared photometric data, \\citet{bh00} demonstrated that the $U$- and $K$-band extinction curve of M31 is consistent with that of the Milky Way, with total-to-selective extinction coefficient $R_V=3.1$. In fact, the former assumption is plausible because in the M31 disc the composition and size distribution of the large normal grains which dominate the dust mass may be similar to those in the Milky Way \\citep[see for details,][]{xh96}.  As an example, we will discuss in some detail the reddening value of the M31 GC B037 (a.k.a. 037-B327), which is known to be an extremely red object. There are a few references that discuss this GC, including \\citet{bk02}, \\citet{Ma06}, \\citet{ma06} and \\citet{cohen06}. \\citet{km60} first noticed an extremely red colour in photographic ($P$) and visual ($V$) bands for B037, and determined its absorption to be $A_V=3.90$ mag. Based on the photometric data for M31 star clusters in $U, B$, and $V$ of \\citet{ve62}, \\citet{ve62b} studied the reddening values for these objects and found that B037 was the most highly reddened in his sample, with $E(B-V)=1.28$ mag ($A_V=4.10$ mag). \\citet{Crampton85} calibrated $(B-V)_{\\rm 0}$ as a function of spectroscopic slope parameter $S$ of the continuum between $\\sim 4000$ and 5000 \\AA, and then determined the intrinsic colours for about 40 GCs and GC candidates, including B037. \\citet{Crampton85} presented a reddening value for B037 of $E(B-V)=1.48$ mag. Armed with a large database of multicolour photometry, \\citet{bh00} determined the reddening value for each individual M31 GC, including B037, using the correlations between optical and infrared colours and metallicity based on various ``reddening-free'' parameters, and derived $E(B-V)=1.38\\pm0.02$ mag for B037. Using spectroscopic metallicities to predict the intrinsic colours, \\citet{bk02} rederived the reddening value for this GC, $E(B-V)=1.30\\pm0.04$ mag. Recently, \\citet{Ma06} determined the reddening and age of the B037 by comparing multicolour photometry with theoretical stellar population synthesis models. The reddening towards B037 determined by \\citet{Ma06} is $E(B-V)=1.360\\pm0.013$ mag. The reddening value for B037 determined in this paper is $E(B-V)=1.21\\pm0.03$ mag. It is clear that the consistent reddening values for B037 from different references confirm that this cluster suffers from very large extinction. In fact, \\citet{ma06} showed the dust lane across the face of the cluster using an {\\sl HST}/ACS image, which may partially account for its very large reddening value \\citep[see also][]{cohen06}."
    },
    "0711/0711.2581_arXiv.txt": {
        "abstract": "The Wide Angle Search for Planets (WASP) survey currently operates two installations, designated SuperWASP-N and SuperWASP-S, located in the northern and southern hemispheres respectively. These installations are designed to provide high time-resolution photometry for the purpose of detecting transiting extra-solar planets, asteroids, and transient events. Here we present results from a transit-hunting observing campaign using SuperWASP-N covering a right ascension range of 06hr $<$ RA $<$ 16hr. This paper represents the fifth and final in the series of transit candidates released from the 2004 observing season. In total, 729,335 stars from 33 fields were monitored with 130,566 having sufficient precision to be scanned for transit signatures. Using a robust transit detection algorithm and selection criteria, 6 stars were found to have events consistent with the signature of a transiting extra-solar planet based upon the photometry, including the known transiting planet XO-1b. These transit candidates are presented here along with discussion of follow-up observations and the expected number of candidates in relation to the overall observing strategy. ",
        "introduction": "More than 30 extra-solar planets are now known to transit their parent stars, most of which were discovered through photometric monitoring. Transit surveys which use a narrow-field and a large magnitude depth have had some success; most notably that of the OGLE transiting planets (e.g., \\citet{kon03}). Surveys which utilise the ``wide and shallow'' technique of monitoring relatively bright field stars have achieved recent rapid success, including the discoveries of TrES-3 \\citep{odo07}, XO-2b \\citep{bur07}, and several new HATNet planets (\\citet{kov07} for example). The transit technique has matured through overcoming serious obstacles which were impeding the data analysis, such as improved optimal photometric methods for wide-field detectors \\citep{har04} and reduction of correlated (red) noise \\citep*{pon06,tam05}.  The Wide Angle Search for Planets (WASP) project currently operates two SuperWASP instruments \\citep{pol06}, one in the northern hemisphere on La Palma (SuperWASP-N) and the other (SuperWASP-S) located at South African Astronomical Observatory (SAAO). The relatively large field-of-view (FOV) of the SuperWASP design allows each instrument to monitor a substantial amount of the visible sky with relatively high time resolution. SuperWASP-N has been acquiring data since late-2003, achieving photometric precision of 1\\% for stars brighter than $V \\sim 11.5$.  The process of extracting reliable transit signatures from the SuperWASP data involves a systematic approach of removing trends from the photometry, employing an efficient transit detection algorithm \\citep{col06}, and performing spectroscopic follow-up using both low-medium resolution and high resolution spectrographs to identify false-positives and confirm planetary candidates. This method has already been used successfully on SuperWASP data to detect the planets WASP-1b and WASP-2b \\citep{col07a}. However, this process typically begins with millions of lightcurves from which many transit candidates are found and provides an extremely useful reference source for future surveys which monitor the same fields. For this reason, the transit candidates detected in several RA ranges have been published for the benefit of the transit survey community \\citep{chr06,cla07,lis07,str07}.  We present the results of a photometric search for exoplanetary transits using data from SuperWASP-N, covering the fields in the range 06hr $<$ RA $<$ 16hr. Section 2 describes the 2003/2004 SuperWASP-N observing campaign, the predicted transit recovery rate, and the data reduction techniques. Section 3 discusses the transit detection algorithm and the candidate selection criteria. The results of this search are presented in section 4 with discussion of the candidates which have been followed up and those which are considered high priority. Section 5 discusses the results from the analysis and the number of final candidates in relation to the limitations of the survey for this particular RA range. ",
        "conclusions": "This paper has described the acquisition, analysis, and results from the transit-hunting SuperWASP-N 2004 observing campaign covering a right ascension range of 06hr $<$ RA $<$ 16hr. Even though 56 fields were monitored during this campaign, the lack of baseline coverage resulted in the rejection of 23 fields from which 130,566 stars were extracted to be searched for transit signatures. The selection criteria which were applied to the resulting candidates proved to be exceptionally powerful at removing eclipsing binaries from the candidate list due to the ellipsoidal variations and the estimated planet size based upon the transit depth. Amongst the final candidates is the known transiting planet XO-1b which further demonstrates the strength of the transit sifting process.  From the final list of 6 transit candidates, the photometry alone indicates that 1SWASP J160242.43+290850.1 is in fact likely to be the signature of a eclipsing binary rather than a planet. Furthermore, follow-up spectroscopy of 1SWASP J161732.90+242119.0 shows that there is no significant radial velocity variation, resulting in its rejection as a planet candidate. Further follow-up observations, both photometric and spectroscopic, are being undertaken for these and other fields from the 2004 SuperWASP observing campaign. However, the real strength of the 06--16h range will be realised when the 2004 data is combined with that of subsequent years to create an exceptional baseline for transit hunting."
    },
    "0711/0711.2548_arXiv.txt": {
        "abstract": "The differential energy spectrum of cosmic rays from Cherenkov radiation measurements in EAS in the energy range of $10^{15} - 10^{20}$~eV has been compared with an anomalous diffusion model for the particles in interstellar space having fractal properties (Lagutin et al., 2001). The close association between experimental data and calculated ``all particle'' spectra in form at $E_{0} \\sim (10^{15} - 10^{18})$~eV is found. In this case, the average mass composition of cosmic rays calculated by five components does not contradict the average mass composition from experimental data which was obtained by several of EAS characteristics in that energy region. ",
        "introduction": "The discovery of irregularities in the cosmic ray energy spectrum at the energy of  $\\sim 3 \\times 10^{15}$~eV (Khristiansen et al., 1956~\\cite{1}) and $\\sim 8 \\times 10^{18}$~eV (Krasilnikov et al., 1978~\\cite{2,3,4}), the detection of sharp decreases in the cosmic ray intensity at $E_{0} > 5 \\times 10^{19}$~eV (Greisen-Zatsepin-Kuzmin effect, 1966~\\cite{5, 6}) at the EAS arrays in Yakutsk, HiRes (USA), AUGER (Argentina) are the most important achievement in the investigation of the superhigh and ultrahigh energy cosmic rays in recent years. Such a character of spectrum turn out to be associated directly with processes in interstellar space, namely, with the origin, acceleration and propagation of cosmic rays in the Galaxy and beyond. The interpretation of these experimental facts using the different models of cosmic ray origin still remains to be answered.  In this paper the comparison of the cosmic ray energy spectrum by EAS Cherenkov light measurements at the Yakutsk array~\\cite{7, 8} with the calculations according to an anomalous diffusion model of cosmic rays in interstellar space~\\cite{9} is performed. ",
        "conclusions": "The differential energy spectrum of primary cosmic rays in the interval of $10^{15} - 5 \\times 10^{19}$~eV obtained from a totality of the all Cherenkov detector measurement data at the Yakutsk EAS array is shown on the Fig.~\\ref{fig01}. Our data confirm an irregularity of the spectrum of ``knee'' type in the energy range of $(2-5) \\times 10^{15}$~eV discovered in~\\cite{1}, and the irregularity of ``ankle'' type at $E_{0} \\sim 8 \\times 10^{18}$~eV. It is established that in the first case the spectrum index is $\\gamma = 2.7 \\pm 0.1$ below the break and $\\gamma = 3.03 \\pm 0.05 at E_{0} > 3 \\times 10^{15}$~eV, and in the second case, the more sloping spectrum with $\\gamma = 2.6 \\pm 0.3$ at $E_{0} > 8 \\times 10^{18}$~eV is observed.  \\begin{figure}[ht] \\centering \\includegraphics[width=0.42\\textwidth,clip]{fig01} \\caption{Energy spectrum of primary cosmic rays by measurement data of Cherenkov light at the Yakutsk complex EAS array.} \\label{fig01} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=0.42\\textwidth,clip]{fig02} \\caption{Differential cosmic ray intensity versus the energy. The points are Yakutsk array data, curves are the calculation from~\\cite{9}.} \\label{fig02} \\end{figure}  For the period of continuous observations of Cherenkov radiation more than 30 years (10\\% relative to one year time of EAS registration with the scintillation detectors), the showers with $E_{0} > 6 \\times 10^{19}$~eV did not detect. This fact confirms once more the GZK hypothesis~\\cite{5, 6} about the sharp break in the cosmic ray energy spectrum at $E_{0} > 5 \\times 10^{19}$~eV.  \\subsection*{The galactic model}  The attempt to explain a form of obtained spectrum from the point of view of cosmic ray anomalous diffusion model and fractality of the Galaxy's magnetic field was made by Lagutin et al~\\cite{9}. The basis for the cosmic ray propagation in the Galaxy is the following assumptions: a) after the generation in the sources, the particles move in fractal interstellar medium by means of two ways: the first way is ``Levy flights'', the second way is the motion along a spiral in the nonhomogeneous magnetic field, $b$) the particles exist during anomalous long time. The lifetime of particles is of a wide distribution and its tail is described by a power law $q(t) \\propto B \\cdot t^{-\\beta -1}, t \\to \\infty, \\beta < 1$ (so-called ``Levy trapping time''). Calculations of the spectrum were separately made for each of following groups of nuclei: p, He, CNO, N-Si, Fe. The resulting sum spectrum for the all particles is shown by a solid curve in Fig.~\\ref{fig01}. From the calculations it follows that the suggested model reproduces the irregularity in the energy spectrum of the ``knee'' type at $E_{0} \\simeq 3 \\times 10^{15}$~eV and also the irregularity of the ``ankle'' type at $E_{0} \\simeq 8 \\times 10^{18}$~eV. This model does not explain the behavior of a spectrum in the energy region of $10^{17} - 10^{18}$~eV and the break of the spectrum at $E_{0} > 6 \\times 10^{19}$~eV in more detail. The mass composition in the energy region of $5 \\times 10^{15} - 5 \\times 10^{18}$~eV is some heavier than at $E_{0} \\simeq 10^{19}$~eV, but this change is not very significant that is expected from an experiment (see Fig.~\\ref{fig03}).  \\subsection*{The galactic model with the sources of two types}  In contrast to~\\cite{9}, in the paper~\\cite{15} a scenario is considered, in which supernovae are the main sources of cosmic rays and the acceleration up to $E_{\\text{max}} \\simeq 105 \\cdot Z$~GeV takes place in the shock fronts. The particle spectrum formed in this case can be presented in the form of  $S_{\\text{SN}} \\sim E^{-2} \\theta (E_{\\text{max}} - E)$, where the Heaviside function $\\theta(x)$ reflects qualitatively the presence of a sharp cut-off in the spectrum at $E > E_{\\text{max}}$~\\cite{16, 17}. The new calculations fulfilled by the above scenario of particle generation in the sources of two different types under the assumption of anomalous diffusion of particles in inhomogeneous medium show that at some parameters the anomalous diffusion model describes satisfactorily the features of cosmic ray energy spectrum and mass composition up to $E_{0} \\sim 10^{18}$~eV observed in an experiment. First of all, it refers to the fine structure of cosmic ray intensity change depending on the energy (see Fig.~\\ref{fig01}). By using these calculations, the sharp peaks in the mass composition depending on energy are also observed (see Fig.~\\ref{fig03}). In this connection, it is of interest to compare calculations in mass composition with experimental data obtained at the Yakutsk EAS big and small Cherenkov sub-arrays in recent years.  \\subsection*{Mass composition of primary cosmic rays}  Fig.~\\ref{fig03} presents the results in mass composition of primary cosmic rays of the Yakutsk array. The data were obtained in the framework of QGSJET-01 model and two-component mass composition (proton-iron nucleus). The several characteristics corresponding to the radial and longitudinal development of EAS are used in the analysis~\\cite{18, 19, 20, 21, 22}.  \\begin{figure} \\centering \\includegraphics[width=0.42\\textwidth, clip]{fig03} \\caption{Mass composition of cosmic rays at superhigh and ultrahigh energies. The curve is a calculation by Lagutin et al (2001) according to the anomalous diffusion model for cosmic ray propogation.} \\label{fig03} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=0.42\\textwidth, clip]{fig04} \\caption{Comparison of the experimental spectrum with the calculated spectrum from~\\cite{23} for the metagalactic protons ($E_{0} > 5 \\times 10^{17}$~eV)  and galactic iron nuclei ($E_{0} = 10^{15} - 5 \\times 10^{17}$~eV).} \\label{fig04} \\end{figure}  The value $\\left<\\ln{A}\\right>$ in each case is determined by using the interpolation method~\\cite{24}. It is seen from Fig.~\\ref{fig03} that the mass composition is varied up to heavy elements in the energy region of $(2-5) \\times 10^{17}$~eV and becomes more light beginning with $E_{0} \\sim 3 \\times 10^{18}$~eV.  The lines are the calculations according to the anomalous diffusion model for the propagation of cosmic rays in the Galaxy (Fig.~\\ref{fig03}, two sources) in the case of inhomogeneous galactic medium. In the first case, the monotone change in the mass composition up to $E_{0} \\ge 3 \\times 10^{18}$~eV is observed, after of which the mass composition becomes more light. In the second case, the complicated structure in the dependence of mass composition on the energy is observed, peaks for the nuclei of different mass in the energy region of $10^{15} - 10^{17}$~eV are noticeable. According to a hypothesis~\\cite{25} and calculations from \\cite{15}, such an inhomogeneous structure can be formed by a near supernova. Our data (Fig.~\\ref{fig03}) do not contradict to this hypothesis.  Such a sharp change of the mass composition in the energy region of $5 \\times 10^{16} - 5 \\times 10^{17}$~eV is not explained in the framework of the galactic model and is likely associated with the existence of a transition boundary from galactic to metagalactic cosmic rays. This conclusion is confirmed by calculations from~\\cite{23}, where a scenario of galactic and metagalactic origin of cosmic rays is considered. These calculations are shown in Fig.4 together with our experimental data. It can be seen from Fig.~\\ref{fig04} that cosmic rays in the energy region of $5 \\times 10^{16} - 5 \\times 10^{17}$~eV are most likely of galactic origin with the noticeable portion of heavy nuclei in the total flux.  It should be noted the estimations of cosmic ray mass composition in the region after the ``knee'', obtained at the compact arrays, agree well with each other. The same cannot be said of the energy region of $\\sim 10^{18}$~eV (see Fig.~\\ref{fig03}) where HiRes array data point to more speedy enrichment of primary radiation by the light nuclei and protons as compared with the Yakutsk array data. The Yakutsk EAS array data, on the contrary, show the gradual change from the heavy to light composition (protons and He nuclei) at $E_{0} \\sim 10^{19}$~eV. In both cases, data point to the existence of the tendency of ``protonization'' of primary cosmic rays at $E_{0} >3 \\times 10^{18}$~eV.  \\subsection*{Conclusions}  Direct  measurements of the cosmic ray energy spectrum in the region of ultrahigh energies (in energy scattered by EAS particles in the atmosphere) have confirmed the complicate form of spectrum. The spectrum becomes steeper at $E_{0} \\ge 3 \\times 10^{15}$~eV and more sloping at $E0 \\ge 8 \\times 10^{18}$~eV. A character of energy dependence of $\\left<\\ln{A}\\right>$ by the Yakutsk EAS data point to the change of the mass composition of primary particles at singular points of cosmic ray energy spectrum. The value $\\left<\\ln{A}\\right>$ rises with the energy after the \u201cknee\u201d up to its maximum equal to $3.5$ at $(2-5) \\times 10^{17}$~eV and then it begins to decrease. Such an energy dependence of $\\left<\\ln{A}\\right>$ does not contradict a hypothesis of cosmic rays propagation according to laws of the anomalous diffusion model in fractal interstellar medium (Lagutin et al., 2001). The value $\\left<\\ln{A}\\right>$ at $E_{0} > 10^{18}$~eV decreases gradually and at $E_{0} \\sim 10^{19}$~eV the mass composition consists of He nuclei and protons. The cosmic ray intensity beyond $E_{\\text{thr.}} > 6 \\times 10^{19}$~eV decreases sharply and this effect is not described in the framework of the galactic model only. Such a character of spectrum does not contradict to the calculations by Berezinsky et al~\\cite{23} for the metagalactic model, in which the ``ankle'', observed in the experiments on ultrahigh energy cosmic ray registration, can be produced by the proton component only arriving from the Metagalaxy. Thereby, the details of experimental spectrum form, for example, ``dip'', i.e. the decrease of intensity at $E_{0} \\times 10^{19}$~eV, are caused by, most likely, the interaction of extragalactic protons with a relic radiation photons ($p + \\gamma_{\\text{\\footnotesize{}CMB}} \\to \\text{p} + e^{+} + e^{-}$). As a direct argument of this hypothesis, the anisotropy can be used which is related to the origin and sources of cosmic rays. Based on data of~\\cite{26, 27, 28}, at $E_{0} \\ge 8 \\times 10^{18}$~eV the weak correlation in the arrival directions of EAS with the Galaxy plane and the close correlation with the Supergalaxy plane are observed and that the quasars can be the possible sources of cosmic rays."
    },
    "0711/0711.2254_arXiv.txt": {
        "abstract": "{The ANTARES neutrino telescope is currently being constructed in the Mediterranean Sea. The complete detector will consist of 12~strings, supplemented by an additional instrumentation line. Nine strings are at present deployed of which five are already connected to the shore and operating. Each string is equipped with 75~Optical Modules (OMs) housing the photomultipliers to detect the Cherenkov light induced  by the charged particles produced in neutrino reactions.  An accurate measurement of the Cherenkov photon arrival times as well as the positions and orientations of the OMs is required for a precise reconstruction of the direction of the detected neutrinos. For this purpose the ANTARES detector is provided with several systems to facilitate the calibration of the detector. The time calibration is performed using light pulses emitted from LED and laser devices. The positioning is done via acoustic triangulation using hydrophones. Additionally, local tilt angles and the orientations of the modules are measured with a set of tiltmeters and compasses.  In this paper, it is demonstrated that the ANTARES time and alignment calibration systems operate successfully in situ. In particular, it is shown that the ANTARES read-out electronics is capable of reaching a sub-nanosecond time resolution.}  \\begin{document} ",
        "introduction": "The ANTARES neutrino telescope \\cite{ref0:status} is being constructed at a depth of 2.5~km in the Mediterranean Sea 40~km off shore from Toulon (France).  The complete detector will consist of 12~strings (also referred to as lines) of 450~m length with about 70~m spacing, anchored on the sea bed and maintained vertically by buoys. Each string is equipped with 75~Optical Modules \\cite{ref1:ompaper}, pressure resistant glass spheres housing a $10\"$ photomultiplier tube (PMT), which detects the Cherenkov light emitted by charged particles produced in neutrino reactions with the surrounding matter. The OMs of a string are arranged into triplets on 25~storeys of 14.5~m vertical distance to each other, with the PMTs looking downwards at an angle of $45^\\circ$ with respect to the horizontal.  In 2006, the first ANTARES lines have been deployed and put into operation \\cite{ref2:firstmuons}. Presently, nine strings have been deployed of which five are already connected to shore and taking data continuously. The completion of the ANTARES telescope is foreseen for the beginning of 2008.  The high angular resolution for neutrino astronomy ANTARES aims at ($0.3^\\circ$ for muon events above 10~TeV) is intimately connected with a precise resolution of the relative arrival times of Cherenkov photons at the PMTs. The relative time resolution between OMs is therefore of importance. The time resolution is intrinsically limited by the transit time spread of the signal in the PMTs ($\\sigma \\approx 1.3\\,\\rm{ns}$) and the optical properties of the seawater \\cite{ref3:opticalprops} such as light scattering and chromatic dispersion ($\\sigma \\approx 1.5\\,\\rm{ns}$ at a distance of 40 meters).  The ANTARES read-out electronics and calibrations systems are designed to contribute less than 0.5~ns to the overall timing resolution. In addition, a precise measurement of the relative positions of the OMs with an accuracy of 10-20~cm is required. ",
        "conclusions": "In situ calibration for all five operating strings have been performed regularly. The time calibration as well as the positioning systems work well within the specifications.  It has been demonstrated that the ANTARES read-out electronics is capable of reaching a sub-nanosecond time resolution in situ.  \\textit{This work is supported by the German BMBF Grant No. 05 CN5WE1/7.}"
    },
    "0711/0711.2312_arXiv.txt": {
        "abstract": "The high-redshift Universe contains luminous Ly$\\alpha$ emitting sources such as galaxies and quasars. The emitted Ly$\\alpha$ radiation is often scattered by surrounding neutral hydrogen atoms. We show that the scattered Ly$\\alpha$ radiation obtains a high level of polarization for a wide range of likely environments of high-redshift galaxies. For example, the back-scattered Ly$\\alpha$ flux observed from galaxies surrounded by a superwind-driven outflow may reach a fractional polarization as high as $\\sim 40\\%$. Equal levels of polarization may be observed from neutral collapsing protogalaxies. Resonant scattering in the diffuse intergalactic medium typically results in a lower polarization amplitude ($\\lsim 7\\%$), which depends on the flux of the ionizing background. Spectral polarimetry can differentiate between Ly$\\alpha$ scattering off infalling gas and outflowing gas; for an outflow the polarization should increase towards longer wavelengths while for infall the opposite is true. Our numerical results suggest that Ly$\\alpha$ polarimetry is feasible with existing instruments, and may provide a new diagnostic of the distribution and kinematics of neutral hydrogen around high-redshift galaxies. Moreover, polarimetry may help suppress infrared lines originating in the Earth's atmosphere, and thus improve the sensitivity of ground-based observations to high-redshift Ly$\\alpha$ emitting galaxies outside the currently available redshift windows. ",
        "introduction": "\\label{sec:intro} The Lyman-$\\alpha$ (Ly$\\alpha$) line serves as an excellent tracer of high-redshift galaxies and quasars \\citep[e.g.][]{Partridge,Hu96,Steidel96,Trager97,Rhoads00,Hu02,Ko03,Ka06,Tapken06,Westra06,Stanway07,Ouchi07,Rauch07}. Searches for redshifted Ly$\\alpha$ emission have discovered galaxies robustly out to redshifts $z=6.96$ \\citep{Iye06}, and potentially out to $z=10$ \\citep{Stark07}. These observations provide a unique glimpse into galaxy formation at high-redshift \\citep[][]{MR02,DW07}, and into the Epoch of Reionization \\citep[EoR, e.g.][]{Loeb99,Haiman99,MR04,HC05,FL06,LF,Stark06,McQuinn07,Ota07,Ko07,F07,Mesinger07}. Future surveys intend to exploit this window further by pushing to even higher redshifts and fainter flux levels \\citep[e.g.][and references therein]{Stark06,N07}.  The interpretation of existing and future observations is complicated by the fact that Ly$\\alpha$ photons are typically scattered both in the interstellar medium (ISM, e.g. Hansen \\& Oh 2006) as well as in the intergalactic medium (IGM, Loeb \\& Rybicki 1999; Santos 2004; Dijkstra et al 2007b). The impact of scattering on the observed Ly$\\alpha$ flux and spectrum may be derived by careful modeling of the observed Ly$\\alpha$ line profile, combined with constraints on the galaxy's stellar population derived from its broad band colors \\citep[][]{Verhamme06,Verhamme07}. Also, observations of the Ly$\\alpha$ line in local starburst galaxies can shed light on the processes that regulate the transfer and escape of Ly$\\alpha$ from star forming galaxies \\citep[][]{Kunth98,Hayes07}. Of course, it is not obvious whether local starburst galaxies are representative of high-redshift Ly$\\alpha$ emitters (LAEs), and detailed modeling of the line profile may be hampered by the quality of the data of high-redshift LAEs.  In this paper, we describe a complementary approach to studying the environments of high-redshift LAEs, which involves the polarization properties of the scattered Ly$\\alpha$ radiation. In the context of the solar system, the polarization of resonantly scattered Ly$\\alpha$ photons from the sun by neutral hydrogen in the interplanetary medium and in the Earth's atmosphere has been studied extensively \\citep[e.g.][]{Brandt59}, occasionally by using a Monte-Carlo method to compute the Ly$\\alpha$ radiative transfer \\citep[][]{Modali72,Keller81}. In the cosmological context, \\citet{Lee98} showed that Ly$\\alpha$ emerging from an unresolved star-bursting galaxy may be polarized to a considerable level ($\\sim 5\\%)$, and \\citet{RL99} showed that scattering of Ly$\\alpha$ photons from a point source embedded in the Hubble flow of a neutral IGM would produce a polarized Ly$\\alpha$ halo around the source. Despite the emergence of several Monte-Carlo Ly$\\alpha$ radiative transfer codes in recent years \\citep[e.g.][]{Zheng02,Cantalupo05,Iro06,D06,Verhamme06,Laursen07, Semelin07}, the polarization properties of scattered Ly$\\alpha$ photons around galaxies have not been investigated in a broader context beyond this early work.  The goal of this paper is to demonstrate that Ly$\\alpha$ radiation may be highly polarized around high-redshift galaxies in a broad range of cosmological circumstances. We will show that Ly$\\alpha$ polarimetry may place constraints on the kinematics of the gas through which the photons propagate. This is important: understanding the impact of scattering in both the ISM and IGM on the observed Ly$\\alpha$ properties is required to fully exploit Ly$\\alpha$ observations as a probe of the high-redshift Universe. Furthermore, we discuss the possibility of separating Ly$\\alpha$ sources from low-redshift line emitters (such as [OII], [OIII], H$\\alpha$, H$\\beta$ emitters, and OH-skylines) based on polarimetry.  The outline of this paper is as follows: In \\S~\\ref{sec:lya} we describe the basic principles of Ly$\\alpha$ radiative transfer and polarization, and how polarization is incorporated in our radiative transfer calculations. In \\S~\\ref{sec:result} we present our main numerical results. In \\S~\\ref{sec:discuss} we discuss our results and their implications, before presenting our final conclusions in \\S~\\ref{sec:conc}. The parameters for the background cosmology used throughout this paper are $(\\Omega_m,\\Omega_{\\Lambda},h)=(0.24,0.76,0.73)$ \\citep{wmap}. ",
        "conclusions": "\\label{sec:conc}  The high-redshift Universe is known to contain luminous Ly$\\alpha$ emitting sources such as galaxies and quasars. The Ly$\\alpha$ photons that are emitted by these sources are typically scattered both in the interstellar medium and in the intergalactic medium. In this paper we have calculated polarization properties of this scattered Ly$\\alpha$ radiation.  We used a Monte-Carlo Ly$\\alpha$ radiative transfer code, and endowed each Ly$\\alpha$ photon with a polarization vector which is perpendicular to the photon's wave vector. In this formulation the Stokes parameters resulted from binning together multiple independent photons. We differentiated between resonant and wing scattering, as these are described by different scattering matrices (see Appendix~\\ref{app:super}). Wing scattering is described by classical Rayleigh scattering and in this case we advanced the photon's polarization vector in a scattering event following the well-tested approach of Rybicki \\& Loeb (1999). However, resonant scattering is described by a superposition of Rayleigh scattering and isotropic scattering with corresponding weights of $1/3$ and $2/3$ (\\S~\\ref{sec:RT2}). In this case, minor modifications to the approach of \\citet{RL99} are required when advancing the photon's polarization vector (\\S~\\ref{sec:RT3}).  We have applied our code to three classes of models which represent the diverse sets of environments around high-redshift galaxies:  \\begin{itemize}  \\item First, we computed the polarization of the back-scattered Ly$\\alpha$ radiation observed from galaxies surrounded by a superwind-driven outflow. We have found that the fractional polarization may reach values as high as $\\mathcal{P}_{\\rm max}\\sim 40\\%$, where the maximum polarization depends on parameters such as the speed of the outflow, $v_{\\rm exp}$, and the column density of neutral hydrogen atoms, $N_{\\rm HI}$ (\\S~\\ref{sec:wind}). In this case we have found the polarization to increase towards longer wavelengths.  \\item Second, we considered resonant scattering in the intergalactic medium (IGM) {\\it after} reionization. Residual intergalactic hydrogen can scatter up to $\\gsim 90\\%$ of the Ly$\\alpha$ emitted by galaxies out of the line of sight.  (This high fraction does not apply to galaxies with superwind-driven outflows where the Ly$\\alpha$ photons are scattered out of resonance before they encounter the IGM.) We have found the polarization of the scattered Ly$\\alpha$ halo around galaxies to be polarized at lower levels ($\\mathcal{P}_{\\rm max}\\lsim 7\\%$) than the halos around galaxies with outflows. This follows from the fact that Ly$\\alpha$ photons scatter multiple times when they enter resonance, which isotropizes the local Ly$\\alpha$ radiation field and reduces its net polarization. The polarization decreases with the mean number of scattering events the photons experience (i.e. the polarization decreases with an increasing effective optical depth in the Ly$\\alpha$ line). For this reason, we expect the polarization of resonantly scattered Ly$\\alpha$ photons in the IGM to be largest around low redshift sources or around sources where the IGM is more highly ionized than average (e.g. in the vicinity of a bright quasar, see \\S~\\ref{sec:igm}).  The overall lower levels of polarization of resonantly scattered Ly$\\alpha$ radiation in a reionized IGM are significantly lower than that of pre-reionization Ly$\\alpha$ halos around galaxies embedded in a fully neutral IGM \\citep{Loeb99}. This suggests that the polarization properties of Ly$\\alpha$ halos may help constrain the neutral fraction of the IGM during the epoch of reionization.  \\item Third, we considered neutral collapsing protogalaxies that are emitting Ly$\\alpha$ cooling radiation. Despite the fact that these clouds are extremely optically thick to Ly$\\alpha$ photons, and photons typically scatter $\\sim \\tau$ times, we found the polarization to linearly increase toward the edge of the cloud, reaching a maximum amplitude of $\\mathcal{P}_{\\rm max}\\sim 35\\%$. (\\S~\\ref{sec:cool}). The resulting polarization increases towards shorter wavelengths in contrast to the trend found in outflow models. \\end{itemize} Our results indicate that resolved high-redshift Ly$\\alpha$ emission may be highly polarized under a variety of likely circumstances. High-redshift Ly$\\alpha$ emitters are usually distinguished from low-redshift line emitters on the basis of their broad-band colors and their asymmetric spectral line shape. The work presented here implies that polarization may provide an additional diagnostic. Moreover, polarimetry has the potential to better remove the glow of infrared lines in the Earth's atmosphere, which would improve the sensitivity of ground-based observations to high-redshift Ly$\\alpha$ emitting galaxies outside the currently available redshift windows.  The polarization properties of Ly$\\alpha$ radiation encode information about the distribution and kinematics of neutral gas in and around galaxies. Polarimetry therefore complements the constraints that are derivable from spectroscopy. Hence, from a theorist's perspective it is well worth to include polarization in Monte-Carlo calculations of Ly$\\alpha$ radiative transfer.  The redshifted Ly$\\alpha$ line has provided us with an important window into the high-redshift Universe. This work suggests that in order to fully exploit the observations, one should focus on both Stokes parameters $I$ and $Q$, rather than just $I$. Polarization measurements to better than $1\\%$ accuracy can be carried out by existing facilities such as the FOcal Reducer/low dispersion Spectrograph (FORS1)\\footnote{http://www.eso.org/instruments/fors/} on the VLT, the LRIS imaging spectropolarimeter at the W.M. Keck Observatory \\citep{Goodrich03}, and the CIAO polarimeter on Subaru \\citep[][]{Tamura}. The results of this paper imply that Ly$\\alpha$ emitting sources would provide excellent future targets for these instruments.  {\\bf Acknowledgments} We thank George Rybicki and Adam Lidz for useful discussions. We thank an anonymous referee for a prompt helpful report that improved the content of this paper. This research was supported by Harvard University funds. \\newpage"
    },
    "0711/0711.3913_arXiv.txt": {
        "abstract": "We present the first clustering results of X-ray selected AGN at $z\\sim3$. Using Chandra X-ray imaging and UVR optical colors from MUSYC photometry in the ECDF-S field, we selected a sample of 58 $z\\sim3$ AGN candidates. From the optical data we also selected 1385 LBG at 2.8 $< z <$ 3.8 with R $<$ 25.5.  We performed auto-correlation and cross-correlation analyses, and here we present results for the clustering amplitudes and dark matter halo masses of each sample.  For the LBG we find a correlation length of $r_{0,LBG} = 6.7 \\pm 0.5$ Mpc, implying a bias value of 3.5$\\pm$0.3 and dark matter (DM) halo masses of log($M_{min}$/$M_{\\odot}$) = $11.8\\pm0.1$. The AGN-LBG cross-correlation yields $r_{0,AGN-LBG} = 8.7\\pm1.9$ Mpc, implying for AGN at 2.8 $<z<$ 3.8 a bias value of 5.5$\\pm$2.0 and DM halo masses of log($M_{min}$/$M_{\\odot}$) = $12.6_{-0.8}^{+0.5}$.  Evolution of dark matter halos in the Lambda CDM cosmology implies that today these $z\\sim3$ AGN are found in high mass galaxies with a typical luminosity $7_{-2}^{+4} L^{*}$. ",
        "introduction": "\\label{sec:intro} There is a wealth of evidence that nuclear supermassive black holes play a significant role in the process of galaxy formation and evolution.  This has become evident in the past few years with the discovery of correlations between the properties of the massive black holes, the stellar systems that host them, and their dark matter halos (\\citealp{magorrianetal98},\\citealp{gebhardtetal00}, \\citealp{ferrarese02}).  Luminous quasars have been studied in great detail, with systematic spectroscopic studies by the Sloan Digital Sky Survey and the 2dF Galaxy Redshift Survey.  These objects are rare and represent the bright end of the active galactic nuclei (AGN) luminosity function. In order to fully understand the link between the growth of supermassive black holes (SMBH) and galaxy evolution, we need to study young, high-redshift galaxies hosting AGN with more typical luminosities.  One of the basic properties of galaxy populations is their clustering strength, but there are few constraints on this quantity for fainter, high-redshift AGN. Furthermore, there are important disagreements in the literature: e.g. galaxy-AGN cross-correlation measurements at $z\\sim3$ by \\citet{adelberger05cons} \\citepalias{adelberger05cons}, imply a bias factor of $3.9\\pm3.0$ for luminous AGN (UV luminosities between -30$<M_{1350}<$-25), while recent work by \\citet{shenetal2007} \\citepalias{shenetal2007} suggests a value of $9.1\\pm0.9$ for sources with similar luminosities at the same redshift.  For galaxies, e.g., measurements indicate a strong luminosity dependence in the clustering length \\citep{giavalisco2001}. On the other hand, for AGN there are a handful of claims that this is not the case (\\citealp{croometal2005}, \\citealp{myersetal2006}, \\citealp{adelbergers05}). Determining the bias with accuracy puts important constraints on models of AGN formation and evolution, (e.g., \\citealp{lidzetal06}, \\citealp{hopkinsetal07}) but requires spanning a broad range in luminosity and obscuration level. This is challenging using optical color selection plus spectroscopy, because the low fraction of AGN (3\\%) found among these candidates \\citep{steideletal02} demands huge amounts of spectroscopy time.  Using X-ray detection plus spectroscopy (e.g., \\citealp{szokolyetal2004}) provides a more efficient and unbiased selection of AGN, since their surface density is higher in X-rays than in optical images and obscuration effects are much less important. However, the unrestricted redshift range sampled by this method makes spectroscopic follow-up highly inefficient, and often prevents confirmation of dimmer, high-z AGN (R$>$24). Hence it is difficult to obtain AGN samples suitable for clustering studies. There have been several measurements of the spatial correlation function of X-ray selected AGN at $z<1$ (\\citealp{mullisetal2004}, \\citealp{gillietal04}, \\citealp{basilakosetal2004}, \\citealp{basilakosetal2005}, \\citealp{yangetal2006}, \\citealp{miyajietal2007}), with samples ranging from 200-500 sources. At higher redshifts, the statistics are much poorer, and consist of purely optically selected AGN. In this work, we constrain the clustering strength of an AGN sample at $z\\sim3$ jointly selected by optical and X-ray photometry.  We determine whether these sources cluster more or less than non-active galaxies at this redshift, and discuss their present-day descendants.  We assume a $\\Lambda$CDM cosmology consistent with WMAP results \\citep{spergeletal2007} with $\\Omega_m=0.3, \\Omega_\\Lambda=0.7$, $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$ and $\\sigma_{8}=0.8$. All quantities are comoving: correlation lengths scale as $h_{70}^{-1}$, number densities as $h_{70}^{3}$ and halo masses as $h_{70}^{-1}$. ",
        "conclusions": "\\label{sec:sumdisc}  We have measured the autocorrelation strength of z=3 LBG and the cross-correlation between these galaxies and AGN selected using both X-ray and optical data. For the LBG sample, we obtained a bias factor of 3.5$\\pm$0.3. This is somewhat higher than the value 2.8$\\pm$0.3 found by \\citet{adelbergeretal05a} at $z=3$. \\citet{hildebrandtetal07} obtained a bias value of 3.2$\\pm$0.2 for an equivalent LBG population, consistent with ours.  From the AGN-LBG cross-correlation and the bias calculated for the LBG, we have deduced a bias factor of 5.5$\\pm$2 for our AGN sample. The active galaxies targeted in this study appear to cluster more than star-forming galaxies with similar restframe-UV colors. This is consistent with cosmic downsizing of AGN, implying that typical SMBH tend to sit in more massive galaxies than the ``normal'' galaxy population. We need greater statistics to confirm this result, since the values are consistent within the uncertainties.  \\citetalias{adelberger05cons} performed this same calculation in an optically selected and spectroscopically confirmed sample of 79 AGN between 1.6 $< z <$ 3.7, dividing the AGN by UV luminosity into bright (25 sources with -30$<M_{1350}<$-25) and dim (54 sources with -25$<M_{1350}<$-19) samples. Although \\citetalias{adelberger05cons} does not report the autocorrelation length of the galaxy sample used to compute these values, we approximate it by averaging the result for LBG and BX galaxies presented in \\citet{adelbergeretal05a}, obtaining a bias factor of 2.6$\\pm$0.3. From this we infer bias factors of 3.9$\\pm$3.0 and 4.7$\\pm$1.7 for the \\citetalias{adelberger05cons} bright and dim AGN samples. Our AGN sample has an UV magnitude range between -26 and -20, almost identical to the \\citetalias{adelberger05cons} faint AGN set, and showing the same clustering strength. In Fig.~4 these results are compared to the present estimate and to that obtained by \\citetalias{shenetal2007} for $\\sim$2250 SDSS quasars at 2.9$<z<$3.5.  To estimate the dark matter halo mass of the typical descendant of the halos that host these $z \\sim 3$ galaxy sets at the present time, we calculated the mode and width of the conditional probability distribution of the expected mass $z=0$ (see \\citealp{hamanaetal06} and references therein). We took the median mass of the halo populations at $z=3.2$ as the representative value, and we report the bias values corresponding to the maximum likelihood z=0 progenitor, with $1-\\sigma$ uncertainties corresponding to halo masses whose likelihood is reduced by a factor $exp(-1/2)=0.607$. For the LBG bias of 3.5 at $z = 3.2$, we obtain a bias of $1.3_{-0.1}^{+0.3}$ at $z=0$. In the nearby universe \\citep{zehavietal2005}, this corresponds to the clustering of a somewhat massive galaxy, with $L=2.7_{-0.6}^{+1.9} L^{*}$. On the other hand, our AGN sample, with a bias of 5.5 at $z=3.2$ would have a typical halo with a bias factor of $2.0_{-0.3}^{+0.6}$ at present. Extrapolating from \\citet{zehavietal2005}, this corresponds to the clustering of galaxies having $7_{-2}^{+4} L^{*}$, which at the present time are most likely located in groups and galaxy clusters.  \\begin{figure} \\plotone{f4.eps} \\caption{ Comparison of AGN bias factors at $z\\sim3$ between this work (square), \\citetalias{shenetal2007} (triangle) and \\citetalias{adelberger05cons} (stars). \\label{fig:biascomp} } \\end{figure}  This confirms the result found by \\citetalias{adelberger05cons} for faint, optically selected AGN. Improving our understanding of the AGN-galaxy connection at $z\\sim3$ and constraining models such as \\citet{lidzetal06} and \\citet{hopkinsetal07} for AGN clustering, requires resolving the discrepancy in the bias estimates for bright AGN seen in Fig.~4 improving the statistics. We have show that targeting high-redshift AGN for clustering analyses can be done very efficiently by means of deep optical and X-ray imaging, and for that reason, future surveys with deep X-ray coverage will be ideal for obtaining large samples of active galaxies in restricted redshift ranges, suitable for clustering studies. Our method can be applied to these surveys to obtain enough X-ray-selected AGN to reduce the current uncertainties.  H.~F. was supported by MECESUP project UCH0118, Andes Fundation fellowship C-13798, ALMA fellowship 31060003 and Fondecyt project 1040719."
    },
    "0711/0711.4158_arXiv.txt": {
        "abstract": "We address the ability of broad iron emission lines from black hole accretion disks to diagnose the spin of the black hole.  Using a high-resolution 3-dimensional MHD simulation of a geometrically-thin accretion disk in a Pseudo-Newtonian potential, we show that both the midplane density and the vertical column density of the accretion flow drop dramatically over a narrow range of radii close to the innermost stable circular orbit (ISCO).  We argue that this drop of density is accompanied by a sharp increase in the ionization parameter of the X-ray photosphere, and that the resulting imprint of the ISCO on the X-ray reflection spectrum can be used to constrain spin.  Motivated by this simulation, we construct a simplified toy-model of the accretion flow within the ISCO of a Kerr black hole, and use this model to estimate the systematic error on inferred black hole spin that may result from slight bleeding of the iron line emission to the region inside of the ISCO.  We find that these systematic errors can be significant for slowly spinning black holes but become appreciably smaller as one considers more rapidly rotating black holes. ",
        "introduction": "Recent years have brought an increasing realization of the astrophysical importance of black hole spin.  Ever since the seminal work of Penrose (1969) and Blandford \\& Znajek (1977), it has been recognized that black hole spin may be an important source of energy, especially for the powerful relativistic jets seen emerging from many black hole systems.  However, the importance of black hole spin goes beyond its role as a possible power source.  The spins of stellar-mass black holes found in accreting Galactic Black Hole Binaries (GBHBs) are expected to be natal (King \\& Kolb 1999) and give us a window into the black hole forming core-collapse supernova.  The spin distribution of the supermassive black hole (SMBH) population (and its dependence on SMBH mass) encodes the black hole growth history, e.g., the role of accretion versus mergers (Moderski \\& Sikora 1996; Volonteri et al. 2005).  Thus, across the mass scales, black hole spin gives us a fossil record of how black holes form and grow.  Of course, we must have robust observational probes of black hole spin if we are to truly test spin-powered jet models or actually use spin to probe the formation and growth of black holes.  The most direct method imaginable is the characterization of the gravitational radiation emitted as another compact object spirals into the black hole.  However, in the current (pre- gravitational wave observatory) era, we must search for signs of black hole spin in the electromagnetic radiation emitted by the surrounding accretion disk. The natural waveband to use is the X-ray band given the fact that the observed X-rays are thought to be predominately emitted from the inner regions of the accretion flow where the relativistic effects related to black hole spin are strongest.  The most important effect of black hole spin on the accretion disk arises from the spin dependence of the location of the innermost stable circular orbit (ISCO); in Boyer-Lindquist coordinates, $r_{\\rm isco}$ shrinks from $6GM/c^2$ for a non-rotating black hole down to $GM/c^2$ for a hole that is maximally rotating in the same sense as the accretion disk.  Thus, for a given black hole mass, the characteristic temporal frequency, energy release (and hence temperature), and gravitational redshift of the inner accretion flow all increase as one considers progressively higher black hole spin.  These considerations lead to the three most widely discussed techniques for using X-ray data to determine black hole spin.  On the timing front, the frequency stability of the high-frequency quasi-periodic oscillations (HFQPOs) seen in the soft intermediate state of GBHBs strongly suggest a connection to the gravitational potential and black hole spin (Strohmayer 2001; also see McClintock \\& Remillard [2003] for a general review of HFQPO phenomenology).  While HFQPOs may eventually provide a robust way of measuring spins in GBHBs, the lack of a compelling theoretical framework with which to interpret the frequencies prevents any robust conclusions from being drawn at the present time.  For example, different and contradictory constraints on the mass and spin of the black hole result from assuming that the HFQPOs are manifestations of trapped global g-modes (Nowak \\& Wagoner 1992; Nowak et al. 1997), parametric resonances (Abramowicz et al. 2003), or Lens-Thirring precession (Merloni et al. 1999).  The second technique involves modeling the thermal continuum spectrum from the disk.  Provided that one selects systems that have optically-thick, geometrically-thin inner accretion disks, observations can be compared with model thermal continuum spectra (which include vertical radiation transfer in the disk atmosphere as well as Doppler and gravitational red/blue-shifts; Davis, Done \\& Blaes 2006).  Spin results from this technique have been reported by McClintock et al. (2006) and Shafee et al. (2006), although the contrary results of Middleton et al. (2006) demonstrate the current fragility of this technique to modeling the non-thermal emission, particularly when applying it to data from the Proportional Counter Array (PCA) on the {\\it Rossi X-ray Timing Explorer (RXTE)} in which one only sees the Wien tail of the thermal disk spectrum.  While this is a powerful method for determining the spin of accreting stellar-mass black holes (especially when applied to broad-band X-ray data that extends into the soft X-ray band), the fact that one needs both an estimate of the black hole mass and a high quality measurement of the thermal continuum shape makes it difficult to apply to the SMBH in active galactic nuclei (AGN).  The thermal continuum of AGN disks is in the UV/EUV region of the spectrum, making a determination of its shape and normalization extremely susceptible to any errors in the (large) correction that must be done for dust extinction and photoelectric absorption.  The third technique for constraining spin uses relativistic broadening and gravitational redshifting of the fluorescent iron emission line seen in many GBHBs and AGN (Fabian et al, 1989; Laor 1991; Fabian et al. 2000; Reynolds \\& Nowak 2003; Fabian \\& Miniutti 2005).  As one considers progressively more rapidly rotating black holes, the primary X-ray emission and hence the iron line emission will be dominated by regions with higher gravitational redshift, leading to broader and more highly skewed iron emission lines.  A major advantage of this technique is that the iron line profiles are completely independent of black hole mass and hence one can apply this to an AGN in which the mass is often extremely uncertain.  Although one actually proceeds via formal fitting of relativistically smeared theoretical disk reflection spectra to data, the black hole spin constraint essentially results from a characterization of the maximum redshift of the iron line. Results using this technique on the Seyfert galaxy MCG--6-30-15 have been reported by Brenneman \\& Reynolds (2006) who report rather extreme constraints on the dimensionless black hole spin of $a>0.987$ (also see Dabrowski et al. 1997).  Both the thermal continuum and the iron line techniques for determining black hole spin rely on the assumption that the ISCO defines an inner edge for the thermally or line emitting region of the disk (see Krolik \\& Hawley 2002 for a general discussion of the various ``inner edges'' that one can define for an accretion disk). The robustness of the thermal continuum technique has been explicitly addressed by Shafee, Narayan \\& McClintock (2007) using a global $\\alpha$-viscosity model that explicitly tracks the flow properties within the ISCO.  Although MHD effects close to the ISCO may invalidate the assumptions of such $\\alpha$-models (e.g., Reynolds \\& Armitage 2001), the results of Shafee et al. (2007) do indeed suggest that modeling of the thermal continuum can produce an accurate measure of black hole spin.  The subject of the current {\\it Paper} is to address the robustness of the iron line fitting technique.  In Section 2, we present a high resolution 3-dimensional magnetohydrodynamic (MHD) simulation of a geometrically-thin ($h/r\\sim 0.05$) accretion disk that we use to explore the nature of the flow close to the ISCO.  In Section 3, we show that both the midplane density and vertically-integrated surface density of the accretion flow drops precipitously over a small range of radii close to the ISCO.  We then argue in Section 4 that photoionization of the flow within the ISCO will effectively suppress the iron line emission (and all other X-ray reflection signatures) across most of the region within the ISCO.  We discuss the implications and limitations of our results in Section 5, including an estimate of the uncertainties in the inferred black hole spin that might arise from iron line emission in the region immediately within the ISCO.  Section 6 presents our conclusions. ",
        "conclusions": "It is a fundamental fact that one can, in principle, observe arbitrarily large gravitational redshifts from around any black hole, irrespective of its spin parameter.  Indeed this is very closely tied to the definition of a black hole.  At first glance, this fact would appear to render impotent attempts to determine spin using the gravitational redshift of iron emission lines.  However, it is the {\\it astrophysics} of black holes accretion and the change in the character of the flow around the ISCO that gives observations of broad iron lines (and X-ray reflection in general) the power to diagnose black hole spin.  To understand the behavior of the accretion flow close to the ISCO, one must capture the angular momentum transport as fully as possible. Models based on parameterized anomalous viscosity (e.g., $\\alpha$-models) can become acausal within the plunge region and cannot capture the non-local MHD stresses that may operate around the ISCO.  Fully 3-d MHD simulations provide the most powerful tool to examine these complex flows.  All previously published global 3-d MHD simulations of black hole accretion disks have revealed rather gradual transitions of flow properties around the ISCO, but have modeled rather thick ($h/r\\sim 0.1$ or greater) accretion flows.  Motivated by the fact that luminous Seyfert nuclei and high-state GBHBs (i.e., the systems in which we actually observe broad iron lines) are thought to possess geometrically-thin disks, we have examined the nature of the flow around the ISCO using a high-resolution pseudo-Newtonian MHD simulation of a disk with $h/r\\sim 0.05$.  We find a rather abrupt transition of flow properties across the ISCO; the radial length-scale of these changes is comparable to the vertical scaleheight of the disk beyond the ISCO. In the language of Krolik \\& Hawley (2002), the ``stress edge'' and the ``turbulence edge'' of the disk are at $r\\approx 5.5r_g$ compared with $r_{\\rm isco}=6r_g$.  Most importantly for the present discussion, we find that both the midplane density and column density of the disk drop dramatically as the flow accelerates inwards. Although there is no radiation in the simulation, we relate the simulation quantities to the expected ionization state of the accretion flow and find that (even under conservative assumptions aimed at minimizing the ionization state of the plunging flow) significant iron line emission cannot be produced from inside $r=5r_g$.  In fact, for canonical parameters, this ``reflection edge'' is at $r\\approx 5.8r_g$, very close to the ISCO.  The fact that the iron line emission does not truncate precisely at the ISCO will, however, introduce a systematic error into iron line based spin measurements that assume strictly ISCO-truncated line profiles.  We make a first attempt to quantify this systematic error in Section~5, although we note that this approach neglects the spin information which is carried by the full line profile (as opposed to just the maximum redshift).  Further study is required to assess whether, in practice, the information carried by the full line profile dramatically decreases the systematic error estimates of Fig.~\\ref{fig:spinerror}."
    },
    "0711/0711.1171_arXiv.txt": {
        "abstract": "We use the photometric information contained in individual pixels of 44,964 ($0.019<z<0.125$ and $-23.5<M_r<-20.5$) galaxies in the Fourth Data Release (DR4) of the Sloan Digital Sky Survey to investigate the effects of environment on galaxy star formation (SF). We use the pixel-z technique, which combines stellar population synthesis models with photometric redshift template fitting on the scale of individual pixels in galaxy images. Spectral energy distributions are constructed, sampling a wide range of properties such as age, star formation rate (SFR), dust obscuration and metallicity. By summing the SFRs in the pixels, we demonstrate that the distribution of total galaxy SFR shifts to lower values as the local density of surrounding galaxies increases, as found in other studies. The effect is most prominent in the galaxies with the highest star formation, and we see the break in the SFR-density relation  at a local galaxy density of $\\approx 0.05\\,$(Mpc/h)$^{-3}$. Since our method allows us to spatially resolve the SF distribution within galaxies, we can calculate the mean SFR of each galaxy as a function of radius. We find that on average the mean SFR is dominated by SF in the central regions of galaxies, and that the trend for suppression of SFR in high density environments is driven by a reduction in this nuclear SF. We also find that the mean SFR in the outskirts is largely independent of environmental effects. This trend in the mean SFR is shared by galaxies which are highly star forming, while those which are weakly star forming show no statistically significant correlation between their environment and the mean SFR at any radius. ",
        "introduction": "}  \\subsection{Environmental dependences of galaxy properties and dependence on morphology \\label{subsec:intro1}}  According to models of hierarchical formation \\citep{Kau:93,Som:99,Col:00}, galaxies form in less dense environments and are then accreted into larger halos (e.g., falling into clusters or groups), having their hot gas reservoir removed as this occurs. This predicts that galaxy properties, such as color, luminosity and SFR, are correlated with their environment. Galaxies in highly dense regions are predicted to be more luminous and more red than those in lesser dense regions. In fact, the mean galaxy environment as a function of color and luminosity has been explored using the Sloan Digital Sky Survey (SDSS) \\citep{Hog:03,Bla:03}, as was the dependence of the color-magnitude relation of bulge-dominated galaxies on their environment \\citep{Hog:04}. The models also predict a gradual decrease in star formation (SF) activity in more dense regions, as seen in recent observational studies exploring the variation of star formation rate (SFR) with environment \\citep{Lew:02,Gom:03}. Numerous physical mechanisms for how such SFR suppression might occur have been suggested, including ram pressure stripping of gas \\citep{GG:72}, gravitational interactions between galaxies \\citep{BV:90} and between galaxies and a non-uniform cluster potential \\citep[``galaxy harassment,\" ][]{Moo:99}.  Recent N-body simulations and semi-analytical models have extended our insight into a number of the important issues with the hierarchical formation scenario. \\citet{Bal:00}, investigating the origin of cluster-centric gradients in SFRs and colors of rich cluster galaxies, used a model where clusters are built through the ongoing accretion of field galaxies. These models assume that after galaxies enter the cluster their SFRs decline on a timescale of a few Gyr which is the typical gas consumption timescale for disk galaxies in the field. They combined these timescales with mass accretion histories from N-body simulations of cluster formation in a $\\Lambda$CDM universe to show that there is an expected strong suppression of SF in cluster galaxies. The simulations also show that a significant fraction of galaxies beyond the virial radius of the cluster may have been within the main body of the cluster in the past. This would explain why star formation in the outskirts of clusters (and as far out as two virial radii) is systematically suppressed relative to the field. The agreement with the data beyond the cluster virial radius is further improved by assuming that gas-stripping happens within lower mass systems, before the galaxy is actually accreted into the cluster. The suggestion is that the SFRs of cluster galaxies depend primarily on the the time since their accretion onto the cluster, and that the SFR suppression happens gradually over a few Gyr.  Here we investigate the relation between environment and SFR for galaxies in the low redshift universe. \\citet{Gom:03}, using H$\\alpha$ equivalent widths (EW) as an indicator of SFR in the Early Data Release (EDR) of the SDSS, established an SFR-density relation for the SDSS, confirming a similar result seen in the 2 degree Field Galaxy Redshift Survey (2dFGRS) by \\citet{Lew:02}. They found that the overall distribution of SFRs is shifted toward lower values in more dense environments. The effect is most noticeable for the strongly star-forming galaxies (EW(H$\\alpha$)$ > 5\\,$\\AA) in the 75th percentile of the SFR distribution. They also found a characteristic ``break''  (or characteristic density) in the density-SFR relation at a local galaxy density of $\\approx 1\\,$(Mpc/h)$^{-2}$. \\citet{Gom:03} explored whether the density-morphology relation \\citep{Dre:80} alone could explain the density-SFR relation, and concluded that it could not. Using the concentration index of SDSS galaxies as a morphology indicator, they showed that SFRs for galaxies of the same type were suppressed in dense regions. Higher redshift samples ($z>0.2$) have also been used to suggest a suppression in the SFR of galaxies in the cores of distant clusters compared to those in the field \\citep{Bal:97,Has:98,Pog:99,Cou:01,Pos:01}. Together this provides strong evidence for a decrease in SFR of galaxies in dense environments, spanning a wide range of densities ($0.08-10\\,$(Mpc/h)$^{-2}$) and redshift (out to $z\\approx 0.5$). In addition, \\citet{Bal:97} and \\citet{Has:98} found that cluster galaxies have a reduced SFR compared with the field, independent of morphology.  Not all studies agree with the general SFR-density trend though. \\citet{Bal:04} analyzed H$\\alpha$ emission strength as a function of galaxy environment using galaxies selected from the SDSS and 2dFGRS. They found that the distribution of H$\\alpha$ EWs is bimodal, consisting of actively star-forming populations with EW(H$\\alpha$)$>4$\\,\\AA\\ and a quiescent population with little current SF. They showed that the distribution of EW(H$\\alpha$) for the star forming population does not itself depend on environment, and concluded that it was unlikely that SFRs are gradually decreasing in a substantial number of star-forming galaxies in or near dense regions today. They did find, however, that the fraction of galaxies with EW(H$\\alpha$)$>4\\,$\\AA\\ decreases steadily with increasing local density.  \\subsection{Galaxy Evolution and environment \\label{subsec:evol_vs_env}}  The most massive galaxies known today are giant ellipticals that reside in the most dense galaxy environments \\citep[e.g.,][]{Dre:80}, which are also known to have formed their stars rapidly at an early stage in cosmic history. Recent studies of the evolution in the global space density of galaxy SFRs find that the majority of the observed SF occurs in the highest mass galaxies at high redshift, moving to lower mass galaxies at lower redshifts \\citep{Pan:04,Jun:05,Pan:06,Sey:07}. This is referred to as ``downsizing'' \\citep{Cow:96}. Together with the trend for massive systems to occur predominantly in dense environments, this results in a scenario where the SFRs in dense environments at low redshift are lower than in less dense regions, consistent with the observations. The implication is that at sufficiently high redshifts, ($z\\gtrsim 2$) where the SFR is dominated by the most massive galaxies, the SFR-density relation should invert, with the most dense environments hosting elevated SFRs compared to the field.  This trend is beginning to be observed too. \\citet{Pog:06} show that at $0.4<z<0.8$ the suppression of SFR in dense environments is weaker than seen locally. \\citet{Ilb:06} and \\citet{Elb:07} show environmental dependent evolution in the galaxy luminosity function from $0.25<z<1.5$, suggesting an increase in the density of faint red galaxies in overdense regions as cosmic time increases. More confirmation is clearly needed, of course, but these examples are consistent with the scenario that galaxies in dense environments form stars rapidly at early times, quickly building up mass and becoming quiescent, while galaxies in less dense environments form stars at a more sedate pace but over longer timescales.  This ``environmentally governed evolution\" scenario makes a distinct prediction compared to the ``infall and quench\" models. The latter suggests that galaxies in dense environments should show an SFR distribution that is progressively suppressed from the outside in, as the outer regions are those which will be affected first by their rapidly changing environment. The former, on the contrary, suggests that the suppression should either happen uniformly as a galaxy ages, or that the inner regions should be suppressed first, since nuclear SF seems to occur more rapidly than disk SF given the elevated gas densities present. Thus by studying the spatial distribution of SFR in star-forming galaxies as a function of environment we should be able to distinguish clearly between these two scenarios.  \\subsection{Radial Variation in Galaxy SFRs \\label{subsec:radial_var}}  Previous studies on the radial dependence of SF have focused on individual galaxies. \\citet{Per:06} studied the recent SF in the early-type galaxy M81 using imaging observations from the far-ultraviolet (UV) to the far-infrared (IR). The data was then compared to models of stellar, gas and dust emission, with results from different sub-galactic regions, including individual HII regions (around 0.1 kpc). They were able to confirm the existence of a diffuse dust emission not directly linked to the SF. Using the H$\\alpha$ emission that probes the unobscured SF, and the IR luminosity (especially the $24\\,\\mu$m emission) that probes the obscured SF in the galaxy, they found a decrease in the ratio of obscured SF to total SF with radius. This fraction varies from an obscured SF of 60\\% in the inner regions of the galaxy to 30\\% in the outer regions. \\citet{Joh:05} used pixel-based spectral energy distribution (SED) fitting to a merging system hosting a compact steep spectrum radio source, in order to explore the connection between the nuclear radio emission and the distribution of star formation. \\citet{Kas:03} used pixel-based colors to explore stellar populations and obscuration in the Antennae, and \\cite{deG:03} used the same technique to explore the Mice and the Tadpole interacting galaxy systems. \\citet{Boi:06} showed, using GALEX and Spitzer data, that for disk galaxies the attenuation varies radially, being highest in the nuclear regions, and is correlated with metallicity. They also found that the Schmidt law connecting the SF and gas surface densities continues beyond the traditional ``threshold\" radius. \\citet{Lan:07} used a study of pixel Color Magnitude Diagrams for a sample of 69 nearby galaxies to study stellar populations and structure of galaxies. They found that these Color Magnitude Diagrams of each galaxy type have distinct trends. In addition, they performed a pixel-by-pixel analysis to show that there is a steady progression in average pixel color along the Hubble sequence. Finally they compared pixel colors to the Bruzual and Charlot stellar population models and used these to map the stellar mass distribution and M/L ratio in galaxies.  The current analysis has the advantage of a very large sample, accessible as a result of the SDSS, which allows a study of the environmental dependence of spatially distributed star formation. Also, since the inferred SFRs come from fitting SEDs, derived from stellar population synthesis models, to broadband photometry, there is no aperture effect such as that affecting fiber-based spectroscopy.  By fitting SEDs to individual pixels within resolved galaxy images, we can spatially resolve the star formation and explore how this distribution varies for galaxies as a function of local galaxy density.  \\subsection{Using photometric redshifts to study (spatially-resolved) properties of stellar populations in galaxies \\label{subsec:photoz}}  Photometric redshifts have become an efficient way of measuring redshifts of galaxies, particularly when the galaxies are too faint for spectroscopic studies. This typically involves fitting a library of SEDs to the observed colors of galaxies, where the SED templates are functions of galaxy type. We can, however, shift the focus from the inferred photometric redshifts themselves, to what the best fitting SEDs can reveal about the properties of the galaxy sample. This highlights instead the reliability and appropriateness of the template spectra used in the fitting. More specifically, as shown by \\citet{Abr:98}, the {\\em spatially resolved\\/} colors of galaxies can be used to understand the relative ages of bulge and disks and the formation histories of galaxies. This can be done by applying the photometric redshift technique to individual pixels of resolved galaxy images, having measured fluxes and flux errors. Given a library of SED templates constructed from population synthesis codes, for which the underlying physical parameters are defined, the spatial distribution of those parameters can then be established. \\citet{Con:03} applied such SED fitting, a technique referred to as ``pixel-z,\" to individual pixels in $\\approx 150$ galaxies with measured spectroscopic redshifts, and $\\approx 1500$ galaxies with measured photometric redshifts, in the Hubble Deep Field-North (HDF-N). The aim was to decompose the internal photometric structure of galaxies into intrinsic properties of the stellar populations like stellar ages and star formation rates. They used the pixels to calculate the comoving density of star formation and metallicity enrichment, as a function of redshift. With the sample of galaxies available, they were able to directly assess the drivers behind the current understanding of the global star formation history.  Here we apply the pixel-z technique to 44,964 SDSS galaxies, to study the dependence on environment of the total galaxy SFR and the spatial distribution of star formation within those galaxies. In \\S\\,\\ref{sec:data}, we describe our SDSS sample. We describe the pixel-z method and its application to the SDSS in \\S\\,\\ref{sec:method}, together with a description of the SEDs and the range of physical properties spanned. \\S\\,\\ref{sec:implementation} demonstrates features of the pixel-z implementation. \\S\\,\\ref{sec:results} presents our results, including the relation between total SFR and local density of galaxies, along with the spatial distribution of SFR as a function of this density. Our conclusions are described in \\S\\,\\ref{sec:conclusions}. We assume throughout that $\\Omega_{\\Lambda}=0.7$, $\\Omega_{\\rm M}=0.3$ , and $H_{0}=75\\,{\\rm km\\,s^{-1}\\,Mpc^{-1}}$. ",
        "conclusions": "}  Pixel-z is a useful technique for exploring spatially distributed galaxy properties, including the star formation rate. Although it has inherent limitations (detailed in this paper and by \\citet{Con:03}), it can provide insights to direct more detailed exploration, and when applied to large samples, many of those limitations can be overcome. We draw several conclusions about the environmental dependence of spatially resolved galaxy star formation:  \\begin{enumerate} \\item By summing the SFRs in individual pixels, a total SFR-density relation is established. The SFR-density relation measured from emission lines (\\citet{Gom:03}, \\citet{Lew:02}) is confirmed based on SFRs inferred from SED template-fitting to pixel fluxes across five bands. \\item The total SFR-density relation is found to be independent of morphological type of the galaxy. Therefore, the SFR-density relation of galaxies in the SDSS is not solely a result of having a larger fraction of early types as we go to higher densities (the morphology-density relation). So there is a SFR suppression in more dense environments. \\item Within galaxies, star formation, averaged over all morphological types and inclinations, is highest in a region spanning $0.125\\lesssim r/R_p \\lesssim 0.5$, with the specific location of the peak depending on local density. It is lower in both the nucleus and the outskirts than in the circum-nuclear region. In denser environments, the low mean SFR in the nuclear regions appears lower still. The higher mean SFR in the circum-nuclear region is also lessened in more dense regions, with the peak SFR moving towards larger radii. The SFR beyond $r/R_p\\approx 0.25$ and in the outskirts of galaxies is not affected by a changing environment. It is thus a depression of SF in the innermost annuli that accounts for the dependence of total galaxy SFR on density.  \\item When the sample is split based on the total SFR of galaxies, it is the highly star-forming systems that are found to be responsible for the dependence of radial SFR on local density. The low star-forming galaxies show little dependence of radial SFR on environment. \\end{enumerate}"
    },
    "0711/0711.1201_arXiv.txt": {
        "abstract": "We report a serendipitous detection of an intense X-ray flare from the Tycho reference source HD\\,161084 during a Suzaku observation of the Galactic Center region for $\\sim$20~ks. The X-ray Imaging Spectrometer (XIS) recorded a flare from this A1-type dwarf or subgiant star with a flux of $\\sim$1.4$\\times$10$^{-12}$~erg~s$^{-1}$~cm$^{-2}$ (0.5--10~keV) and a decay time scale of $\\sim$0.5~hr. The spectrum is hard with a prominent Fe\\emissiontype{XXV} K$\\alpha$ emission line at 6.7~keV, which is explained by a $\\sim$5~keV thin-thermal plasma model attenuated by a $\\sim$1.4$\\times$10$^{21}$~cm$^{-2}$ extinction. The low extinction, which is consistent with the optical reddening, indicates that the source is a foreground star toward the Galactic Center region. Based on the spectroscopic parallax distance of $\\sim$530~pc, the peak X-ray luminosity amounts to $\\sim$1$\\times$10$^{32}$~erg~s$^{-1}$ (0.5--10~keV). This is much larger than the X-ray luminosity of ordinary late-type main-sequence stars, and the X-ray emission is unattributable to a hidden late-type companion that comprises a wide binary system with the A-star. We discuss possible natures of HD\\,161084 and suggest that it is most likely an interacting binary with elevated magnetic activity in the companion such as the Algol-type system. The flux detected by Suzaku during the burst is $\\sim$100 times larger than the quiescent level measured using the archived XMM-Newton and Chandra data. The large flux amplification makes this star a unique example among sources of this class. ",
        "introduction": "Main-sequence stars with intermediate (late B to early A) spectral types are considered intrinsically X-ray inactive. These stars have no X-ray production mechanisms such as shocks in unstable stellar winds causing the X-ray emission in early-type stars (earlier than B2; \\cite{lucy80,berghoefer97}), or the magnetic activity as a consequence of surface convection layer and differential rotation responsible for the X-ray emission in late-type stars. Indeed, most X-ray surveys of A-type stars in the field and in open clusters show a consistent paucity of their X-ray detections (e.g. \\cite{schmitt85,schmitt90,micela90}).  X-ray emission from A-type stars, if detected, is usually attributed to its magnetically-active late-type companion (e.g. \\cite{huelamo00,briggs03}) that comprises a binary system. The idea is consistent with the following observational facts: (1) There is a lack of any correlation between the X-ray luminosity and various stellar properties of A-type stars \\citep{simon95,panzera99}. (2) The X-ray emission from A-type stars is characterized by X-ray luminosities of 10$^{28}$--$10^{30}$~erg~s$^{-1}$ with averaged soft emission, which is similar to those seen in late-type main-sequence stars.  In some high-resolution X-ray imaging studies, the position of the X-ray emission is located precisely enough to claim that the emission is from the late-type companion and not from the A-type star itself \\citep{stelzer03a,stelzer03b}. In some unresolved systems, too, the X-ray emission is confirmed to be from the late-type companion by the eclipsing observation \\citep{schmitt93} or by the Doppler shift measurements of emission lines in the X-ray spectrum \\citep{chung04}.  Nevertheless, these pieces observational evidence do not entirely rule out the possibility of the intrinsic X-ray emission from A-type stars. A-type stars with a debris disk (e.g., $\\beta$ Pic) may radiate thermal X-rays fueled by accretion \\citep{hempel05} and those with strong magnetic activity (e.g, IQ Aur) may cause X-ray emission caused by magnetically-confined winds \\citep{babel97}.  The long-term X-ray behavior of A-type stars is an important property to characterize their emission and to understand its origin. Monitoring observations on particular A-type stars are necessary to reveal their long-term flux and spectral variation, but the number of such studies is only a few (e.g. \\cite{briggs03}). This is because systematic X-ray studies of A-type stars have been conducted by utilizing all-sky survey data \\citep{simon95,panzera99} or by snapshot observations of open clusters including A-type stars \\citep{briggs03}.  \\medskip  Here, we study the long-term X-ray behavior of an A-type star by exploiting the wealth of data obtained in the Galactic Center region. Every X-ray satellites repeatedly observed this region, which contains many A-type stars as foreground objects. We surveyed the Suzaku images of the Galactic Center and serendipitously found that the A-type star HD\\,161084 showed an intense flare during one of the mapping observations. The star is a Tycho reference source classified as A1\\,V or A1\\,IV \\citep{wright03}. In the X-ray band, the star has been soft and faint at $\\sim$10$^{29.5}$~erg~s$^{-1}$ in the 0.5--10~keV luminosity in the archived X-ray data, but a sudden flaring made this source as bright as $\\sim$10$^{32}$~erg~s$^{-1}$, which is hard to achieve by a hidden ordinary late-type companion. We report the result obtained by the X-ray Imaging Spectrometer (XIS) on-board Suzaku and discuss the nature of this source. ",
        "conclusions": "\\subsection{Extinction and Distance to HD\\,161084}\\label{sect:4-1} HD\\,161084 has been a largely ignored source and nothing beyond the catalog information is known. We start with estimating the distance to HD\\,161084. From the SIMBAD database, the source has a \\textit{B}-band and \\textit{V}-band magnitudes of 10.29~mag and 10.11~mag, respectively. It is classified as a star of A1\\,VI or A1\\,V.  Assuming the source is an A1\\,V star, the intrinsic (\\textit{B}--\\textit{V}) color is $\\sim$0.015~mag \\citep{drilling00}. The source is estimated to have a reddening of (\\textit{B}--\\textit{V}) $\\sim$~0.165~mag to be compatible with the observed color, which can be converted to an extinction of $A_{V} \\sim$~0.51~mag. This is consistent with the independent measurement of the X-ray absorption of $\\sim$1.4$\\times$10$^{21}$~cm$^{-2}$ assuming the conversion factor of $N_{\\rm{H}}$/$A_{V}=$1.79$\\times$10$^{21}$~cm$^{-2}$~mag$^{-1}$ \\citep{predehl95}. The agreement gives another support for the identification of the Suzaku source with HD\\,161084. From the dereddended \\textit{V}-band magnitude of HD\\,161084 and the intrinsic \\textit{V}-band magnitude of A1\\,V stars, we estimate a distance of $\\sim$530~pc.  We repeated the same procedure assuming that HD\\,161084 is an A0\\,V or A2\\,V star to see how the uncertainty of the spectral-type measurement affects our distance estimate. The intrinsic (\\textit{B}--\\textit{V}) color differs by $\\sim$0.035~mag and the absolute \\textit{V}-band magnitude by $\\sim$0.33~mag among A0\\,V, A1\\,V, and A1\\,V types \\citep{drilling00}. The derived distance is $\\sim$586~pc for A0\\,V and $\\sim$480~pc for A2\\,V. This converts to the uncertainty of the X-ray luminosity by $\\sim$20\\%.   \\subsection{Flaring and Quiescent X-rays}\\label{sect:4-2} We detected intense X-ray emission from HD\\,161084 with a flare-like light curve and hard emission accountable by a plasma temperature of $\\sim$5~keV. The peak luminosity at the maximum count rate bin in figure~\\ref{fg:f3} is $\\sim$1$\\times$10$^{32}$~erg~s$^{-1}$. The $e$-folding time of the star is similar to those of other flare stars such as the Sun and RS CVn-type binaries. Based on the plot of X-ray volume emission measure versus the plasma temperature \\citep{shibata99}, we estimate the magnetic field strength to be $\\sim$50~gauss and the flare loop length to be $\\sim$0.05~AU. The flare was detected once in a total duration of $\\sim$9 days (table~\\ref{tb:t1}), which indicates that the flare frequency is once in 2--113 days (90\\% confidence; \\cite{kraft91}).  The XMM-Newton spectrum was obtained during a quiescent state of this source. We could not obtain any meaningful X-ray spectra from other observations and from the first 10~ks (phase 1) of the Suzaku observation. We speculate that HD\\,161084 was at a quiescent state in these observations with a similar X-ray brightness and hardness recorded by XMM-Newton.   \\subsection{Nature of HD\\,161084}\\label{sect:4-3} In the quiescent phase, the X-ray emission from HD\\,161084 appears to be attributable to an uncovered normal low-mass main-sequence companion or the intrinsic emission from an A-type star. However, the intense flare with an amplified X-ray flux and a hard spectrum makes both interpretations unlikely. The flare luminosity is beyond the reach of normal late-type stars. The X-ray emission from the debris disk and magnetic A-type stars is weak ($\\sim$10$^{28}$~erg~s$^{-1}$), soft, and constant \\citep{hempel05,cash82,golub83}, which contradicts the behavior of HD\\,161084 during the flare.  We speculate that the X-ray emission from HD\\,161084 is from its hidden late-type companion, and the secondary star is exceptionally bright among stars of this type. The level of X-ray activity in late-type stars is known to be an increasing function of rotational velocity (e.g. \\cite{pallavicini81}). This is a natural consequence that the X-ray--emitting plasma is fueled by the energy release of magnetic fields generated by the turbulence in the surface convection layer and the differential rotation. Two major classes of late-type stars are known to rotate faster and emit exceptionally brighter X-rays than ordinary late-type stars.  One is pre--main-sequence sources, whose faster rotation stems from the accumulation of angular momentum via accretion and contraction during their formation. These sources can be bright up to $\\sim$10$^{32}$~erg~s$^{-1}$ with hard X-ray emission (e.g, \\cite{imanishi01}). A bright and hard X-ray flare of $\\sim$10$^{31.6}$~erg~s$^{-1}$ (0.5--8.0~keV) was observed from an intermediate-type star HD\\,38563S (B3--B5; \\cite{yanagida04}). The X-ray emission from this source can be from a hidden late-type pre--main-sequence companion around the B star, which is quite conceivable considering that HD\\,38563S is in a young star-forming region NGC\\,2068. We consider that HD\\,161084, in contrast, is not located in a star-forming region and thus is unlikely to have a pre--main-sequence companion. We examined the near-infrared image by SIRIUS, but no clustering of red sources was present, which would indicate star-forming regions. Moreover, the extinction toward HD\\,161084 is consistent with the integration of interstellar hydrogen density of 1~cm$^{-3}$ along the line of sight. Most sources in star-forming regions suffer an additional extinction by natal molecular cloud and circumstellar matter. The lack of such extra extinction further supports that HD\\,161084 is a field star.  The other class of fast rotators is interacting binaries. The tidal force and mass transfer between the two components of a close binary synchronizes the stellar rotation with the orbital motion, giving rise to a faster rotation and elevated X-ray activity in the late-type constituent. A handful of classes of such interacting binaries are known \\citep{hall89,richards93}. Among them, the classes that may contain an intermediate-type main-sequence star are the Algol-type (semi-detached) and the W UMa-type (contact) binaries.  The Algol-type binary is a well-established class of intense X-ray emitters with an intermediate-type dwarf and late-type giants or sub-giants. The X-ray emission from Algol-type sources are attributed to fast-rotating late-type star \\citep{white83,schmitt99,chung04}. They emit hard and luminous X-rays with occasional flares \\citep{singh95,singh96,ottmann96}, which is comparable to those observed from HD\\,161084 in the flare state. We suggest that HD\\,161084 is an interacting binary of the Algol-type based on the observed X-ray properties.  W UMa-type stars also cause X-ray flares with the luminosity up to $\\sim$10$^{31}$~erg~s$^{-1}$ \\citep{mcgale96,choi98}. In almost all these systems, the spectral type of the two stars are nearly equal \\citep{maceroni96}. The secondary star of HD\\,161084 is thus likely to be another A1-type star. The origin of the X-ray emission boils down to the intrinsic X-ray emission from A-type stars, which we concluded unlikely to account for the emission from HD\\,161084. We thus speculate that HD\\,161084 is not a W UMa-type binary.  What makes this source unique among these classes is the degree of amplification of the flare flux from the quiescent flux. Both the Algol and W UMa binaries show occasional flares, but their flux increase from the quiescent state is commonly a few times \\citep{choi98,gondoin04,gondoin04_2,mcgale96,nordon07,white86}. It may be the case that the X-ray emitting late-type companion in the HD\\,161084 system was eclipsed during all the previous X-ray observations except for the phases 2 and 3 of the Suzaku observation. The interacting binary nature of HD\\,161084 is conclusively identified by established methods, such as photometric eclipse and spectroscopic monitoring. We encourage follow-up observation of this source."
    },
    "0711/0711.2458_arXiv.txt": {
        "abstract": "{} {We analyze the distribution of the interstellar matter in the environs of the \\hii\\ region Gum\\,31, excited by the open cluster NGC\\,3324, located in the complex Carina region, with the aim of investigating the action of the massive stars on the surrounding neutral material.} {We use neutral hydrogen 21-cm line data, radio continuum images at 0.843, 2.4 and 4.9 GHz, $^{12}$CO(1-0)  observations, and IRAS and MSX infrared data.} {Adopting a distance of 3 kpc for the \\hii\\ region and the ionizing cluster, we derived an electron density of 33$\\pm$3 cm$^{-3}$ and an ionized mass of (3.3$\\pm$1.1)$\\times$10$^3$  M$_{\\odot}$ based on the radio continuum data at 4.9 GHz. The \\hi\\ 21-cm line images revealed an \\hi\\ shell surrounding the H\\,{\\sc ii} region. The \\hi\\ structure is 10.0$\\pm$1.7 pc in radius, has a neutral mass of 1500$\\pm$500 M$_{\\odot}$, and is expanding at 11 km\\,s$^{-1}$. The associated molecular gas amounts to (1.1$\\pm$0.5)$\\times$10$^5$ M$_{\\odot}$, being its volume density of about 350 \\cmtres. This molecular shell could represent the remains of the cloud where the young open cluster NGC\\,3324 was born or could have originated by the shock front associated with the \\hii\\ region. The difference between the ambient density and the electron density of the \\hii\\ region suggests that the \\hii\\ region is expanding.  The distributions of the ionized and molecular material, along with that of the emission in the MSX band A, suggest that a photodissociation region has developed at the interface between the ionized and molecular gas. The copious UV photon flux from the early type stars in NGC\\,3324 keeps the \\hii\\ region ionized.  The characteristics of a relatively large number of the IRAS, MSX, and 2MASS point sources projected onto the molecular envelope are compatible with  protostellar candidates, showing the presence of active star forming regions. Very probably, the expansion of the \\hii\\ region has triggered stellar formation in the molecular shell. } {} ",
        "introduction": "The interstellar medium associated with the birth place of massive stars, such as O or early B-type stars, is made up of dense giant molecular clouds. Massive stars are characterized by intense photon fluxes which ionize and photodissociate the surrounding material. An \\hii\\ region is a direct consequence of the high rate of Lyman continuum luminosity. Initially, the \\hii\\ region is a small, high density region, commonly named ultra-compact \\hii\\ region (UC \\hii, Wood $\\&$ Churchwell 1989). If the photon flux rate of the massive star is sufficiently high, the \\hii\\ region evolves into a normal \\hii\\ region.  A neutral shell encircles the \\hii\\ region during the expanding phase (e.g. Spitzer 1978). The presence of these neutral shells is observed in the \\hi\\ 21-cm line emission distribution (e.g. Deharveng et al. 2003). Molecular line studies have allowed the identification of molecular gas following the outer borders of \\hii\\ regions, indicating the presence of photodissociation regions (PDRs). Deharveng et al. (2005) have detected these PDRs in a number of \\hii\\ regions.  In the present study, we analyze the distribution of the ionized and neutral material associated with the \\hii\\ region Gum\\,31 (Gum 1955) based on \\hi\\ 21-cm line emission data, radio continuum information at different frequencies, and IR and molecular data.  The \\hii\\ region Gum\\,31 is about 15\\arcmin\\ in size and approximately circular in shape (Fig. 1). It is located at  {\\it (l,b)} = (286\\deg 12\\arcmin, --0\\deg 12\\arcmin) in the complex region of Carina and is considered a member of the Car OB1 association. The SuperCOSMOS image (Parker et al. 2005) shows a quite inhomogeneous \\hii\\ region, with a sharp and bright rim towards lower galactic longitudes and lower galactic latitudes, looking fainter and more diffuse towards higher galactic longitudes.  Based on data of the radio recombination lines (RRLs) H109$\\alpha$ and H110$\\alpha$ at 5 GHz, Caswell \\& Haynes (1987) found that the LSR velocity of the ionized gas in Gum\\,31 is --18 \\kms, similar to the velocities of other \\hii\\ regions in the area of the Car OB1 association (Georgelin et al. 1986), and derived an electron temperature $T_e$ = 7100 K. They estimated a flux density $S_{5GHz}$ = 35 Jy.  The excitation sources of the \\hii\\ region Gum\\,31 are the OB star members of the open cluster NGC\\,3324. The brightest star in this cluster is HD\\,92206, which is the visual double star IDS\\,10336-5806 in the Index catalogue (Jeffers et al. 1963) with a 1 mag fainter companion (HD\\,92206B) placed 5\\arcsec\\ to the East. Another bright cluster member is located 35\\arcsec\\ to the SW. This star is CD--57\\deg 3378, also referred to as HD\\,92206C in the literature.  The three brightest stars in NGC\\,3324 have published spectral types. Both HD\\,92206A and B are classified as O6.5V, and HD\\,92206C as O9.5V by Mathys (1988). Walborn (1982) classifies HD\\,92206A as O6.5V(n) and the component C as O8.5Vp.  Moffat $\\&$ Vogt (1975) first carried out photometric observations of about 12 stars in the cluster and estimated a color excess $E(B-V)\\sim$ 0.45$\\pm$0.05 mag and a distance $d$ = 3.3 kpc. Clari\\'{a} (1977), using UBV photometry, confirmed previous results by Moffat $\\&$ Vogt and estimated $d$ = 3.1 kpc. Vazquez \\& Feinstein (1990) derived a distance $d$ = 3.6 kpc for the cluster from UBVRI photometry. More recently, Carraro et al. (2001) found about 25 new possible cluster members and derived  $d$ = 3.0$\\pm$0.1 kpc. On the other hand, distance estimates for Car\\,OB1 are in the range 1.8-2.8 kpc (Walborn 1995). Bearing in mind these results we adopted  $d$ = 3.0$\\pm$0.5 kpc for both Gum\\,31 and the ionizing cluster.  Carraro et al. (2001) find evidence for pre-main sequence members beginning at about late B spectral type, which suggests an extremely young age for NGC\\,3324 ($\\leq$ (2-3)\\por 10$^6 $ yr). The O-type members would be stars recently arrived on the Zero Age Main Sequence.  \\begin{figure} \\resizebox{8.8cm}{!}{\\includegraphics{7028fi1.eps}} \\caption{SuperCOSMOS image showing the \\hii\\ region Gum\\,31. The cross  marks the position of the multiple system HD\\,92206. The intensity units are arbitrary. } \\label{optical image} \\end{figure}  \\begin{table} \\centering \\caption[]{Radio  data: relevant parameters.} \\begin{tabular}{lc} \\hline Radio continuum at 4.85 GHz \\\\ \\hline Angular resolution             &  5\\farcm 0 \\\\ rms noise                      &  10 \\mjyb \\\\ \\hline Radio continuum at 2.4 GHz \\\\ \\hline Angular resolution             &  10\\farcm 4 \\\\ rms noise                      &  12 \\mjyb \\\\ \\hline Radio continuum at 0.843 GHz \\\\ \\hline Angular resolution      &  43\\arcsec\\ $\\times$ 51\\arcsec \\\\ rms noise                      &  1 \\mjyb \\\\ \\hline \\hi\\ data \\\\ \\hline Synthesized beam               &   2\\farcm 4 $\\times$ 2\\farcm 1 \\\\ Number of channels             & 256 \\\\ Velocity coverage              &   (--190,+230) \\kms \\\\ Velocity resolution            &   1.64  \\kms  \\\\ RMS noise level                &  1.6 K \\\\ \\hline $^{12}$CO(1-0) NANTEN data \\\\ \\hline Angular resolution             &  2\\farcm 7 \\\\ Velocity coverage              &  (--50,+50) \\kms \\\\ Velocity resolution            &   0.2 \\kms  \\\\ RMS noise level                &   1.0 K \\\\ \\hline \\end{tabular} \\end{table} ",
        "conclusions": "\\begin{table*} \\centering \\caption{Main parameters of the ionized and neutral gas in Gum\\,31} \\begin{tabular}{lc} \\hline Adopted distance (kpc)              &  3.0$\\pm$0.5  \\\\ \\hline {\\it \\hii\\ region} \\\\ Flux density at 4.85 GHz (Jy)       &  37.7$\\pm$2.5  \\\\ Angular radius (arcmin)             &  7.5$\\pm$0.2      \\\\ Linear radius (pc)                  &  6.5$\\pm$1.0  \\\\ Emission measure (pc cm$^{-6}$)     &  (1.5$\\pm$0.2)\\por 10$^4$  \\\\ rms electron density [$n_e$](\\cmtres)      &  33$\\pm$3  \\\\ Used Lyman UV photons [$ log N_{Ly-c}$] (s$^{-1}$) &  49.0$\\pm$0.1  \\\\ Ionized mass (\\msun)                & 3300$\\pm$1100  \\\\ \\hline {\\it Neutral atomic shell} \\\\ {\\it (l,b)} centroid                & 286\\deg 15\\arcmin,--0\\deg 10\\arcmin  \\\\ Velocity interval v$_1$,v$_2$ (\\kms)  & $-$13,$-$32  \\\\ \\hi\\ systemic velocity (\\kms)       & $-$23        \\\\ Expansion velocity $v_{exp}$ (\\kms) & 11           \\\\ Radius of the \\hi\\ structure (arcmin)  & 11.5       \\\\ Radius of the \\hi\\ structure $R$ (pc)  & 10.0$\\pm$1.7 \\\\ Atomic mass in the shell (\\msun)       &  1500$\\pm$500         \\\\ \\hline {\\it Molecular shell}\\\\ Velocity interval v$_1$,v$_2$ (\\kms)  & --27.2,--14.0 \\\\ Mean H$_{2}$ column density (\\cmdos) &  1.2\\por 10$^{22}$ \\\\ Molecular mass (\\msun)              &   (1.1$\\pm$0.5)\\por 10$^5$ \\\\ \\hline {\\it Dust related to the \\hii\\ region and the neutral shell} \\\\ Total dust mass (\\msun)                   & 60$\\pm$20        \\\\ Dust color temperature (K)       & 34$\\pm$7       \\\\ \\hline \\end{tabular} \\end{table*}  \\subsection{Physical parameters}  The main parameters of the dust and the ionized and neutral gas related to Gum\\,31 are summarized in Table 2. The parameters of the ionized gas were derived from the image at 4.85 GHz. The uncertainty in the flux density corresponds  to an error of 0.1 \\jyb\\ in the estimate of the radio continuum background. The electron density and the \\hii\\ mass were obtained from the expressions by Mezger \\& Henderson (1967) for a spherical \\hii\\ region of constant electron density (rms electron density $n_e$). The presence of He\\,{\\sc ii} was considered by multiplying the \\hi\\ mass by 1.27. The number of UV photons necessary to ionize the gas $log \\ N_{Ly-c}$ was derived from the radio continuum results. Errors in the linear radius, in the rms electron density, and in the excitation parameter come from the distance uncertainty. The high electron density of the \\hii\\ region is compatible with the relatively short lifetime for the cluster.  The parameters of the neutral atomic gas includes: the {\\it (l,b)} position of the centroid of the \\hi\\ shell, the velocity interval spanned by the structure, the systemic and expansion velocities, the radius of the neutral gas structure, and the associated atomic mass.  The expansion velocity was estimated as in previous papers (see Cappa et al. 2005 and references therein) as $v_{exp} = (v_2 - v_1$)/2 + 1.6 \\kms. The extra 1.6 \\kms\\ allows for the presence of \\hi\\ in the caps of the expanding shell, which should be present in the images at v$_1$+1.6 km/s and v$_2$--1.6 km/s. These caps are not detected in the present case, as in most of the cases, probably because of their weak brightness temperature.  The radius of the \\hi\\ structure was estimated from Fig. 4 and corresponds to the position of the maxima in the shell. The neutral atomic mass corresponds to the mass excess in the shell, assuming that the gas is optically thin and including a He abundance of 10\\%.  The $H_2$ column density ($N_{H2}$) and the molecular mass  were estimated from the $^{12}$CO data, making use of the empirical relation between the integrated emission $W_{CO}$ (= $\\int T dv)$ and $N_{H2}$. We adopted $N_{H2}$ = (1.06$\\pm$0.14) \\por W$_{CO}$ $\\times$ 10$^{20}$ cm$^{-2}$ (K \\kms)$^{-1}$, obtained from $\\gamma$-ray studies of molecular clouds in the Orion region (Digel \\etal\\ 1995). To derive the molecular mass, we integrated the CO emission within a circle of about 16\\arcmin\\ in radius (= 15 pc at 3.0 kpc), centered at {\\it(l,b)} = (286\\deg16\\arcmin,--0\\deg8\\arcmin).  The ambient density derived by distributing the molecular mass over a sphere of 15.0 pc in radius is  $\\simeq$350 \\cmtres, reinforcing the idea that the molecular gas represents the remains of the original material where the open cluster NGC\\,3324 was born. The difference between the electron density and the ambient density indicates that the \\hii\\ region is expanding.  It is known that massive stars have stellar winds with high mass loss rates and terminal velocities (Chlebovski \\& Garmany 1991; Prinja et al. 1990), which generally affect the surrounding interstellar medium creating {\\it interstellar bubbles} (Weaver et al. 1977). In the optical domain, they appear as ring-shaped structures (Lozinskaya 1982). The optical appearance of the \\hii\\ region, almost spherical without clear evidence for a central cavity, suggests that the massive stars in the cluster have weak stellar winds (which is supported by the short age of the cluster) and/or have existed during a very short period of time to create an interstellar bubble in an interstellar medium as dense as observed here.  \\begin{figure} \\resizebox{8.0cm}{!}{\\includegraphics{7028fi8.eps}} \\caption{Schematic view of Gum\\,31 and the respective locations of the different gas components and of interstellar dust.} \\end{figure}  Figure 8 shows schematically the distribution of the different gas components. Although \\hi\\ and molecular clumps anticorrelate in position, the large scale \\hi\\ and molecular gas coincide in position around the nebula but in the region at higher galactic latitudes, where the CO emission is observed outside the \\hi\\ region. The large amount of molecular gas compared to the neutral atomic gas supports the idea that the CO emission which encompasses the ionized nebula represents the remains of the molecular material where the open cluster was born. Part of this \\hi\\ gas may have  originated in the photodissociation of the molecular gas.  \\begin{table*} \\centering \\caption[]{IR Point sources and YSO candidates from the IRAS, MSX and 2MASS catalogs} \\begin{tabular}{rllcrrrrc} \\hline \\hline $\\#$ &    $l$[$\\circ\\ \\arcmin$]    &  $b$[$\\circ\\ \\arcmin$]    & {\\it IRAS} source &     \\multicolumn{4}{c}{Fluxes[Jy]}  & $L_{IRAS}[10^3 L_{\\sun}]$ \\\\ & & &  &  12$\\mu$m  &    25$\\mu$m &   60$\\mu$m &  100$\\mu$m &  \\\\ \\hline 1 &  285\\deg 59\\farcm 82 & +0\\deg 05\\farcm 82 & 10349-5801 &   1.6 &   2.5 &   27.9 &  104  &  1.5\\\\ 2 &  286\\deg 04\\farcm 98 & --0\\deg 25\\farcm 08 & 10335-5830 &   0.8 &   3.5 &   22.8 & 111  &  1.5 \\\\ 3 &  286\\deg 08\\farcm 58 & --0\\deg 06\\farcm 96 & 10351-5816 &   6.3 &   6.8 &  315   & 1480  & 18\\\\ 4  &  286\\deg 08\\farcm 64 & --0\\deg 18\\farcm 78 & 10343-5826 &   5.5 &   4.3 &  103   & 340  &  5 \\\\ 5  &  286\\deg 10\\farcm 92  & --0\\deg 14\\farcm 7 &  10349-5824 & 5.9 &  9.0 & 163 & 1660  &  18\\\\ 6 &  286\\deg 12\\farcm 18 & +0\\deg 10\\farcm 20 & 10365-5803 &   7.2 &  86   & 1170   & 2780  & 43  \\\\ 7 &  286\\deg 15\\farcm 12 & --0\\deg 25\\farcm 44 & 10346-5835 &   1.1 &   3.3 &   11.7 & 1430  &  14\\\\ 8 &  286\\deg 17\\farcm 34 & +0\\deg 00\\farcm 42 & 10365-5814 &   2.3 &   1.7 &   85   & 246  &  4\\\\ 9 &  286\\deg 22\\farcm 5 & --0\\deg 15\\farcm 3 & 10361-5830 &  12.4 &  38.4 &  626   & 2160  &  30\\\\ 10 &  286\\deg 26\\farcm 94 & --0\\deg 22\\farcm 92 & 10361-5839 &   2.8 &   5.4 &   51.6 & 307   & 4 \\\\ 11 &  286\\deg 29\\farcm 34 & --0\\deg 16\\farcm 68 & 10368-5835 &   2.1 &   2.3 &   84   & 271   & 4 \\\\ 12 &  286\\deg 33\\farcm 72 & --0\\deg 07\\farcm 08 & 10379-5828 &   1.0 &   2.0 &   13.8 &  540   & 5 \\\\ \\hline \\hline $\\#$ &   $l$[$\\circ \\arcmin$] &  $b$[$\\circ\\ \\arcmin$] &   {\\it MSX} source &     \\multicolumn{4}{c}{Fluxes[Jy]} \\\\ &  &  &  &  8$\\mu$m &  12$\\mu$m & 14$\\mu$m & 21$\\mu$m & Class. \\\\ \\hline 13 & 286\\deg 09\\farcm 78 & --0\\deg 11\\farcm 28 & G286.1626-00.1877 & 0.7311 & 1.298 & 1.256 & 2.806 &  C\\hii\\ \\\\ 14 & 286\\deg 12\\farcm 36 & --0\\deg 09\\farcm 66 & G286.2056-00.1611 & 0.1585 & 0.9313 & 2.151 & 6.904  & MYSO\\\\ 15 & 286\\deg 12\\farcm 48 & --0\\deg 10\\farcm 32 & G286.2077-00.1720 & 0.0870 & 0.9481 & 1.764 & 2.9 &  MYSO \\\\ 16 & 286\\deg 12\\farcm 54 & +0\\deg 10\\farcm 14  & G286.2086+00.1694 & 1.353  & 2.882  & 7.182 & 40.57  & MYSO\\\\ 17 & 286\\deg 12\\farcm 6  & --0\\deg 10\\farcm 68 & G286.2096-00.1775 & 0.2202 & 0.9538 & 1.368 & 7.206  & MYSO \\\\ 18 & 286\\deg 15\\farcm 42 & --0\\deg 19\\farcm 44 & G286.2566-00.3236 & 2.04   & 2.151  & 1.34 & 4.328  & C\\hii \\\\ 19 & 286\\deg 21\\farcm 48 & --0\\deg 17\\farcm 58 & G286.3579-00.2933 & 0.7126 & 1.815 & 2.677 & 6.065  & MYSO\\\\ 20 & 286\\deg 22\\farcm 5 & --0\\deg 15\\farcm 78 & G286.3747-00.2630 & 3.591  & 4.756 & 2.409 & 7.577  & C\\hii\\ \\\\ 21 & 286\\deg 22\\farcm 62 & --0\\deg 15\\farcm 36 & G286.3773-00.2563 & 1.628 & 2.918 & 3.855 & 12.1  & MYSO \\\\ \\hline \\hline $\\#$ &  $l$[$\\circ \\arcmin$]   &  $b$[$\\circ \\arcmin$]  &   {\\it 2MASS} source & $J$[mag] & $H$[mag] & $K_s$[mag] & $(J-H)$ & $(H-K)$  \\\\ \\hline 22 & 285\\deg 54\\farcm 42 & --0\\deg 13\\farcm 74 &  10350210-5831039 & 10.584 & 10.631 & 10.564 & -0.047 & 0.067\\\\ 23 & 286\\deg 04\\farcm 44 & --0\\deg 00\\farcm 24 & 10365972-5824186 & 12.042 & 10.51  & 9.284  & 1.532  & 1.226\\\\ 24 & 286\\deg 06\\farcm 36 & --0\\deg 08\\farcm 64 & 10364112-5832326 & 11.367 & 11.385 & 11.28  & -0.018 & 0.105\\\\ 25 & 286\\deg 06\\farcm 54 & --0\\deg 08\\farcm 40 & 10364296-5832267 & 10.465 & 10.513 & 10.439 & -0.048 & 0.074\\\\ 26 & 286\\deg 09\\farcm 60 & --0\\deg 00\\farcm 30 & 10373406-5826540 & 12.487 & 11.454 & 10.513 & 1.033  & 0.941\\\\ 27 & 286\\deg 09\\farcm 78 & --0\\deg 11\\farcm 22 & 10365396-5836293 & 12.242 & 10.99  & 10.103 & 1.252  & 0.887\\\\ 28 & 286\\deg 10\\farcm 08 & --0\\deg 09\\farcm 12 & 10370395-5834489 & 10.772 & 9.712  & 8.685  & 1.06   & 1.027\\\\ 29 & 286\\deg 10\\farcm 20 & --0\\deg 11\\farcm 10 & 10365749-5836366 & 13.99  & 12.815 & 11.984 & 1.175  & 0.831\\\\ 30 & 286\\deg 11\\farcm 16 & --0\\deg 02\\farcm 64 & 10373574-5829405 & 15.318 & 13.343 & 11.614 & 1.975  & 1.729\\\\ 31 & 286\\deg 12\\farcm 96 & --0\\deg 04\\farcm 86 & 10373956-5832311 & 10.926 & 10.752 & 10.554 & 0.174  & 0.198\\\\ 32 & 286\\deg 13\\farcm 02 & --0\\deg 10\\farcm 86 & 10371717-5837460 & 11.793 & 11.773 & 11.665 & 0.02   & 0.108\\\\ 33 & 286\\deg 13\\farcm 32 & --0\\deg 08\\farcm 46 & 10372824-5835492 & 12.191 & 11.301 & 10.448 & 0.89   & 0.853\\\\ 34 & 286\\deg 13\\farcm 38 & --0\\deg 10\\farcm 20 & 10372226-5837229 & 7.563  & 7.588  & 7.479  & -0.025 & 0.109\\\\ 35 & 286\\deg 13\\farcm 92 & --0\\deg 17\\farcm 87 & 10365763-5844052 & 12.22  & 11.882 & 11.51  & 0.338  & 0.372\\\\ 36 & 286\\deg 19\\farcm 26 & --0\\deg 06\\farcm 84 & 10381421-5837192 & 12.635 & 11.886 & 11.349 & 0.749  & 0.537\\\\ 37 & 286\\deg 20\\farcm 10 & --0\\deg 19\\farcm 80 & 10373105-5849026 & 15.155 & 12.8   & 11.127 & 2.355  & 1.673\\\\ 38 & 286\\deg 21\\farcm 06 & --0\\deg 16\\farcm 86 & 10375219-5847133 & 12.271 & 11.363 & 10.675 & 0.908  & 0.688\\\\ 39 & 286\\deg 21\\farcm 90 & --0\\deg 04\\farcm 92 & 10383875-5836566 & 12.116 & 11.636 & 11.187 & 0.48   & 0.449\\\\ 40 & 286\\deg 22\\farcm 92 & --0\\deg 13\\farcm 20 & 10381461-5844416 & 12.367 & 11.815 & 11.398 & 0.552  & 0.417\\\\ 41 & 286\\deg 24\\farcm 92 & --0\\deg 18\\farcm 66 & 10380736-5850240 & 12.477 & 11.538 & 10.78  & 0.939  & 0.758\\\\ 42 & 286\\deg 28\\farcm 11 & --0\\deg 24\\farcm 50 & 10380702-5857039 & 12.724 & 11.599 & 10.726 & 1.125  & 0.873 \\\\ 43 & 286\\deg 28\\farcm 44 & --0\\deg 23\\farcm 71 & 10381226-5856318 & 13.348 & 12.274 & 11.355 & 1.074  & 0.919\\\\ 44 & 286\\deg 29\\farcm 46 & --0\\deg 06\\farcm 30 & 10392451-5841486 & 13.836 & 12.561 & 11.619 & 1.275  & 0.942\\\\ 45 & 286\\deg 29\\farcm 46 & --0\\deg 08\\farcm 10 & 10391799-5843257 & 13.195 & 12.322 & 11.624 & 0.873  & 0.698\\\\ 46 & 286\\deg 30\\farcm 90 & --0\\deg 06\\farcm 30 & 10393410-5842321 & 11.701 & 10.423 & 9.334  & 1.278  & 1.089\\\\ 47 & 286\\deg 31\\farcm 27 & --0\\deg 28\\farcm 40 & 10381363-5902003 & 13.368 & 12.557 & 11.982 & 0.811  & 0.575\\\\ \\hline \\hline \\end{tabular} \\end{table*}  \\subsection{Energy budget}  Bearing in mind that massive stars have a copious UV flux capable of ionizing the surrounding \\hi\\ gas, we investigate in this section whether the massive stars in NGC\\,3324 can provide the energy to ionize the gas.  As described in Sect. 1, the only O-type stars in the open cluster are the multiple system HD\\,92206, which contains three O-type stars classified as O6.5V, O6.5V and O8.5 (Walborn 1982, Mathys 1988). Considering that HD\\,92206A is almost 1 mag brighter than component B (Clari\\'{a} 1977), the similar spectral types found by Mathys (1988) suggest that HD\\,92206A is probably a spectroscopic binary O6.5V + O6.5V. Taking into account the UV photon fluxes emitted by the stars, $N_{*}$ (s$^{-1}$), derived by Martins et al. (2005) from stellar atmosphere models, a group of four massive stars having the spectral types indicated above have a  total UV photon flux corresponding to $log \\ N_*$ =  49.4. By comparing $N_{*}$ with the UV photons used to ionize the gas ( $N_{Ly-c}$)  listed in Table 2, we conclude that the massive O-type stars in NGC\\,3324 can maintain the \\hii\\ region ionized.  The radius of the Str\\\"omgren sphere formed in a region with a volumetric ambient density of 350  cm$^{-3}$ is $\\simeq$ 1.7 pc, which is lower than the radius of the ionized region. Following Spitzer (1978), we estimated that the \\hii\\ region has been expanding for $\\approx$ 1\\por 10$^6$ yr.  The ambient density we  derived from the  $^{12}$CO(1-0) data is larger than the rms electron density, 33 cm$^{-3}$ (see Table 2), giving additional support for the interpretation that the \\hii\\ region is expanding.  \\begin{figure*} \\resizebox{13cm}{!}{\\includegraphics{7028fi9.eps}} \\caption{{\\it Top:} Point sources from the IRAS (triangles), MSX (circles), and 2MASS catalogs overlaid onto the SuperCOSMOS image and the $^{12}$CO and C$^{18}$O contours (see Fig. 5 for details). } \\label{ir} \\end{figure*}    We have analyzed the interstellar medium in the environs of the \\hii\\ region Gum\\,31 to investigate the action of the massive stars in the exciting open cluster NGC\\,3324 on the surrounding neutral material.  We based our study on \\hi\\ 21-cm line emission data belonging to the SGPS, radio continuum data at  0.843, 2.4 and 4.85 GHz from the PMN Southern Radio Survey, $^{12}$CO data from Yonekura et al. (2005), and IRAS (HIRES) and MSX data.  Adopting a distance of 3.0$\\pm$0.5 kpc, we have derived an ionized mass of 3300$\\pm$1100 \\msun\\ and an electron density of 33$\\pm$3 \\cmtres. The four O-type stars in the HD\\,92206 multiple system can provide the necessary UV photon flux to maintain the \\hii\\ region ionized.  The \\hi\\ emission distribution in the environs of Gum\\,31 shows the presence of an \\hi\\ shell approximately centered at the position of the multiple system HD\\,92206. The \\hi\\ shell closely encircles the optical nebula. It is detected within the velocity range --32 to --13 \\kms\\ and its systemic velocity of  --23 \\kms\\ is coincident, within errors, with the velocity of the ionized gas in the nebula (--18 \\kms). The \\hi\\ structure is 10.0$\\pm$1.7 pc in radius and expands at about 11 \\kms. The associated atomic mass is 1500$\\pm$500 \\msun.  Molecular gas with velocities in the range --27.2 and --14.0 \\kms\\ surrounds the brightest parts of Gum\\,31. The sharp interface between the ionized and molecular material indicates that these gas components are interacting. We have estimated a molecular mass of  (1.1$\\pm$0.5)\\por 10$^5$ \\msun. Two possible origins can be suggested for the molecular shell: 1) it could have originated in the {\\it collect and collapse} process driven by the expansion of the \\hii\\ region; or 2) it could represent the remains of the natal cloud where the open cluster NGC\\,3324 formed. In both cases, a large amount of the neutral hydrogen in the \\hi\\ shell may have originated in the photodissociation of the molecular gas.  The volume density of the molecular cloud, higher than the electron density (33 \\cmtres), implies that the \\hii\\ region is expanding.  The emission in the far infrared correlates with  the \\hii\\ region and the \\hi\\ envelope, indicating that the observed emission is probably related  to the neutral and molecular envelopes.  The distribution of the emission in the MSX band A, which closely delineates the \\hii\\ region and correlates with the molecular emission, suggests the presence of a PDR at the interface between the ionized and molecular gas.  A number of MSX, IRAS, and 2MASS point sources with IR spectra compatible with protostellar objects appear projected onto the molecular envelope, implying that stellar formation is active in the higher density cores of the molecular envelope around Gum\\,31, where massive stars or star clusters are probably being formed.  The optical image of the nebula does not show clear evidence of a central cavity, as expected in a stellar wind bubble, suggesting that the massive stars in the cluster have weak stellar winds or existed for a short period of time to develop an interstellar bubble in a high density interstellar medium."
    },
    "0711/0711.0177_arXiv.txt": {
        "abstract": "We generalize the Maximum Likelihood-type method used to study cross correlations between a catalog of candidate astrophysical sources and Ultrahigh Energy Cosmic Rays (UHECRs), to allow for differing source luminosities.  The new method is applicable to any sparse data set such as UHE gamma rays or astrophysical neutrinos.  Performance of the original and generalized techniques is evaluated in simulations of various scenarios.   Applying the new technique to data, we find an excess correlation of about 9 events between HiRes UHECRs and known BLLacs, with a $6\\times 10^{-5}$ probability of such a correlation arising by chance. ",
        "introduction": "Correlation studies play a fundamental role in establishing or ruling out candidate sources for rare events such as the highest energy cosmic rays and UHE gamma rays. When the number of events is low, it is necessary to use sensitive methods to be able to identify the sources and quantitatively assess the possibility of an incorrect identification. The purpose of the present work is to evaluate and improve these methods.  If the hypothetical sources are commonplace rather than rare, so that the typical separation between candidate sources is not large compared to the directional uncertainty of individual events, then an entirely different approach must be used, requiring much larger datasets.  Fortunately, the GZK horizon allows the surface density of source candidates to be reduced by increasing the energy threshold; this played an essential role in the recent Auger discovery \\citep{AugerAGN:2007} of a correlation between UHECRs and galaxies in the Veron-Cetty Veron Catalog of Quasars and AGNs.   Furthermore, it is plausible that only unusual astrophysical objects can produce very rare events, in which case the candidate source catalog may naturally be sparse.  Hints of excess correlations were reported earlier between ultrahigh energy cosmic rays (UHECRs) and BLLacs by \\cite{TT:BLLacs01}, \\cite{TT:BLLacs04}, \\cite{HR:BLLacs} and between UHECRs and x-ray clusters by \\cite{pf05}. The analyses in  \\cite{TT:BLLacs01}, \\cite{TT:BLLacs04}, and \\cite{pf05} determine the number of correlated events within some given angular separation between UHECR and source, and calculate the ``chance probability'' of finding a correlation at the observed level by doing a large number of simulations with no true correlations.  The \\citep{AugerAGN:2007} analysis in addition ``scans\" to find the optimal UHECR energy threshold and maximum source redshift.  In order to incorporate the experimental resolution on an event-by-event basis, \\cite{HR:BLLacs} proposed a Maximum Likelihood-type procedure and applied it to studying correlations between UHECRs and BLLacs. This procedure (denoted the HiRes procedure, below) is motivated under the unphysical assumption that every candidate BLLac source has the same apparent luminosity. Even if BLLacs were standard candles with respect to UHECR emission, the BLLacs in the catalog have a large range of distances which would imply an even larger range of apparent luminosities, so one does not want to rely on such an assumption. Furthermore the validity of the method was not demonstrated via simulations.  In the present paper we introduce a ML prescription which avoids the assumption of equal apparent source luminosity and allows the potential sources to be ranked according to the probability that they have emitted the correlated UHECRs. We test and compare both methods with simulations in a variety of situations. We find that the HiRes method gives the correct total number of correlated events even when the sources do not have equal apparent luminosities, as long as the numbers of events are sufficiently low and the candidate sources are not too dense or clustered themselves. We find that our new method performs better in these more challenging cases. In a final section, we apply the new procedure to BLLacs and x-ray clusters. ",
        "conclusions": "\\label{summary}  We have introduced a generalization of the HiRes Maximum Likelihood method, which allows the most likely sources of individual events to be identified and ranked.  Using simulations we have tested the two Maximum Likelihood methods and find that they complement each other well: the HiRes method allows a fast way to estimate the number of true correlated events, while the new method gives the quality of correlation between individual sources and cosmic rays rather than just the total number of correlated events.  Furthermore, the new method is less sensitive to the validity of the estimated angular resolution, and has less systematic bias when candidate sources are clustered (as astrophysical sources such as BLLacs are). Of course, for any given data set, mere statistical fluctuations can result in conclusions that would not be borne out with a larger data set. We conclude that both methods should be used: if they disagree markedly on the total number of correlations the data may have some statistical anomaly and be difficult to interpret.  Although we have tested both Maximum Likelihood methods' ability to reproduce the true number of correlations under many different conditions, it is impossible to test for every conceivable statistical property of future data sets to which these methods could be applied. Thus we recommend, when applying these methods on new data samples, that the methods first be tested by simulations using the candidate source positions but randomly generated cosmic ray directions with some fraction of them Gaussianly aligned to some randomly chosen sources.  By repeating this procedure, the variance and possible systematic bias in the extracted number of correlations can be determined for that particular detector exposure, source and event statistics, resolution, etc.  We applied the ML methods to find correlations between x-ray clusters and AGASA events with $E>4\\times 10^{19}$ eV, as done in \\cite{pf05} but without imposing the GZK horizon.  We find $\\approx 10$ excess correlations, with a chance probability of $6\\times 10^{-3}$.  The same analysis for HiRes data with $E>10^{19}$ indicates $3-5$ excess correlations, which is expected by chance about 30\\% of the time.   Imposing the restriction that the correlated x-ray cluster not be farther than the GZK distance cutoff was found in \\cite{pf05} to reduce the number of chance correlations without significantly decreasing the number of true correlations.  Since in this work we did not impose the GZK restriction, we cannot compare the present analysis with that of \\cite{pf05}.  Finally, we applied the generalized Maximum Likelihood method to check previously claimed correlations between UHECRs and BLLacs.  We corroborate that there is a significant correlation between BLLacs of the Veron 10th catalog and HiRes cosmic rays with $E>10^{19}$ eV \\citep{HR:BLLacs}.  The generalized ML method yields a slightly greater number ($\\approx 9$ instead of $\\approx 8$) excess correlations and a lower chance probability ($6 \\times 10^{-5}$ instead of $2 \\times 10^{-4}$).  Simulations show that the  HiRes ML method systematically underestimates the number of correlations by 0.7 when the source candidates are clustered like the BLLacs, suggesting that the present result is the more reliable. The striking difference between the apparently significant correlation observed in the Northern sky (HiRes) and Auger's null result \\citep{Harari:2007up} is a mystery. It may be due to greater average magnetic deflection in the Southern sky, or greater incompleteness in the Southern BLLac catalog, or a statistical fluctuation or a combination of the above.   \\ack We thank S. Westerhoff, C. Finley, E. Pierpaoli and I. Zaw for information and discussions.  This research has been supported in part by NSF-PHY-0401232."
    },
    "0711/0711.0341_arXiv.txt": {
        "abstract": "% ASPID stands for the \"Archive of Spectral, Photometric, and Interferometric Data\".  The world largest collection of raw 3D spectroscopic observations of galactic and extragalactic sources is provided. ASPID-SR is a prototype of an archive of heterogeneous science ready data, fed by ASPID, where we try to exploit all the power of the IVOA Characterisation Data Model. Multi-level Characterisation metadata is provided for every dataset. The archive provides powerful metadata query mechanism with access to every data model element, vital for the efficient scientific usage of a complex informational system. We provide a set of access interfaces: SIAP/SSAP, HTTP-based characterisation metadata query, Web-service accepting ADQL/x. ",
        "introduction": "Interoperability between various data archives and data discovery, retrieval and analysis tools is a cornerstone for a success of the International Virtual Observatory. Data model, providing full and self-sufficient description of a dataset, is an essential for the interoperability. Characterisation DM of IVOA (McDowell et al. 2006) is one of the most general data models, positioning the dataset in the space of physical parameters.  ASPID stands for the \"Archive of Spectral, Photometric, and Interferometric Data\". It is operated by the Laboratory of Spectroscopy and Photometry of Extragalactic Objects and contains over 72000 raw observational datasets coming from selected instruments at the 6-m telescope of SAO RAS since 1992: direct images; 1D, long-slit, multi-object, and IFU spectroscopy; Fabry-Perot datacubes; spectropolarimetric observations. The world largest collection of 3D spectroscopic observations of galactic and extragalactic sources is provided. ",
        "conclusions": "ASPID-SR is one of the first implementations of the Characterisation Data Model, and we believe it is an important step toward development of VO-aware complex astronomy informational systems."
    },
    "0711/0711.3612_arXiv.txt": {
        "abstract": "{} {X-ray properties of the stellar population in the Carina\\, OB1 association are examined with special emphasis on early-type stars. Their spectral characteristics provide some clues to understanding the nature of X-ray formation mechanisms in the winds of single and binary early-type stars.} {A timing and spectral analysis of five observations with XMM-Newton is performed using various statistical tests and thermal spectral models.} {235 point sources have been detected within the field of view. Several of these sources are probably pre-main sequence stars with characteristic short-term variability. Seven sources are possible background AGNs. Spectral analysis of twenty three sources of type OB and WR\\,25 was performed. We derived spectral parameters of the sources and their fluxes in three energy bands. Estimating the interstellar absorption for every source and the distance to the nebula, we derived X-ray luminosities of these stars and compared them to their bolometric luminosities. We discuss possible reasons for the fact that, on average, the observed X-ray properties of binary and single early type stars are not very different, and give several possible explanations.} {} ",
        "introduction": "Although \\mbox{X-ray} emission by early-type stars is nowadays well established, its origin is still not fully understood. Observations with the {\\sl Einstein}, {\\sl ROSAT} and {\\sl ASCA} satellites indicated that hot luminous stars have rather soft thermal {\\mbox X-ray} spectra. This picture contrasts with the expected spectral properties if the {\\sl X-rays} were produced in a hot corona at the base of the stellar wind (e.g.\\ Waldron \\cite{Waldr84}). Indeed, in the latter case one would expect to observe substantial absorption of the softest emission. The lack of this absorption in the observed spectra triggered the elaboration of an alternative scenario. In fact, according to the phenomenological model proposed by Lucy \\& White (\\cite{Lucy80}) and further elaborated by Lucy (\\cite{Lucy82}), hydrodynamic shocks are generated throughout a radiation driven stellar wind as the consequence of the intrinsic instabilities of radiative driving. The velocity jumps in such shocks heat the post-shock plasma to temperatures of a few million degrees. These shocks would be distributed throughout the stellar wind with the result that even soft {\\mbox X-rays} could escape the wind without substantial absorption. Recent theoretical studies of such instabilities have been presented for instance by Feldmeier et al.\\ (\\cite{Feld97}), Owocki \\& Cohen (\\cite{OwockiC99}), and Dessart \\& Owocki (\\cite{Dess02}). High resolution spectra obtained with the new generation of {\\mbox X-ray} observatories have confirmed the thermal origin of the bulk of the {\\mbox X-ray} emission from early-type stars (e.g.\\ Kahn et al.\\ \\cite{Kahn01}), but in some cases, the properties of the observed line profiles led also to new challenges for the shock model (see e.g.\\ Rauw \\cite{ElEscorial} for a review).  Another well established, though yet unexplained, property of the {\\mbox X-ray} emission of hot stars is the linear scaling between the observed {\\mbox X-ray} luminosity ($L_{\\rm X}$) and the bolometric luminosity ($L_ {\\rm bol}$). Using {\\sl Einstein} data, Long \\& White (\\cite{Long80}), Pallavicini et al.\\ (\\cite{Pall81}) and Chlebowski et al.\\ (\\cite{Chle89}) found that $L_{\\rm X} \\simeq 10^{-7} \\times L_{\\rm bol}$ for OB stars. {\\sl ROSAT} observations broadly confirmed this relation. Bergh\\\"ofer et al.\\ (\\cite{Berg97}) and Kudritzki et al.\\ (\\cite{Kudrit96}) found that including a weak dependence on the characteristic wind density leads to a somewhat tighter relationship. However, considering the scaling of the X-ray emission with various wind and stellar properties in the framework of the shock model, Owocki \\& Cohen (\\cite{OwockiC99}) found that a delicate equilibrium between emission and absorption is required to reproduce the empirically derived $L_{\\rm X}$ -- $L_{\\rm bol}$ relation. With the broader energy range and improved sensitivity of the new generation of {\\mbox X-ray} satellites it is important to re-address the question of the empirical $L_{\\rm X}$ -- $L_{\\rm bol}$ scaling for OB stars.  When doing this it is crucial to pay attention to the multiplicity of the stars in the sample. Indeed, in OB binaries, the collision of the two stellar winds is expected to generate a strong {\\mbox X-ray} bright shock between the stars (see e.g.\\ Stevens et al.\\ \\cite{Stev92}) and this feature is believed to manifest itself as an ``excess'' in the $L_{\\rm X}/L_{\\rm bol}$ ratio (see e.g. Pollock\\ \\cite{Pollock95}). Moreover, since the relative velocities of colliding winds can be much higher than the shock velocity jumps in winds of single stars, the {\\mbox X-ray} spectra of colliding wind early type binaries can be quite hard. In fact the hardness of the observed spectra of some early-type stars, and, in particular, the presence of the Fe\\,XXV and Fe\\,XXVI lines at $~6.7$\\,keV indicative of very hot material has been suggested as an indication of possible binarity (e.g. Raassen et al., \\cite{Raas03}).  This latter suggestion leads to an interesting question. Clearly in a general case it is very difficult to ``hide'' such a hard secondary component of a colliding wind binary. Indeed, to produce strong hard \\mbox{X-ray} emission, the companion must be a massive hot star, and as such, should manifest itself in other wavelength domains (e.g.\\ optical spectra) and also via orbital spectral and photometric variability in the \\mbox{X-rays}. Traditionally, a lack of variations was attributed to a ``pole-on'' orientation of the binary orbit. However, lately, a growing number of examples of apparently single O and WR stars with hard {\\mbox X-ray} spectra (e.g. Skinner et al.\\ \\cite{Skinn02}, Raassen et al. \\cite{Raas03}, the present study) has become available. It seems unlikely that all these objects are binaries seen from their orbital poles. Thus the question is whether there exists another, yet unknown, mechanism which could produce hard {\\mbox X-rays} in single star winds? Exceptions to this rule are stars such as $\\theta^1$\\,Ori\\,C with strong enough magnetic fields to confine the stellar wind into the plane near the magnetic equator (Cassinelli et al., \\cite{Cassin02}, Brown et al., \\cite{Brown03}, Maheswaran \\cite{Mahes03}, Gagn\\'e et al.\\ \\cite{Gagne}, Li et al., \\cite{Li07}). In such cases, the head-on collision of the winds from the two hemispheres of a single star produces a rather hard X-ray emission. The fundamental problem here is that the only source of energy in the classical picture of a hot star wind which is able to produce high temperatures and hence hard thermal {\\mbox X-rays}, is the wind kinetic energy. But to release most of this energy through radiation one needs a ``wall'' able to stop the supersonic flow and heat the material in the formed shock. In the winds of single stars we do not expect to find such a wall except in the magnetised case. On the other hand, there is a number of known early-type binaries which show neither significant X-ray excess nor particularly hard X-ray spectra. What could be the mechanism which dumps the X-ray production in these systems?  The Carina OB\\,1 association with its impressive number of early-type stars -- including some of the youngest, hottest and most massive objects known in our Galaxy -- is an ideal place to investigate the {\\mbox X-ray} properties of these objects. The history of X-ray studies of the region goes back to the {\\em Einstein} era. Seward et al. (\\cite{Seward79}) presented {\\em Einstein} X-ray (0.2\\,--\\,4.0\\,keV) observations of the Carina open cluster Tr\\,16 and its environment, including six O-type stars and one WR star, WR\\,25. Subsequent {\\em Einstein} observations by Seward \\& Chlebowski (\\cite{Seward82}) of the same region showed X-rays from 15 O-type and WR stars. More studies were performed with the use of the {\\em ROSAT} and {\\em ASCA} satellites. Some of them were focused on individual objects like {\\em the LBV} $\\eta$\\,Car (Corcoran et al. \\cite{Corcoran00} and references therein) or WR\\,25 (e.g. Pollock \\& Corcoran \\cite{Pollock06}, Skinner et al. \\cite{Skinner95}). Others were devoted to investigating general X-ray properties of the early-type stars population (Corcoran \\cite{Corc99} and references therein).  Recently, four studies of the region based on X-ray observations with {\\em Chandra} and {\\sl XMM-Newton} satellites were published. Evans et al. (\\cite{Evans03}) (hereafter EV03) and Evans et al. (\\cite{Evans04}) (hereafter {\\em EV04}) using {\\em Chandra} data, detected 154 point sources, 23 of which are of O, B types. They present luminosities and hardness ratios for the detected sources and confirm the $L_{\\rm X}$ -- $L_{\\rm bol}$ relation for early-type stars. They also discuss the low-resolution spectra of the 14 brightest sources. Sanchawala et al. (\\cite{Sanchawala07}), using the same {\\em Chandra} data as EV03 plus another archival {\\em Chandra} data set, detected 454 sources in the region, among which 38 are known OB stars. The latter paper is focused on cross-identification of the X-ray sources with optical and near-infrared imaging observations specifically obtained for this study. This allowed the authors to conclude that about 300 sources are likely late-type pre-main-sequence stars. The X-ray luminosities of the detected OB stars were estimated based on their count rate and using a single-temperature Raymond-Smith model of thermal plasma with $\\log\\,T=6.65$ (0.384\\,keV). Albacete-Colombo et al. (\\cite{AC03}) (hereafter AC03) detected 80 discrete sources in the region using two observations by {\\sl XMM-Newton} (revolution numbers 115 and 116). They discuss X-ray properties of the sources, including the $L_{\\rm X}$ -- $L_{\\rm bol}$ relation for early-type stars. In the present paper, we used three additional {\\sl XMM-Newton} observations, which allow us to increase the signal to noise ratio and to detect more than 200 discrete sources. We shall compare our results with those of AC03 throughout the current paper. We do not present a detailed spectral analysis of individual objects in this paper as many such objects deserve dedicated papers. A detailed study of WR\\,25 based on the same {\\sl XMM-Newton} data as in the current study was published by Raassen et al. (\\cite{Raas03}), whilst $\\eta$\\,Car has been the subject of many studies including those utilizing extensive data sets obtained with various X-ray observatories (see, e.g., Leutenegger et al. \\cite{Leut03}; Weis et al. \\cite{Weis04}; Viotti et al. \\cite{Viotti04}; Corcoran \\cite{Corc05}; Hamaguchi et al. \\cite{Hamaguchi07}). Analysis of other interesting objects (e.g. the early type binary HD\\,93205) will be the subject of forthcoming papers.  The goals of this paper are: \\vspace*{-2mm} \\begin{itemize} \\item To describe general X-ray properties of the detected sources. \\item To look for their variability, both short and long term. \\item To investigate the $L_{\\rm X}/L_{\\rm bol}$ relation for early type stars. \\item To compare X-ray properties of single versus binary early-type stars and to try to get some clues to their differences and similarities. \\end{itemize}    The structure of the paper follows these outlined priorities. In the second section, the observational details are given. Source detection, their identification with the optical and IR catalogues, X-ray count rates, identification of possible extra-galactic sources are given in Section 3. Section 4 provides general information about open clusters in the Carina nebula, along with the discussion of their distance from the Earth. The results of the variability study are presented in Section 5. The X-ray properties of early-type stars detected in the region are presented in Section 6. We discuss our results in section 7 and a short summary is given in Section 8. ",
        "conclusions": "High quality and large amounts of {\\sl XMM-Newton} data obtained for the Carina Nebula, allowed us to detect 235 discrete sources in the field of view. Several of these sources are probably pre-main sequence stars with a characteristic short-term variability, while seven sources are recognized as possible background AGNs.  We concentrated in our study on the properties of early-type stars in the region. Spectral analyses of twenty three sources of type OB and WR\\,25 were performed. We derived the spectral parameters of the sources and their fluxes in three energy bands. Estimating a value of interstellar absorption towards every source and assuming a distance of 2.5\\,kpc to the Nebula, we derived the X-ray luminosities of these stars. Comparison of these luminosities with bolometric luminosities has confirmed the known scaling relation $L_{\\rm X}\\,\\sim\\,10^{-7}\\,\\times\\,L_{\\rm bol}$ for early type stars and allowed us to make a few more detailed statements about the characteristics of this relation. We pointed out that, {\\em on average}, the observed X-ray properties of binary and single early type stars are similar, and give several possible explanations for this fact.  Further work will be aimed at the detailed analysis of the data of particular sources."
    },
    "0711/0711.4204_arXiv.txt": {
        "abstract": "{We present an investigation of the clustering of the faint ($i'_{AB}<24.5$) field galaxy population in the redshift range $0.2<z<1.2$. Using 100,000 precise photometric redshifts extracted from galaxies in the four ultra-deep fields of the Canada-France Legacy Survey, we construct a set of volume-limited galaxy samples. We use these catalogues to study in detail the dependence of the amplitude $A_w$ and slope $\\delta$ of the galaxy correlation function $w$ on absolute $M_B$ rest-frame luminosity, redshift, and best-fitting spectral type (or, equivalently, rest-frame colour). Our derived comoving correlation lengths for magnitude-limited samples are in excellent agreement with measurements made in spectroscopic surveys. Our principal conclusions are as follows: 1. The comoving correlation length for all galaxies with $-19<M_B-5\\log h<-22$ declines steadily from $z\\sim0.3$ to $z\\sim1$.  2. At all redshifts and luminosity ranges, galaxies with redder rest-frame colours have clustering amplitudes between two and three times higher than bluer ones. 3. For both the red and blue galaxy populations, the clustering amplitude is invariant with redshift for bright galaxies ($-19<M_B-5\\log h<-22$). 4. At $z\\sim0.5$ for less luminous galaxies with $M_B-5\\log h\\sim-19$ we find higher clustering amplitudes of $\\sim 6h^{-1}$~Mpc. 5. The relative bias between redder and bluer rest-frame populations increases gradually towards fainter magnitudes. Among the principal implications of these results is that although the full galaxy population traces the underlying dark matter distribution quite well (and is therefore quite weakly biased), redder, older galaxies have clustering lengths which are almost invariant with redshift, and by $z\\sim1$ are quite strongly biased. ",
        "introduction": "In the cold dark matter model structures grow hierarchically under the influence of gravity. Galaxies form inside ``haloes'' of dark matter \\citep{1978MNRAS.183..341W}. Because these haloes can only form at the densest regions of the dark matter distribution, the distribution of galaxies and dark matter is not the same; the more strongly clustered galaxies are said to be ``biased'' \\citep{1984ApJ...284L...9K,1986ApJ...304...15B} with respect to the dark matter distribution. This relationship between dark and luminous matter provides important information concerning the galaxy formation process and tracing the evolution of bias as a function of scale and mass of the hosting dark matter halo is one of the key objectives of observational cosmology. On large scales, ($>10 h^{-1}$~Mpc) structure growth is largely driven by gravitation (where we measure the correlations between separate haloes of dark matter); however on smaller scales ($<1~h^{-1}$) non-linear effects generally associated with galaxy formation dominate the structure formation process. In this paper we must bear in mind that although we measure a clustering signal to around $0.1\\deg$ at $z\\sim1$ this corresponds to around $\\sim 3h^{-1}$~Mpc and $\\sim 2h^{-1}$~Mpc at $z\\sim0.5$, which means that our observations are mostly in non-linear to strongly non-linear regimes where environmental effects play an important role in the evolution of structure.  On linear scales, as theory and simulations have shown (for example \\cite{1998ApJ...499...20J} or \\cite{2004ApJ...601....1W}), the clustering amplitude of dark matter decreases steadily with redshift. If galaxies perfectly traced the dark matter component, then their clustering amplitudes would decrease at each redshift slice, in step with the underlying dark matter. However, as the galaxy distribution is biased, stellar evolution intervenes to complicate this picture; in effect, the actual measured clustering amplitudes are a complicated interplay between the underlying dark matter component and how well the luminous matter traces this galaxy distribution, or how efficiently galaxies form. Understanding fully the evolution of galaxy clustering requires, therefore, some insights into the galaxy formation process. Thanks to large spectroscopic redshift surveys we now have a much more complete picture of the evolution of the galaxy luminosity function with redshift \\citep{2005A&A...439..863I} and how the fraction of galaxy types evolves with redshift \\citep{2006A&A...455..879Z}. For example, \\cite{2005A&A...439..863I}, using first-epoch data from the VVDS redshift survey have shown that the luminosity function brightens considerably between $z=0.3$ and $z=1$, with $M*$ increasing by one or two magnitudes at $z\\sim1$. We must take this into account when comparing clustering amplitudes measured at the same absolute luminosities in different redshift ranges.  In the local Universe, million-galaxy redshift surveys have greatly expanded our knowledge of galaxy clustering at low redshift. We now have a broad idea how the distribution of galaxies depends on their intrinsic luminosity and spectral type \\citep{2001MNRAS.328...64N,2002MNRAS.332..827N,2005ApJ...630....1Z}. In general, these works have shown to a high precision that at the current epoch more luminous galaxies are more clustered than faint ones, and that similarly redder objects have higher clustering amplitudes than bluer ones. Other works have shown that slope of the galaxy correlation depends on spectral type \\citep{2003MNRAS.344..847M}. These studies have indicated that, in general, more luminous, redder, objects are more strongly clustered than bluer, fainter galaxies. Some studies have also related physical galaxy properties, such a total mass in stars, with the clustering properties \\citep{2006MNRAS.368...21L}. But how do these relationships change with look back time?  At intermediate redshifts ($z\\sim1$), however, our knowledge is still incomplete. Multi-object spectrographs mounted on ten-metre class telescopes have made it possible to construct samples of a few thousand galaxies.  The first studies investigating galaxy clustering as a function of the object's rest frame luminosity and colour for large galaxy samples at $z\\sim1$ have now appeared \\citep{2006A&A...452..387M,2006A&A...451..409P,2006ApJ...644..671C}. Unfortunately, these surveys typically contain $\\sim10^3$ galaxies, which are enough to select objects either by type and absolute luminosity, but not, for instance, to apply both cuts simultaneously.  These works confirm some of the broad trends seen at lower redshift and with magnitude-limited samples \\citep{2005A&A...439..877L} but are still not quite large enough to investigate in detail how galaxy clustering depends simultaneously on more than one galaxy property. For example, one may investigate the dependence of clustering amplitude within a volume-limited sample \\citep{2006A&A...452..387M}, but one may not, as yet, investigate simultaneously samples selected by type and absolute luminosity. Unfortunately, even with efficient wide-field multi-object spectrographs, gathering redshift samples of thousands of galaxies at redshift of one or so requires a significant investment of telescope time.  Photometric redshifts offer an exit from this impasse, and represent a middle ground between simple studies using imaging data with magnitude or colour-selected samples and spectroscopic surveys. Several attempts have been made in the past to carry out galaxy clustering studies with photometric redshifts, mostly using the Hubble deep field data sets \\citep{2002MNRAS.329..355A,2001ApJ...548..127T,1999MNRAS.306..988M}. However, such works either suffered from sampling and cosmic variance issues or poorly-controlled systematic errors. The advent of wide-field mosaic cameras like MegaCam \\citep{2000SPIE.4008..657B} has made it feasible for the first time to construct samples of tens to hundreds of thousands of galaxies from $z\\sim0.2$ all the way to $z\\sim1$ and beyond. Two key advances have made this possible; firstly, rigorous quality control of photometric data, and secondly, the availability of much larger, reliable training samples reaching to faint ($i'\\sim24$) magnitudes  In this paper we will describe measurements of galaxy clustering derived from a large sample of galaxies with accurate photometric redshifts in Canada-France legacy survey (CFHTLS) deep fields. These fields have been observed repeatedly since the start of survey operations in June 2003 as part of the on-going SNLS project \\citep{2006A&A...447...31A} and consequently each filter has very long integration times (for $r$ and $i$ bands the total integration time in certain fields is over 100 hours).  A full description of our photometric redshift catalogue can be found in \\cite{2006A&A...457..841I}.  Containing almost 100,000 galaxies to $i'<24.5$ we are able to divide our sample by redshift, absolute luminosity and spectral type.  These photometric redshifts have been calibrated using 8,000 spectra from the VIRMOS-VLT deep survey (VVDS; \\citet{2005A&A...439..845L}). In addition, our sample has sufficient volume to provide reliable measurements of galaxy clustering amplitudes at redshifts as low as $z\\sim0.2$; and we are thus able to follow the evolution of galaxy correlation lengths over a wide redshift interval.  In the lower redshift bins, the extremely deep CFHTLS photometry means it is possible to measure clustering properties of a complete sample of objects as faint as $M_B-5 \\log(h)\\sim-18$ (at $z\\sim0.2$ we have large numbers of very faint objects with $M_B-5\\log(h)\\sim-15$ although we do not consider them here).  Moreover, by using all four independent deep fields of the Canada-France legacy survey we are able to robustly estimate the amplitude of cosmic variance for each of our samples.  Our objective in this paper is to determine, first of all, how the observed properties of galaxies determines their clustering.  We are able to carry out such an investigation of galaxy clustering strength for samples of galaxies selected independently in absolute luminosity, rest-frame colour and redshift.  Our paper is organised as follows: in Section~\\ref{sec:prep-catal-comp} we describe how our catalogues were prepared and how we computed our photometric redshifts; in Section~\\ref{sec:meas-galaxy-clust} we describe how we measure galaxy clustering in our data; our results are presented in Section~\\ref{sec:results}. Finally, our discussions and conclusions are presented in Section~\\ref{sec:disc-results-comp}.  In this work we divide the CFHTLS galaxy samples in three principal ways: first of all, we consider simple magnitude limited samples, divided by bins of redshift (described in Section~\\ref{sec:magn-limit-sampl}); next, at two fixed redshift ranges, we consider galaxy samples selected by absolute luminosity and type (Section \\ref{sec:lumin-deped-clust}); and lastly, at a range of redshift bins and for the same slice in absolute luminosity, we consider galaxies selected by type (presented in Section~\\ref{sec:volume-type-depend}).  Throughout the paper, we use a flat lambda cosmology ($\\Omega_m~=~0.3$, $\\Omega_\\Lambda~=~0.7$) and we define $h~=~H_{\\rm0}/100$~km~s$^{-1}$~Mpc$^{-1}$. Magnitudes are given in the AB system unless otherwise noted. ",
        "conclusions": "\\label{sec:conclusions}  We have presented an investigation of the clustering of the faint ($i'<24.5$) field galaxy population in the redshift range $0.2<z<1.2$. Using 100,000 precise photometric redshifts extracted in the four ultra-deep fields of the Canada-France Legacy Survey, we construct a series of volume-limited galaxy samples and use these to study in detail the dependence of the amplitude $A\\_w$ and slope $\\delta$ of the galaxy correlation function $w$ on absolute $M_B$ rest-frame luminosity, redshift, and best-fitting spectral type (or, equivalently, rest-frame colour). Our principal conclusions are as follows:  1. The comoving correlation length for all galaxies with $-19<M_B-5\\log h<-22$ declines steadily from $z\\sim0.3$ to $z\\sim1$.  2. At all redshifts and luminosity ranges, galaxies with redder rest-frame colours have clustering amplitudes between two and three times higher than bluer ones.  3. For both the red and blue galaxy populations, the clustering amplitude is invariant with redshift for bright galaxies with $-22<M_B-5\\log h<-19$.  4. At $z\\sim0.5$, less luminous galaxies with $M_B-5\\log h\\sim-19$, we find higher clustering amplitudes of $\\sim 6h^{-1}$~Mpc.  5. The relative bias between populations of redder and bluer rest-frame populations increases gradually towards fainter magnitudes.  The main implications of these results is that although the full bright galaxy population traces the underlying dark matter distribution quite well (and is therefore quite weakly biased), redder, older galaxies have clustering lengths which are almost invariant with redshift, and must therefore by $z\\sim1$ be quite strongly biased. In addition, at $z\\sim0.5$ there is some evidence that fainter red objects are more strongly clustered than $\\sim L^*$ galaxies at these redshifts. This is consistent with a picture in which fainter red objects exist primarily as satellite galaxies in galaxy clusters.  It is tempting to interpret our results in terms of studies which show that the number density of massive, luminous galaxies evolves little from $z\\sim1$ to the present day \\citep{2006A&A...453L..29C,2006A&A...455..879Z,CSH2}. In our survey, the clustering amplitudes of bright ellipticals are already 'fixed in' at $z\\sim1$. Most of the changes in the clustering amplitude occur in the fainter galaxy population. However, a full understanding of the processes at work here will require mass-selected samples covering a larger interval in redshift. Such samples will become possible in the near future with the addition of near-infrared data to the CFHTLS survey fields."
    },
    "0711/0711.1617_arXiv.txt": {
        "abstract": "We present a new sample of $4634$ eclipsing binary stars in the Large Magellanic Cloud (LMC), expanding on a previous sample of 611 objects and a new sample of $1509$ eclipsing binary stars in the Small Magellanic Cloud (SMC), that were identified in the light curve database of the MACHO project. We perform a cross correlation with the OGLE-II LMC sample, finding 1236 matches. A cross correlation with the OGLE-II SMC sample finds 698 matches. We then compare the LMC subsamples corresponding to center and the periphery of the LMC and find only minor differences between the two populations. These samples are sufficiently large and complete that statistical studies of the binary star populations are possible. ",
        "introduction": "Eclipsing binary stars (EBs) are important for astrophysical research in many ways. They may be used to obtain accurate estimates of star masses and radii \\cite[and references therein]{andersen91}. Precise determination of stellar parameters can in turn be used to put theories of stellar structure and evolution to a stringent test by comparing measured parameters with theoretical predictions \\citep[and references therein]{lvg02,lvg03}. \\par EBs may also be used for distance determination and this use goes back several decades; its history is reviewed by \\citet{kruszewski99}. Since \\citet{stebbins11} used an estimate of the parallax to $\\beta$ Aurig\\ae~to infer the surface brightness of both its components, it has been known that a good photometric light curve plus a double line spectroscopic orbit admits a simple geometric relationship between the surface brightnesses of the stars and the distance to the EB; \\citet{stebbins11} however had no way at the time to make the reverse ``surface brightness to distance'' inference and his paper does no mention this possibility. After \\citet{stebbins11}, other early papers \\citep{gaposchkin33,woolley34,pilowski36,kopal39,gaposchkin38,gaposchkin40} used parallaxes obtained independently to estimate surface brightnesses, but, as remarked by \\citet{kruszewski99}, these pioneers surely knew of the potential of this technique to estimate distances. Modern analyses of EBs have usually focussed on this technique (e.g. \\cite{andersen91}). The method affords high precision due to its purely geometrical nature and has been applied by a number of authors to determine the distance to the Large Magellanic Cloud (LMC) using HV2274 \\citep{ud98a,guinan98,nelson99, groen01}, HV982 \\citep{fitzp02}, EROS 1044 \\citep{ribas02} and HV 5936 \\citep{fitzp03}; an attempt to use EBs to determine the distance to M31 is currently under way \\citep{ribas03} and the DIRECT project is attempting to measure the distance to M31 and M33 via EBs and Cepheids \\citep{kaluzny98,bonanos03,bonanos05}; other recent examples include \\citet{michalska05} who present a sample of detached binaries in the LMC for distance determination and \\citet{ribas05} who present the first determination of the distance and properties of an EB in M31; \\citet{north05} presents a sample of EBs with total eclipses in the LMC suitable for spectroscopic studies. In general it is important that distances be determined using a large sample of EBs to minimize the impact of systematic errors. A recent collection of references on extragalactic binaries can be found in \\citet{ribas04}. \\par Large-scale surveys to detect gravitational microlensing events have identified and collected light curves for large numbers of variable stars in the bulge of the Milky Way and in the Magellanic Clouds. Eclipsing binary stars comprise a significant fraction of these collections. The MACHO collaboration\\footnote{\\url{http://www.macho.mcmaster.ca/}} has presented a sample of 611 EBs in the LMC with preliminary analyses of their orbits \\citep{lacy97}. A catalogue of $3031$ EBs in the LMC found in the MACHO database has been just published by \\citet{derekas07}; this catalogue was compiled by analyzing a list of $6835$ stars classified as possible EBs in the MACHO database; a cross correlation between these $6835$ stars and our sample finds just $1987$ matches, thus at least about $2700$ EBs in our catalogue are new identifications. The $6835$ classified as possible EBs were found in regions of parameter space such as color, magnitude, and period, where one does not expect to find pulsating variables and therefore the detected variability of these stars was \\emph{tentatively} ascribed to eclipses. Regions where pulsating variables could exist were not considered while making this preliminary classification and EBs there were therefore not included in the list. In our search we did not rely primarily on cuts in parameter space and and we did not exclude a priori regions of this space where pulsating variables are present; therefore we were able to classify many EBs in these regions, that were not included in the preliminary classification. The OGLE collaboration\\footnote{\\url{http://sirius.astrouw.edu.pl/\\~{}ogle/}} has introduced a sample of $2580$ EBs in the LMC \\citep{wyr03} and of $1351$ EBs in the Small Magellanic Cloud \\citep[SMC:][]{wyr04}. Both samples were selected from their catalogue of variable stars in the Magellanic Clouds \\citep{zebrun01a} compiled from observations taken during the second part of the project \\citep[OGLE II:][]{udalski97} and reduced via Difference Image Analysis \\citep[DIA:][]{zebrun01b}. An earlier sample of $79$ EBs in the bar of the LMC was presented by the EROS collaboration\\footnote{\\url{http://eros.in2p3.fr/}}\\citep{gri95}. Other large variable star data sets are being produced by surveys not specifically designed to detect gravitational microlensing, such as the All Sky Automated Survey \\citep[ASAS:][]{pojmanski97} \\footnote{\\url{http://www.astrouw.edu.pl/\\~{}gp/asas/asas.html}}. \\par The availability of large samples of EBs (and the even larger ones that can be found by future surveys such as Pan-STARRS\\footnote{\\url{http://pan-starrs.ifa.hawaii.edu/public/}} and LSST\\footnote{\\url{http://www.lsst.org/lsst\\_home.shtml/}}) can have an important impact on stellar astrophysics. This impact can arise in two qualitatively different approaches. First, a large catalogue allows the discerning researcher to select carefully a few EBs for detailed follow-up study; the distance estimation described above is an example of this. Second, statistical analyses of an entire population become possible when a large collection is assembled; such analyses of EBs have not previously been possible. To fulfill this promise there are challenges to overcome, including finding EBs in large data sets and automating their analysis. With regard to the first task, the \\textit{discovery} problem is complicated by the fact that EBs do not have clear relationships between their parameters (period, luminosity, colors) as do the major classes of pulsating variables. This makes it difficult to find them via simple and well understood cuts in parameter space. The first step toward automated discovery is thus to have a large sample of data on which to experiment with search techniques. This non-trivial exercise in mining large data sets can be useful for future surveys that are not necessarily aimed at binary star research. An example is given by \\citet{wyr03} and \\citet{wyr04} who employ an artificial neural network to identify EBs in the OGLE-II LMC and SMC samples, but more needs to be done. With regard to \\textit{analysis} of EBs, the traditional approach has been to carefully analyze individual systems with the help of dedicated computer codes such as the Wilson-Devinney code \\citep[WD:][]{wilson71,wilson79}. This becomes impracticable when many thousands of stars are involved and an automated approach is required. The light curves in a previous sample of 1459 EBs in the SMC found by OGLE-II \\citep{ud98b} were systematically solved by \\citet{wilson01,wilson02} using an automated version of the WD code; the ASAS collaboration has developed an automated classification algorithm for variable stars based on Fourier decomposition \\citep{pojmanski02}; \\citet{devor05} found and analyzed 10000 Bulge EBs from OGLE-II using DEBiL\\footnote{\\url{http://www.cfa.harvard.edu/\\~{}jdevor/DEBiL.html}}, an EB analysis code that allows automated solutions of large EB data sets and works best for detached EBs; a genetic algorithm based approach to finding good initial parameters for WD is described in \\citet{metcalfe99}. \\par This paper is the first of a series of papers aimed at describing the EB samples in the MACHO database and is organized as follows: Section \\ref{sec:sample} introduces the LMC and SMC samples; Section \\ref{sec:diagrams} describes the Color Magnitude Diagram (CMD) and the Color Period Diagram, pointing out significant features in them, Section \\ref{sec:comparison} compares the LMC and SMC samples, Section \\ref{sec:crossogle} describes the results of the cross correlation with the OGLE LMC and SMC samples, and Section \\ref{sec:online} reports where and in what form the data presented in the paper can be accessed on line. ",
        "conclusions": ""
    },
    "0711/0711.4518_arXiv.txt": {
        "abstract": "In the fluorescence detection of ultra high energy ($\\gtrsim 10^{18}$ eV) cosmic rays, the number of emitted fluorescence photons is assumed to be proportional to the energy deposited in air by shower particles. We have performed measurements of the fluorescence yield in atmospheric gases excited by electrons over energies ranging from keV to hundreds of MeV in several accelerators. We found that within the measured energy ranges the proportionality holds at the level of few \\%. ",
        "introduction": "\\label{} The detection of ultra high energy  ($\\gtrsim 10^{18}$eV) cosmic rays using nitrogen fluorescence emission induced by extensive air showers (EAS) is a well established technique \\cite{flyseye}. Atmospheric nitrogen molecules, excited by EAS charged particles (mainly $e^{\\pm}$), emit fluorescence light in the $\\approx$ 300-400 nm range. The fluorescence detection of UHECR is based on the assumption that the number of fluorescence photons of wavelength $\\lambda$ emitted at a given stage of a cosmic ray shower development, {\\it i.e.} at a given altitude $h$ in the atmosphere, is proportional to the energy $E_{dep}^{shower}(h)$ deposited by the shower particles in the air volume \\cite{airflyAP}: \\begin{equation} N_{\\lambda}^{shower}(h) = E_{dep}^{shower}(h) Y_{air}(\\lambda,p_0,T_0) F(\\lambda,p,T), \\label{eq:nlambdash} \\end{equation} where $ Y_{air}(\\lambda,p_0,T_0)$ is the absolute yield (in number of photons per MeV) at a reference pressure $p_0$ and temperature $T_0$, $F(\\lambda,p,T)$ accounts for quenching effects, and $p$ and $T$ are the air pressure and temperature at the altitude $h$. Since a typical cosmic ray shower extends up to about 15 km altitude, the fluorescence yield must be known over a wide range of air pressure and temperature. Measurements of the fluorescence yield dependence on atmospheric parameters ($F(\\lambda,p,T)$) by AIRFLY are presented in a separate contribution \\cite{airflyAP,paolo,paolo2}.  Simple considerations suggest that fluorescence emission should indeed be proportional to the energy deposited. In fact, the cross sections for electron excitation of the 2P and 1N nitrogen systems, which are the most relevant in the 300-400 nm range, are peaked at very low energies (tens of eV) and  decrease rapidly with energy of the electron ($\\approx E^{-2}$ for the 2P and $\\approx \\log E/E$ for the 1N). Therefore the fluorescence light induced by a high energy electron ($>$ keV) will be mainly produced by the secondary electrons of eV energies. Since the total number of secondary electrons produced by the passage of the primary electron in the air volume is roughly proportional to the energy deposited, the fluorescence light is also expected to be proportional to the energy deposited. The constant of proportionality should not depend on the primary electron energy.  In fact, $E_{dep}^{shower}(h)$ in Eq. (\\ref{eq:nlambdash}) is the sum of the energies deposited by EAS particles with a spectrum spanning from keV to GeV. It is thus important to verify the proportionality of the fluorescence emission to the energy deposit over a wide range of electron energies. Available measurements are limited to a few energies \\cite{kakimoto} or used indirect methods \\cite{flash2}. The AIRFLY (AIR FLuorescence Yield) collaboration has performed measurements of the energy dependence of the fluorescence yield at several accelerators covering a range of electron kinetic energy from keV to hundreds of MeV.  Results of these studies are reported in the following. ",
        "conclusions": "We presented measurements of the energy dependence of the fluorescence yield performed at several accelerators. We tested the proportionality of the fluorescence light to the energy deposited at a level of few \\% over the energy ranges 0.5 to 15 MeV, 50 to 420 MeV and 6 to 30 keV. Notice that we performed only relative measurements within each range, and absolute measurements of the fluorescence yield are in principle needed to verify that the proportionality constant is the same in the three measured energy ranges \\cite{paolo2}.  Work in this direction is ongoing. On the other hand, given that the basic mechanism for the fluorescence yield is excitation by very low energy secondary electrons, it is hard to find any physical mechanism which could change the proportionality constant between 15 and 50 MeV. The AIRFLY data presented here would then indicate that the fluorescence yield is indeed proportional to the energy deposit for electron energies at least between 0.5 and 420 MeV. Most of the EAS energy is deposited by shower particles within this energy range."
    },
    "0711/0711.0427_arXiv.txt": {
        "abstract": "{ We propose a new scenario for compound chondrule formation named as ``fragment-collision model,\" in the framework of the shock-wave heating model. A molten cm-sized dust particle (parent) is disrupted in the high-velocity gas flow. The extracted fragments (ejectors) are scattered behind the parent and the mutual collisions between them will occur. We modeled the disruption event by analytic considerations in order to estimate the probability of the mutual collisions assuming that all ejectors have the same radius. In the typical case, the molten thin ($\\sim 1 \\, {\\rm mm}$) layer of the parent surface will be stripped by the gas flow. The stripped layer is divided into about 200 molten ejectors (assuming that the radius of ejectors is $300 \\, {\\rm \\mu m}$) and then they are blown away by the gas flow in a short period of time ($\\sim 0.01 \\, {\\rm s}$). The stripped layer is leaving from the parent with the velocity of $\\sim 4 \\, {\\rm cm \\, s^{-1}}$ depending on the viscosity, and we assumed that the extracted ejectors have a random velocity $\\Delta v$ of the same order of magnitude. Using above values, we can estimate the number density of ejectors behind the parent as $n_{\\rm e} \\sim 800 \\, {\\rm cm^{-3}}$. These ejectors occupy $\\sim 9\\%$ of the space behind the parent in volume. Considering that the collision rate (number of collisions per unit time experienced by an ejector) is given by $R_{\\rm coll} = \\sigma_{\\rm coll} n_{\\rm e} \\Delta v$, where $\\sigma_{\\rm coll}$ is the cross-section of collision (e.g., Gooding \\& Keil 1981, Meteoritics 16, 17), we obtain $R_{\\rm coll} \\sim 36 \\, {\\rm collisions/s}$ by substituting above values. Since most collisions occur within the short duration ($\\sim 0.01 \\, {\\rm s}$) before the ejectors are blown away, we obtain the collision probability of $P_{\\rm coll} \\sim 0.36$, which is the probability of collisions experienced by an ejector in one disruption event. The estimated collision probability is about one order of magnitude larger than the observed fraction of compound chondrules. In addition, the model predictions are qualitatively consistent with other observational data (oxygen isotopic composition, textural types, and size ratios of constituents). Based on these results, we concluded that this new model can be one of the strongest candidates for the compound chondrule formation.  It should be noted that all collisions do not necessarily lead to the compound chondrule formation. The formation efficiency and the future works which should be investigated in the forthcoming paper are also discussed. }\\\\   \\noindent {\\it Keywords:} meteorites, Solar System origin, Solar Nebula     \\clearpage ",
        "introduction": "\\label{sec:introduction} Chondrules are millimeter-sized, once-molten, spherical-shaped grains mainly composed of silicate material. They are abundant in chondritic meteorites, which are the majority of meteorites falling onto the Earth. They are considered to have formed from chondrule precursor dust particles about $4.56\\times10^9 \\, {\\rm yr}$ ago in the solar nebula (Amelin et al. 2002); they were heated and melted through flash heating events in the solar nebula and cooled again to solidify in a short period of time (e.g., Jones et al. 2000, and references therein). Typical chondrules are single spherical objects, while compound chondrules are composed of two or more chondrules fused together. They are rare in all chondrules ($\\sim 4 \\%$), but occur in many classes of chondrites, so they offer crucial information regarding the physical and chemical state of solid materials during chondrule formation (Gooding \\& Keil 1981, Wasson et al. 1995, Sekiya \\& Nakamura 1996, Ciesla et al. 2004, Akaki \\& Nakamura 2005).  Wasson et al. (1995) defined compound chondrules based on textures and assigned each constituent chondrule as primary or secondary. The primary chondrule was rigid enough to retain its original shape at the time of compound chondrule formation. On the contrary, the secondary chondrule had a low viscosity enough to allow it to conform the shape of the primary. Akaki \\& Nakamura (2005) considered four types of compound chondrules as follows: (1) the enveloping type, wherein a secondary chondrule encloses a primary one, (2) the adhering type, wherein a small secondary chondrule forms a hemispherical bump on the surface of a larger primary chondrule, (3) the consorting type, wherein the conjoined chondrules are of similar size, and (4) the blurred type, wherein compound chondrules have a blurred boundary between two chondrules, which makes it difficult to define the primary or the secondary. They classified 30 compound chondrule sets found in two CV3 chondrites into 7 blurred, 13 adhering, 4 consorting, and 6 enveloping.  At least four models have been proposed for compound chondrule formation. (1) Random collision model; totally or partially molten particles randomly collided in the solar nebula (Gooding \\& Keil 1981, Sekiya and Nakamura 1996). (2) Collision between individual chondrules in shock-wave heating (Ciesla 2006). (3) Eruption model; immediately after the formation of the solidified shell in the primary chondrule, the inner residual melt was vacuumed out through a crack of the shell, and formed the secondary chondrule on the surface of the primary (Sanders \\& Hill 1994). (4) Relict grain model; fine-grained dust particles accreted on the surface of already-formed primary were melted in second heating event (Wasson 1993, Wasson et al. 1995). The relict grain model is good for enveloping type if considering that fine-grains accreted uniformly on the primary surface, but it seems to be difficult to explain other types. In the eruption model, the secondary was formed from the inner residual melt after the formation of the solidified shell on the primary surface, however, the chemical compositions of some sets of compounds are inconsistent with this scenario (Akaki \\& Nakamura 2005). The random collision model is considered not to account for the observed fraction of compound chondrules because of the low density of matter in the nebula (Gooding \\& Keil 1981, Wasson et al. 1995, Sekiya and Nakamura 1996). Collision between individual chondrules in shock-wave heating can account for the observed fraction of compound chondrules if the enhancement of the dust particle in the pre-shock region is $\\sim 400$ times that expected under canonical conditions (Ciesla 2006). However, such a highly dust-rich region is unfavorable to explain the scarcity of isotopic fractionation of sulfur if the shock wave has a large spatial extent in solar nebula (Tachibana \\& Huss 2005, Miura \\& Nakamoto 2006).  In this paper, we propose a new scenario for compound chondrule formation. The shock-wave heating model is one of the most plausible models for chondrule formation (e.g., Connolly \\& Love 1998). In this model, the dust particles are exposed to a high-velocity gas flow and heated by the gas frictional heating. It has been suggested that the maximum sizes of chondrules are regulated by the gas flow because large dust particles should be disrupted by the strong gas ram pressure when they melt (Susa \\& Nakamoto 2002, Kato et al. 2006). Recently, we carried out three-dimensional hydrodynamics simulations of molten dust particle exposed to the gas flow and showed that molten cm-sized dust particle is disrupted into many small pieces in a typical setting of nebula shocks (Miura \\& Nakamoto 2007). These pieces have many chances of mutual collisions to form compound chondrules because the local number density of them behind the disrupted particle is enhanced. We name this scenario ``fragment-collision model\" and think that it can be a strong candidate for compound chondrule formation model. This model seems very similar to the model of collision between individual chondrules in the shock-wave heating at the point that compound chondrules are formed in the gas flow (Ciesla 2006). The difference is that in our model, compound chondrules are formed from a single large dust particle, so the dust enhancement in the pre-shock region is not necessarily required.  The purpose of this paper is to estimate the collision probability, i.e., the number of collisions which a single piece will experience in a disruption event, by using a simple formulation. We call the disrupted dust particle as ``parent\" and small pieces as ``ejectors.\" For simplicity, we assume that all ejectors have the same radius of $r_{\\rm e}$ in this paper. We describe the formulations for estimating the collision probability in \\S \\ref{sec:formulation}. The expected probability is discussed in \\S \\ref{sec:predicted_collision_frequency}. We consider appropriate situations in which our model can be applied in \\S \\ref{sec:parameter_range}. We compare our model with observational data of compound chondrules in \\S \\ref{sec:observation}. Finally, we make conclusions in \\S \\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary} We proposed a new scenario for compound chondrule formation named as ``fragment-collision model,\" in the framework of the shock-wave heating model. We modeled the disruption of molten cm-sized parent dust particle exposed to a high-velocity gas flow in order to estimate the efficiency of mutual collisions between small fragments assuming that all of them have the same radius. We obtained collision probability $P_{\\rm coll}$ for a wide range of parameters (parent radius $r_{\\rm p} = 0.2 - 5 \\, {\\rm cm}$, ejector radius $r_{\\rm e} = 100 - 1000 \\, {\\rm \\mu m}$, ram pressure of the gas flow $p_{\\rm fm} = 10^4 - 10^5 \\, {\\rm dyne \\, cm^{-2}}$, and viscosity of molten parent $\\mu = 10 - 10^3 \\, {\\rm g \\, cm^{-1} \\, s^{-1}}$). The estimated collision probability was $\\sim 0.1 - 1$ or more, which is about one order of magnitude larger than the observed fraction of compound chondrules. If we assume that all collisions lead to compound chondrule formation, this result suggests that our model can account for the observed fraction of compound chondrules when about $10\\%$ or less of all chondrules formed via the fragmentation events. Since the fraction of chondrules which have undergone the fragmentation event depends on a size distribution of precursor dust particles in chondrule-forming region, it would be difficult to make a conclusion about the fraction of compound chondrules. However, numerical results of dust coagulation equations in the solar nebula seem to match well with our estimation (Nomura et al. 2007). In addition, our model does not require the dust enhancement in the pre-shock region because compound chondrules are formed from a single large dust particle. This is advantageous to explain the scarcity of isotopic fractionation of sulfur (Tachibana \\& Huss 2005, Miura \\& Nakamoto 2006).  We also compared our model with other observational data. Akaki \\& Nakamura (2005) measured the oxygen isotopic compositions of constituents of compound chondrules and found that in a three-isotopic diagram, all compound chondrules except for enveloping-types fall in the typical range obtained for single chondrules from the same chondrites. These observations are consistent with our model because two constituent chondrules are expected to originate from the same dust reservoirs. Gooding \\& Keil (1981) and Akaki \\& Nakamura (2005) reported that compound chondrules with non-porphyritic pairs are more frequent than that with porphyritic pairs. It is thought that these observations can be also explained by our model because the collision probability depends on the viscosity of molten parent dust particle. The dependence on the viscosity is consistent with the experimental results that porphyritic and non-porphyritic textures have been formed from partially- and completely-molten dust particles, respectively (Lofgren \\& Russell 1986). Finally, the size ratios of secondary to primary in each set of compound chondrule have the mean value about 0.3 (Wasson et al. 1995). This result might be explained by our model because two fragments with different sizes are accelerated by the gas flow with different accelerations. As a result, these two fragments obtain large relative velocity and it would enhance the collision probability. Therefore, the compound chondrule of different-sized pair tends to be formed more frequently than that of same-sized one.  However, it should be noted that all collisions do not necessarily lead to compound chondrule formation. For example, undesirably fast collisions cause disruption of ejectors upon impact. Assuming that the upper limit of the collisional velocity for coalescence is $10^4 \\, {\\rm cm \\, s^{-1}}$ (Ciesla 2006), the appropriate size ratio of secondary (small) chondrule to primary (large) one is $r_{\\rm s} / r_{\\rm l} \\ga 0.5$, which does not account for the observations (mean value of $r_{\\rm s} / r_{\\rm l} \\sim 0.3$, Wasson et al. 1995). In addition, ejectors should cool moderately before collision to make compound chondrule not to fuse into a single droplet. The collision probability after the moderate cooling, however, is much smaller than that of total collisions. In order to overcome these difficulties, we consider other physics that we did not take into account in this paper (e.g., the gas flow is blocked by the parent and/or numerous numbers of ejectors). These issues will be discussed in the forthcoming paper. In addition, our model should be tested by other methods, e.g., three-dimensional (thermo-)hydrodynamics simulation or aerodynamic disruption experiment in the future."
    },
    "0711/0711.2939_arXiv.txt": {
        "abstract": "We calculated profiles of CIV\\,1550, Si\\,IV 1400, NV\\,1240 and OVI\\,1035 doublet lines using results of 3D MHD simulations of disc accretion onto young stars with dipole magnetic field. It appeared that our calculations can not reproduce profiles of these lines observed (HST/GHRS-STIS and FUSE) in CTTSs's spectra. We also found that the theory predicts much larger C\\,IV 1550 line flux than observed (up to two orders of magnitude in some cases) and argue that the main portion of accretion energy in CTTSs is liberated outside accretion shock. We conclude that the reason of disagreement between the theory and observation is strongly non-dipole character of CTTS's magnetic field near its surface. \\footnote{ Published in Proc. of IAU Symp. 243: \"Star-Disk Interaction in Young Stars\", Grenoble, France, 2007, Eds. J. Bouvier \\& I. Appenzeller, p.115 } ",
        "introduction": "Since the beginning of the 1990s, there has been a consensus that the line and continuum emission observed in the spectra of classical T Tauri stars (CTTSs) results from the magnetospheric accretion of circumstellar material. More precisely, the magnetic field of the star is believed to stop the accretion disk from reaching the stellar surface. In some way the disk material becomes frozen in the magnetospheric field lines and slides along them toward the stellar surface, eventually being accelerated to velocities $V_0 \\sim 300$ km/s. The gas is then decelerated in an accretion shock (AS), whose radiation presumbly gives rise to the observed line and continuum emission.  Radial extention of pre- and post-shock radiating regions of CTTS's AS is much smaller than stellar radius -- $Z_{pre}, Z_{pst} \\ll R_*$ in figure~\\ref{fig:as-scheme}, -- making it possible to calculate the structure and spectrum of the AS in 1-D approximation (Lamzin, 1995). Calculations of Lamzin (1998) and Calvet \\& Gullbring (1998) indicated that the structure of the flow can be specified nearly unambiguously by two parameters: the velocity $V_0$ and density $\\rho_0$ (or particle number density $N_0)$ of the gas far in front of the shock. \\begin{figure} \\begin{center} \\includegraphics[height=7cm]{fig1.epsi} \\caption{Schematic structure of CTTS's AS.} \\label{fig:as-scheme} \\end{center} \\end{figure}  Calvet \\& Gullbring (1998) used results of their calculations to derive the parameters of the accretion shock via modeling of the continuum spectral energy distributions of classical T Tauri stars. However, they did not took into account limb darkening effect and what is more the agreement between the calculated and observed spectra of the veiling continuum cannot be considered as a decisive support for the magnetospheric model, since boundary-layer models provide equally good agreement -- see e.g. Basri \\& Bertout (1989). The line spectrum is far more informative, and comparisons of the calculated and observed intensities and profiles of emission lines enable detailed studies of the accretion processes.  Optically thin lines are best suited for this purpose: their intensity ratios can be used to derive physical conditions independent of the geometry of the region where they are formed, and the line profiles provide information about both the velocity field and geometry of this region. The calculations of Lamzin \\& Gomez de Castro (1999) demonstrated that the O\\,III] 1663, Si\\,III] 1892 and C\\,III] 1909 lines should display the highest intensities among the optically thin lines, and just these lines were used to determine the accretion-shock parameters for several young stars in that paper. However, Gomez de Castro \\& Verdugo (2001, 2007) questioned whether these lines form in the AS. Spectral lines of neutral or singly ionized atoms apparantly form not only in AS but in magnetospheric flow and wind as well -- see e.g. Edwards (2007) -- and thus likewise cannot be used for diagnostic of AS.  Resonant UV lines of the C\\,IV, Si\\,IV, N\\,V and O\\,VI uv1 doublets looks as the most suitable lines for diagnostic of CTTSs's AS, especially lines of C\\,IV 1550 doublet: they are strong in CTTS's spectra and calculation of their intensities is relatively simple. ",
        "conclusions": "We demonstrated that observed intensity and profiles of C\\,IV 1550 doublet lines significantly differ from theoretical predictions based on the assumption that magnetic field of CTTSs near stellar surface is close to dipole. We conclude therefore that geometry of CTTS's magnetic field near stellar surface is strongly non-dipole. Multipole components of global magnetic field of young star or/and small-scale magnetic fields of active regions probably produce large divergency of accreted gas stream lines within accretion zone that presumbly can explain disagreement between the theory and observations.  \\medskip  {\\it We thank the LOC of the Symposium for the invitation, financial support and hospitality.}  \\medskip"
    },
    "0711/0711.0611_arXiv.txt": {
        "abstract": "{Numerical simulations predict that minor mergers are an important channel for the mass assembly of galaxies. However, minor mergers are relatively difficult to detect using imaging, especially at high redshift. While such events are much less violent than major mergers, they can nevertheless leave several features on the kinematical structures of remnant galaxies which could be detected using 3D spectroscopy. }{We present the first direct detection of a minor merger in a z$\\sim$0.6 galaxy. Such events could indeed be good candidates to explain the kinematics of perturbed rotating disks observed with GIRAFFE at z$\\sim$0.6.}{We present photometric and kinematical evidence of such an event in a combined analysis of three-band HST/ACS imaging and VLT/GIRAFFE 2D-kinematics.}{Using these data, we are able to demonstrate that a minor merger of a relatively small satellite (mass ratio $\\sim$1:18) is occurring in this galaxy. We also derive a total SFR of $\\sim$21$M_\\odot/yr$.}{Minor mergers could be one of the physical processes explaining the kinematics of perturbed rotating disks, which represent $\\sim$25\\% of emission line intermediate mass galaxies at z$\\sim$0.6. 3D spectroscopy appears to be a very good tool to identify minor mergers in distant (and local) galaxies.} ",
        "introduction": "Minor mergers are usually defined as the merging between two galaxies with a mass ratio smaller than 1:3 to 1:5. Hence, during a single minor merger event, the amount of stellar mass added to the main progenitor remains negligible. However, cosmological simulations predict that minor mergers are one order of magnitude more frequent than major mergers (e.g., \\citealt{Khochfar06}); Moreover, they can be responsible for \\emph{in situ} star formation episodes \\citep{Woods06,Woods07} during the interaction. Therefore, minor mergers are thought to be an important channel for the galaxy mass assembly, even possibly the dominant contributor for intermediate mass galaxies \\citep{Guo07}.   Contrary to major mergers, minor mergers are much less violent dynamical processes since they do not destroy the disk of the main progenitor. Nevertheless, minor mergers can result in several important imprints on the morphological and kinematical structure of the remnant: they can thicken the disk of the main progenitor by increasing its vertical scale-length and velocity dispersion, as found both in numerical simulations \\citep{Toth92,Walker96,Velazquez99,Bournaud05,Robertson06} and observations \\citep{Schwarzkopf00}. They are believed to play an important role during the formation of thick disks (e.g., \\citealt{Yoachim06}), and they could also explain several other structural properties of spiral galaxies, such as the formation of anti-truncated surface brightness profiles \\citep{Younger07}, or activation of bars \\citep{Laine99}.    In spite of their possible cosmological importance and influence on galaxy structure, relatively few studies have focused on minor mergers. Statistical studies aiming at identifying an important number of minor mergers in large photometric surveys rely on counts of galaxies in pairs which have a large difference in a red band magnitude, used as a proxy for the stellar mass. One of the main difficulty encountered is then the high frequency of false pairs due to the large number of background galaxies, which makes it necessary to have spectroscopic confirmations for both progenitors \\citep{Woods06}. Moreover, once the minor progenitor gets too close to the main progenitor, it becomes difficult to separate both objects using imaging alone, especially at high redshift. Therefore, detailed studies of minor mergers are quite rare \\citep{Laine99}.   Another route for identifying on-going minor mergers may be to use 3D spectroscopy. Using GIRAFFE at the VLT, \\cite{Yang07} studied the 2D-kinematical properties of a complete sample of 63 z$\\sim$0.6 galaxies, which are representative of $M_{stellar}\\geq$1.5 10$^{10}M_{\\odot}$ emission line ($EW_0([OII])\\geq$ 15\\AA) galaxies. Interestingly, they have classified 25\\% of these objects as ``perturbed rotating'' galaxies (PR): their kinematics shows all the features of a rotating disk, but with a peak in the velocity dispersion map shifted away from the dynamical center. These local perturbations cannot be reproduced from the rotation patterns, and thus cannot be attributed to rotation \\citep{Yang07}. They might illustrate a possible minor merger event, which does not affect the disk stability, but has an expected signature in the $\\sigma$ map.   Other studies have also revealed the existence of morphologically and/or kinematically disturbed galaxies at higher redshifts. \\cite{Forster06} have obtained 3D spectroscopy of several z$\\sim$2 galaxies among which some also show off-centered $\\sigma$ peaks in relatively well-ordered VFs. However, all these studies (including GIRAFFE) are seeing limited, and the physical process responsible for such off-centered $\\sigma$ peaks remains unclear, even when using adaptive optics \\citep{Law07}. In this \\emph{letter}, we present what is, to our knowledge, the first direct evidence for the detection of a minor merger at high redshift. This object (J033226.23-274222.8) has been classified as a PR by \\cite{Yang07} and is thus an ideal case to test the possibility that minor mergers could be a viable process explaining PR kinematics. This \\emph{letter} is organized as follows: Sect. 1 presents the general properties of J033226.23-274222.8; Sect. 2 details the kinematical evidences for a minor merger event in this galaxy; Sect. 3 gives some estimates for the mass and the star formation rate; A conclusion is given in Sect. 4. We adopt $H_0$=70 km/s/Mpc, $\\Omega _M$=0.3, and $\\Omega _\\Lambda$=0.7, and the $AB$ magnitude system. ",
        "conclusions": "In this \\emph{letter}, we have presented the first detection of a minor merger at high redshift (z=0.6679). Using ACS imaging, we identified an infalling satellite onto a late-type spiral. Gas kinematics reveals a velocity field with no large-scale perturbations, except a twist in the isovelocities which might indicate a possible warp of the gaseous disk. Such a warp is also seen in the stellar disk, using photometric isophotes. The gas $\\sigma$-map shows a peak off-centered from the dynamical center, and close to the infalling satellite. Such a peak cannot be attributed to rotation and is interpreted as the signature of shock-heated gas during the collision. We derive a total SFR$\\sim$21$M_\\odot/yr$ and a rough stellar mass ratio estimate $\\sim$1:18, in agreement with this kind of events. Therefore, J033226.23-274222.8 seems to be the first detection of a minor merger at z$\\sim0.6$, and more generally at high redshift. This study demonstrates that using 3D spectroscopy appears to be a promising tool to identify minor mergers, both in nearby and distant galaxies. Finally, minor mergers are a good candidate to explain the kinematics observed in perturbed rotating galaxies, which account for 25\\% of emission line intermediate mass galaxies at z$\\sim$0.6."
    },
    "0711/0711.2752_arXiv.txt": {
        "abstract": "We present a detailed analysis of observations of the high mass X-ray binary Cen~X-3 spanning two consecutive binary orbits performed with the \\textsl{RXTE} satellite in early March 1997. During this time Cen~X-3 had a luminosity of $L_{2-10\\,{\\mathrm{keV}}} \\sim 4$--$5 \\times 10^{37} {\\mathrm{erg\\,s}}^{-1}$ and a pulse period of $4.814$\\,s. The PCA and HEXTE light curves both show a clear reduction in count rate after mid-orbit for both binary revolutions. We therefore analyze two broad band spectra for each orbit, before and after mid-orbit. Consistent with earlier observations these four joint PCA and HEXTE spectra can be well described using a phenomenological pulsar continuum model, including an iron emission line and a cyclotron resonance scattering feature.  While no strong spectral variations were detected, the second half of orbit 2 shows a tendency toward being softer and more strongly absorbed. In order to follow the orbital phase-dependent evolution of the spectrum in greater detail, we model spectra for shorter exposures, confirming that most spectral parameters show either a gradual or sudden change for the second half of the second orbit. A comparison with a simple wind model indicates the existence of an accretion wake in this system. We also present and discuss high resolution pulse profiles for several different energy bands, as well as their hardness ratios. PCA and HEXTE spectra were created for 24 phase bins and fitted using the same model as in the phase averaged case. Systematic pulse phase-dependent variations of several continuum and cyclotron line parameters were detected, most notably a significant increase of the cyclotron line energy during the early rise of the main peak, followed by a gradual decrease. We show that applying a simple dipole model for the magnetic field is not sufficient to describe our data. ",
        "introduction": "\\label{sec:intro} The pulsar nature of the accreting high mass X-ray binary (HMXB) Centaurus~X-3 was first discovered by \\textsl{UHURU} \\citep{schreier72,giacconi71}, measuring a spin period of $\\sim$ 4.8\\,s and an orbital period of $\\sim$ 2.1\\,days. The binary system consists of a neutron star with a mass of $1.21\\pm 0.21\\,{\\mathrm{M}}_\\odot$ accompanied by an O\\,6--8~III supergiant star with a mass of $20.5\\pm0.7 {\\mathrm{\\,M}}_\\odot$ \\citep{hutchings79,ash99}. The distance of the binary system has been estimated to be roughly 8\\,kpc with a lower limit of 6.2\\,kpc \\citep{krzeminski74} and an eccentricity of $\\leq 0.0016$ \\citep{bildsten97}. \\citet{day91} used dust scattering measurements to derive a distance of $5.4\\pm0.3$\\,kpc, a result only marginally consistent with \\citet{krzeminski74} within given uncertainties.  In order to allow for easier comparison with earlier publications we will adopt the widely used distance of 8\\,kpc. HMXBs generally show a strong stellar wind which can supply a continuous tidal stream to the compact object when the companion star is close to filling its Roche lobe \\citep{petterson78, stevens88, blondin91}. \\citet{day93} discussed that in a system like Cen~X-3 an X-ray excited wind would be sufficient as a possible mechanism for mass transfer able to sustain a constant mass flux.  A system with a combined mass transfer mechanism is, e.g., GX~301$-$2 \\citep{leahy02, leahy07}, where the folded \\textsl{RXTE}/ASM light curve was best modeled by a stellar wind with an additional gas stream as described by \\citet{stevens88}. In addition evidence for an accretion disk in Cen~X-3 can be found: \\citet{tjemkes86} compared the observed light curve of Cen~X-3 with modeled binary light curves and achieved best results when including an accretion disk. The observed overall spin-up trend of 1.135\\,ms/yr with fluctuations on timescales of a few years \\citep{tsunemi96} also indicates the presence of an accretion disk, attributed to Roche lobe overflow. A realistic picture of the mass transfer mechanism in Cen~X-3 would therefore be a combination of disk accretion with an excited stellar wind, as proposed by \\citet{petterson78}.  An absorption feature in the X-ray spectrum at $\\sim$30\\,keV, interpreted as a cyclotron resonance scattering feature (CRSF) also known as cyclotron line, was first observed in Cen~X-3 with \\textsl{Ginga} by \\citet{nagase92}. \\textsl{BeppoSAX} and \\textsl{RXTE}/HEXTE confirmed this feature at 28\\,keV and 30\\,keV, respectively \\citep{santa98,heindl99}. It is assumed to be the fundamental line and so far no higher line harmonics have been observed for this source.  CRSFs are generated due to the interaction of photons with electrons which are quantized in Landau levels, the discrete energy states populated by electrons in strong magnetic fields. Photons of energies corresponding to the separation of the Landau levels are absorbed and re-emitted by the electrons. The net effect is that photons from the accretion column above the magnetic poles that have energies near the CRSF centroid energy are scattered out of the line of sight and/or redistributed in energy.  These features can be used to measure directly the magnetic field strength in the line-forming region since the fundamental cyclotron line energy is given by \\begin{equation} E_{\\mathrm{cyc}} = \\frac{1}{1+z}\\frac{\\hbar e B}{m_{\\mathrm{e}}} = \\frac{11.6 {\\mathrm{keV}}}{1+z} B_{12}, \\end{equation} where $B_{12}$ is the magnetic field strength near the neutron star surface, in units of $10^{12}$\\,Gauss and $z$ is the gravitational redshift \\citep{canuto77}. Assuming a typical neutron star mass of $1.4 {\\mathrm{\\,M}}_\\odot$ and neutron star radius of 10\\,km gives $z =0.3$. Cen~X-3's cyclotron line at $\\sim$ 30\\,keV then corresponds to a magnetic field of $\\sim 3.4\\times 10^{12}$\\,G.  \\cite{kohmura01} analyzed the PCA data of this observation in the energy band $2-13$\\,keV which has a high time resolution of 15.625\\,ms. They saw a delay in the time variability of the iron K band by $0.39\\pm 0.10$\\,ms, and concluded that these photons were reprocessed away from the neutron star and are probably part of the matter accreting onto the NS.  A preliminary analysis by \\citet{heindl99} of the HEXTE data presented in this work concentrated on basic pulse phase resolved spectroscopy, dividing the pulse profile into four broad phase bins, and provided a first indication of a possible increase of the cyclotron line energy during the rise of the main pulse. Systematic variations of the CRSF centroid energy over pulse phase have also been observed in other sources, e.g., Her X-1 \\citep{soong90,gruber01}, Vela~X-1 \\citep{labarbera03,kreykenbohm02}, 4U\\,0352$+$309 \\citep{coburn01}, 4U\\,1538$-$52 \\citep{clark90}, and GX~301$-$2 \\citep{kreykenbohm04}.  The origin of these changes is not yet well understood.  Straightforward mechanisms like the angular dependence of relativistic corrections to the cyclotron energy may not be sufficient to explain the observed variations.  Higher harmonics rarely, if ever, are integer multiples of the fundamental line energy. In some sources this can be explained by invoking the proper non-integer line spacing expected from the relativistic cross section \\cite[e.g.,][]{pottschmidt05}. However, more complicated mechanisms such as a B-field gradient at the magnetic pole, have also been invoked \\citep{nishimura05}. In addition a correlation between CRSF energy and luminosity can be observed in several sources \\citep[e.g.,][]{mihara95,mowlavi06,labarbera05,staubert07, nakajima06}.  In the following we further characterize the phase dependence of the line parameters in Cen~X-3 using the exceptionally long -- spanning $\\sim$4\\,days, i.e., two binary revolutions -- \\textsl{RXTE} observation of the source in 1997.  \\citet{burderi00}, in the following referred to as B00, observed Cen~X-3 simultaneously with \\textsl{BeppoSAX} overlapping 11\\,hours of \\textsl{RXTE} observations during the first binary orbit. We will compare our results with B00 whenever possible.  We use combined PCA and HEXTE data for the analysis which allows us to constrain the continuum shape of the spectrum and the CRSF parameters much more strongly than with HEXTE alone. We compare our analysis with other cyclotron line sources and apply a newly developed physical model for the CRSF to the data. For the analysis over the orbit we compare with a simplified wind model to discuss the evolution of $N_{\\mathrm{H}}$ through the observation.  The remainder of this paper is organized as follows: basic observation parameters as well as the data reduction are described in \\S\\ref{sec:data}. In \\S\\ref{sec:light} we describe the PCA and HEXTE light curves observed over the two binary orbits.  Phase averaged spectra for the first and second halves of each orbit and their best fit models are presented in \\S\\ref{sec:phase_avg}, as well as the temporal evolution of the spectral parameters with a resolution of roughly one \\textsl{RXTE} orbit throughout the two Cen X-3 binary orbits.  \\S\\ref{sec:phase_res} is dedicated to the pulse phase resolved analysis, including the presentation of pulse profiles in multiple energy bands and modeling of phase resolved spectra. The results are discussed in \\S\\ref{sec:discussion}. ",
        "conclusions": ""
    },
    "0711/0711.0743_arXiv.txt": {
        "abstract": "Gravitational wave standard sirens have been considered as precision distance indicators to high redshift; however, at high redshift standard sirens or standard candles such as supernovae suffer from lensing noise. We investigate lensing noise as a signal instead and show how measurements of the maximum demagnification (minimum convergence) probe cosmology in a highly complementary manner to the distance itself.  Revisiting the original form for minimum convergence we quantify the bias arising from the commonly used approximation.  Furthermore, after presenting a new lensing probability function we discuss how the width of the lensed standard siren amplitude distribution also probes growth of structure.  Thus standard sirens and candles can serve as triple probes of dark energy, measuring both the cosmic expansion history and growth history. ",
        "introduction": "}  Observations of distant standardized sources carry information not only on the global cosmology but the signal propagation through the intervening universe.  In particular, gravity affects the propagation on all scales, from the global curvature of spacetime to the local inhomogeneities of the mass distribution.  In this sense, detecting the acceleration of the cosmic expansion through measuring distances from the luminosity, size, or amplitude of a source is as much gravitational lensing as is the image distortion pattern around massive structures.  While the global spacetime properties determine the distances in a manner fixed by the cosmology, inhomogeneities spread the derived distances about the ``truth'' or smooth universe value according to a lensing probability distribution function depending on the energy-momentum distribution and the convergence and shear it induces in the image. An interesting feature of this probability distribution is the presence of a minimum convergence or maximum demagnification corresponding to an ``empty beam'' \\cite{dr1}.  This cutoff carries cosmological information.  Thus the mean distance measured to standardized sources such as Type Ia supernovae in luminosity or inspiraling supermassive black hole binaries in gravitational wave amplitude (standard sirens \\cite{holzhughes,dalal}) can probe cosmology, but so can the lensed amplitude distribution through the minimum convergence and the width of the distribution. Gravitational wave sources in particular may be measured to high redshift $z\\gg1$ where lensing effects are quite evident.  Standard sirens and candles thus have the potential to be triple probes of cosmology.  In \\S\\ref{sec:beam} we derive a general expression for the minimum convergence and maximum demagnification, correcting an often used expression in the literature.  Implications for gravitational wave standard sirens and supernova standard candles are examined in \\S\\ref{sec:stds}. The use of the minimum convergence as a cosmological probe in its own right is considered in \\S\\ref{sec:lever}, along with its complementarity with distance measurements, and the potential for the lensed distribution width to serve as a growth probe.  In the Appendix we briefly discuss effective averaging procedures in an inhomogeneous universe and why the standard weak gravitational lensing formula remains valid. ",
        "conclusions": "}  Standardized siren or candle measurements of the mean distance-redshift relation provide precise cosmological parameter estimation, but their full information is even richer.  The distribution in the binary black hole gravitational wave strain amplitude or supernova luminosity carries with it characteristics reflecting the propagation along the line of sight.  In particular, this distribution has a physically precise lower boundary due to the minimum convergence, which involves a distance ratio, providing a further geometric measurement of the cosmology.  From the full, long known definition of the minimum convergence or maximum demagnification we show that a common approximation induces a substantial bias as measurements move toward higher redshifts and better precision. We explore several observational issues with using minimum amplification as a practical cosmological probe, suggesting that gravitational wave standard sirens might enable this method if sufficient numbers of sirens can be detected with precise redshifts.  Along with a new probability distribution function for lensing amplification, we discuss how the width of the distribution might be used as a measure of the mass power spectrum or linear growth factor.  From a Fisher analysis it appears that the method would need to be developed to allow 2\\% accuracy on the lensing dispersion for a significant gain in cosmological constraints, but weaker precision could still provide useful crosschecks or consistency tests of gravity.  Substantial work needs to be done to realize the promises of the lensing distribution information for precision cosmology.  In particular, large ray tracing simulations are needed to verify and accurately calibrate the minimum detectable convergence and distribution width.  While challenging, nevertheless it is exciting that distance measurements of gravitational wave standard sirens and possibly supernova standardized candles can act as a triple probe of cosmology and dark energy: 1) the mean amplitude (unaffected by lensing) gives a geometric measure of the distance, 2) the minimum amplification gives a complementary geometric measure consisting of distance ratios, and 3) the dispersion encodes the growth history through the linear growth factor. If this potential can be realized, distance-redshift measurements of sirens or candles can provide probes of both the cosmic expansion and growth histories."
    },
    "0711/0711.2570_arXiv.txt": {
        "abstract": "We report the discovery of a light echo (LE) from the Type Ia supernova (SN) 2006X in the nearby galaxy M100. The presence of the LE is supported by analysis of both the {\\it Hubble Space Telescope (HST)} Advanced Camera for Surveys (ACS) images and the Keck optical spectrum that we obtained at $\\sim$300~d after maximum brightness. In the image procedure, both the radial-profile analysis and the point-spread function (PSF) subtraction method resolve significant excess emission at 2--5 ACS pixels ($\\sim0.05''-0.13''$) from the center. In particular, the PSF-subtracted ACS images distinctly appear to have an extended, ring-like echo. Due to limitations of the image resolution, we cannot confirm any structure or flux within 2 ACS pixels from the SN. The late-time spectrum of SN 2006X can be reasonably fit with two components: a nebular spectrum of a normal SN~Ia and a synthetic LE spectrum. Both image and spectral analysis show a rather blue color for the emission of the LE, suggestive of a small average grain size for the scattering dust. Using the Cepheid distance to M100 of 15.2 Mpc, we find that the dust illuminated by the resolved LE is $\\sim$27--170~pc from the SN. The echo inferred from the nebular spectrum appears to be more luminous than that resolved in the images (at the $\\sim$2$\\sigma$ level), perhaps suggesting the presence of an inner echo at $<$2 ACS pixels ($\\sim0.05''$). It is not clear, however, whether this possible local echo was produced by a distinct dust component (i.e., the local circumstellar dust) or by a continuous, larger distribution of dust as with the outer component. Nevertheless, our detection of a significant echo in SN 2006X confirms that this supernova was produced in a dusty environment having small dust particles. ",
        "introduction": "Light echoes (LEs) are produced when light emitted by the explosive outburst of some objects is scattered toward the observer by the foreground or surrounding dust, with delayed arrival time due to the longer light path. This phenomenon is rare, having been observed only around a few variable stars in the Galaxy, and around several extragalactic supernovae (SNe). The best-studied events are SN 1987A (Schaefer 1987; Gouiffes et al. 1988; Chevalier \\& Emmering 1988; Crotts 1988; Crotts, Kunkel, \\& McCarthy 1989; Bond et al. 1990; Xu et al. 1995) and the peculiar star V838 Mon (Bond et al. 2003). Other SNe with LEs include the Type II SNe 1993J (Liu et al. 2002; Sugerman 2003), 2002hh (Meikle et~al. 2006; Welch et al. 2007), and 2003gd (Sugerman 2005; Van Dyk et al. 2006), as well as the Type Ia SNe 1991T (Schmidt et al. 1994; Sparks et al. 1999), 1998bu (Cappellaro et al. 2001; Garnavich et al. 2001), and possibly 1995E (Quinn et al. 2006). Besides their spectacular appearance, LEs offer a unique means to diagnose the composition, distribution, and particle size of the scattering dust. In particular, LEs from the circumstellar environments might provide constraints on SN progenitors.  The Type Ia SN 2006X was discovered on 2006 February 7.10 (UT dates are used throughout this paper) by S. Suzuki and M. Migliardi (IAUC 8667, CBET 393) in the nearby spiral galaxy NGC 4321 (M100). Extensive photometric and spectroscopic coverage is presented by Wang et al. (2007, hereafter W07). They suggest that SN 2006X is highly reddened [$E(B - V)_{\\rm host} = 1.42 \\pm 0.04$ mag] by abnormal dust with $\\Re_{V} = 1.48 \\pm 0.06$. Its early-epoch spectra are characterized by strong, high-velocity features of both intermediate-mass and iron-group elements. In addition to the anomalous extinction and the very rapid expansion, SN 2006X exhibits a continuum bluer than that of normal SNe~Ia. Moreover, its late-time decline rate in the $B$ band is slow, $\\beta = 0.92 \\pm 0.05$ mag (100~{\\rm d})$^{-1}$, significantly below the 1.4 mag (100~{\\rm d})$^{-1}$ rate observed in normal SNe~Ia and comparable to the decay rate of 1.0 mag (100~{\\rm d})$^{-1}$ expected from $^{56}$Co $\\rightarrow$ $^{56}$Fe decay. This may suggest additional energy sources besides radioactive decay, such as the interaction of the supernova ejecta with circumstellar material (CSM) and/or a LE.  Attempts to detect the CSM in SNe~Ia in different wavebands were unsuccessful before SN 2006X, and only some upper limits could be placed (see Patat et al. 2007a, and references therein) except for the peculiar SNe~Ia/IIn 2002ic (Hamuy et al. 2003; Deng et al. 2004; Wang et al. 2004; Wood-Vasey et~al. 2004) and 2005gj (Aldering et~al. 2006; Prieto et~al. 2007). Recent progress in this respect was made from high-resolution spectroscopy by Patat et~al. (2007b, hereafter P07), who find time-variable Na~I~D absorption lines in spectra of SN 2006X. This has been interpreted as the detection of CSM within a few 10$^{16}$ cm ($\\sim 0.01$ pc) from the explosion site of the supernova. With the inferred velocity, density, and location of the CSM, P07 proposed that the companion star of the progenitor of SN 2006X is most likely to be a red giant (but see Hachisu et~al. 2007, who present a main-sequence star model with mass stripping). Note, however, that SN 2006X exhibited somewhat abnormal features in spectra and photometry; it may not represent a typical SN~Ia. Multi-epoch, high-resolution spectral observations of SN 2007af, a normal SN~Ia, do not reveal any significant signature of CSM absorption (Simon et~al. 2007).  In this paper we report the discovery of an optical LE around SN 2006X, with evidence from {\\it Hubble Space Telescope (HST)} Advanced Camera for Surveys (ACS) images and Keck optical spectra. The paper is organized as follows. In \\S 2 we briefly describe the late-epoch data available for SN 2006X, while the data analysis and the interpretation are presented in \\S 3. We discuss the properties of the light echo and the underlying dust in \\S 4. Our conclusions are given in \\S 5. ",
        "conclusions": "The emergence of a LE in SN 2006X has been confirmed with PSF-subtracted {\\it HST} ACS images which show a ring-like, but rather extended, echo 2--5 pixels ($0.05''$--$0.13''$) from the SN site at $t = 308$~d past maximum brightness. A Keck nebular spectrum of the SN taken at a similar phase provides additional evidence for the LE scenario; it can be decomposed into a nebular spectrum of a normal SN~Ia and a reflection spectrum consisting of the SN light emitted at early phases.  From the resolved echo image, we derive that the intervening dust is $\\sim$27--170~pc from the supernova. Based on the quite blue color of the echo, we suggest that the mean grain size of the scattering dust is substantially smaller than Galactic dust. Smaller dust particles are also consistent with the low $\\Re_{V}$ value obtained from the SN photometry. Our detection of a LE in SN 2006X confirms that this SN~Ia occurred in a dusty environment with atypical dust properties, as suggested by the photometry (W07).  Analysis of the nebular spectrum might also suggest a local echo at $<$27 pc (or at $<$2 pixels) that is not resolved in the PSF-subtracted image. This possible local echo is likely associated with the CSM dust produced by the progenitors, though detailed modeling of the echo spectrum and/or further high-resolution imaging are required to test for the other possibilities, such as very forward scattering by a distant cloud or CSM-ejecta interaction."
    },
    "0711/0711.0433_arXiv.txt": {
        "abstract": "% The \\swift\\ spacecraft detects and autonomously observes \\til100 Gamma Ray Bursts (GRBs) per year, \\til96\\%\\ of which are detected by the X-ray telescope (XRT). GRBs are accompanied by optical transients and the field of ground-based follow-up of GRBs has expanded significantly over the last few years, with rapid response instruments capable of responding to \\swift\\ triggers on timescales of minutes. To make the most efficient use of limited telescope time, follow-up astronomers need accurate positions of GRBs as soon as possible after the trigger. Additionally, information such as the X-ray light curve, is of interest when considering observing strategy. The \\swift\\ team at Leicester University have developed techniques to improve the accuracy of the GRB positions available from the XRT, and to produce science-grade X-ray light curves of GRBs. These techniques are fully automated, and are executed as soon as data are available. ",
        "introduction": "The \\swift\\ satellite triggers on \\til100 Gamma Ray Bursts (GRBs) per year, and the X-ray telescope (XRT, Burrows \\etal2005) provides localisations accurate to \\til5 arcsec for $>90\\%$ of these. However, GRBs, and their optical counterparts, fade rapidly and ground-based observers have limited telescope time available to use for observations. It is thus desirable to produce precise, accurate positions rapidly. Furthermore, a key indication of the scientific interest of a GRB comes from the \\swift-XRT light curve. These are non-trivial to produce correctly, but it is desirable to generate them accurately and rapidly.  We describe how these two challenges are being met by the \\swift\\ team at Leicester, providing automatic data analysis which gives results of publication-grade quality. While our work is specific to \\swift, rapid response time-domain astrophysics is a fast-growing field, and this type of software will be increasingly useful to astronomers. ",
        "conclusions": ""
    },
    "0711/0711.4951_arXiv.txt": {
        "abstract": "{The faint-end slopes of galaxy luminosity functions (LFs) in clusters of galaxies have been observed in some cases to vary with clustercentric distance and should be influenced by physical processes (mergers, tides) affecting cluster galaxies. However, there is a wide disagreement on the values of the faint-end LF slopes, ranging from $-1$ to $-2.3$ in the magnitude range $-18 < M_r < -14$.} {We investigate the LF in the very relaxed cluster Abell~496.} {Our analysis is based on deep images obtained at CFHT with MegaPrime/MegaCam in four bands ($u^*$$g'$$r'$$i'$) covering a 1$\\times$1 deg$^2$ region, which is centered on the cluster Abell 496 and extends to near its virial radius.  The LFs are estimated by statistically subtracting a reference field taken as the mean of the 4 Deep fields of the CFHTLS survey.  Background contamination is minimized by cutting out galaxies redder than the observed Red Sequence in the $g'-i'$ versus $i'$ colour-magnitude diagram.  } {In Abell~496, the global LFs show a faint-end slope of $-1.55\\pm0.06$ and vary little with observing band. Without colour cuts, the LFs are much noisier but not significantly steeper. The faint-end slopes show a statistically significant steepening from $\\alpha=-1.4\\pm0.1$ in the central region (extending to half a virial radius) to $-1.8\\pm0.1$ in the Southern envelope of the cluster. Cosmic variance and uncertain star-galaxy separation are our main limiting factors in measuring the faint-end of the LFs. The large-scale environment of Abell~496, probed with the fairly complete 6dFGS catalogue, shows a statistically significant 36 Mpc long filament at $\\rm PA=137^\\circ$, but we do not find an enhanced LF along this axis. } {Our LFs do not display the large number of dwarf galaxies ($\\alpha \\approx -2$) inferred by several authors, whose analyses may suffer from field contamination caused by inexistent or inadequate colour cuts. Alternatively, different clusters may have different faint-end slopes, but this is hard to reconcile with the wide range of slopes found for given clusters and for wide sets of clusters. } ",
        "introduction": "\\label{sec:intro}  Clusters of galaxies represent an extreme environment for galaxy evolution, either in situ or through the accretion of galaxies within groups, which are situated in the filamentary network of our hierarchical Universe.  The analysis of the galaxy luminosity function (LF) in several wavebands is a good way to sample the history of the faint galaxy population (e.g. \\citealp{Adami+07}) including star formation history, evolutionary processes and environmental effects.  In particular, the slope of the faint-end of the LF is a direct indicator of the importance of dwarf galaxies, which are expected to be more fragile in the environment of clusters.  The great majority of studies of the LF indicate faint-end slopes in the range $-0.9$ to $-1.5$, but these mostly did not reach very faint magnitudes (see Table~1 in \\citealp{DePropris+03} and references therein).  Recent deep imaging has shed more light on the LF at faint luminosities. Table~\\ref{alphalit} shows deep (fainter than absolute magnitude $M=-16$) estimations of the faint-end slope: most studies conclude to fairly shallow slopes $\\alpha \\simeq -1.3$ (with typical uncertainties of 0.1  to 0.2), while several point to faint-end slopes as steep as $\\alpha \\approx -2.3$, which diverges in luminosity (unless the LF becomes shallower or is cut off at some very faint luminosity).   Part of the wide range of faint-end slopes may be caused by cosmic variance of the background counts. The range of faint-end slopes of the LF may also be due to different mass buildup histories of clusters, through spherical and filamentary infall and major cluster-cluster mergers. However, some of this dispersion in slopes could be caused by systematic uncertainties such as in the star/galaxy separation or through different surface brightness cuts.  Spectroscopic-based LFs would alleviate this problem, but until recently, such spectroscopic-based LFs do not extend fainter than $M=-16$. The exceptions are studies by \\cite{RG08} and \\cite{MMB08} who both find shallow slopes for Virgo cluster galaxies down to $M_r = -14$, as well as \\cite{PC07} for the Perseus cluster core.   Unfortunately, galaxy formation simulations do not yet probe the galaxy LF down to sufficiently faint luminosities: the deepest study, by \\cite{Lanzoni+05}, probed the LF with the GALICS semi-analytical galaxy formation model only down to $M_B < -16.5$ (they found $\\alpha \\simeq -1.3$ in the $B$ band and $\\simeq -1.4$ in the $K$ band within the virial radii of clusters, and $\\alpha \\simeq -1.0$ between 1 and typically 3 virial radii).  The present paper aims at clarifying the debate on the faint-end of the LF and on understanding the different effects of spherical and filamentary infall on the LF. Indeed, the influence of infall on the galaxy population, in particular on the LF is not always well understood, except for a few clusters, such as Coma (e.g. \\citealp{Adami+07}). To achieve this aim, we chose to analyse the very relaxed cluster, Abell~496, known to be very regular both at X-ray and optical wavelengths, and also from a dynamical point of view \\citep{Durret+00}. Our data cover a field of view that is wide enough to reach the virial radius and thus probe a variety of environments. The LFs are computed after a better filtering of artefacts through a minimum galaxy width, a better filtering of background galaxies through the rejection of galaxies redder than the Red Sequence of early-type galaxies, and make use of a thorough analysis of the uncertainties due to cosmic variance, photometric errors and imperfect star/galaxy separation. Our general motivation is to obtain LFs in various regions of Abell~496 and compare them with previous works.   The paper is organised as follows. We present our MegaCam data and data reduction in Section 2. In Section 3, we describe how we compute LFs using large comparison fields from the CFHTLS to statistically subtract the fore- and background galaxy population. In Section 4, we present our results obtained for the LFs of Abell~496 in various regions. In Section 5, we briefly discuss our results concerning the LFs in terms of large scale environmental effects on the cluster galaxy populations.  Finally, in Section 6 we compare our LFs to other determinations of the Abell~496 LF, and we discuss the discrepancy between our moderate faint-end slope and the steep faint-end slopes recently found by several authors in several clusters.  With a mean heliocentric velocity of $9885 \\, \\rm km \\, s^{-1}$ \\citep{Durret+00}, Abell~496 has a (luminosity) distance modulus of 35.73, and the scale is 0.636 kpc arcsec$^{-1}$ (including cosmological corrections\\footnote{We used the cosmological corrections in the CMB frame as  provided in the NASA/IPAC Extragalactic Database (NED).}). We give magnitudes in the AB system. ",
        "conclusions": "\\subsection{Comparison with previous analyses of Abell~496}  The LF of Abell~496 was previously measured by \\cite{MCMdG98}, who analyzed the cluster in 4 small fields, one including the cluster centre, and by \\cite{DAL02}, who measured th LF in a $42'\\times28'$ field in the $I$ band. The faint-end slopes of $-1.69\\pm0.04$ in $r$ and $-1.49\\pm0.04$ in $i$ found by \\citeauthor{MCMdG98} are consistent with our slopes of $-1.62\\pm0.05$ and $-1.52\\pm0.05$ in these bands. Moreover, their faint-end slope of $-1.34\\pm0.04$ in $g$ is probably consistent with our slope ($-1.53\\pm0.08$), given that their estimated error neglects cosmic variance, which, as we show in Fig.~\\ref{fig:allsig}, is many times greater than the Poisson variance. \\cite{DAL02} find a slope of $-1.79\\pm0.01$ in $I$, whereas we find here a slope of $-1.52\\pm0.05$. Given our analysis of cosmic variance, we estimate the error from cosmic variance on the slope of \\cite{DAL02} to be roughly 0.08. If this error estimate is correct, then the difference is faint-end slopes would be statistically significant. This difference may be caused by the different regions probed in the two studies and the different reference fields used.      \\subsection{Comparison of radial trends with other clusters}  Our $g'$, $r'$ and $i'$ LFs exhibit slopes that increase significantly from the centre outwards (Table~\\ref{tab:LF}). This agrees with the trend for flatter LFs around cluster cDs found by \\cite{Lobo+97}, with the steeper slope in the outer envelope of Coma found by \\cite{BHvDvdH02}, and with the greater dwarf-to-giant ratio found by \\cite{DCP98}. Note, however, that all three trends are either qualitative or of marginal statistical significance. Although our radial trend is opposite to that found by \\cite{DePropris+03} in 2dFGRS clusters, their shallower slope in the cluster envelopes is not statistically significant.  \\subsection{A ubiquitous dwarf population?} \\label{ubiq} Whereas our faint-end slope for the core of Abell 496 is fairly shallow ($\\alpha \\simeq -1.4$), it is still steeper than most estimates of the field LF: \\cite{Blanton+03} find $\\alpha_r=-1.05\\pm0.01$ for the SDSS galaxies, while \\cite{JPCS06} find $\\alpha$ between $-1.10\\pm0.04$ ($J$-band) and $-1.21\\pm0.04$ ($r$-band) for 6dFGS galaxies. However, \\cite{Blanton+05} estimate that a careful inclusion of SDSS low surface brightness galaxies yields $\\alpha_r < -1.3$ and perhaps as steep as $-1.5$. But several authors found very steep slopes ($\\alpha$ as steep as $-2.2$) for faint galaxy populations in clusters (see Table~\\ref{alphalit}), suggesting an important population of dwarf galaxies in clusters, which is not seen in the field LFs.  It is difficult to understand how dwarf galaxies could survive better in the hostile cluster environment than in the field. Direct galaxy mergers should have little effect on dwarfs, whose cross-sections are small. Mergers of dwarfs after orbital decay by dynamical friction into the central cD cannot be an explanation: the decay times are expected to be proportional to galaxy mass, so the giant galaxies can disappear into the central cD. However, the orbital decay times for the low mass galaxies should be long enough that the faint-end mass function in clusters would be the same as the field, not steeper. Moreover, there is no observational evidence for luminosity segregation of faint galaxies in clusters (e.g., \\citealp{Pracy+05}). Finally, tidal effects from the cluster potential are the same, to first order, on giant and dwarf galaxies. Given the trend that low luminosity ellipticals are less concentrated than more luminous ones \\citep{GTC01}, the latter will survive better the strong tides near the cluster centre, which should make the LF shallower, not steeper. The only possible explanation for a steeper faint-end slope in clusters would be that galaxies of moderately low luminosity are tidally fragmented by the cluster potential or by close encounters.  Now, if the field counts in the cluster field are underestimated (because the reference fields are slightly underdense), the resultant net cluster counts will be highly contaminated by field counts. Writing the field counts as $dN/dm = \\hbox{dex}[\\beta\\,( m-m_0)]$ and the cluster faint-end luminosity function as $\\Phi(L) \\propto L^{\\alpha}$, it is easy to show that if the field dominates the net cluster counts, one will end up measuring $\\alpha=(-\\beta/0.4)-1$. For example, if $\\beta=3/5$ (Euclidean counts), one would measure $\\alpha=-5/2$, while if $\\beta \\simeq 0.36$ (as in the DFs in the magnitude range where we are measuring the faint-end slope of the LF), one should find $\\alpha = -1.9$.  We are confident that we suffer little from a background contamination of our LFs. Indeed, the faint-end slopes that we have computed for the global LF (Table~\\ref{tab:LF}) and for the central region of Abell~496 (Table~\\ref{tab:LF}) are all considerably and significantly shallower than $\\alpha = -1.9$ (except in the less sensitive $u^*$ band).  We also took great care not to underestimate the counts in our field count process by the use of very homogeneous field data covering a large enough field of view in order to treat properly the cosmic variance (as described above).  Finally the selection in colour that we have applied eliminates (both in the cluster fields and in the fields used for statistical subtraction) very red objects that are very unlikely to belong to the cluster, so the field subtraction is quite conservative and therefore secure.  The MegaCam colour redshift diagrams derived by \\cite{Ilbert+06} from the VIMOS VLT Deep Survey (VVDS) followup of the CFHTLS DF show that galaxy spectra are shifted to the red to such an extent that, in the redshift range $0.4 < z < 0.6$, the bulk of field (intrinsically blue spiral) galaxies become redder in $g'-r'$ than nearby ellipticals, and the same happens in the range $0.5 < z < 1.2$ for $r'-i'$ colours.  We are therefore probably contaminated by intrinsically blue (spiral) background galaxies with colours almost as red as our Red Sequence at $z < 0.4$ or 0.5.  Note that two of the studies concluding to steep faint-end LF slopes \\citep{dPPHM95,Milne+07} have negative LFs at some magnitudes, which suggests inadequate background subtraction (Fig.~\\ref{fig:llf496} shows that 4 of our 12 non-central local LFs also display negative LFs at some intermediate magnitudes, which makes these particular LFs suspicious).  Moreover, all steep faint-end slopes were found by authors who made no colour cuts, with the exception of \\cite{PBBR06}.  These authors estimated the LFs of clusters in the SDSS, using the statistical subtraction method, and found $\\alpha \\approx -2.2$.  The range of absolute magnitudes where \\citeauthor{PBBR06} see the rise in the faint-end LF,  [$-17.7, -13.7$], corresponds to apparent magnitudes in the range $18.03 < m < 22.03$, which matches our analyzed range of apparent magnitudes. It is puzzling that we do not find in our deeper images the low surface brightness dwarf population found by \\citeauthor{PBBR06} in the SDSS images.  \\begin{figure}[ht] \\centering \\includegraphics[width=9cm]{boue_fig11.ps} \\caption{Our $u^*-r'$ vs $i'$ colour-magnitude diagram for Abell~496. The \\emph{dashed contours} are for our selected cluster members using equation~(\\ref{coloureq}), while the \\emph{solid (red)} contours (same levels, spaced by factors of 1.6) are the galaxies identified as field galaxies, because they are redder than our Red Sequence of equation~(\\ref{coloureq}).  Also shown are the limits used by Popesso et al. (2006) to separate their red cluster galaxies and field galaxies (\\emph{upper horizontal line}) and their blue and red cluster galaxies (\\emph{lower horizontal line}), after correcting both for $(u^*-r')_{\\rm MegaCam} = (u-r)_{\\rm SDSS} - 0.35$, as found using the matches in a field of Abell~85 for which we had both SDSS catalogues and reduced MegaCam images.  The \\emph{vertical lines} delimit the region where Popesso et al. clearly found a steep faint-end LF slope. \\label{umrpopesso}} \\end{figure}   Fig.~\\ref{umrpopesso} displays the $u^*-r'$ vs $i'$ colour-magnitude diagram of the Abell~496 field. The Red Sequence is clearly visible at bright magnitudes and its red edge is still sharp at magnitudes $18 < i' < 22$.  The galaxies called red cluster members by \\citeauthor{PBBR06} and fainter than $M_i' = -19$ are almost all field galaxies according to our colour-magnitude diagram.  Similarly, at $M_i' = -16.5$, our Red Sequence (or its extrapolation) becomes bluer than the blue-red galaxy cut of \\citeauthor{PBBR06}. This suggests that an increasing fraction of both \\citeauthor{PBBR06}'s red and blue galaxies are in fact field galaxies.  {}From inspection of the $u^*-g'$ and $g'-r'$ colour-redshift diagrams of \\cite{Ilbert+06}, one infers that the $u^*-r'$ colour cut used by \\citeauthor{PBBR06} should vary negligibly with redshift. This explains why the galaxies we assign to the field (solid red contours in Fig.~\\ref{umrpopesso}) --- because they are redder than our sharp $g'-i'$ vs $i'$ Red Sequence --- pollute the $u^*-r'$ vs $i'$ colour-magnitude diagram. Therefore, \\emph{colour cuts based upon $u^*-r'$ are not efficient in rejecting background galaxies} (with $z \\la 0.4$).  Nevertheless, while LF analyses based upon the statistical subtraction method with inappropriate or inexistent colour cuts should lead to noisier LFs, it is not clear why they should be biased to steeper slopes, unless the noise is such that the faint-end of the LF fluctuates from positive to negative (which is not the case for the global LF of \\citeauthor{PBBR06}). Although simulations by \\cite{VML01} showed that clusters selected in 2D will have faint-end slopes much steeper than what they put in their simulations ($-1.4$ instead of $-1$ in roughly the same absolute magnitude range as where \\citeauthor{PBBR06} found their slope of $-2.3$), \\citeauthor{VML01} have also shown that LFs measured in clusters selected in 3D, such as all the clusters reported in this paper (including the X-ray selected clusters studied by \\citeauthor{PBBR06}), show no bias in the faint-end slope. A confirmation of this result is that X-ray selected clusters show shallower faint-end slopes than do non X-ray clusters \\citep{VMML04}, which are expected to be more prone to projection effects.   A possible cause for the discrepancy between different analyses could be the uncertainties from star/galaxy separation. This is a serious issue for our analysis, because Abell~496 lies at lower absolute galactic latitude ($b = -36^\\circ$) than our reference DF fields ($b=-58^\\circ, 42^\\circ, 60^\\circ$ and $-53^\\circ$), so that Abell~496 is more contaminated by stars. Indeed,  our $i'$-band LF has a considerably shallower faint-end slope of $-1.27\\pm0.04$ instead of $-1.52\\pm0.05$ (Table~\\ref{tab:LF}). Moreover, in general, the errors on star/galaxy separation could go either way, leading to either an underestimation or overestimation of the faint-end LF slope.  A possible explanation for the wide range of faint-end slopes (see Table~\\ref{alphalit}) is cosmic variance: some clusters may exhibit steeper slopes than others. However, inspecting the slopes found in the literature (Table~\\ref{alphalit}), one notices a wide range of slopes for the same cluster analyzed in similar wavebands, magnitude ranges and fields of view (contrast the slopes in similar regions of Coma of $-1.7$, $-1.8$ of \\citealp{Lobo+97} and \\citealp{Trentham98a} with the slopes shallower or equal to $-1.4$ of \\citealp{AC02}). Moreover, the two studies that have enough clusters to `beat' cosmic variance have inconsistent slopes ($-2.2$ \\citealp{PBBR06} and $-1.4$ \\citealp{Trentham98b}) in the magnitude range $-18$ to $-14$.  \\subsection{Conclusions}   We have analyzed the galaxy LFs in the relaxed cluster of galaxies Abell~496.  We have shown that the LFs are not only well-defined in the central region (with a faint-end slope $\\alpha = -1.4\\pm0.1$), but also in the South. A concentration of clusters is indeed observed towards the South-East and along a filament extending South-East to North-West (Fig.~\\ref{fig:lss}), suggesting the existence of a cosmological filament linking Abell~496 with various poor clusters and groups. However, such a filament cannot be very dense since no X-ray emission is detected in this direction, contrary to what is observed in Abell~85 \\citep{Durret+03}.  We discuss the disagreement of our fairly shallow faint-end slope for Abell~496 with that found by other authors in this and other clusters. Although it is clear that a careful estimate of the reference field galaxy counts and their cosmic variance are crucial, there may also be a cosmic variance in the faint-end slopes of cluster LFs. We find that uncertainties in the star/galaxy separation can be responsible for some (but probably not all) of the scatter in the faint-end LF slopes given in the literature. We highlight the removal of galaxies redder than the Red Sequence as a means to reduce the noise in the LFs, but it is not clear if the lack of adequate colour cuts causes a bias towards steeper slopes. The advent of very deep spectroscopy in cluster fields should rapidly settle the issue of the faint-end slope of the cluster LF."
    },
    "0711/0711.1059_arXiv.txt": {
        "abstract": "We conduct a detailed analysis of the photometric redshift requirements for the proposed Dark Energy Survey (DES) using two sets of mock galaxy simulations and an artificial neural network code - ANNz. In particular, we examine how optical photometry in the DES $grizY$ bands can be complemented with near infra-red photometry from the planned VISTA Hemisphere Survey (VHS) in the $JHK_s$ bands. We find that the rms scatter on the photometric redshift estimate over $1<z<2$ is $\\sigma_z$=0.2 from DES alone and $\\sigma_z$=0.15 from DES+VISTA, i.e. an improvement of more than 30\\%. We draw attention to the effects of galaxy formation scenarios such as reddening on the photo-z estimate and using our neural network code, calculate the extinction, $A_v$ for these reddened galaxies. We also look at the impact of using different training sets when calculating photometric redshifts. In particular, we find that using the ongoing DEEP2 and VVDS-Deep spectroscopic surveys to calibrate photometric redshifts for DES, will prove effective. However we need to be aware of uncertainties in the photometric redshift bias that arise when using different training sets as these will translate into errors in the dark energy equation of state parameter, $w$. Furthermore, we show that the neural network error estimate on the photometric redshift may be used to remove outliers from our samples before any kind of cosmological analysis, in particular for large-scale structure experiments. By removing all galaxies with a neural network photo-z error estimate of greater than 0.1 from our DES+VHS sample, we can constrain the galaxy power spectrum out to a redshift of 2 and reduce the fractional error on this power spectrum by $\\sim$15-20\\% compared to using the entire catalogue.  Output tables of spectroscopic redshift versus photometric redshift used to produce the results in this paper can be found at $www.star.ucl.ac.uk/\\sim mbanerji/DESdata$. ",
        "introduction": "It is now widely accepted that dark energy is responsible for driving the observed acceleration of the Universe. In recent years, measuring and constraining the nature of this dark energy has become a central focus of current studies in cosmology. Several different methods have been developed and shown to probe the nature of dark energy through its effects on the geometry and structure of the Universe e.g. \\citet{Riess:98,Hu:WL,Blake:03}. Large-scale sky surveys such as the Sloan Digital Sky Survey have no doubt aided this kind of study \\citep{Tegmark:SDSS,Spergel:WMAP3} and more such galaxy surveys are now being planned to exploit the different techniques that will help us better understand the nature of dark energy.  The proposed Dark Energy Survey is one such experiment. It will use four independent probes namely galaxy clusters, galaxy power spectrum measurements, weak lensing studies and a supernova survey to constrain the nature of dark energy. Each of these methods relies on accurate distance measurements extending over cosmological scales. Given the wealth of data that will be available to us from such surveys, measuring distances and redshifts for all the objects using spectroscopic methods clearly becomes unfeasible. Hence the need for photometric redshifts.  Photometric redshift estimation methods have been around since the 1960s but have undergone a recent revival with proposals for a new generation of large-scale photometric surveys such as the Dark Energy Survey. New algorithms have been developed for photo-z estimation and made available to the community e.g. \\citet{Collister:ANNZ, Bolzonella:HyperZ, Benitez:BPZ, Feldmann:ZEBRA, Babbedge:IMPZ}. Furthermore, there is currently a lot of emphasis on optimising the depth and number of bands that will be used to image galaxies in future galaxy surveys so as to obtain accurate photometric redshifts. It is widely known that imaging in more bands can help reduce errors on photometric redshifts but the costs of adding more filters to planned surveys are significant. Given the plethora of data that is now becoming available to us covering the whole range of the electromagnetic bands and a large portion of the observable sky, it is vital that we explore the overlap between different surveys and fully exploit the data sets available to us, in order to achieve the best compromise between cost and science.  In this paper we analyse the prospect of combining optical data from the Dark Energy Survey (DES) with near infra-red data from the Vista Hemisphere Survey (VHS) in order to obtain accurate photometric redshifts that will help us better constrain the nature of dark energy. We begin with a brief description of these two proposed surveys and a description of the method used to generate mock galaxy samples for both of these surveys. We then proceed to a full photometric redshift analysis of simulated data from these two surveys using artificial neural networks. We assess the impact of reddening on our photometric redshift estimate as well as the effects of removing outliers and using different training sets. In each case, we present results obtained for the optical data from DES only and for optical and near-infra red data from DES and VHS. Finally, we look at the implications of our results for cosmological constraints on dark energy. In particular, we concentrate on the impact of photometric redshift errors on constraints on dark energy using galaxy power spectrum measurements. All magnitudes quoted in this paper are in the AB system. ",
        "conclusions": "In this work we have shown the role of near infra-red photometry from the VISTA Hemisphere Survey in constraining photometric redshifts for the Dark Energy Survey. We have examined the effects of galaxy reddening and training sets on the photo-z estimate. We have also studied the biases in the photometric redshift estimate when using training and testing sets with different sizes and levels of completeness and quantified the error in the dark energy equation of state parameter, $w$ that arises from differences in these biases.  A method of clipping our galaxy catalogues by removing outliers based on the ANNz error estimate on the photometric redshift has been presented. By applying a threshold error to our catalogues and rejecting all objects with errors bigger than this threshold error, we can effectively reduce the overall scatter on the photometric redshifts of our sample.  Finally, we have conducted a full galaxy power spectrum analysis using our DES and DES+VHS catalogues and looked at how our clipping method can improve the uncertainties on our galaxy power spectrum measurements. We find that there is an optimum threshold error at which we should clip our catalogues and this error depends on the catalogue being used and the redshift range in which we are evaluating the power spectrum. If we use a high value for the threshold error, the scatter on our photometric redshift estimate is high leading to large positional uncertainties and therefore large errors in the power spectrum due to cosmic variance. However, if we adopt a very low value for our threshold error, we remove most of the galaxies from our sample before calculating the power spectrum and the resulting uncertainties in the power spectrum are dominated by shot noise. We find that the optimum threshold error is smaller for the DES+VHS catalogues compared to the DES only catalogues in the same redshift range and hence more outliers are removed from this sample before analysis. Adding the VHS NIR data thus helps us to compute the galaxy power spectrum more accurately out to higher redshifts than for the DES only case.  In summary, our main conclusions are:  \\begin{itemize}  \\item{NIR data from VISTA VHS helps to reduce the scatter on DES photometric redshifts by $\\sim$30\\% for $z>1$.}  \\item{Reddening the galaxies can increase the photo-z scatter of DES by $\\sim$ 30\\% in some redshift ranges due to the degeneracy between redshift and reddening that exists in these redshift ranges. However, this is unlikely to be a major issue as most of our mock DES galaxies do not suffer from heavy extinction.}  \\item{ANNz can be used to predict the extinction, $A_v$ of DES galaxies to an accuracy of 0.27 and to classify them into six spectral types - E, Sbc, Scd, Im, SB2 and SB3.}  \\item{The VVDS-Deep and DEEP2 spectroscopic surveys, when finished will provide a very complete training set for DES out to a redshift of 2.}  \\item{Using different numbers of training set galaxies can lead to a difference in the photometric redshift bias of the order of $10^{-3}$.}  \\item{If we have an incomplete training set, we can improve the photometric redshift estimate by imposing the same colour cuts on the testing set as are applied to the training set. When this is done, the improvement in the photometric redshift bias is of the order of 10\\% compared to if we used a complete testing set with no imposed colour cuts.}  \\item{The clipping method introduced by \\citet{Abdalla:DUNEphotoz} can be effectively applied to the DES+VHS sample and applying a threshold error of 0.1 at which to cut our sample, reduces the scatter on the photometric redshift by $\\sim$50\\% by removing $\\sim20$\\% of the galaxies.}  \\item{A clipping threshold of 0.1 is optimal for calculating the DES power spectrum out to a redshift of 1.4. Applying this clipping threshold reduces $\\delta P/P$ by $\\sim$20\\%. When the VHS NIR data is added to the DES sample, the optimal clipping threshold in the same redshift range is 0.05 and this reduces the fractional error in the power spectrum by $\\sim$20-30\\%. In order to calculate the power spectrum out to a redshift of 2, the addition of VHS NIR data is crucial. In this redshift range, applying the optimal clipping threshold of 0.1 results in an improvement in $\\delta P/P$ by $\\sim$15-20\\%.}  \\item{Systematic errors on the photometric redshift bias arising from cosmic variance lead to uncertainties in the dark energy equation of state parameter, $w$, of about a percent.}  \\item{However if we use an incomplete training set to determine photometric redshifts on a testing set that hasn't been cut to match the training set, the resulting uncertainties in the photometric redshift bias can lead to errors in $w$ of the order of $\\sim$10\\% if we keep all other cosmological parameters as constant. Note however that this result is model specific and depends to a certain extent on our choice of mock catalogues and the algorithm used to calculate photometric redshifts, in this case, the neural networks.}  \\end{itemize}  In the absence of large spectroscopic surveys, the DES and VHS datasets, when combined, will prove extremely effective in constraining dark energy through large scale structure signals like baryon acoustic oscillations. By clipping photometric redshift catalogues and carefully removing a suitable number of outliers, one can achieve reasonably precise measurements of the galaxy power spectrum out to a redshift of 2."
    },
    "0711/0711.3789_arXiv.txt": {
        "abstract": "Using numerical simulations at moderate magnetic Reynolds numbers up to 220 it is shown that in the kinematic regime, isotropic helical turbulence leads to an alpha effect and a turbulent diffusivity whose values are independent of the magnetic Reynolds number, $\\Rm$, provided $\\Rm$ exceeds unity. These turbulent coefficients are also consistent with expectations from the first order smoothing approximation. For small values of $\\Rm$, alpha and turbulent diffusivity are proportional to $\\Rm$. Over finite time intervals meaningful values of alpha and turbulent diffusivity can be obtained even when there is small-scale dynamo action that produces strong magnetic fluctuations. This suggests that small-scale dynamo-generated fields do not make a correlated contribution to the mean electromotive force. ",
        "introduction": "The generation and maintenance of large-scale magnetic fields in stars and galaxies is often studied within the framework of the mean-field dynamo (MFD); see, e.g., \\cite{Mof78,Par79,KR80}. A particularly important driver of MFDs is the $\\alpha$-effect. For isotropic turbulence and weak magnetic fields, i.e.\\ in the kinematic regime, the $\\alpha$-effect can be expressed purely in terms of the kinetic helicity. Research in recent years has mostly been concerned with clarifying the effects of nonlinearity, but there are serious uncertainties even in the linear (kinematic) regime. In particular, whether or not $\\alpha$ can then be expressed in terms of the kinetic helicity depends on the applicability of the first order smoothing approximation (FOSA) or other closures used to calculate $\\alpha$. Such approaches become questionable when the magnetic Reynolds number, $\\Rm$, is large, i.e.\\ when the magnetic diffusion time is long compared with the turnover time which, in turn, is comparable with the correlation time of the turbulence. In the high conductivity limit FOSA can only be applied if the correlation time for the velocity field is much smaller than the eddy turnover time. This is not the case for high Reynolds number turbulence, where the two time scales are equal, i.e.\\ the Strouhal number is unity \\citep{BS05b,BS07}, so FOSA should in principle break down. In this case all higher order terms need to be taken into account \\citep{Kno78}. Furthermore, high $\\Rm$ random flows typically lead to a fluctuation dynamo which leads to rapidly growing small-scale magnetic fields independent of the mean field. This also breaks the assumption made by FOSA that fluctuating fields are much smaller than the mean field.  The existence of $\\alpha$-effect and turbulent diffusion has been worrying dynamo researchers over several decades. Assuming steady flow patterns, \\cite{Chi79} found that motions that concentrate magnetic fields into thin flux sheets lead to an $\\alpha$-effect whose value diminishes with $\\Rm$ like $\\Rm^{-1/2}$. But for an analogous helical motion which concentrates the field into an axial flux rope, $\\alpha$ tends to a finite limit as $\\Rm \\to \\infty$. He conjectured that the latter estimate may be typical of steady three-dimensional motion.  The validity of turbulent diffusion has been questioned by \\cite{Pid81}. Calculations by \\cite{Kraichnan76} suggest that $\\alpha$ and turbulent diffusion converge to finite values for statistically isotropic velocity fields with gaussian statistics. However, numerical simulations \\citep{Drummond86} using a frozen velocity field suggest that in the limit of large magnetic Reynolds numbers $\\alpha$ tends to zero. Based on specific imposed (kinematic) flow patterns it has been suggested that there is no simple relation between $\\alpha$ and helicity of the flow; see \\cite{CHT06}. In fact, their results may suggest that in the kinematic regime, $\\alpha$ exhibits a strong $\\Rm$ dependence (for large $\\Rm$ up to $2\\times10^5$) and could change sign for $\\Rm\\approx20$ and long correlation times. It is important to emphasize that in the nonlinear regime, i.e.\\ for finite magnetic field strength and including the Lorentz force, a strong $\\Rm$ dependence is now indeed well established \\citep{CH96,B01}; see \\cite{BS05a} for a review. However, in the following we shall be concerned with the purely linear regime.  In order to clarify the $\\Rm$ dependence in the kinematic regime we perform numerical turbulence experiments where we adopt an externally imposed body force to drive the flow. This is a common technique applied in simulations and helps develop homogeneous and isotropic turbulence which is easier to handle analytically. While the calculation of $\\alpha$ from a turbulence simulation is relatively straightforward by imposing a uniform magnetic field, the calculation of turbulent diffusion is more uncertain. One possibility is to determine the decay rate of an initial large-scale magnetic field \\citep{YBR03}. Another more reliable method is to calculate both $\\alpha$ and turbulent diffusion tensors simultaneously by computing the mean electromotive force using a number of different nonuniform test fields in different directions and of different spatial structure \\citep{S05,S07}. This is also the approach used in the present paper. ",
        "conclusions": "The present work has shown that for isotropic turbulence the first order smoothing results give quite accurate expressions for the $\\alpha$ effect and the turbulent magnetic diffusivity in the kinematic regime. This result comes almost as a surprise given that in recent years mean field theory has been seriously challenged based on numerical simulations. However, we can now clearly say that for moderate Reynolds numbers up to about 220 and under the conditions stated (scale separation, isotropy, etc.) there is {\\it no} evidence that the kinematic results obtained using FOSA are flawed, even though its applicability can then no longer be guaranteed and we know of its shortcomings in the nonlinear regime \\citep{BS05a}. As explained above, a possible reason for this might be that the predictions of FOSA and MTA are rather similar, even though FOSA loses its justification whilst MTA hinges on a not well justified closure hypothesis.  The emergence of small-scale dynamo action for supercritical values of $\\Rm$ of about 30 does not seem to affect the average values of $\\alpha$ and $\\eta_{\\rm t}$. This suggests that the exponentially growing part of the small-scale field does not make a contribution to the mean emf $\\meanemf$, correlated with the imposed test fields. However, small-scale dynamo action makes the calculation of reliable average values of $\\alpha$ and $\\eta_{\\rm t}$ more difficult. Of course, in the supercritical case the long time limit of any kinematic problem becomes unphysical. Nevertheless, within a certain time interval the averaged values of $\\alpha$ and $\\eta_{\\rm t}$ match reasonably with theoretical expectations of FOSA."
    },
    "0711/0711.2166_arXiv.txt": {
        "abstract": "The generally adopted value for the solar abundance of indium is over six times higher than the meteoritic value. We address this discrepancy through numerical synthesis of the 451.13~nm line on which all indium abundance studies are based, both for the quiet-sun and the sunspot umbra spectrum, employing standard atmosphere models and accounting for hyperfine structure and Zeeman splitting in detail. The results, as well as a re-appraisal of indium nucleosynthesis, suggest that the solar indium abundance is close to the meteoritic value, and that some unidentified ion line causes the 451.13~nm feature in the quiet-sun spectrum. ",
        "introduction": "\\label{sec:intro}  The solar abundance of indium is controversial because its generally accepted value significantly exceeds the meteoritic value. At a factor of six difference, this remains the largest unexplained discrepancy between meteoritic and solar abundance values. In this paper we address this problem by considering the nucleosynthesis of indium and through indium line synthesis for the quiet solar photosphere and sunspot umbrae including hyperfine structure in detail.  The meteoritic indium abundance is \\mbox{$\\Am = 0.80 \\pm 0.03$} (\\citealt{2003ApJ...591.1220L} and references therein) where $A_{\\rm In} \\equiv \\log(n_{\\rm In}/n_{\\rm H}) + 12$ with $n_{\\rm In}$ and $n_{\\rm H}$ the indium and hydrogen particle densities, respectively. Table~\\ref{tab:abundances} summarises the determinations of the solar indium abundance in the literature. All measurements are based on a single, very weak feature in the quiet-sun spectrum at $\\lambda = 451.1307$~nm which is commonly identified as one of the \\InI\\ resonance lines. The initial result of \\citet{1960ApJS....5....1G} was based on an erroneous oscillator strength. The other three determinations scatter around $\\AInS = 1.6$, the value listed in the compilation of \\citet{2005ASPC..336...25A}. The 0.8~dex discrepancy with the meteoritic value cannot be explained by the usual uncertainties of abundance determination such as line strength measurement, imprecise atomic data and solar modeling deficiencies.  \\begin{table} \\caption{$\\AInS$ determinations} \\label{tab:abundances} \\centering \\begin{tabular}{@{}lccc} \\hline Authors                      &  $\\AInS$ & Specified error \\\\ \\hline \\citet{1960ApJS....5....1G}  & 1.16       &       \\\\ \\citet{1969MNRAS.142...71L}  & 1.71       &       \\\\ \\citet{1998SSRv...85..161G}  & 1.66       & 0.15  \\\\ \\citet{2002SoPh..211....3B}  & 1.56       & 0.20  \\\\ \\hline \\end{tabular} \\end{table}  The origin of elements heavier than \\hbox{Fe} is mostly attributed to neutron-capture processes \\citep[see][for a review]{1994ARA&A..32..153M}. $Slow$ neutron capture (the $s$ process) occurs for relatively low neutron densities ($\\simeq 10^7 $~cm$^{-3}$), while $rapid$ neutron capture (the $r$ process) occurs for relatively high neutron densities ($> 10^{20}$~cm$^{-3}$). \\citet{1978ApJ...219..307A} analysed the Sn/In abundance ratio. He found that `no combination of $r$- or $s$-process products even remotely resembling those which generally predict the solar system abundances very successfully can give Sn/In as low as 1.4' (which results from taking $\\AInS = 1.71$).  Recently, \\citet{2006MNRAS.370L..90G} suggested that because of its low condensation temperature (536~K, \\citealt{2003ApJ...591.1220L}), indium may have been depleted in chondritic meteorites to an abundance much smaller than the solar one. However, we note that, e.g., thallium has a similar condensation temperature (532~K, \\citealt{2003ApJ...591.1220L}) and a similar meteoritic abundance ($0.78 \\pm 0.04$) whereas its well-determined photospheric abundance ($0.9 \\pm 0.2.$) is only slightly higher than the meteoritic one. In a sequel paper \\citet{2006MNRAS.371..781G} studied the indium abundance in a sample of 42 sun-like stars of which five are known to host planets. He found a strong negative correlation between the [\\hbox{In}/\\hbox{Fe}] and [\\hbox{Fe}/\\hbox{H}] logarithmic abundance ratios. This trend is much steeper than the comparable relation for europium, which is a pure $r$-process element. However, one would expect a less steep trend for indium because it received contributions both from the $s$- and the $r$-processes.  In this paper we once again scrutinize the solar indium abundance, paying specific attention to indium nucleosynthesis (Sect.~\\ref{sec:ratio}), indium line identification and appearance in the solar spectrum (Sect.~\\ref{sec:solarindium}), and indium line synthesis accounting for hyperfine structure both for quiet sun and sunspot umbrae (Sect.~\\ref{sec:synthesis}). The conclusion is that, after all, the solar indium abundance is likely to be close to the meteoritic value. ",
        "conclusions": "\\label{sec:discussion}  Our update in Sect.~\\ref{sec:ratio} of the analysis of \\citet{1978ApJ...219..307A} did not change the conclusion that there is no combination of nucleosynthetic processes that can produce a Sn/In ratio as obtained using the most recent solar values, while the chondritic ratio can be easily accounted for. This is a strong indication that the solar abundance of indium cannot be much higher than the meteoritic value. This also suggested by the absence of the \\InII\\ resonance line at 158.6~nm.  Furthermore, Fig.~\\ref{fig:synthetic_profile} shows that standard line synthesis including hyperfine structure does not fit the observed profile for either the high or the low abundance value. In contrast, its appearance in the sunspot spectrum is well reproduced, including hyperfine Zeeman splitting, using the low meteoritic abundance value (Fig.~\\ref{fig:sunspot_profile}) and assuming a reasonable straylight fraction (60 per cent) and umbral field strength of 3000~G.  These results together suggest that the line at 451.13~nm in the quiet-sun spectrum is not due to indium but to some other species. It should be from an ion with low ionisation energy of the neutral stage since the line disappears in the sunspot spectrum (cf.\\ Fig.~\\ref{fig:iron2}). The observation by \\citet{2006MNRAS.371..781G} of significant trends of the assigned indium abundance with $T_{\\rm{eff}}$ and [Fe/H] in a sample of stars cooler than the sun suggests that the unidentified transition should have high excitation energy. The best possible candidate in the list of \\citet{2006MNRAS.371..781G} would be the \\SmII\\ line, but as noted in Sect.~\\ref{sec:solarindium} this is actually a resonance transition of \\SmI.  Finally, an alternative explanation of the \\InI~451.13~nm line formation might be NLTE effects. In particular, the appearance of the 451.13~nm line as emission line in the spectra of long-period variables during the downward modulation phase was explained already by \\citet{1937ApJ....86..499T} as optical pumping between H$\\delta$ and a \\InI\\ line at 410.1~nm which shares its upper level with \\InI~451.13~nm. However, we this mechanism does not operate at all in the deep solar photosphere where the formation of the hydrogen Balmer lines is close to LTE (cf.\\ Fig.~30 of \\citealt{1981ApJS...45..635V}).  In conclusion, we suggest that the solar indium abundance is close to the meteoritic value as evidenced in the sunspot spectrum, whereas the feature at 451.13~nm in the quiet-sun spectrum remains unidentified and is likely to be an ion line at high excitation energy from a species with low first ionisation energy."
    },
    "0711/0711.0213_arXiv.txt": {
        "abstract": "High-resolution X-ray observations have revealed cavities and `cold fronts' with sharp edges in temperature and density within galaxy clusters.  Their presence poses a puzzle since these features are not expected to be hydrodynamically stable, or to remain sharp in the presence of diffusion.  However, a moving core or bubble in even a very weakly magnetized plasma necessarily sweeps up enough magnetic  field to build up a dynamically important sheath; the layer's strength is set by a competition between `plowing up' and slipping around of field lines, and depends primarily on the ram pressure seen by the moving object.  In this inherently three dimensional problem, our analytic arguments and numerical experiments show that this layer modifies the dynamics of a plunging core, greatly modifying the hydrodynamic instabilities and mixing, changing the geometry of stripped material, and slowing the core through magnetic tension.  We derive an expression for the maximum magnetic field strength and thickness of the layer, as well as for the opening angle of the magnetic wake. The morphology of the magnetic draping layer implies the suppression of thermal conduction across the layer, thus conserving strong temperature gradients.  The intermittent amplification of the magnetic field as well as the injection of MHD turbulence in the wake of the core is identified to be due to vorticity generation within the  magnetic draping layer. These results have important consequences for understanding the complex gasdynamical processes of the intra-cluster medium, and apply quite generally to motions through other magnetized environments, \\eg{} the ISM. ",
        "introduction": "\\label{sec:intro} \\input paper1intro.tex ",
        "conclusions": "\\label{sec:conclusion}   This work aims at understanding the morphology and the dynamical properties of magnetic draping to set a solid ground for its astrophysical applications.  A core, bullet, or bubble that moves super-alfv{\\'e}nically in even a very weakly magnetized plasma necessarily sweeps up enough magnetic field to build up a dynamically important sheath around the object; the layer's strength is set by a competition between `plowing up' of field and field lines slipping around the core, and to first order depends only on the ram pressure seen by the moving object.  This layer is developed very quickly, potentially faster than a crossing time of the projectile. The energy density in the draped layer, at its maximum, exceeds the ram pressure by a factor of two, necessary to anisotropically redirect the flow.  This effect has important implications for galaxy cluster physics as it suppresses hydrodynamic instabilities at the interface of AGN bubbles. It naturally explains so-called `cold fronts' by keeping temperature and density interfaces of merging cores sharp that would otherwise be smoothed out by thermal conduction and diffusion. Other important astrophysical implications of this effect include draping of the solar wind magnetic field at the bow shock of the Earth as well as pulsar wind nebulae.  In this paper, we compare a simplified analytical solution of the problem that neglects the back-reaction of the dynamically important magnetic field on the potential flow with a high-resolution AMR simulation and find very good agreement between both solutions in the region ahead of the bubble.  Non-linear back-reaction of the magnetic field in the draping layer necessarily implies the generation of vorticity in the flow. The induced vortices in the wake deform the core hydrodynamically and eventually cause the magnetic sheath to peel off. There is a strong indication that the vorticity generation is responsible for the intermittent amplification and stretching of the magnetic fields as well as the injection of MHD turbulence in the wake of the core. If this withstands further critical analysis, this mechanism might have profound astrophysical implications for the amplification and generation of large-scale magnetic fields in the inter-galactic medium.  The magnetic layer, once fully developed, has a characteristic geometry which we have shown here to be roughly conical in the plane along the magnetic field lines, with opening angle $\\theta \\sim \\arctan(\\Alfvenvelmag/u)$, and remains contained in the perpendicular plane, with the Kelvin-Helmholtz instability acting on the object interface.  The magnetic tension in the layer significantly decelerates the object, dominating over any hydrodynamic drag.  Over long times, the anisotropy imposed by the field -- despite the fact that the field is initially highly subthermal and with an energy density significantly less than the kinetic energy of the ambient medium in the frame as the projectile -- can significantly distort the projectile, keeping it significantly more constrained in the plane perpendicular to the ambient field, and keeping any stripped material inside the drape.  If a magnetic draping layer such as generated in these simulations were astrophysically observable, it would be possible to get independent measurements of the magnetic field strength in the ambient medium provided the local gas density and the velocity of the moving object is known. The tools are both the opening angle of the drape and the thickness of the magnetic field layer -- but not through the layer's field strength.  Similarly, for a known magnetic field strength, we would have an alternate measure of the velocity of the projectile.    \\bigskip"
    },
    "0711/0711.1812_arXiv.txt": {
        "abstract": "A fully self-consistent model of the neutron star inner crust based upon models of the nucleonic equation of state at zero temperature is constructed. The results nearly match those of previous calculations of the inner crust given the same input equation of state. The extent to which the uncertainties in the symmetry energy, the compressibility, and the equation of state of low-density neutron matter affect the composition of the crust are examined. The composition and pressure of the crust is sensitive to the description of low-density neutron matter and the nuclear symmetry energy, and the latter dependence is non-monotonic, giving larger nuclei for moderate symmetry energies and smaller nuclei for more extreme symmetry energies. Future nuclear experiments may help constrain the crust and future astrophysical observations may constrain the nuclear physics input. ",
        "introduction": "The inner crust of a cold neutron star can be defined as the region between the density where neutrons drip out of nuclei (about 4 $\\times 10^{11}$ g/cm$^{3}$) and the density for the transition to homogeneous nucleonic matter at about half of the nuclear saturation density. This region is sensitive to the nuclear physics input because the nature of the crust is determined by the structure of neutron-rich nuclei and the energetics of the surrounding dripped neutrons. In this work, the dependence of the description of the neutron-rich nuclei and the dripped neutrons on the equation of state (EOS) of homogenous nucleonic matter is examined.  The inner crust is of broad interest because a large variety of astrophysical observations are dependent on and sensitive to the properties of the neutron star crust. One recent motivation is the suggestion that the giant flares in Soft Gamma-Ray repeaters trigger seismic events in the neutron star crust and are sensitive to the shear modulus of the crust crust~\\cite{Thompson95,Strohmayer06,Samuelsson07}. The shear modulus, in turn, is sensitive to the composition of the neutron star crust and the relative magnitude of the proton and neutron numbers of the nuclei in the inner crust. The moment of inertia of rotating neutron stars is also sensitive to the inner crust and depends on the transition density between the crust and the core~\\cite{Link99}.  Neutrino and photon opacities are also sensitive to the properties of the nuclei in the inner crust. For example, neutrino-nucleus scattering, which scales like $A^2$, is the most important neutrino process during the lepton-trapped phase of a Type II supernova (see Ref.~\\cite{Burrows06} for a recent review). Finally, the cooling and evolution of neutron star crusts depends on the both the size of the crust~\\cite{Lattimer94} and by its transport properties \\cite{Potekhin97,Brown00}, which are both related to the composition. These astrophysical connections motivate the study of the magnitude of the uncertainty of the properties of the inner crust which come from present uncertainties in the nuclear physics inputs.  In this article, several models of the neutron star inner crust are constructed systematically using inputs from with the current experimental information while allowing the range of uncertainty allowed due to the uncertainty in the EOS of homogeneous matter. Of particular importance, is that the symmetry energy is varied in {\\em both} the description of the nuclei and the description of the neutron matter at the same time. The composition depends on the symmetry energy, but is nearly independent of the compressibility. This means that astrophysical observations which are connected to properties in the crust can constrain the nuclear symmetry energy.  The inner crust is quite sensitive to the EOS of neutron matter at sub-saturation densities. At sufficiently low densities, neutron matter is somewhat well understood because three-body interactions are small, and the two-body neutron-neutron interaction is strongly constrained by the experimentally measured neutron-neutron scattering phase shifts~\\cite{Carlson03} (see also the review in Ref.~\\cite{Heiselberg00}). Many of the currently available EOSs, however, do not respect this understanding of low-density neutron matter because they are fit to the properties of nuclei which are more sensitive to matter near saturation densities.  The neutron matter EOSs used in this work are designed to have a realistic behavior below the saturation density, within the precision required for the description of the crust. ",
        "conclusions": "The composition of the neutron star crust is still partially unknown, due to uncertainties in the nuclear mass formula and the equation of state. The composition (and to a lesser extent, the overall pressure) is quite sensitive to the equation of state of low-density neutron matter, and the nuclear symmetry energy, both its magnitude and its density dependence. The dependence of the composition on the symmetry energy is not monotonic, as models with moderate symmetry energies can have larger nuclei than models with lower or higher symmetry energies. To the extent to which neutron stars depend on the composition, this means that it is important to explore the full range of variation in the crust allowed by the present knowledge of the input nuclear physics, while ensuring that the EOS is constrained by what is already known about the EOS of low-density neutron matter. Nuclear experiments will continue to provide better constraints on the symmetry energy, including from the PREX experiment~\\cite{Michaels00,Horowitz01b} to measure the neutron skin thickness of lead at Jefferson Lab and from intermediate-energy heavy-ion collisions as has been done in Ref.~\\cite{Tsang04}.  It remains to be seen if these results persist in the more microscopic models which include pairing, corrections beyond the Wigner-Seitz approximation, long-range correlations, and better treatments of the nuclear structure."
    },
    "0711/0711.3210_arXiv.txt": {
        "abstract": " ",
        "introduction": "Cosmological constant problem is the deepest problem of fundamental physics. To date, the attempts to explain a small cosmological constant in effective field theory coupled to General Relativity using dynamical adjustment have not yielded an answer \\cite{Wein}. The perspective may change if our universe is a brane in extra dimensions \\cite{rusha}-\\cite{selftun}. In such setups one may divert the brane vacuum energy into the extrinsic curvature \\cite{selftun}, so that the vacuum energy could remain invisible to long distance $4D$ gravity, yielding a low curvature $4D$ geometry. However, while recovering exact $4D$ General Relativity at large distances, by using additional branes in the bulk one rediscovers a fine-tuning similar to the standard $4D$ one \\cite{selftun,nillest}. A program of \\cite{msledcc} has been initiated to seek the means to relax these fine-tunings in codimension-2 brane setups. Interesting precursors to these ideas however can be found already in \\cite{markus}.  One may try to alleviate the problem by only approximating $4D$ General Relativity at large distances, as in the brane induced gravity (BIG {\\it a.k.a.} DGP) \\cite{DGP}. There the extra-dimensional space may have an infinite volume, so that the $4D$ graviton completely decouples. Thus $4D$ gravity must emerge from the exchange of the continuum of bulk modes. This may occur with the help of the induced curvature terms on the brane \\cite{DGP}, which control the momentum transfer due to the scattering of virtual bulk gravitons and may yield a force which scales as $1/r^{2}$, simulating $4D$ behavior \\cite{DGP}. The exact gravitational shock waves of \\cite{shocks} confirm this nicely.  These mechanisms may operate on higher-codimension defects \\cite{DGP,highercod}. In \\cite{kkland}, with D. Kiley we explored in a very detailed fashion the BIG mechanism on general codimension-2 branes floating in $6D$.  We presented exact background solutions with flat $4D$ geometries for nonvanishing tensions, and regulating the short distance properties by considering hollow cylindrical branes we calculated the various regimes of gravity. The most interesting case, where the theory displayed a remarkably good perturbative behavior, was for near-critical tensions, where the deficit angle around the brane was almost $2\\pi$. In this limit the bulk compactifies  to a very thin sliver \\cite{kalwall}, inside which bulk gravity is $5D$. Because of this gravity along the brane changes dimensions with distance according to $4D \\rightarrow 5D \\rightarrow 6D$. The crossover scale where gravity stops being $4D$ is independent of the tension, and can be quickly obtained by using the naive crossover formula on a codimension-1 brane, realized as a wrapped 4-brane in a $6D$ flat space with a compact circle: by Gauss formulas for Planck masses $M_{4 \\, eff}^2 = M_5^3 r_0$ and $M_{5 \\, eff}^3 = M_6^4 r_0$, we find exactly $r_c \\sim \\frac{M_{4 \\, eff}^2}{M_{5 \\, eff}^3} \\sim \\frac{M^3_5}{M^4_{6 }}$, which is our crossover formula from the exact shock waves, and which is the same as the see-saw scale $r_c \\sim \\frac{M^2_4}{M^4_{6}r_0}$ of \\cite{highercod}. Thus we see precisely {\\it how} the see-saw mechanism emerges. However, BIG is not free of the vacuum energy problem: although the brane is flat, when we fix $M_4$, we must finely tune the tension of the brane to get the bulk sliver geometry, that yields the separation of the $4D$ and $6D$ regimes by a $5D$ one.  From the shock waves we have seen that the theory contains $4D$ General Relativity. We unveil the rest in the linearized perturbation theory, which shows that there are scalar modes with gravitational couplings to matter. In the near-critical limit linearized perturbation theory around the vacuum remains under control in the $4D$ regime below the crossover scale. While the radion is suppressed, the helicity-0 mode mediates an extra long range force around the $4D$ vacuum backgrounds, and where the linearized theory is $4D$ it a Brans-Dicke gravity with $\\omega=0$. This would disagree with the classic tests of General Relativity \\cite{vdvz}, but perhaps non-linearities or dynamics on vacua of more complex structure could come to the rescue. As the tension decreases towards sub-critical values the scalars start to couple strongly, and the perturbativity of the $4D$ regime is lost exactly at the $4D$ crossover scale $r_c$, which is the Vainshtein scale \\cite{veinsh} of the vacuum itself. This happens since the scalar gravitons probe the regulating axion flux, which sources brane bending. On the other hand, we do not find any instabilities in the leading order of perturbation theory.  In this note, we will review the near-critical limits for the sake of clarity. The detailed analysis of the general sub-critical setups with flat $4D$ geometries can be found in \\cite{kkland}. We should note that very recently in \\cite{many} similar features of BIG theory in $6D$ environments were argued to emerge in a framework  with more branes. Also, some ideas along these lines were explored in a setup with codimension-3 branes with a solid deficit angle in \\cite{gia}. ",
        "conclusions": "BIG models may in fact approximate $4D$ worlds between the UV cutoff that resolves the core of the gravity-localizing defect and the crossover scale, if their Scherk-Schwarz sector is carefully tuned to yield $4D$ Minkowski vacua. We have seen this explicitly in the example of regulated codimension-2 defects, which remain perturbatively calculable and have a $4D$ regime for near-critical tensions. We feel that similar conclusions may extend to any setups with codimension $\\ge 2$. The gravitational sector contains additional scalars, but at least at the linearized order there are no ghosts, of the kind that plague the codimension-1 DGP models \\cite{Luty,koyama}-\\cite{minjoon}, or were argued to arise in soft gravity resolutions of higher-codimension cases \\cite{durub}.  The cosmological constant problem changes completely in BIG codimension-2 setups. Putting brane localized terms and changing gravity in the IR by itself does not solve it. Instead, one still has vacuum energy problem, since the tension of the brane needs to be tuned to get a desired low energy theory. But now, vacuum energy does not need to curve the $4D$ geometry at all. Instead, we must tune the vacuum energy to get a weak $4D$ gravity over a large spatial region and maintain perturbativity in the scalar graviton sector. This might offer a different strategy to help with the vacuum energy problem: perhaps the fact that graviscalars go into the strong coupling for detuned tensions can be used as a new tool. One might try to look, for example, for theories where the tuned value of the tension, which ensures perturbativity, marks an IR fixed point of the combined matter and graviscalar theory. In this case, at least some of the premises of the Weinberg's no-go theorem for cosmological constant adjustment \\cite{Wein} would be violated, and it would be interesting to see what the consequences are.  The numerology of the near-critical codimension-2 BIG is rather intriguing. Using the formula for the $4D \\rightarrow 5D$ crossover scale $r_c = M_5^3/M_6^4$, we see that to get $r_c$ to be of the order of the current horizon size, $H_0^{-1} \\sim 10^{28} {\\rm cm}$, and ensure that matter on the brane looks pointlike at energies below a ${\\rm TeV}$, we need $M_6 \\ga {\\rm TeV}$, and therefore $M_5 \\la 10^{19} {\\rm GeV}$. Further to reproduce the $4D$ Newton's constant below the crossover scale, one needs to pick $r_0 \\ga 10^{-19} {\\rm GeV}^{-1}$. So we must tune the theory to make sure that the brane is very thin, with many hidden sector fields, to get such hierarchies. We don't have much to add to the discussion of how to do it, but merely note that this is in line with the current philosophy of the BIG models \\cite{DGP}. While some work along the lines of embedding the theory into string theory, that would provide the framework for its full UV completion, has been pursued \\cite{kiritsis,lowe,ignatios}, the existing constructions are really semiclassical models where gravity is treated classically. At least, we can explore the low energy consequences of such models.  Many questions remain open. Can the helicity-0 modes decouple? Are there dangerous instabilities beyond the linear order? Can BIG theory be UV-completed? Most interestingly, could the new guise of the cosmological constant problem open some new directions to circumvent the Weinberg no-go theorem \\cite{Wein}? While we can't address these issues here, we hope to return to them soon.  \\vskip.4cm  {\\bf \\noindent Acknowledgements}  \\vskip.2cm  We would like to thank Derrick Kiley for collaboration and Alberto Iglesias, Robert Myers, Minjoon Park, Oriol Pujolas, Lorenzo Sorbo and John Terning for interesting discussions. N.K. is grateful to Galileo Galilei Institute, Florence, Italy, for kind hospitality in the course of this work, and to the organizers of the ``Sowers Workshop\", Virginia Tech, May 14-18, 2007, ``Cosmology and Strings\" workshop at ICTP, Trieste, Italy, July 9-13, 2007, ``Dark Energy In the Universe\", Hakone, Japan Sep 1-4, 2007 \\cite{hakone} and ``Zagreb Workshop 2007\", Zagreb, Croatia, Nov 9-11, 2007, for the opportunity to present the results and the kind hospitality. This work has been supported in part by the DOE Grant DE-FG03-91ER40674, in part by the NSF Grant PHY-0332258 and in part by a Research Innovation Award from the Research Corporation."
    },
    "0711/0711.1165_arXiv.txt": {
        "abstract": "We present techniques for long-term, stable, and accurate evolutions of multiple-black-hole spacetimes using the `moving puncture' approach with fourth- and eighth-order finite difference stencils. We use these techniques to explore configurations of three black holes in a hierarchical system consisting of a third black hole approaching a quasi-circular black-hole binary, and find that, depending on the size of the binary, the resulting encounter may lead to a prompt merger of all three black holes, production of a highly elliptical binary (with the third black hole remaining unbound), or disruption of the binary (leading to three free black holes).  We also analyze the classical Burrau three-body problem using full numerical evolutions. In both cases, we find behaviors distinctly different from Newtonian predictions, which has important implications for N-body black-hole simulations.  For our simulations we use approximate analytic initial data.  We find that the eighth-order stencils significantly reduce the numerical errors for our choice of grid sizes, and that the approximate initial data produces the expected waveforms for black-hole binaries with modest initial separations. ",
        "introduction": "The recent dramatic breakthroughs in the numerical techniques to evolve black-hole-binary spacetimes~\\cite{Pretorius:2005gq,Campanelli:2005dd,Baker:2005vv} has led to rapid advancements in our understanding of black-hole physics. Notable among these advancements are developments in mathematical relativity, including systems of PDEs and gauge choices~\\cite{Gundlach:2006tw, vanMeter:2006vi}, the exploration of the validity of the cosmic censorship conjecture~\\cite{Campanelli:2006uy,Sperhake:2007gu}, and the application of isolated horizon formulae~\\cite{Dreyer:2002mx,Schnetter:2006yt,Campanelli:2006fg, Campanelli:2006fy,Cook:2007wr,Krishnan:2007pu}.  There are many exciting new results on recoil velocities~\\cite{Baker:2006vn,Sopuerta:2006wj,Gonzalez:2006md, Sopuerta:2006et,Herrmann:2006cd,Herrmann:2007zz,Herrmann:2007ac, Campanelli:2007ew,Koppitz:2007ev,Choi:2007eu,Gonzalez:2007hi, Baker:2007gi,Campanelli:2007cga,Berti:2007fi,Tichy:2007hk, Herrmann:2007ex,Brugmann:2007zj,Schnittman:2007ij,Krishnan:2007pu, HolleyBockelmann:2007eh,Pollney:2007ss,Lousto:2007db}, post-Newtonian (PN) and numerical waveform comparisons~\\cite{Baker:2006ha,Pan:2007nw,Husa:2007rh,Hannam:2007ik}, modeling of the remnant spin~\\cite{Campanelli:2006uy, Campanelli:2006fg,Campanelli:2006fy,Herrmann:2007ex,Rezzolla:2007xa, Marronetti:2007wz,Rezzolla:2007rd}, new studies of eccentric black hole binaries \\cite{Pretorius:2007jn,Sperhake:2007gu,Hinder:2007qu} and producing waveforms for matched filtering~\\cite{Vaishnav:2006ci,Baker:2006kr, Ajith:2007qp,Vaishnav:2007nm,Buonanno:2007pf,Baker:2007fk, Ajith:2007kx}.  In particular, the recent discovery of very large merger recoil kicks for black-hole binaries with spins in the orbital plane, which was originally inferred from the results in~\\cite{Campanelli:2007ew}, then observed in~\\cite{Gonzalez:2007hi}, and determined to have a maximum value of $4000\\ \\KMS $ in~\\cite{Campanelli:2007cga},  has had a great impact in the astrophysical community, with several groups now seeking for observational traces of such high speed holes as the byproduct of galaxy collisions~\\cite{Bonning:2007vt,HolleyBockelmann:2007eh}.  Three-body and four-body interactions are expected to be common in globular clusters~\\cite{Miller:2002pg,Gultekin:2003xd} and in galactic cores hosting supermassive black holes (when stellar-mass-black-hole-binary systems interact with the supermassive black hole).  Hierarchical triplets of supermassive black holes might also be formed in galactic nuclei undergoing sequential mergers~\\cite{Makino:1990zy,valtonen96}.  The gravitational wave emission from such systems was recently estimated using post-Newtonian techniques~\\cite{Gultekin:2005fd}.  The recent discovery of a probable triple quasar~\\cite{Djorgovski:2007ka} (with estimated masses of $50\\times10^6 M_\\odot$, $100\\times10^6M_\\odot$ and $500\\times10^6 M_\\odot$ and projected separations of between $30$ and $50$ kpcs) indicates that hierarchical supermassive-three-black-hole systems are possible.  Triple stars and black holes are much more common in globular clusters~\\cite{Gultekin:2003xd}, and galactic disks. The closest star to the solar system, Alpha Centauri, is a triple system, as is Polaris and HD 188753.  On the theoretical side, even in the Newtonian theory of gravity, the three-body problem is much more complicated than the two body one, and is generically chaotic. In Ref.~\\cite{Campanelli:2007ea} we established that our numerical formalism is able to handle the evolution of three fully relativistic bodies.  In this paper we continue our quest to expand our understanding of the multiple-black-hole problem by studying a sequences of black-hole-binary---third-black-hole configurations, with the third black hole intersecting the binary along the axis of rotation, to explore the influence of a third black hole on the binary dynamics. We also examine the classical Burrau three-body problem and find that the dynamics can differ dramatically from the Newtonian prediction, with important consequences for N-body black-hole simulations.  This paper is organized as follows: in Sec.~\\ref{sec:tech} we describe the techniques used for multiple black-hole evolutions, including eighth-order techniques and approximate initial data. In Sec.~\\ref{sec:res} we show the results from multiple black-hole evolutions of three families of hierarchical configurations consisting of a third black hole interacting with a black-hole binary along the binary's axis of rotation, as well as the classical Burrau three-body problem. We conclude our analysis in Sec.~\\ref{sec:disc} where we discuss the implications of our numerical results to black-hole astrophysics and N-body simulations. We provide the techniques to generate 2PN quasi-circular orbits of our hierarchical three-black-hole systems in Appendix~\\ref{ap:3-body-H}. ",
        "conclusions": "\\label{sec:disc} Our fully nonlinear evolutions of three-black-hole systems demonstrate dissipative General Relativistic effects due to emission of gravitational radiation and black-hole mergers, which do not have Newtonian (or even 2PN) counterparts, that dramatically change the qualitative behavior of the system. These effects become important when two objects approach each other closer than about 100 times their combined mass~\\cite{Lousto:2007ji} (or $d\\sim 100 G (m_1+m_2)/c^2$ in more conventional units) and become dominant when the two objects approach within about 10 times their combined mass. Even when the effect is small, the sensitivity of three (and more) body encounters to small perturbations makes including PN and GR effects important for obtaining the correct dynamics of an N-body simulation. The GR effects in particular provide a natural regularization of the Newtonian problem on small scales.   An important problem in galactic dynamics is determining how the merger process of the two (initially) central black holes proceeds after the collision of the host galaxies~\\cite{GouldRix:2000, Milosavljevic:2001vi,Yu:2002}.  In our simulations we confirm that the resulting binary will, in general, become highly elliptical if it interacts with a similar mass black hole. These eccentric orbits will drive the binary merger due to emission of gravitational radiation. However, we also found that the interaction can produce an immediate merger of the triple system. If we set physical scale of our simulation by setting the black-hole masses to  $10^9M_\\odot$, then the prompt mergers that we observe occur at separations of the order of milli-parsecs. These three-body interactions provides a possible mechanism for resolving the ``last parsec problem.''~\\cite{Milosavljevic:2002ht} We will present results from the effects of unequal-masses, spins, and further separated binaries in a forthcoming paper, where we will also explore the potential for critical phenomena. It is important to note that these three-black-hole interactions provide a mechanism for producing highly-elliptical close-binaries (which would otherwise have circularized due to emission of gravitational radiation during the inspiral) like those studied in Ref.~\\cite{Hinder:2007qu}."
    },
    "0711/0711.4095_arXiv.txt": {
        "abstract": "Results are presented for a simulation carried out to test the precision with which a detector design (HERON) based on a superfluid helium target material should be able to measure the solar \\pp\\ and \\be7 fluxes.  It is found that precisions of $\\pm 1.68\\%$ and $\\pm 2.97\\%$ for \\pp\\ and \\be7 fluxes, respectively, should be achievable in a 5-year data sample.  The physics motivation to aim for these precisions is outlined as are the detector design, the methods used in the simulation and sensitivity to solar orbit eccentricity. ",
        "introduction": "This paper presents an analysis of the projected capability of a detector design, HERON \\cite{heron}, based on a target material of superfluid helium to make a precise measurement of both the \\pp\\ and \\be7 solar neutrino fluxes (\\phipp{} and \\phibe, resp.) in a single, real-time experiment.  The detection reaction used would be the elastic scattering of neutrinos by electrons (ES).  In addition to the novel use of helium, the detector also includes the novel application of a coded aperture \\cite{dicke,ura} technique for accurate measurement of the location and recoil energy of each elastic scattering event and to aid in background discrimination.  According to models of the Sun \\cite{bethe}, the neutrinos from the \\pp-I and \\pp-II branches of the fusion chain (known as the \\pp\\ and \\be7 neutrinos, respectively) are, when taken together, $> 98\\%$ of the neutrino flux and are associated with the reactions producing a similar fraction of the solar energy.  At the present writing there have been no real-time experiments to measure the flux and spectra of \\phipp{} but recently the Borexino collaboration has made the first real-time spectral detection of \\phibe ($861\\,{\\rm keV}$) \\cite{borexino}.  As we explain in Sec.~\\ref{sec:goal} there are several important physics issues which can be addressed if a detector can be constructed to measure both these fluxes and spectra with sufficient precision.  For the \\pp\\ neutrinos ($<420\\,$keV) there does not yet exist any detector with demonstrated feasibility to measure either their flux or spectra; however, there are a number of efforts \\cite{heron,exps} which aim to do so.  Section~2 of this paper discusses the physics goals motivating the HERON detector.  Sec.~3 briefly presents the requirements for and description of the HERON detector design.  Sec.~4 and 5 provide the details of the HERON capability analysis.  In an Appendix we discuss an application to measuring the solar orbit eccentricity. ",
        "conclusions": "We have described that, on the basis of the new knowledge gained in recent years of neutrino properties and of higher energy solar neutrino fluxes, there are excellent reasons to perform precision real-time measurements of the very low-energy neutrino fluxes from the Sun.  The physics goals outlined in Sec.  2 include determining the luminosity of the Sun in neutrinos, providing checks on some details of the SSM, testing the MSW effect in the LMA solution and improving constraints on the neutrinos mass-mixing parameters as well as providing discovery opportunities in the new low energy regime.  To achieve these goals detectors are required which can measure the \\phipp\\ flux with a precision better than $3\\%$ and the \\phibe\\ flux to better than $5\\%$.  Such detectors must be capable of collecting very large event samples and maintain good control of systematic errors.  We have described the design of such a possible detector, HERON, and have simulated its performance.  Although the HERON detector is not presently scheduled or funded for construction, by experimentation in prototypes of several liters we have measured the details of energy loss processes (scintillation, phonons/rotons, electron bubbles) for low energy electrons in the superfluid \\cite{adams,huang,bandler,baskar}.  The development of wafer calorimeters capable of detecting all three channels has been carried out \\cite{fleischmann,enss,bhaskar,kim}.  We have not tested an array of wafers as a coded aperture; however, given the well tested use of the method in other fields, performance as simulated can be reasonably expected. The superfluid helium target material is itself free of intrinsic internal background and provides two channels (scintillation and drifted electrons) to distinguish and separate externally entering background from point-like neutrino ES signals via an active coded-aperture array.  The simulation has been directed towards establishing the systematic and other errors to be expected for the HERON detector in an exposure of 5 years.  For that purpose large samples of both signal and background events were generated and then fully reconstructed according to the physical processes in helium, the detector geometry and the properties of the coded aperture design.  The expected signals were based on current best understanding of the neutrino mass-mixing parameters and the well-known electroweak scattering cross-sections.  The backgrounds (gamma-ray conversions in the helium) were simulated from radioactive sources distributed throughout the major materials surrounding the helium.  The level and nature of the activities assumed was in line with current best practice in solar and double beta decay neutrino experimentation.  To separate a combined sample of pp, Be7 and backgrounds into their respective flux components, an extended loglikelihood method was used employing probability distribution functions constructed from various samples of the above simulation prescription.  By design, the method included all correlations among variables imposed by event reconstruction, various cuts on the data as well as those arising from the properties of the PDF's themselves.  The results are quite promising as can be seen in Tables 2-5.  It appears that should the detector be built and perform as modeled it would be capable of satisfying the criteria necessary for the precision pp and Be7 flux measurements.  For example, to take the particular choice of the so-called ``standard threshold and fiducial cut'' on a 5-year exposure and combining all errors except energy scale (statistical, high-level cuts, likelihood method) a precision on pp flux of ${\\pm 1.35\\%}$ (or ${\\pm 1.68 \\%}$ including energy scale uncertainty) results.  Similarly for \\phibe, we find errors of ${\\pm 2.96\\%}$ and ${\\pm 2.97\\%}$ without and with, respectively, the energy scale uncertainty included.  The full neutrino flux obtained without attempting to employ the separation of \\phipp{} and \\phibe\\ individually would present a combined error of ${\\pm 1.31\\%}$.  The validity of the background model used in these simulations is a key issue.  The composition and relative magnitude of the background sources assumed were based on current experience in the field and should therefore be realizable in practice; nonetheless, it is important to ask to what extent are the simulation results dependent upon the assumed model and to what extent can the model be checked in practice.  The decay modes of the sources, branching ratios and energies of the decay products are well known and the method of their propagation through the simulation programs are well established. Perhaps most important is the question of whether we may have mis-estimated the total background level and if so how strongly the result would be affected.  A mistake in the magnitude would enter principally through the separation procedure using the PDF's; consequently we have tested this effect by varying the assumed size of the background over a wide range.  The effect is not drastic; for example, should the background be 5 times larger, the \\phipp\\ error would double while a factor 50 larger background would raise the \\phipp\\ error by six times (however, this latter rate would introduce prohibitive wafer deadtime).  In contrast, reducing the background by a factor of $0.5$ only improves the \\phipp\\ error by $15\\%$.  In practice, there are some checks available on the model.  Due to the good position and energy resolution for the point-like signal events, fiducial volumes of various sizes can be made and the stability of the flux results checked.  As we have seen in Tables 2, 3 and 5, varying the high-level cuts by setting different fiducial volumes and thresholds does not have a strong effect upon the expected errors.  Similarly, the nature and dependence of the observed spatial and energy distributions as a function of these cuts can also be compared directly to the model.  In conclusion we believe that a detector of the HERON design utilizing superfluid helium and a coded aperture array could provide the capability to carry out the multiple physics goals achievable through precision, real-time, simultaneous measurements of \\phipp{} and \\phibe."
    },
    "0711/0711.1353_arXiv.txt": {
        "abstract": "In this contribution we investigate how to describe the results and usage of evolutionary synthesis models. In particular, we look for an explicit and quantitative description of the parameter space of synthesis models and the evaluation of their associated uncertainties and dispersion. First, we need to understand what synthesis models actually compute: we show that a synthetic stellar population with fixed physical parameters (age, metallicity, star formation history, initial mass function and {\\it size} of the system) can only be described in terms of probability distributions (i.e. there is an intrinsic dispersion in any model). Second, we need to identify and characterize the coverage in the parameter space of the models (i.e. the combinations of input parameters that yield meaningful models) and the different sources of systematic errors. Third, we need a way to describe quantitatively the intrinsic dispersion, the systematic error and the parameter space coverage of the models.  Up to now, the parameter space coverage and uncertainties have been described qualitatively in the models' reference papers, with potential misinterpretations by models' users. We show how Virtual Observatory developments enable a correct use of synthesis models to obtain accurate (and not simply precise) results. ",
        "introduction": "The aim of the Virtual Observatory (VO) is to provide an environment to produce more accurate science. Not only does the VO provide an opportunity for technical developments, it also offers unique possibilities for scientific research: The VO framework requires an explicit description of the data that it provides with no implicit assumptions about the data themselves. This requirement is fundamental since interoperability is performed by machine to machine communications without supervision by human beings. Such a description is performed by the data model working group in the International Virtual Observatory Alliance (IVOA).  The rationale of the data model working group is to provide an abstract description of concepts and their interrelationships, used to fix both the names and meanings of concepts in the VO context and also their internal structure and cross-connections (IVOA Spectral Data Model 1.01, 2007)\\footnote{\\tiny{\\tt http://www.ivoa.net/Documents/WD/DM/SpectrumDM-20070515.html}}, that is to describe all the possible universes of classes of data. Note that such description and interrelationships are fundamental in scientific research also, with the only difference that, in the last case, researchers are focused on {\\it particular} types of objects.   \\begin{figure} \\centering \\includegraphics[width=0.8\\linewidth]{cervinof1.eps} \\caption{IVOA spectrum data model.} \\label{fig:IVOA-SDM} \\end{figure}  Figure \\ref{fig:IVOA-SDM} shows the spectrum data model generated by the data model working group. This data model aims to describe all the universe of tables in the VO that contain a spectrum, hence the data model must not only be valid for observed data, but also for theoretical spectra.  This data model (the generic description of a spectrum) allows to formulate new questions for researchers that either produce or use theoretical spectra: Does this data model describe theoretical spectra? Can all the ``boxes'' in the data model be filled? And, if this is the case, what is the accuracy of a theoretical model? Do theoretical models have systematic errors? And, do they have statistical errors?  In addition to the spectrum data model, this working group has also proposed a data model for astrophysical dataset characterization (IVOA Data Model for Astronomical DataSet Characterization 1.11, 2007)\\footnote{\\tiny{\\tt http://www.ivoa.net/Documents/PR/DM/CharacterisationDM-20070530.html}}. This data model defines the information required to describe the parameter space of observed or simulated astronomical datasets. Fig. \\ref{fig:IVOA-CDM} is a draft illustration of the description of an atmosphere library like those used in population synthesis.  \\begin{figure} \\centering \\includegraphics[width=0.8\\linewidth]{cervinof2.eps} \\caption{Draft illustration of the concepts of the IVOA characterization data model for a typical atmosphere library like those used in populations synthesis showing the resolution of the library ({\\it precision} in the VO data model) and its coverage ({\\it support} in the VO data model). Isochrones are also drawn for an intuitive comparison of the support of both sets of data.} \\label{fig:IVOA-CDM} \\end{figure}  Again, a correct characterization of models is fundamental for their application to observations. All theoretical models have a limited range of validity, given by the computational algorithms, the resolution of the input data and other factors. Usually, these limitations are described qualitatively in the reference papers, but unfortunately, it is not so common for these limitations to be introduced in the code that computes the models \\cite[Fortunately, there are some significant exceptions, like CLOUDY by][]{cloudy}.  We stress again that an answer of the previous issues is fundamental in scientific research, since the basic way to obtain the physical parameters of an observed system is to compare observational data with theoretical models. Only if these questions are solved in a quantitative way, will it be possible to evaluate the accuracy (not the precision) of our knowledge of the observed system.  In the case of stellar populations, theoretical models are provided by stellar population synthesis \\citep{fanelli,tinsley}. The general case of population synthesis aims to decompose the observed light of the system under study in different stellar populations components, that is the contribution of classes of stars, each class with different effective temperature and gravity \\citep{fanelli}. In this case, no further information about the observed system can be obtained unless additional a priori information is considered. In the case of {\\it evolutionary} population synthesis \\citep{tinsley}, the contribution of the different stellar classes is defined by the evolutionary status of the system, hence the knowledge of the contribution of the different stellar classes can be transformed into knowledge about the age, metallicity, stellar birth-rate (i.e. star formation history, initial mass function and the amount of formed stars) of the system. Although evolutionary synthesis models are powerful, it is necessary to take into account that: \\begin{itemize} \\item The transformation between the contribution of stellar classes and the associated physical parameters requires some a priori hypotheses: a set of atmosphere models, evolutionary tracks and defined star formation properties (i.e. stellar birth rate). \\item There are different evolutionary conditions that result in the same integrated spectra, so the physical parameters obtained from the analysis of observations are not necessarily unique. \\end{itemize}  These a priori hypotheses, together with the computational algorithms and approximations used by different synthesis codes, define the range of use (i.e. characterize) of the results. They are also the principal source of systematic errors in model computations. On the other hand, the degeneracy of the possible solutions produces an intrinsic precision loss in any fit of observational data and theoretical models: the best-fit solutions are not necessarily the correct ones and the distribution of the goodness of fit is also needed for a correct interpretation of the observed data. In order to analyze these issues in detail, we need to understand what a synthesis model is. ",
        "conclusions": "Synthesis models, as any other theoretical model, are affected by systematic errors. Although the evaluation of such errors is a difficult task that only can be performed by synthesis models makers, the current IVOA description of spectrum data model allows to add these errors to the output.  In the case of statistical errors, the spectral data model does not apply directly to theoretical results since the goal of statistics is inferring the underlying distribution from which the data have been drawn, which is assumed to be known in the computation of theoretical models. However, the {\\it StatError} box in the current spectral data model can be used to describe the {\\it intrinsic dispersion} of theoretical models. The dispersion of synthesis models, although not recognized as such, is computed by several codes under the label of surface brightness fluctuations, so this piece of information is already available.  Finally, a proper characterization of isochrones and atmosphere libraries coverage allows to obtain a realistic characterization of synthesis models, and its range of application."
    },
    "0711/0711.0281_arXiv.txt": {
        "abstract": "Models describing dust-driven winds are important for understanding the physical mechanism and properties of mass loss on the asymptotic giant branch. These models are becoming increasingly realistic with more detailed physics included, but also more computationally demanding. The purpose of this study is to clarify to what extent the applied numerical approach affects resulting physical structures of modelled winds, and to discuss resulting changes. Following the previously developed radiation hydrodynamic model -- which includes descriptions for time-dependent dust formation and gas-dust drift -- and using its physical assumptions and parameters, numerical improvements are introduced. Impacts of the so-called adaptive grid equation and advection schemes are assessed from models calculated with different numerical setups. Results show that wind models are strongly influenced by numerical imprecision, displaying differences in calculated physical properties of up to one hundred per cent. Using a non-adaptive grid, models become periodic (in multiples of stellar pulsation periods), instead of irregular, as obtained previously. Furthermore, the numerical improvements reveal changes in physical structures. The influence of gas-dust drift is confirmed to be highly important, in particular for the dust component. Gas and dust are less tightly coupled than previously, and drastically larger amounts of dust form assuming drift. ",
        "introduction": "\\label{sec:introduction} Dust-driven winds off of asymptotic giant branch (AGB) stars are complex physical phenomena that form in a highly non-linear interaction between a strong radiation field, a cool gas and conditions suitable for efficient dust formation. It is the radiation pressure on dust grains that with the support of levitated regions of an enclosed pulsating star is believed to form and drive more massive winds. Starting with \\citet{Bo:88} a variety of time-dependent wind models, of increasing complexity, have been created in order to understand AGB star winds physically. An overview of properties and capabilities of recent wind models is given by, e.g., \\citet[and references therein]{Ho:05}. \\citet{Wi:00} reviews mass loss from cool stars in more general terms.  An important feature of self-consistent wind models, as presented here, is an adaptive grid equation. This equation distributes a fixed number of gridpoints across the model domain according to those physical properties that are required to be spatially resolved \\citep[henceforth {\\rDD}]{DoDr:87}. The intent to use an adaptive grid is sound because shocks form in radially steep density gradients of long period variables. Shocks that assist in providing appropriate conditions for dust formation and hence wind acceleration. It has earlier been found that a vast majority of all wind models using such a grid equation have formed irregular outflow patterns; for examples see: \\citet[fig.~1]{HoDo:97}, \\citet[henceforth {\\rSaHoa}, fig.~2]{SaHo:03}, \\citet[fig.~2]{HoGaArJo:03}, \\citet[henceforth {\\rSaHob}, fig.~5]{SaHo:03b}, and \\citet[figs.~5--7]{MaHoHe:07}. Such irregularities could be understood as a result of physically (mildly) chaotic interactions, but they might also be a consequence of insufficient numerical accuracy.  Few studies address the use of the adaptive grid equation. To mention two \\citet{Ku:94} compares results between models using first and second order accurate advection schemes; concluding that a second order advection is important. \\citet[henceforth {\\rDPSG}]{DoPiSt.:06} introduce both an improved discretization scheme and an improved advection scheme. In addition to the advection and discretization schemes, other ``numerical properties'' also have the potential of affecting the physical outcome of models. Such examples are how the adaptive grid equation is used, and the adopted grid resolution. There is no thorough study of wind models available in the literature, where the influence on results due to the numerical approach is explored. For models using the adaptive grid equation there is no such study available at all. This is an important issue of concern, since a model interpretation is uncertain when it is unclear to what degree the outcome is affected by numerical inaccuracy.  This paper is based on the work presented in the previous papers in the series, i.e., {\\rSaHoa}, {\\rSaHob}, and \\citet[henceforth {\\rSaHoc}]{SaHo:04}; a comprehensive summary is given in \\citet{CSa:03}. The purpose of this paper is twofold. Firstly to fill the gap of a numerical study by exploring and determining a modelling approach by which numerical effects on results are separated and minimized. All physical parameters are thus fixed to previously used values. Secondly, the improved numerical approach leads to significant changes in the outcome, which are new and well worth addressing. For instance, when the adaptive grid does not track shocks outflow structures turn out to be periodic, instead of irregular as before. Assuming drift the wind also becomes more tenuous, simultaneously drastically larger amounts of dust are formed. Issues of the numerical method and physical results are treated together in order to emphasise how strongly they depend on each other. Results should, finally, also be valuable when improving related studies working with mass loss rates and yields of dust from AGB stars further, such as, e.g., \\citet{FeGa:06}.  Numerical modifications and improvements are introduced, after first summarising the physical and numerical setup of the current AGB star wind model, in Sect.~\\ref{sec:numerical}. The modelling procedure and model setup are then described in Sect.~\\ref{sec:moprores}, where results and code consistency checks are also presented. The improved numerics in several cases gives rise to drastic changes to the physical structure. These changes and the revised outcome are discussed, after differences of the numerical modifications are treated, in Sect.~\\ref{sec:discussion}. The paper is closed with conclusions in Sect.~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions} In recent years models of dust-driven winds from AGB stars have been made increasingly realistic with the refinement and addition of more physics. These models must handle many complex non-linear and time-dependent processes that describe a combined evolution of gas, dust, and a strong radiation field. Solving the resulting radiation hydrodynamic system of equations puts large demands on the numerical method used, which has to correctly treat, e.g., narrow shocks, and time scales covering many orders of magnitude. Results of previous studies (using the same modelling approach as in earlier papers of this series) show a majority of wind structures that are irregular (i.e., mildly chaotic). There is no thorough study where the occurrence and extent of such irregularities are clarified.  The purpose of this paper was to fill this gap by studying the influence of the numerical method on wind models that use the so-called adaptive grid equation \\citep[][]{DoDr:87}. The study was based on the model description introduced in \\citet[{\\rSaHoa}--{\\rSaHoc}]{SaHo:03,SaHo:03b,SaHo:04}. Clarifying the numerical influence all physical assumptions and parameters were kept unchanged. As before the emphasis was on improving the understanding of effects of gas-dust drift. In order to reassess previous results models were calculated both allowing drift (drift models), and not allowing drift (position coupled [PC] models).   An important result of this study is that an adaptive grid equation is not in general capable of tracking a large number of shocks without introducing numerical errors into the solution. All presented and revised models using a non-adaptive grid (which does not track shocks) were found to form periodic structures with a stationary outflow showing small amplitude fluctuations (originating in the piston). This was found to be the case both for drift and PC models. The same structures were all irregular when the grid was instead adaptive (tracking shocks). Although shocks are well resolved with an adaptive grid, regions between shocks simultaneously become unadequately resolved; because, there is only a fixed amount of gridpoints. In drift models dust is not bound to the gas and might require an appropriate resolution also between shocks. Using such a powerful tool as the adaptive grid equation without first assuring that its numerical influence does not dominate over the physical structure it is impossible to tell whether a modelled wind is physically irregular, or if variations are of numerical origin.  In attaining a high accuracy it is, moreover, necessary to use a precise advection scheme and a sufficient number of gridpoints. Detailed studies showed that smaller amounts of dust form with unresolved shocks. Incorrect structures also formed with a less accurate advection scheme. Using the PPM scheme \\citep{CoWo:84}, numerical errors were found to be small; at the per cent level. Compared to previously calculated models, differences in individual model properties were overall as large as 100\\,per cent.  It is important to point out that absolute values of results in this study are highly dependent on adopted physical assumptions and parameters, as well as the range of stellar parameters used. Qualitatively, the results are comparable with, or applicable to, models that also turn out to form periodic structures. An application to models which instead form irregular structures is not meaningful, but requires a separate study.  Many physical results of previous papers in this series were confirmed, but there were also new findings. The gas component typically showed small variations. Unlike earlier there were no shocks in the outer envelope. The dust component, however, showed large variations in drift models, confirming previous results. Although the periodic variation pattern diffused away at larger radii; and it consequently appears that also the outflow of the dust component becomes stationary. Moreover, unlike previous results the dust component was less tightly coupled to the gas. Typically the drift velocity was fairly low in the entire envelope, $\\vdri\\!\\la\\!10\\,\\kms$. The highest values were assumed in the inner envelope, decreasing outwards to $3\\!\\la\\!\\vdri\\!\\la\\!5\\,\\kms$ at the outer boundary (at about 50 stellar radii).  Drastically larger amounts of dust formed in drift models compared to PC models; the increase was 160--250\\,per cent in every case but one (for which the increase was 40\\,per cent). It was, furthermore, found that mass loss rates typically were lower in drift models, by $\\la30\\,$per cent. Terminal velocities were, to the contrary, larger, by 10--50\\,per cent. Winds of drift models were consequently more tenuous than those of PC models. Altogether results (again) show that neglecting gas-dust drift creates imprecise outcome.  The modelled winds were neither the least, nor the most massive. Partly this is a consequence of the Planck mean absorption coefficients used, which result in less dense wind structures, than would be the case had instead frequency-dependent opacities been used. Partly because weak winds are difficult to calculate with drift. It is highly likely that the range of mass loss rates where drift (and dust) is important is limited. Finally, there are issues for model improvement where a closer study would make the understanding of the wind formation mechanism more complete. Three such issues addressed were: using a grain size distribution instead of a mean size, using frequency-dependent opacities, and the role of a two- or three-dimensional modelling approach."
    },
    "0711/0711.0248_arXiv.txt": {
        "abstract": "We find a strong correlation between the extension of the Na-O anticorrelation observed in  red giant branch (RGB) stars and the high temperature extension of the  horizontal branch (HB)  blue tails of Galactic globular clusters (GCs).  The longer is the O-depleted tail of the Na-O anticorrelation  observed in the  RGB stars, the higher is the maximum temperature reached by  the bluest HB stars in the GC.  This result provides a clear, empirical evidence of a link between the extension of the HB and the presence of  star-to-star abundance variations of proton-capture elements in GC stars. We discuss the possible interpretation of this correlation. ",
        "introduction": "Since more than thirty years, we know that globular cluster (GC) stars show a dispersion in the content of light elements like C, N, O, Na, Mg, Al. The lightest C and N are anticorrelated both in field and cluster stars (Smith  \\& Martell 2003). The environment of GCs must be responsible for the bulk of chemical anomalies observed among stars in GCs (at odds with field stars, Gratton et al. 2000), where the heavier proton-capture elements are anticorrelated (Na vs O and Mg vs Al, respectively; see Gratton, Sneden, \\& Carretta 2004  for extensive references).  The theoretical background (e.g. Denisenkov \\& Denisenkova 1990, Langer et al. 1993) points to the simultaneous run of the CNO, NeNa and MgAl cycles of H-burning at high temperature. In turn, this  implies that the observations of the same anticorrelations in unevolved cluster stars (Gratton et al. 2001; Ramirez \\& Cohen 2002; Carretta et al. 2004, 2005; Cohen \\& Melendez 2005) can be explained only if the involved proton-reactions occurred in more massive stars  of a previous generation than the presently observed one, and the anomalies concern a large fraction of the stellar structure, non simply a surface layer (Cohen et al. 2002). Multiple star formation episodes in GCs have eventually been directly observed in the form of a main sequence split in $\\omega$ Centauri (Bedin et al. 2004) and NGC 2808 (Piotto et al. 2007), and in a subgiant branch split in NGC 1851 (Milone et al. 2007).  Red giant branch (RGB) stars with both normal and anomalous composition must eventually end up on the Zero Age Horizontal Branch (ZAHB) after the onset of core He-burning, and we might expect to track their different chemical signatures in the ensueing evolution.  On one hand, the  horizontal branch (HB) phase is an ideal locus to investigate the effects of chemical anomalies on the stellar structure, as it acts as an amplifier of the  physical conditions consequent to the star's previous evolution. On the other hand, the location of stars onto the ZAHB is sensitive to a number of secondary parameters (age, CNO composition, helium content, rotation etc, see e.g., Fusi Pecci et al. 1993) beside the main parameter, metallicity,  to make it difficult to disentangle the contribution of the single one.  A combination of these parameters actually conspire to give a variety of HB morphologies.  Recently, Recio-Blanco et al. (2006;  hereafter RB06) provided a quantitative estimate of the temperature extent of the HB by using the homogeneous set of color-magnitude diagrams (CMDs) of the HST snapshot program by Piotto et al. (2002). RB06 found a  significant correlation between the HB extension and  the cluster total mass. They interpreted this correlation as an observational evidence of self-pollution, assuming, as suggested by D'Antona et al. (2002), that the factor responsible for the HB extension might be an He enhancement predicted to accompany  the observed CNONaMgAl star-to star abundance variations on the RGB.  In this Letter we present a tight and direct empirical correlation between the maximum temperature along the ZAHB and the extent of the Na-O and Mg-Al anticorrelations, as defined by Carretta (2006), and discuss its implications. ",
        "conclusions": ""
    },
    "0711/0711.0412_arXiv.txt": {
        "abstract": "Three cornerstones for the 3D data support in the Virtual Observatory are: (1) data model to describe them, (2) data access services providing access to fully-reduced datasets, and (3) client applications which can deal with 3D data. Presently all these components became available in the VO. We demonstrate an application of the IVOA Characterisation data model to description of IFU and Fabry-Perot datasets. Two services providing SSA-like access to 3D-spectral data and Characterisation metadata have been implemented by us: ASPID-SR at SAO RAS for accessing IFU and Fabry-Perot data from the Russian 6-m telescope, and the Giraffe Archive at the VO Paris portal for the VLT FLAMES-Giraffe datasets. We have implemented VO Paris Euro3D Client, handling Euro3D FITS format, that interacts with CDS Aladin and ESA VOSpec using PLASTIC to display spatial and spectral cutouts of 3D datasets. Though the prototype we are presenting is yet rather simple, it demonstrates how 3D spectroscopic data can be fully integrated into the VO infrastructure. ",
        "introduction": "Integral field (or 3D) spectroscopy is a modern technique in astrophysical observing that was proposed by Georges Court\\'es in the late 60's. The idea is to get a spectrum for every point in the field of view of a spectrograph. Several instrumental approaches in the optical domain (as well as NIR and near-UV) exist: scanning Fabry-Perot interferometry, image slicing and transforming two-dimensional field of view into a slit using Integral-Field Unit (see review in P\\'econtal-Rousset et al., 2004 for a description of different image slicing techniques).  At present, nearly all large telescopes in the world are equipped with 3D spectroscopic devices, and rapidly growing volume of data produced by them pose a number of questions regarding the data discovery and retrieval. In this paper we demonstrate how 3D data are handled in a framework of the International Virtual Observatory.  All 3D spectroscopic observations result in datasets having both spatial and spectral information. They are usually referred as ``datacubes'', though sometimes (in case of IFU) they are not regularly gridded in spatial dimensions.  There are three cornerstones for the 3D data support in the Virtual Observatory: \\begin{enumerate} \\item data model -- an abstract, self-sufficient and standardised description of the data \\item data access services -- archives, providing access to fully reduced science-ready datasets \\item client applications -- data-model aware software that is able to search, retrieve, and display 3D data, as well as to give a possibility for sophisticated scientific data analysis \\end{enumerate}  All these blocks became available, and we will review them in the forthcoming sections. ",
        "conclusions": "In Chilingarian et al. (2006) we concluded that \"all the necessary infrastructural components exist for building VO-compliant archives of science-ready 3D data and tools for dealing with them\". Since that time there was a substantial progress of VO standards and protocols. Now we are able to provide access to first two such VO-compliant archives. This not only a ``proof-of-concept'', but the services that can be used for real scientific purposes. Another important conclusion that can be drawn is that the present state of VO standards (including PLASTIC -- a prototype of the VO application messaging protocol) is totally sufficient for dealing with complex datasets in a VO framework without need to develop new client applications for every particular kind of data."
    },
    "0711/0711.4141_arXiv.txt": {
        "abstract": "The idea of this work is to compare a new positive and entropy stable approximate Riemann solver by Francois Bouchut with a state-of the-art algorithm for astrophysical fluid dynamics. We implemented the new Riemann solver into an astrophysical PPM-code, the Prometheus code, and also made a version with a different, more theoretically grounded higher order algorithm than PPM. We present shock tube tests, two-dimensional instability tests and forced turbulence simulations in three dimensions. We find subtle differences between the codes in the shock tube tests, and in the statistics of the turbulence simulations. The new Riemann solver increases the computational speed without significant loss of accuracy. ",
        "introduction": "In modern astrophysics the interplay between observations and numerical experiments plays a central role. Typically hydrodynamical flows with high Reynolds numbers and Mach numbers are studied, and they are modelled by the Euler equations \\begin{align} \\rho_t + \\text{div}(\\rho {\\bf u})=0\\nn\\\\ (\\rho u)_t +  \\text{div}(\\rho u{\\bf u}) + p_x=\\rho f_1\\nn\\\\ (\\rho v)_t +  \\text{div}(\\rho v{\\bf u}) + p_y=\\rho f_2\\nn\\\\ (\\rho w)_t +  \\text{div}(\\rho w{\\bf u}) + p_z=\\rho f_3\\nn\\\\ E_t + \\text{div}((E+p){\\bf u})=\\rho {\\bf f }\\cdot{\\bf u}. \\label{euler} \\end{align} Here $\\rho$ is the mass density, ${\\bf u}=(u,v,w)$ is the velocity field, $p$ is the pressure, and $E$ is the energy density $E=\\half\\rho {\\bf u}^2 + \\rho e$ with $e$ the specific internal energy. External forces are given in units of acceleration by ${\\bf f }=(f_1,f_2,f_3)$. The system is closed by the equation of state that relates $p$ to $\\rho$ and $e$. In this work we consider ideal gases where $\\rho e=p/(\\gamma-1)$ for some $\\gamma>1$, and isothermal gases. An isothermal gas has constant temperature $T$, which implies $p=a^2\\rho$ for $a=\\frac{RT}{\\mu}$, with $R$ the gas constant and $\\mu$ the mean molecular weight. In a real astrophysical flow additional physical phenomena such as magnetic fields, gravitational forces and electromagnetic radiation may be important, but this paper is only concerned with the hydrodynamics. The specific physical entropy $s$ is defined by the relation \\beq {\\text d}e + p {\\text d}\\frac{1}{\\rho}=T{\\text d}s \\eeq with $T=T(\\rho,e)>0$ the temperature. The second law of thermodynamics implies that \\beq (\\rho \\phi(s))_t + \\text{div}(\\rho {\\bf u} \\phi(s))\\leq 0 \\label{entropyinequality} \\eeq for any smooth, nonincreasing and convex $\\phi$. In high Mach number flows this condition is needed to ensure the dissipativity of shocks, since the viscous forces are ignored in \\eqref{euler}.  To numerically solve \\eqref{euler}, shock-capturing finite volume schemes are widely used. In astrophysics it is often done with the PPM algorithm described in \\cite{ppmpaper}, often with an iterative method to approximate the exact midpoint value of the Riemann fan. This was implemented in the Prometheus code in 1989, using the iterative Riemann solver of \\cite{ppmriemann} with a fixed number of iterations, see \\cite{prometheus}. An efficient, parallelised version was then implemented in 2001, see \\cite{Reinecke1} and \\cite{Reinecke2}. Stochastic forcing for turbulence simulations was added later, see \\cite{schmidtphd}, \\cite{schmidtnum}. Results produced by the Prometheus code have been presented in many astrophysical publications, for example in \\cite{Schmidtsupernova2}, \\cite{Reinecke1}, \\cite{Reinecke2} and \\cite{Schmidtsupernova}. We used this code from 2001 as the basis for this work, and implemented our changes into it.  First, we switched the Riemann solver to an HLLC solver with the signal speeds of Bouchut (\\cite{BouchutBGK}, \\cite{Bouchutbook}). As far as we know, this is the first implementation of this advancement into an astrophysics code. It is well known that approximate Riemann solvers like HLLC are computationally efficient and easy to implement. In addition, the Riemann solver of Bouchut has two good properties: a) It automatically ensures that a discrete version of the entropy inequality \\eqref{entropyinequality} holds. b) The density and pressure stay positive. Often, finite volume codes need to check for each cell update whether the new data are physically reasonable, but with these two a priori estimates, no checks are necessary apart from underflow treatment. The two estimates are true in the first order case.  When using this Riemann solver in a higher order scheme, these two properties are not automatically inherited. Hence, we introduced a piecewise linear reconstruction, and replaced the characteristic backtracing with Runge-Kutta time integration. This was done in such a way that positivity is preserved at one half the CFL-number required in the first order case. Such a second-order scheme which also satisfies the entropy inequality is however impractical, see \\cite{secordentropic}, but entropy stability for first order schemes has so far seemed to be a good condition in practice. A different notion of stability comes from scalar conservation laws, which have solutions with nonincreasing total variation. This notion is also important in the design of higher order methods for systems. The reconstruction algorithm and the Runge-Kutta integration we use ensure a nonincreasing total variation when applied to scalar equations.  This gave rise to four codes as summarized in Figure \\ref{codestable}. \\begin{figure} \\begin{tabular}[c]{|l||c|c|} \\hline & \\multicolumn{2}{|c|} {}\\\\[-8pt] & \\multicolumn{2}{|c|}{Riemann solver:}\\\\[3pt] \\cline{2-3} & &\\\\[-8pt] Higher order algorithm:& Iterative of Prometheus& HLLC-Bouchut\\\\[3pt] \\hline \\hline &&\\\\ PPM reconstruction and back-tracing& PPM & PPM-HLLC\\\\[6pt] \\hline &&\\\\ Piecewise linear in space, Runge-Kutta in time&RK-exact&RK-HLLC\\\\[6pt] \\hline \\end{tabular} \\caption{The table summarises the four codes we tested. Along the horizontal axis the Riemann solver changes, while vertically the higher order algorithm varies} \\label{codestable} \\end{figure} The codes RK-HLLC and PPM-HLLC both run about 20\\% faster than PPM on the same data with the same resolution. However, the difference between the algorithms might be larger, since only the original PPM-code was optimised. The RK-exact code is the slowest, but it was of fundamental interest in this project to compare its accuracy to the RK-HLLC code's.  In the remainder of this introduction we will first sum up the main ideas of the underlying PPM algorithm in the Prometheus code. Then we will describe the new Bouchut-HLLC solver and its theoretical advantages. The second order algorithm is outlined next, and we show how it preserves positivity. In sections 2,3 and 4 the one-, two- and three-dimensional test comparisons are presented. At the end there is a conclusion.  \\subsection{PPM and the Prometheus code} The basis of the Prometheus code is the one-dimensional PPM-method of \\cite{ppmpaper} with the iterative Riemann solver from \\cite{ppmriemann}. Strang splitting is then used to handle multidimensions. The crucial point of PPM is the so-called characteristic back-tracing. This technique produces a second order approximation to the states at the cell interfaces at the half time step, allowing the use of the midpoint method in time. These approximate states are then used as input to the Riemann solver. Although the overall accuracy is second order, the spatial reconstruction is piecewise parabolic, which is reported to give better resolution than piecewise linear reconstruction. Furthermore, the accuracy at contact discontinuities is improved by an algorithm that detects them, and then steepens the reconstructed density. There is also an algorithm that adds artificial diffusion in order to avoid oscillations behind strong, slow-moving shocks without smearing out the solution much. The reconstruction is required to be monotocity preserving. This means that the order of the scheme may drop locally at extremal points of a reconstructed quantity, which means a primitive quantity in the case of PPM as in \\cite{ppmpaper}. This drop in order is also a feature of the second order reconstruction we use with the Runge-Kutta time integration.  In order to resolve shocks and shock interactions a Riemann problem is solved with the data from the back-tracing operation as input. For the Euler equations there is no general explicit formula for the solution of the Riemann problem, so in Prometheus an iterative procedure provides instead approximate values of the fluxes at cell interfaces. For efficiency the number of iterations is limited to a fixed number.  \\subsection{The HLLC solver of Bouchut and stability} The notion of an approximate Riemann solver goes back to the Roe solver, see \\cite{roescheme}, which is based on a local linearization of the fluxes at the cell interface. We refer to \\cite{Torobook} and \\cite{Bouchutbook} for a modern presentation of the basic ideas. The basic idea is to replace the exact Riemann fans in Godunov's method with something simpler that still gives a numerical flux that is consistent and conservative. In addition entropy consistency and positivity of density must be somehow ensured. For this linearized solvers require additional treatment, a so-called entropy fix.  The simplest approximate Riemann solver is the HLL solver, where the Riemann fan is replaced by a constant state separated by two discontinuities moving with constant speeds $C_l$ and $C_r$. A sufficient condition for stability is that the exact Riemann solution does not have waves with speed outside the interval $[C_l,C_r]$. The main weaknesses of this approach is that material contact waves are smeared out, and that the signal velocities $C_l$ and $C_r$ have to be guessed. A solution to the first problem was hinted at already in \\cite{HLL}, and was carried out in \\cite{HLLC}, see also \\cite{Torobook}, with the so called HLLC-solver. The HLLC-solver consists of three discontinuities traveling with speeds  $C_l$, $u^*$ and $C_r$, where velocity and pressure are held constant across the middle wave, and $u^*$ is this intermediate value of the velocity.  In \\cite{Bouchutbook} the HLLC-solver was improved by showing that it results from a relaxation system, which established its entropy stability. Furthermore, sharp explicit formulas for the signal speeds that ensure positivity and entropy stability could be given. We refer to these as Bouchut speeds, and they are given by formula (2.133) and Proposition 2.18 of \\cite{Bouchutbook}.  \\subsection{The new second order algorithm} When going to higher order, requiring the reconstruction to be entropy dissipative leads to impractical methods, but there is a way to preserve positivity in a rigorous manner. In the rewritten version of Prometheus RK-HLLC and RK-exact we used the following reconstruction, based on \\cite{Bouchutbook}  As a limiter we use \\beq \\text{minmod}(a_l,a_r)=\\bca 0, \\quad\\quad a_la_r\\leq 0\\\\ \\text{sign}(a_l)\\min\\(\\alpha |a_l|,\\,\\, \\half(|a_l|+|a_r|),\\,\\,\\alpha |a_r|\\), \\quad a_l a_r>0 \\eca \\eeq with $\\alpha$ set to 1.4. This is applied to produce the discrete differential \\beq D\\rho_i=\\frac{1}{h}\\text{minmod}(\\rho_{i+1}-\\rho_i, \\rho_i-\\rho_{i-1}), \\eeq and $Du_i$, $D(u_\\n)_i$, $D(\\rho e)_i$ similarily. The positivity of the reconstructed density is guaranteed since $\\rho_i-\\frac{h}{2}|D\\rho_i|>\\min_k(\\rho_k)$. Conservation of momentum dictates that we take the reconstruction slope \\beq D(\\rho u)_i=\\rho_i Du_i+u D\\rho_i-\\frac{h^2}{4}D\\rho_i Du_i, \\eeq and similarily for $D(\\rho u_\\n)_i$. Energy is conserved by replacing $e_i$ with \\beq \\tilde{e}_i=e_i-\\frac{h^2}{8}\\(1-\\frac{h^2}{4\\rho_i^2}D\\rho_i^2\\)\\(Du_i^2+(Du_\\n)_i^2\\) \\eeq when computing the reconstructed internal energy. The extra terms cancel out the conservation errors in kinetic energy caused by the linear reconstruction. Hence positivity means that $\\rho_i\\tilde{e}_i-\\frac{h}{2}|D(\\rho e)_i|>0$, or in other words \\beq \\frac{h^2}{8}\\(Du_i^2+(Du_\\n)_i^2\\)< \\frac{1}{1-\\frac{h^2}{4\\rho_i^2}D\\rho_i^2}\\(\\rho_i e_i-\\frac{h}{2}|D(\\rho e)_i|\\)\\hat{=}\\Lambda_i^2 \\label{posreconstr} \\eeq To ensure \\eqref{posreconstr} in a consistent way we first restrict $|D(u_\\n)_i|$ to less than or equal to $\\Lambda_i$, and then set $Du_i^2$ less than or equal to $\\Lambda_i^2-|D(u_n)_i|^2$. Note that in practice we multiply $\\Lambda_i $ with a number slightly less than one to ensure that the inequality \\eqref{posreconstr} is strict.  We did not apply any special treatment of material contact waves in this code version, and no artificial diffusion was added at shocks.  The numerical time integration is a second order Runge-Kutta method. That is, one does two full time steps, and then averages the resulting cell average with the initial one. This procedure preserves positivity, and is total variation diminishing. Multidimensionality is taken care of by Strang splitting just as in the PPM-codes. ",
        "conclusions": "From our numerical experiments we have made the following observations: \\begin{itemize} \\item There is slightly more smearing of stationary shocks with HLLC-Bouchut compared to the exact solver. \\item The RK-codes smear out most features more than the PPM-codes, and especially contact discontinuities. % \\item RK-HLLC handles small densities better than the other codes. \\item All codes exhibit spurious oscillations. We see more of them with the PPM-codes, except at the near stationary shock, where PPM has a specialised 'flattening' procedure. \\item The growth of Kelvin-Helmholtz and Richtmeyer-Meshkov instabilities appears to be little affected by which of the two Riemann solvers are used. \\item In turbulence simulations of adiabatic gas, the dissipativity of RK-HLLC seems to be less for higher than for lower Mach numbers, while the dissipation with PPM is independent of the Mach number. \\item For turbulence in an adiabatic gas with an RMS (root mean squared) Mach number less than about 5, RK-HLLC seems to be more dissipative than PPM. \\item For turbulence with an RMS (root mean squared) Mach number of 2.5 and higher in an isothermal gas, there were no significant differences between Riemann solvers or higher algorithms. \\end{itemize} The widespread use of PPM in the astrophysics community has lead to concern about how much the results depend on this algorithm. We conclude that with respect to the Riemann solver their results are accurate. However the efficiency of the HLLC Riemann solver of Bouchut suggests that it may be used instead.   {\\bf Acknowledgements} The authors would like to thank Prof. Jens C. Niemeyer at the University of W\\\"urzburg for valuable comments and suggestions. We also thank Christoph Federrath for performing the post-processing of the turbulence simulation data."
    },
    "0711/0711.1880_arXiv.txt": {
        "abstract": "We describe in this paper a new binary millisecond pulsar, PSR J1738+0333. Using Arecibo, we have achieved good timing accuracy for this object, about 220 ns for 1-hour integrations over 100~MHz. This allowed us to measure a precise proper motion, parallax and orbital parameters for this system. We highlight the system's potential for constraining alternative theories of gravitation. ",
        "introduction": "PSR~J1738+0333 is a 5.85-ms pulsar in a binary system with an orbital period of 8.5 hours and a companion white dwarf (WD) with a mass of about 0.2 M$_{\\odot}$. This millisecond pulsar (MSP) was found with the Parkes 64-m Radio Telescope in a 20-cm Multi-Beam search for pulsars in intermediate Galactic latitudes ( $5^\\circ < | b | < 30^\\circ$) \\cite{jbo+07}; we have been timing it with Arecibo for the last 4 years using the Wide-band Arecibo Pulsar Processors (WAPPs, \\cite{dsh00}). We have obtained a TOA residual rms of 220 ns per WAPP per hour. This pulsar will be used in the array that is being used to search for nano-Hertz gravitational waves. ",
        "conclusions": ""
    },
    "0711/0711.0973_arXiv.txt": {
        "abstract": "We investigate magnetic field amplification in a turbulent velocity field with nonzero helicity, in the framework of the kinematic Kazantsev-Kraichnan model. We present the numerical solution of the model for the practically important case of Kolmogorov distribution of velocity fluctuations, with a large magnetic Reynolds number. We find that in contrast to the nonhelical case where growing magnetic fields are described by a few bound eigenmodes concentrated inside the inertial interval of the velocity field, in the helical case the number of bound eigenmodes considerably increases; moreover, new unbound eigenmodes appear. Both bound and unbound eigenmodes contribute to the large-scale magnetic field. This indicates a limited applicability of the conventional alpha model of a large-scale dynamo action, which captures only unbound modes. ",
        "introduction": "Magnetic fields in planets and stars, protogalaxies and galaxies, and possibly intergalactic medium are generated due to random stretching of magnetic field lines by turbulent motion of highly conducting fluids or plasmas in which these lines are frozen \\citep[e.g.,][]{lynden-bell,parker,moffatt78,kulsrud1,zweibel,schekochihin,kulsrud2}. It is natural to expect that such a dynamo mechanism can amplify magnetic fields at scales smaller than the correlation scales of the velocity fields. However, magnetic fields observed in astrophysical systems often appear to be correlated at larger scales. Such magnetic fields can be explained if one assumes that the velocity field ${\\bf v}({\\bf x, t})$ possesses nonzero kinetic helicity, that is,  $H=\\int {\\bf v}\\cdot ({\\nabla \\times {\\bf v}})\\, d^3 x\\neq 0$. To describe the large-scale fields, one generally applies the alpha model~\\citep{skr,moffatt78,kulsrud1}. This model is obtained if one averages the induction equation \\begin{eqnarray} \\partial_t {\\bf B}=\\nabla \\times ({\\bf v}\\times {\\bf B})+\\eta \\nabla^2 {\\bf B}, \\label{induction} \\end{eqnarray} where $\\eta$ is resistivity or magnetic diffusivity,  over small-scale fluctuations of the velocity and magnetic fields, assuming that these fluctuations are much weaker and concentrated at the scales much smaller than the scales of the growing large-scale field. As a result one obtains the alpha model equation for the large-scale magnetic field $\\overline{\\bf B}({\\bf x}, t)$, that is, the magnetic field averaged over the scales larger than the scales of velocity fluctuations, \\begin{eqnarray} \\partial_t \\overline{\\bf B}=\\alpha \\nabla \\times \\overline{\\bf B}+\\beta \\nabla^2 \\overline{\\bf B}. \\label{alpha} \\end{eqnarray} In this equation, $\\alpha\\sim \\overline{{\\bf v}\\cdot ({\\nabla \\times {\\bf v}})} \\tau_0$, and $\\beta \\sim v_0 l_0$, where $v_0$ is a characteristic velocity, $l_0$ is the characteristic scale, and $\\tau_0\\sim l_0/v_0$ is a characteristic eddy turnover time (decorrelation time) of fluid fluctuations. In astrophysical systems resistivity is very small, so that turbulent velocity field effectively generates small-scale magnetic fluctuations, which grow much faster than the large-scale field. In this case, the applicability of the mean-field equation~(\\ref{alpha}) is questionable \\citep[see, e.g., the discussion in][]{vainshtein1,diamond,blackman}.  To investigate this question, we use a solvable model of kinematic dynamo, introduced by Kazantsev~(1968) and Kraichnan~(1968). In this model, velocity is assumed to be Gaussian, with zero mean, $\\langle{\\bf v}\\rangle=0$, and the covariance tensor \\beq \\langle {v^i}({\\bf x},t){v^j}({\\bf x}',t') \\rangle \\!=\\! \\kappa^{ij}(|{\\bf x}-{\\bf x}'|)\\delta(t-t'), \\label{V_V_TENSOR} \\eeq where $\\kappa^{ij}$ is an isotropic tensor, \\beq \\kappa^{ij}({\\bf x})\\!=\\!\\kappa_N \\left(\\delta^{ij}-\\frac{x^ix^j}{x^2}\\right)+ \\kappa_L \\frac{x^ix^j}{x^2}+g\\epsilon^{ijk}x^{k}. \\label{KAPPA} \\eeq Here $\\langle \\rangle$ denotes ensemble average, $\\epsilon^{ijk}$ is the unit anti-symmetric pseudo-tensor and summation over repeated indices is assumed. The first two terms on the right-hand side of (\\ref{KAPPA}) represent the mirror-symmetric, nonhelical part, while function $g(x)$ describes the helical part of the velocity fluctuations. For an incompressible velocity field (the only case we are considering here), we have $\\kappa_N(x)=\\kappa_L(x)+x\\kappa'_L(x)/2$, where the prime denotes  derivative with respect to~$x=|{\\bf x}|$. Therefore, to describe the velocity field, we need to specify only two independent functions, say, $\\kappa_L(x)$ and $g(x)$. The magnetic field correlator can similarly be expressed as \\begin{eqnarray} \\langle B^i({\\bf x}, t)B^j(0,t)\\rangle = M_N\\left(\\delta^{ij}-\\frac{x^ix^j}{x^2}\\right) + M_L\\frac{x^ix^j}{x^2}+K\\epsilon^{ijk}x^k, \\label{B_B_TENSOR} \\end{eqnarray} where the corresponding solenoidality constraint implies $M_N(x,t)=M_L(x,t)+(x/2)M'_L(x,t)$. To describe the magnetic correlator we therefore need only two functions, $M_L(x, t)$ and $K(x, t)$, corresponding to magnetic energy and magnetic helicity.  Suppose that the velocity field~(\\ref{V_V_TENSOR}) is given. The problem is then to find the correlation function~(\\ref{B_B_TENSOR}) of the magnetic field. In the nonhelical case, $g(x)\\equiv 0$, it was established by \\citet{kazantsev68} that the problem is reduced to a quantum mechanical problem with imaginary time: \\begin{eqnarray} -\\partial_t \\psi ={\\hat{\\cal H}}\\psi. \\label{quantum} \\end{eqnarray} More precisely, given the velocity correlator~(\\ref{V_V_TENSOR}) and magnetic resistivity $\\eta$, one constructs the self-adjoint Hamiltonian~${\\hat{\\cal H}}$. The magnetic correlator $M_L(x, t)$ is then mapped to the ``wave function''~$\\psi(x, t)$. If the equation~(\\ref{quantum}) has growing solutions $\\psi(x, t)$, corresponding to negative eigenvalues of ${\\hat{\\cal H}}$, then dynamo action is possible. One can show that the smaller the resistivity~$\\eta$, the deeper the effective potential in the quantum mechanical problem~(\\ref{quantum}), and, therefore, nonhelical dynamo is always possible when the magnetic Reynolds number, ${\\rm Rm}\\propto \\eta^{-1}$, is large enough~\\citep{vainshtein,boldyrev-cattaneo,iskakov}. The quantum mechanical representation~(\\ref{quantum}) is important since it ensures that the eigenvalues are real and the eigenfunctions are mutually orthogonal, which allows one to apply a variational principle for estimating dynamo growth rates. The Kazantsev-Kraichnan model of nonhelical dynamo action has been well investigated in the literature.  The situation is more complex when the velocity field possesses nonzero helicity, i.e., $g(x)\\neq 0$. In this case, given kinetic energy~$\\kappa_L(x)$ and kinetic helicity~$g(x)$, one needs to solve {\\em two} coupled partial differential equations for functions $M_L(x,t)$ and $K(x, t)$ related to magnetic energy and magnetic helicity. Such equations were first derived by \\citet{vainshtein}. Due to their complexity, there have been relatively fewer results obtained for the helical case  \\citep[e.g.,][]{pouquet,kulsrud,kim,brandenburg,blackman,brandenburg2}. However, it is the helical case that is practically more important since astrophysical systems generally possess nonzero helicity. Moreover, while magnetic fields at small (velocity) scales are naturally expected on the base of equation~(\\ref{quantum}), explanation of astrophysically observed large-scale magnetic fields commonly requires helicity effects, as, e.g., in equation~(\\ref{alpha}).  Recently, it has been established in \\citet{boldyrev05} that the \\citet{vainshtein} equations also possess a self-adjoint structure, and can be reduced to a quantum mechanical ``spinor'' form: \\begin{eqnarray} -\\partial_t \\psi^{\\alpha} ={\\hat{\\cal H}}^{\\alpha \\beta}\\psi^{\\beta}, \\label{spinor} \\end{eqnarray} where $\\alpha=\\{1,2\\}$ and summation over repeated indices is assumed. Similarly to equation~(\\ref{quantum}), the self-adjoint Hamiltonian ${\\hat{\\cal H}}^{\\alpha\\beta}$ depends on kinetic energy and kinetic helicity, $\\kappa_L(x)$ and $g(x)$, and on magnetic resistivity~$\\eta$. The two components of the $\\psi^{\\alpha}(x, t)$ function are then related to functions $M_L(x,t)$ and $K(x, t)$ in the magnetic correlator~(\\ref{B_B_TENSOR}).  Similar structures of the dynamo equations~(\\ref{quantum}) and~(\\ref{spinor}) allow one to investigate them on the same footing. In particular, one can compare the growth rates and the eigenvalues associated with helical and nonhelical dynamo action, and address the important question of whether the large-scale magnetic field generated due to helical dynamo action is described by the alpha model~(\\ref{alpha}). This is the goal of the present work. We assume that the velocity field has the Kolmogorov spectrum and that the Reynolds number and the magnetic Reynolds number are large. Then we present the full numerical solution of the helical dynamo model~(\\ref{spinor}), i.e., we find its eigenvalues and eigenfunctions. We demonstrate that at least at the kinematic stage, the alpha model~(\\ref{alpha}) may provide a nonadequate description of large-scale fields, since it misses the rapidly growing large-scale eigenmodes. In the next section we present our main results; a detailed discussion will be presented elsewhere. nonhelical ",
        "conclusions": "We find that when the Reynolds number is of order unity (a single-scale velocity field), the magnitude of the kinetic helicity parameter $h$ does not have much effect on the bound magnetic eigenmodes (plots~A and~B in Fig.~\\ref{FIGURE_1}). However, in the case when the Reynolds number is large and the velocity fluctuations extend over a large range of scales, the bound eigenmodes are significantly affected by kinetic helicity at scales larger than the viscous scale. (Magnetic fluctiations at the scales much smaller than the viscous scale, which are excited in the case ${\\rm Pr}>>1$, are not significantly affected by magnetic helicity, which is consistent with previous considerations; e.g.,~\\citet{kulsrud}.) The number of bound eigenmodes increases considerably when the kinetic helicity increases, and the corresponding eigenvalues, $\\lambda_n$, become strongly concentrated near the eigenvalue of the fastest unbound eigenmode, $\\lambda_0$. This result follows from a nearly uniform distribution of $\\lambda_n$ on the logarithmic-scale plots in Fig.~\\ref{FIGURE_1}.  Moreover, we observe an important fact that in all the cases in Fig.~\\ref{FIGURE_1} the growth rate of the first bound eigenmode,~$\\lambda_1$, happens to be very close to the growth rate of the fastest growing unbound eigenmode,~$\\lambda_0$.\\footnote{We found same result for all other high-Reynolds cases that we investigated (not reported here), with different helicity parameters.} We conjecture that for high Reynolds numbers the potential in~(\\ref{W2_AND_W3}) always has a  shallow bound state~$\\lambda_1$ (such that $\\lambda_1 -\\lambda_0\\ll \\lambda_0$) whose characteristic scale is much larger than the velocity correlation scale, due to~(\\ref{W_LARGE_X_TRAPPED}). This has an important consequence for the dynamo mechanism. To understand it, we note that the alpha dynamo model~(\\ref{alpha}), describing the large-scale magnetic field, does not capture the bound eigenmodes. However, our analysis shows that shallow bound eigenmodes have faster growth rates and large correlation lengths, so at a given scale~$x$ they may rapidly become dominant over the unbound modes. In practical applications, this means that such modes rather than the modes described by~(\\ref{alpha}) become essential in the large-scale magnetic field configurations. In this case the conventional alpha-dynamo model~(\\ref{alpha}) leads to an inadequate description of the large-scale magnetic field even at the kinematic stage of dynamo action.  We are grateful to Fausto Cattaneo and Alexander Obabko for many fruitful discussions. This work was supported by the NSF Center for Magnetic Self-Organization in Laboratory and Astrophysical Plasmas at the Universities of Chicago and Wisconsin--Madison, and by the US Department of Energy under grant DE-FG02-07ER54932. We thank the Aspen Center for Physics, where a part of this work was done,  for hospitality and support."
    },
    "0711/0711.2976_arXiv.txt": {
        "abstract": "{Iron line emission is common in the X-ray spectra of accreting black holes. When the line emission is broad or variable then it is likely to originate from close to the black hole. X-ray irradiation of the accretion flow by the power-law X-ray continuum produces the X-ray 'reflection' spectrum which includes the iron line. The shape and variability of the iron lines and reflection can be used as a diagnostic of the radius, velocity and nature of the flow. The inner radius of the dense flow corresponds to the innermost stable circular orbit and thus can be used to determine the spin of the black hole. Studies of broad iron lines and reflection spectra offer much promise for understanding how the inner parts of accretion flows (and outflows) around black holes operate. There remains great potential for XMM-Newton to continue to make significant progress in this work. The need for high quality spectra and thus for long exposure times is paramount.} ",
        "introduction": "Most of the radiation from luminous accreting black holes is released within the innermost 20 gravitational radii (i.e. $20 r_{\\rm g}=20 GM/c^2$). In such an energetic environment, iron is a major source of line emission, with strong lines in the 6.4--6.9~keV band depending on ionization state. Observations of such line emission then provides us with a diagnostic of the accretion flow, conditions and the strong gravity in this regime (Fabian et al 2000; Reynolds \\& Nowak 2003; Matt 2006; Miller 2007).  The rapid X-ray variability found in many Seyfert galaxies requires that  the emission orginates at small radii. The high frequency break in their power spectra, for example, corresponds to orbital periods at $\\sim20 r_{\\rm g}$ and variability is seen at still higher frequencies (Uttley \\& McHardy 2004; Vaughan et al 2005). Key evidence that the very innermost radii are involved comes from Soltan's (1982) argument relating the energy density in radiation from active galactic nuclei (AGN) to the local mean mass density in massive black holes, which are presumed to have grown by accretion which liberated that radiation. The agreement found between these quantities requires that the radiative efficiency of accretion be 10 per cent or more (Yu \\& Tremaine 2002; Marconi et al 2005). This exceeds the 6 per cent for accretion onto a non-spinning Schwarzschild black hole and inevitably implies that most massive black holes are rapidly spinning with accretion flows extending down to a just few $r_{\\rm g}$. Moreover, this is where most of the radiation in such accretion flows originates.  \\begin{figure}[] \\includegraphics[width=100mm]{profiles.ps} \\caption{The profile of an intrinsically narrow emission line is modified by the interplay of Doppler/gravitational energy shifts, relativistic beaming, and gravitational light bending occurring in the accretion disc (from Fabian et al 2000). The upper panel shows the symmetric double--peaked profile from two annuli on a non--relativistic Newtonian disc. In the second panel, the effects of transverse Doppler shifts (making the profiles redder) and of relativistic beaming (enhancing the blue peak with respect to the red) are included. In the third panel, gravitational redshift is turned on, shifting the overall profile to the red side and reducing the blue peak strength. The disc inclination fixes the maximum energy at which the line can still be seen, mainly because of the angular dependence of relativistic beaming and of gravitational light bending effects.  All these effects combined give rise to a broad, skewed line profile which is shown in the last panel, after integrating over the contributions from all the different annuli on the accretion disc. Detailed computations are given by Fabian et al (1989), Laor (1991), Dovciak et al (2004) and Beckwith \\& Done (2004). } \\end{figure}  \\begin{figure}[] \\includegraphics[height=80mm,angle=-90]{xrayspec1.ps} \\caption{ The main components of the X--ray spectra of an unobscured accreting BH are shown: soft X--ray emission from the accretion disc; power law from Comptonization of the soft X--rays in a corona above the disc; reflection continuum and narrow Fe line due to reflection of the hard X--ray emission from dense gas.  } \\end{figure}  The X-ray spectra of AGN are characterized by several components: a hard power-law which may turnover at a few hundred keV, a soft excess and a reflection component (Fig.~1). This last component is produced from surrounding material by irradiation by the power-law. It consists of backscattered X-rays, fluorescence and other line photons, bremsstrahlung and other continua from the irradiated surfaces. Examples of reflection spectra from photoionized slabs are shown in Fig.~2. At moderate ionization parameters ($\\xi=F/n\\sim 100$~erg~cm~s$^{-1}$, where $F$ is the ionizing flux and $n$ the density of the surface) the main components of the reflection spectrum are the Compton hump peaking at $\\sim 30$~keV, the iron line at 6.4-6.9~keV (depending on ionization state) and a collection of lines and reradiated continuum below 1~keV. When such a spectrum is produced from the innermost parts of an accretion disk around a spinning black hole then the outside observer sees it smeared and redshifted (Fig.~3) due to doppler and gravitational redshifts into the soft excess, the broad iron line and the Compton hump.  \\begin{figure}[] \\includegraphics[height=80mm,angle=-90]{rossreflion.ps} \\includegraphics[height=80mm,angle=-90]{refblur.ps} \\caption{{Upper panel}: Computed X--ray reflection spectra as a function of the ionization parameter $\\xi$ ( Ross \\& Fabian 2005). The illuminating continuum has a photon index of $\\Gamma=2$ and the reflector is assumed to have cosmic (solar) abundances. {\\it Lower panel}: Relativistic effects on the observed X--ray reflection spectrum (solid line).  We assume that the intrinsic rest--frame spectrum (dotted) is emitted in an accretion disc and suffers all the relativistic effects shown in Fig.~3. } \\end{figure}    All three main parts of the reflection spectrum have now been seen from some AGN and Galactic Black Holes (GBH). The broad iron line and reflection hump are clearly seen in the Seyfert galaxy MCG--6-30-15 and in the GBH J1650-400 (Fig.~4). More recently it has been realised that the soft excess in many AGN can be well explained by smeared reflection (Crummy et al 2005).  \\begin{figure}[] \\includegraphics[height=80mm,angle=-90]{diskline.ps} \\caption{ The line profile dependence on the inner disc radius is shown for the two extreme cases of a Schwarzschild BH (red, with inner disc radius at $6~r_g$) and of a Maximal Kerr BH (blue, with inner disc radius at $\\simeq 1.24~r_g$) } \\end{figure} \\begin{figure}[] \\includegraphics[height=80mm,angle=-90]{mcg6_linebw.ps} \\caption{ The broad iron line in MCG--6-30-15 from the XMM observation in 2001 (Fabian et al.\\ 2002a) is shown as a ratio to the continuum model. } \\end{figure}  The extent of the blurring of the reflection spectrum is determined by the innermost radius of the disk (Fig.~4). Assuming that this is the radius of marginal stability then enables the spin parameter $a$ of the hole to be measured. Objects with a very broad iron line like MCG--6-30-15 are inferred to have high spin $a>0.95$ (Dabrowski et al 1997, Brennemann et al 2006). Some (Krolik \\& Hawley 2002) have argued that magnetic fields in the disk can blur the separation between innermost edge of the disk and the inner plunge region so that the above assumption is invalid. This probably makes little difference however since the reflection comes from very dense matter. Matter in the plunge region very rapidly drops to a low density (Reynolds \\& Begelman 1997) and only very strong fields, much larger than are inferred in disks, can stop this steep decline in density. ",
        "conclusions": "Clear examples of relativistically-broadened iron lines are seen in some AGN and GBH in some states. Such objects must have dense inner accretion disks in order that the gas is not overionized. Detection of a line is helped greatly if the iron abundance is super-Solar and if there is little extra absorption due to very strong warm absorbers or winds. Where broad lines are seen and can be modelled satisfactorally then the spin of the central black hole can be reliably determined.  The study of absorption and emission variability of iron-K lines is in its infancy, with some interesting and tantalising results produced so far. The inner regions of accretion flows are bound to be structured and so give rise to variations (e.g. Armitage \\& Reynolds 2003). Some may be due to motion or transience in the corona or primary power-law source while others may reflect structure, e.g. spiral waves, on the disk itself, intercepting primary radiation from much smaller radii.  \\begin{figure}[] \\includegraphics[height=80mm,angle=-90]{waterfall.ps} \\caption{ Line variability expected from a turbulent accretion disc (Armitage \\& Reynolds 2003).  } \\end{figure} \\begin{figure}[] \\includegraphics[width=80mm]{3516data.ps} \\caption{Iron line variability, suggestive of periodic behaviour, see in an XMM observation of NGC\\,3516 (Iwasawa et al 2004).  } \\end{figure}  What XMM-Newton, and Suzaku, have established is the reality and relative ubiquity of broad iron lines and other blurred reflection features. We have a tool with which to study the inner workings of the ultimate powerhouse in the Universe, the accreting black hole. To examine variability on the light crossing time of the innermost orbits of the black holes in the bright AGN is beyond XMM, since square-metre-class collecting area is required, but there are enough hints from the existing patchy observations for a) the problem to be tackled in a statistical way, b) more black hole spins to be well established, c) the inner regions of jetted sources to be explored, and d) the position and/or motion of the power-law source to be tracked. We need to map how the energy is extracted from the accretion flow and deposited into the power-law generating corona, presumably by magnetic fields. This is not a minor process but accounts for at least 12 per cent of the accretion power in high Eddington ratio sources and 50 per cent or so in the low Eddington ones ($L_{\\rm Bol}/L_{\\rm Edd}<0.1$; Vasudevan \\& Fabian 2007).  This requires us to make much deeper exposures on many more suitable AGN, to push them over the 150,000 ct boundary, an approach being pursued by Matteo Guainazzi and colleagues. This does of course mean long exposures on the next fainter level of sources. I favour very long exposures on the brightest sources. Ms exposures on MCG--6-30-15, for instance, or of the brightest NLS1s such as 1H0707 or IRAS\\,13224. We are only just beginning to explore these objects and parameter space is ripe for serendipitous discovery as much as nailing down the parameters we expect. As a community, we owe it to the legacy of XMM-Newton to propose, accept and follow through such work which will be invaluable for understanding the most powerful quasi-steady sources in the Universe. It cannot be done any other way!"
    },
    "0711/0711.0697_arXiv.txt": {
        "abstract": "We review the status of a certain (infinite) class of four-dimensional generally covariant theories propagating two degrees of freedom that are formulated without any direct mention of the metric. General relativity itself (in its Pleba\\'nski formulation) belongs to the class, so these theories are examples of modified gravity. We summarize the current understanding of the nature of the modification, of the renormalizability properties of these theories, of their coupling to matter fields, and describe some of their physical properties. ",
        "introduction": "General relativity (GR) was formulated by Einstein as a theory that dynamically determines the spacetime metric. This theory is famously non-renormalizable, which means that it gives only the ``low-energy'' description of gravity. Because the Newton's constant is so small, what is low energy for gravity may be very high energy by the standards of particle physics, so for all practical purposes we are quite happy with the description of gravity given by GR. The problem of quantum gravity, which is to find a satisfactory description of gravity at ``high'' (Planckian) energy is then a very hard one as it is not possible, and is unlikely to be possible, to probe the relevant range of energies by direct experiments. In spite of our almost absolute ignorance as to what happens at Planckian energy, the common consensus is that it is unlikely that the metric description of gravity survives there. For this reason many alternative descriptions have been developed, some of them rather radical, e.g., even calling for abandoning of the notion of the spacetime manifold. The most popular approach to the problem is given by string theory, which is well defined as a quantum theory at the expense of introducing extra symmetries and dimensions that so far are not observed.  It is not universally appreciated, however, that general relativity itself admits a rather radical reformulation that almost entirely eliminates the spacetime metric from the picture. This formulation was discovered a long ago by Pleba\\'nski \\cite{Plebanski}, and was rediscovered more recently in \\cite{CDJ} after the work on the new Hamiltonian formulation for gravity by Ashtekar \\cite{Ashtekar:1987gu}. Ashtekar's new formulation of GR can be seen \\cite{CDJ} to be just the Hamiltonian formulation of the Pleba\\'nski theory. The Plebanski formulation, together with the idea of describing gravity by something else than the spacetime metric, leads to a class of gravity theories much larger than GR, which is the subject of this review.  In Pleba\\'nski formulation of gravity, the notion of the dynamical spacetime metric is replaced by that of the dynamical Hodge operator. Recall that, given a metric $g_{\\mu\\nu}$, the Hodge operator can be defined as the map acting on the space of two-forms and mapping a two-form $B_{\\mu\\nu}$ into its ``dual'' two-form $(*B)_{\\mu\\nu} = (1/2) \\epsilon_{\\mu\\nu}^{\\,\\,\\,\\,\\rho\\sigma}B_{\\rho\\sigma}$. Here $\\epsilon_{\\mu\\nu\\rho\\sigma}$ is the volume four-form compatible with the metric. It is easy to check that the Hodge dual operator is invariant under conformal rescalings of the metric $g_{\\mu\\nu}\\to\\Omega g_{\\mu\\nu}$. The key fact is that the converse is also true and two metrics that define the same Hodge operator are related by a conformal transformation, see, e.g., \\cite{DKS} for a simple proof. This means that the Hodge dual operator defines a metric up to conformal transformations, and that it can be used as the main dynamical object in a theory of gravity, instead of the metric.  This is essentially the idea that was realized in \\cite{Plebanski} (even though it is only implicit in this work). More precisely, to describe the Hodge operator it is sufficient to specify which two-forms are self-dual. The self-dual forms are those satisfying $*B=\\sqrt{\\sigma} B$, where $\\sigma=*^2$, so $\\sigma=\\pm 1$ for the Euclidean and Lorentzian signatures, respectively. Knowing the subspace $W^+$ of self-dual two-forms, the anti-self-dual forms $*B=-\\sqrt{\\sigma} B$ are found as those orthogonal to self-dual ones with respect to the scalar product in the space of two-forms given by the wedge product (the scalar product itself depends on the choice of a volume form; the notion of the orthogonality, however, does not). The knowledge of the subspaces $W^{\\pm}$ is equivalent to the knowledge of the Hodge operator, as is not hard to see.  The description then proceeds as follows. One introduces three two-form fields $B^i$. Declaring the subspace (in the space of two-forms) spanned by the fields $B^i$ to be that of self-dual two-forms determines the metric up to a conformal factor. This last is fixed by a choice of the volume form, for which there is a natural choice $(vol)=(1/3)\\delta^{ij}B^i\\wedge B^j$, where $\\delta_{ij}$ is the Kronecker delta. The metric defined by this data is then invariant under ${\\rm SO}(3)$ rotations of the triple $B^i$. This suggests that we should think of the triple $B^i$ as of a section of a vector bundle $V\\to M$ associated with a principal ${\\rm SO}(3)$ bundle over the spacetime manifold $M$. Therefore, in addition to $B^i$, we should introduce a connection $A^i$ on $V$. It is now not hard to write the field equations for the theory. First, it is natural to require the connection $A^i$ to be compatible with the triple $B^i$, i.e., require $(D_A B)^i=0$. Given a triple $B^i$, this set of $4\\times 3$ equations can be solved for $4\\times 3$ components of $A^i$ in terms of derivatives of $B^i$. It remains to write an equation that would allow one to find $B^i$. It is natural to require that this equation be second order in derivatives and ${\\rm SO}(3)$ covariant. This points to some equation that involves $F^i$(A), the curvature of the connection $A^i$. The obvious choice $F^i(A)=0$ does give $6\\times 3$ equations for $6\\times 3$ components of $B^i$. However, the theory that one gets this way is not very interesting as it is void of any physics (does not have local degrees of freedom). A much more interesting choice is: \\be\\label{int-1} F^i(A) = \\Lambda^i_j B^j, \\ee where $\\Lambda^i_j$ is some (undetermined) matrix. This equation says that the curvature is purely self-dual as a two-form. This gives us $3\\times 3$ equations $(F^i(A))_{asd}=0$ for $6\\times 3$ components of $B^i$ so the theory is not well-defined as we cannot solve for the two-form fields. One can make it into a well-defined theory by writing an action principle that leads to (\\ref{int-1}), namely: \\be\\label{int-2} S[B,A,\\Lambda]= \\int B_i \\wedge F^i(A) - \\frac{1}{2} \\Lambda_{ij} B^i\\wedge B^j, \\ee where we raise and lower indices using $\\delta^{ij}$. Varying this action with respect to $A^i$ we get the equation $(D_A B)^i=0$, while varying it with respect to $B^i$ gives (\\ref{int-1}). Note that only the symmetric part of $\\Lambda_{ij}$ enters equation (\\ref{int-2}), so this matrix should be assumed symmetric also in (\\ref{int-1}). The action (\\ref{int-2}) tells us that we should treat $\\Lambda_{ij}$ as a dynamical object and vary with respect to it. We get: $B^i\\wedge B^j=0$, which gives six equations for $B^i$. Together with (\\ref{int-1}) this gives 15 equations. Noting that three components of $B^i$ are pure gauge we thus get the right number of equations to determine $B^i$ completely. The components of the symmetric matrix $\\Lambda_{ij}$ are determined from the unused equations in (\\ref{int-1}). Note, however, that the equation $B^i\\wedge B^j=0$ implies, in particular, that the volume form $(vol)$ constructed above is zero. Also, it is rather obvious that there is no Hodge dual operator for which the space $W^+$ of self-dual two-forms is null, $W^+\\wedge W^+=0$. So, this theory is not of any physical interest.\\footnote{Its Hamiltonian analysis shows that this is simply GR with the Hamiltonian constraint removed.}  To get an interesting theory, we will impose one constraint on the symmetric matrix $\\Lambda_{ij}$ that appears on the right-hand-side of (\\ref{int-1}). If we do so, then only five of the six components of $\\Lambda_{ij}$ remain independent. Equation (\\ref{int-1}) then gives ten equations for the components of $B^i$ (this is the number of equations in GR) plus five equations for the independent components of $\\Lambda_{ij}$ in terms of the second derivatives of $B^i$. Variation of the action (\\ref{int-2}) with respect to the independent components of $\\Lambda_{ij}$ gives an additional five equations on $B^i$, which gives just enough equations to determine the two-form fields.  What constraint can be imposed on $\\Lambda_{ij}$? It is clear that the five independent components of $\\Lambda_{ij}$ can only be chosen to be those of its {\\it traceless\\/} part, which we denote by $\\Lambda_{ij}-(1/3)\\delta_{ij} {\\rm Tr}(\\Lambda):=\\Psi_{ij}$. Indeed, this is the part that forms an irreducible representation with respect to the action of ${\\rm SO}(3)$. The trace part ${\\rm Tr}(\\Lambda)$ then becomes a function of its traceless part $\\Psi_{ij}$ that we denote by $-\\phi(\\Psi)$ (the minus sign is to agree with certain conventions; see below). Thus, we have: \\be\\label{lambda} \\Lambda_{ij} = \\Psi_{ij} - \\frac{1}{3} \\delta_{ij} \\phi(\\Psi). \\ee  The theory defined by the action (\\ref{int-2}) where only the traceless part $\\Psi_{ij}$ of $\\Lambda_{ij}$ is an independent variable (\\ref{lambda}) gives the right number of equations to solve for the components of the two-form field $B^i$ and thus determine the Hodge dual operator, which in turn (together with the volume form) determines the metric. We also note that the volume form $(1/3)\\delta_{ij}B^i\\wedge B^j$ is no longer trivial: the equation one obtains by varying the action (\\ref{int-2}) with respect to $\\Psi_{ij}$, \\be\\label{metricity} B^i\\wedge B^j - \\frac{1}{3} \\delta^{ij} (B^k\\wedge B_k) = \\frac{1}{3} \\frac{\\partial\\phi}{\\partial \\Psi_{ij}} (B^k\\wedge B_k), \\ee no longer implies any relation for the volume form. We thus get a theory with potentially interesting physical implications. It becomes even more interesting after one observes that, for $\\phi(\\Psi)=\\Lambda=\\mbox{const}$, the above theory is nothing else but the Einstein's GR in disguise (i.e., in its Pleba\\'nski formulation), with $\\Lambda$ being the cosmological constant (our choice of the sign in (\\ref{lambda}) was motivated precisely by the desire to agree with the usual convention for the sign of $\\Lambda$). Therefore, what we have obtained is an infinite family of generalizations of general relativity, where to fix a theory one has to specify the function $\\phi(\\Psi)=\\phi({\\rm Tr}(\\Psi^2),{\\rm Tr}(\\Psi^3))$ of two invariants that can be constructed from the traceless matrix $\\Psi$. The above counting of the field equations suggests that this theory makes sense for an arbitrary $\\phi(\\Psi)$. A more detailed analysis, whose results we will review in the next section, shows that this is indeed the case.  Our above discussion was motivated by the idea to describe gravity as a dynamical theory of an object other than the spacetime metric. We have seen how GR itself can be reformulated in these terms, and how an infinite class of gravity theories different from GR and parametrized by a function $\\phi(\\Psi)$ is obtained this way. In the next section we will review the (classical) properties of this class of theories without worrying too much about the principle that can fix $\\phi$. We will see that, for an arbitrary $\\phi$, the gravity theory (\\ref{int-2}) resembles GR in many ways. Still, there are some new physical effects predicted by this theory. It is important to emphasize that, at the level of the classical physics, nothing forces us to depart from the familiar ground of GR, and the new class of theories described above may be viewed as not more than a mathematical curiosity. However, when one considers the quantum theory of usual GR in Pleba\\'nski (Hodge operator) formulation, one {\\it is\\/} forced to consider theories more general than GR, similar to what happens in quantum gravity in the metric formulation. In the usual metric formulation of quantum gravity, one has to introduce theories with higher derivatives and, thus, with rather unpleasant properties (instability etc.). In the Hodge operator formulation of gravity, many (if not all; see Section \\ref{sec:quantum}) terms that have to be added to the action do not change the character of the theory so drastically, i.e., do not introduce higher derivatives and new DOF. The quantum corrected theory is one of the class (\\ref{int-2}) with (\\ref{lambda}). Thus, one does have to take this class of gravity theories seriously, as the modifications they describe will be induced by quantum corrections. It is a different question whether any of these quantum corrections can survive and be of relevance at low (astrophysical) energy scales. It can only be answered after the quantum mechanical behavior of this class of theories is understood. But even in the absence of such understanding, it is interesting to quest what kind of new physical effects can be expected.  The theory (\\ref{int-2}) with (\\ref{lambda}) was first proposed in \\cite{Krasnov:2006du}, where renormalizability properties of GR in Pleba\\'nski formulation were considered. It was only later appreciated that it is also quite interesting as a purely classical theory, the point of view which was developed in \\cite{Krasnov:2007uu}, \\cite{Krasnov:2007ky}. The Hamiltonian analysis of this theory (in a slightly different version) was done in \\cite{Bengtsson:2007zx} and more recently in \\cite{Krasnov-recent}. We will describe the main results of the classical analysis in the next section. The current understanding of the quantum mechanical behavior of this class of theories is reviewed in Section \\ref{sec:quantum}. ",
        "conclusions": "A class of gravitational theories propagating two degrees of freedom and formulated without any direct reference to the spacetime metric was described. A theory from this class is specified by an (arbitrary) function $\\phi$ of two arguments. General relativity belongs to the class considered and corresponds to $\\phi=\\mbox{const}=\\Lambda$, $\\Lambda$ being the cosmological constant. One way to describe the nature of the modification is to say that the cosmological constant became a non-trivial function of the ``curvature''. However, the theories we consider cannot be obtained by simply inserting a scalar function of the Weyl curvature into the Lagrangian. The construction is much more subtle, and involves, in particular, replacing the paradigm of gravity being a dynamical theory of the metric by a new paradigm of gravity being about the dynamics of the Hodge dual operator. We have described some simple physical consequences of this modification of gravity, as well as reviewed the status of the question of coupling these theories to matter. Thus, gauge fields couple seamlessly, but no descriptions for fermions in this framework is as of yet known. Work is currently in progress on this very important issue.  It is tempting to try to apply the new class of gravity theories to the fundamental problems of cosmology, namely those of ``dark matter'' and ``dark energy''. We have indicated how this might be possible, but much more work is needed, in particular on coupling of these theories to massive matter fields, to see if any realistic model can be built along these lines.  We have reviewed a conjecture to the effect that the described class of theories (with varying $\\phi$) is closed under the renormalization. Some arguments in support of this conjecture were given, but its status remains open."
    },
    "0711/0711.1234_arXiv.txt": {
        "abstract": "We have measured the HI power spectrum of the nearly face-on spiral galaxy NGC628 (M74) using a visibility based estimator. The power spectrum is well fitted by a power law $P(U)=AU^{\\alpha}$, with $\\alpha\\ =-\\ 1.6\\pm0.2$ over  the length scale $800 \\, {\\rm pc} \\, {\\rm to} \\, 8 \\, {\\rm kpc}$. The slope is found to be independent of the width of the velocity channel. This value of the slope is a little more than one in excess of what has been seen at considerably smaller length scales  in  the Milky-Way, Small Magellanic Cloud (LMC), Large Magellanic Cloud (SMC) and the dwarf galaxy  DDO210. We interpret this difference as indicating a transition from three dimensional turbulence at  small scales to two dimensional turbulence in the plane of the galaxy's disk at  length scales larger than galaxy's HI scale height. The slope measured here is similar to that found at large scales in the  LMC. Our analysis also places an upper limit to the galaxy's scale height  at $800\\ {\\rm pc}$ . ",
        "introduction": "\\label{sec:intro} Evidence has been mounting in recent years that turbulence plays an important role in the physics of the ISM as well as in governing star formation. It is believed that turbulence is responsible for generating the hierarchy of structures present across  a range of spatial scales in the ISM (e.g. \\citealt{ES04I}; \\citealt{ES04II}). In such models the ISM has a fractal structure and the power spectrum of  density fluctuations is a power law, indicating that there is no preferred ``cloud\" size.  On the observational front, the power spectrum analysis of  HI intensity fluctuations is  an important technique to probe the structure of the neutral ISM in galaxies.   The power spectrum analysis in our own Galaxy, Large Magellanic Cloud(LMC) and Small Magellanic Cloud(SMC) finds  the  power spectrum of the HI intensity fluctuation to be a power law(\\citealt{CD83}; \\citealt{GR93}; \\citealt{DD00};  \\citealt{EK01}; \\citealt{SSD99}) which is a characteristic of a turbulent medium.  Similarly, Westpfahl et al.(1999)  showed that the HI distribution in several galaxies  in the M81 group is a fractal. Further, Willett et al.(2005) used the Fourier transform  power spectra of  the V and H$\\alpha$ images of a sample of irregular galaxies to show that the power spectra in optical and H$\\alpha$ pass-bands are also  power law, indicating that  there is no characteristic mass or luminosity scale for OB associations and star complexes.  Recently \\citet{AJS06} presented a visibility based formalism for determining the power spectrum of  HI intensity fluctuations in galaxies  whose emission is extremely weak. This was applied to a dwarf galaxy, DDO 210. Interestingly, the HI power spectrum of this extremely faint, largely quiescent galaxy  was found to be a power law with the same slope as that observed in much brighter galaxies. In this paper the same  formalism is used to  measure the power spectrum of HI intensity fluctuation in the nearby spiral galaxy NGC628.  NGC628(M74) is a nearly face-on  SA(s)c spiral galaxy with an inclination angle in the range  $6^{\\circ}$ to $ 13^{\\circ}$ \\citep{KB92}. It has a very large HI disk extending out to more than 3 times the Holmberg diameter \\citep{KB92}.  Elmegreen et al. (2006) have found a scale-free size and luminosity  distribution of  star forming regions in this galaxy, indicating turbulence to be functional here. The distance to this galaxy is uncertain with previous estimates ranging from $6.5 \\, {\\rm Mpc}$ to $ 10 \\, {\\rm Mpc}$. \\citealt{BW80} and \\citealt{KB92} have used a Hubble flow distance of $\\sim$10 Mpc. On the other hand, \\cite{SK96} estimated a distance of $7.8\\pm 0.9$ Mpc from  the brightest blue star in the galaxy. This distance estimate matches with an independent photometric distance estimate by \\citet{SD96}.  In a recent study \\citet{VI04} inferred the  distance to be $6.7 \\pm  4.5 $ Mpc by applying  the expanding  photosphere method to the hyper novae SN2002ap. Throughout this paper we adopt the photometric distance of 8 Mpc  for NGC628. At this distance 1$^{\\prime\\prime}$ corresponds to 38.8 pc.  In this paper we  present the power spectrum of HI intensity fluctuations of NGC628 derived using the visibility based formalism developed by  \\citet{AJS06}. Studies of external galaxies  like NGC628 which has a very extended HI   disk holds the potential of probing the ISM and  its power spectrum at length scales much larger  than  has been possible  in earlier studies(\\citealt{OY02}).   \\begin{figure} \\begin{center} \\mbox{\\epsfig{file=fig1.ps,width=2.5in,angle=0}} \\end{center} \\caption{The 38$^{\\prime\\prime} \\times 36^{\\prime \\prime}$ resolution integrated HI column density map of NGC628 (contours) overlayed on the optical DSS image (grey scale). The contours levels are 0.24, 1.96, 3.68, 5.41, 7.13, 8.85, 10.57 and 12.29 $\\times 10^{20}$ cm$^{-2}$. } \\label{fig:mom0} \\end{figure} ",
        "conclusions": "\\label{ref:discuss}   Figure \\ref{figure:f2}A shows the real and imaginary parts of $P_{\\rm HI}(U)$, which is the observed value  of the estimator $\\hat{P}_{\\rm HI}(U)$  for the $64$ channels which have HI emission.  As expected from the theoretical considerations mentioned earlier, the imaginary part is well suppressed compared to  the real part. To test for  a possible contribution from residual continuum, we also show the real part of  $P_{\\rm HI}(U)$  using $64$ line free channels. This is found to be much smaller than the signal.  For the channels with HI emission the observed $P_{\\rm HI}(U)$ may be directly interpreted as the HI power spectrum at $U$ values that are considerably larger than $D=0.4 \\, {\\rm K} \\lambda$. We find that a power law $P_{\\rm HI}(U)=A U^{\\alpha}$ with slope  $\\alpha=-1.6 \\pm 0.2$ provides a good fit to the results over the $U$  range $1.0 \\, {\\rm K} \\, \\lambda$ to  $10.0 \\, {\\rm K} \\, \\lambda$ (Figure \\ref{figure:f2}B) which corresponds to spatial scales  of $800 {\\rm pc}$ to $8 \\, {\\rm kpc}$.  Both HI density fluctuations as well as spatial fluctuations in the velocity of the HI gas contribute to fluctuations in the HI specific intensity. Considering a turbulent ISM, \\citet{LP00} have shown that it is possible to disentangle these two contributions  by studying the behavior of the HI power spectrum as the thickness of the frequency  channel is varied. If the observed HI power spectrum is due to the gas velocities, the slope of the  power spectrum is predicted to change as the frequency  channel thickness is increased. We have tested  this  by determining  the HI power spectrum for different values of  the channel width in the range $1.29 \\,  {\\rm   km  \\, s^{-1}}$  to $82.6 \\,  {\\rm   km \\, s^{-1}}$  (Figure \\ref{figure:f2}B)  and do not find any change in the slope of the HI power spectrum (Table \\ref{table:t3}). As the  thickest channel that we have used is considerably wider than the  typical HI  velocity dispersion of $7-10 \\, {\\rm  km s^{-1}}$ seen in spiral galaxies \\citep{SK84}, we conclude that    the observed HI power spectrum of NGC628 is purely due to density fluctuations. Our finding is similar to that  of  \\citet{AJS06} who noticed no change of the slope with channel width for the dwarf  galaxy DDO210.  Further, \\citet{EK01}  also reported a similar behavior for  LMC.  \\begin{table} \\centering \\begin{tabular}{|c|c|c|c|} \\hline N & $\\Delta V$ &$\\alpha $ & $1 \\sigma $\\\\ & (kms$^{-1}$) & &\\\\ \\hline \\hline 1 & 1.3 &-1.7 &  $\\pm 0.2$\\\\ 32 & 42.3 &-1.6 &  $\\pm 0.2$\\\\ 64 & 82.5 &-1.6 &  $\\pm 0.2$\\\\ \\hline \\end{tabular} \\caption{NGC628 power spectrum are well fitted by a power law of $P_{HI}=AU^{\\alpha}$. This table sunrises the result for different channel width.} \\label{table:t3} \\end{table}   Earlier studies of the Milky-Way,  and also of the dwarf galaxies LMC, SMC and DDO210 \\citep{CD83,GR93,SSD99,DD00,EK01,AJS06}  have all found a power law   HI power spectrum with slope $\\sim -3$. On the contrary, we find a  slope $-1.6 \\pm 0.2$ for NGC628. This is a little more than one in excess of the earlier values. However, when comparing these values it should be noted that the earlier works have all measured the HI power spectrum at much smaller length scales in the range  $10\\ \\mathrm{to}\\ 500\\ {\\rm pc}$ [MW ($5 \\pc - 200 \\pc$), SMC($4 \\pc - 30 \\pc$),  DDO210($80 \\pc - 500 \\pc$)] whereas the current measurement probes much larger length scales from $800 \\pc$  to $8.0 \\, {\\rm kpc}$. The typical HI scale heights within the Milky-Way \\citep{LH84,WB90} and external galaxies (e.g. \\citealt{narayan}) are well within $1.5 \\, {\\rm  kpc}$ . This implies that on the  largest length scales which we have probed, the turbulence is definitely confined to the plane of the galaxy's disk and is therefore two dimensional.  \\citet{EK01} have found that the HI power spectrum of LMC flattens  at large length scales, which was interpreted as a transition from three dimensional to   two dimensional turbulence. We conclude that the slope is  different in our observations because it probes two dimensional turbulence, whereas the earlier observations were on length scales smaller than the scale height where we can expect three dimensional turbulence. To the best of our knowledge our results  are the first observational determination of the HI power spectrum of  an external spiral galaxy at such  large length scales which are comparable to the radius of the galaxy's disk.  \\citet{WC99} have performed a fractal analysis using the perimeter-area dimension  of  intensity contours in  HI images of several galaxies in the  M81 group. Of particular interest is the galaxy M81, a spiral galaxy for which the perimeter-area dimension was found to be $\\sim 1.5$ at a length scale $\\sim 10 \\, {\\rm kpc}$.  This observation is consistent with a power law power spectrum of slope $-1.6 \\pm 0.2$ provided  the assumption  that the local dimension has the same value as the perimeter-area dimension is valid.  It is difficult to probe the HI scale height of external face-on galaxies. \\citet{PK01} present a method to probe the scale height from a change in the slope of the Spectral Correlation Function(SCF)[\\citet{RG99}], and applied it to HI data for the LMC to estimate the scale height to be $\\sim 180$pc. \\citet{EK01} suggested that one could use a change in the slope of the power spectrum of the density fluctuations to measure the scale height of face on gas disks. Applying this method to HI data for the LMC they measure scale height of $100$ pc. To the best of our knowledge, prior to this work, there has been no observational constraint on the HI scale height of NGC628. Since the scale height is definitely less than $8 \\, {\\rm kpc}$ and the power spectrum is found to have the same slope from $800 \\pc$ to $8 \\, {\\rm kpc}$, we conclude that the scale height must be less than $800 \\pc$. \\cite{kregel04} present HI images of a large sample of edge on intermediate to late type spirals; from their data the ratio of the HI disk height to the radius of the HI disk (at a column density of 1$M_\\odot$pc${-2}$) is $\\sim 0.06 \\pm 0.15$. From Fig.~\\ref{fig:mom0} the disk of NGC628 has a diameter of $\\sim 28$kpc at this column density. From the average thickness to radius ratio for edge on galaxies, one would expect NGC628 to have a scale height of $\\sim 840$pc, consistent with our result. Future observations of NGC628 with higher angular resolution should be able to put a tighter constraint or even determine the scale height."
    },
    "0711/0711.1002_arXiv.txt": {
        "abstract": "The synchrotron-emitting nebulae formed by energetic winds from young pulsars provide information on a wide range phenomena that contribute to their structure. High resolution X-ray observations reveal jets and toroidal structures in many systems, along with knot-like structures whose emission is observed to be time-variable. Large-scale filaments seen in optical and radio images mark instability regions where the expanding nebulae interact with the surrounding ejecta, and spectral studies reveal the presence of these ejecta in the form of thermal X-ray emission. Infrared studies probe the frequency region where evolutionary and magnetic field effects conspire to change the broadband synchrotron spectrum dramatically, and studies of the innermost regions of the nebulae provide constraints on the spectra of particles entering the nebula. At the highest energies, TeV gamma-ray observations provide a probe of the spectral region that, for low magnetic fields, corresponds to particles with energies just below the X-ray-emitting regime.  Here I summarize the structure of pulsar wind nebulae and review several new observations that have helped drive a recent resurgence in theoretical modeling of these systems. ",
        "introduction": "It has long been known that the Crab Nebula is produced by the wind from a young, energetic pulsar whose spin-down power manifests itself as a synchrotron-emitting bubble of energetic particles.  Optical observations reveal toroidal wisps and jet-like outflows near the pulsar. Faint portions of a ring are observed, marking the region where the pulsar wind undergoes a shock upon entering the surrounding nebula, the outer portions of which are rich in filamentary structure from swept-up ejecta. Though clearly formed in a supernova (SN) event in 1054~CE, the Crab shows no trace of an associated supernova remnant (SNR), presumably indicating that the SN blast wave is propagating through a very low density medium; the nebula results entirely from the pulsar input (and its confinement by surrounding material -- in this case, stellar ejecta) and is the prototypical example of a pulsar wind nebula (PWN).  The diagram shown in Figure~1 illustrates the main points of the most basic picture for a PWN: in an inner zone the wind flows away from the neutron star (NS) with Lorentz factor $\\gamma \\sim 10^6$; at a distance $R_w$ from the NS the wind passes through a termination shock, decelerating the flow while boosting particle energies by another factor of $\\ga 10^3$; and beyond $R_w$ energetic electrons in the wind radiate synchrotron emission in the wound-up toroidal magnetic field to form the PWN which is confined at a radius $R_{PWN}$ by the inertia of the SN ejecta or the pressure of the interior of a surrounding SNR. A detailed theoretical framework, incorporating particle injection and diffusion, magnetic field evolution and radiative and adiabatic losses, has been constructed within this picture, allowing us to successfully predict and explain some basic PWN properties (Reynolds \\& Chevalier 1984; Kennel \\& Coroniti 1984).  For a more detailed review on the structure and evolution of PWNe, see Gaensler \\& Slane (2006).  \\begin{figure}[t] \\includegraphics[width=\\textwidth]{f1.ps} \\caption{ Schematic diagram of a PWN within an SNR (see text for description). \\chandra\\ images of the Crab Nebula, PSR~B1509$-$58, and Vela are shown in the left panels. On the right are corresponding images of the jets in these nebula, with indications of the relative sizes. Here and throughout, north is up and east is to the left. } \\end{figure}  Most models of PWNe have been built and tailored to explain the properties of the Crab, but questions have lingered as to whether or not other PWNe were really well-described by this ``prototype.'' Recent observations, particularly with high resolution X-ray telescopes, have now shown that despite considerable differences in age, environmental conditions, and pulsar properties, a large number of PWNe share many of the Crab's traits.  In particular, it has become startlingly apparent that pulsars release their energy not in an isotropic fashion, but rather in equatorial winds and polar outflows defined by the spin axis of the system. In Figure 1 (left panels) we show \\chandra\\ images of the Crab Nebula (Weisskopf et al.  2000), PSR~B1509$-$58 (Gaensler et al. 2002), and the inner nebula around the Vela pulsar (Helfand et al. 2001). Each shows a well-defined axis of symmetry associated with the spin axis of the pulsar, along with arc-like structures in the equatorial plane. Similar structures have now been identified around a large number of young pulsars.  The symmetry extends to large scales as well, with enhanced confinement by the toroidal field producing nebulae whose extent along the spin axis is larger than at the equator (see Figure 2).  In the extended nebulae, particles suffer synchrotron losses which gradually steepen the spectra with radius, as the higher energy particles are burned off as they diffuse toward the outer boundary -- an effect readily observed in X-ray spectra of PWNe (e.g. Slane et al. 2001, Slane et al.  2004) that directly probes the efficiency of the synchrotron cooling (Chevalier 2000). ",
        "conclusions": "Recent advances in observations of PWNe have uncovered a wealth of information regarding the axisymmetric nature of the wind and the large- and small-scale structure of the nebulae. It is now clear that the Crab Nebula is indeed the prototypical member of this class -- albeit a particularly bright and energetic example; the elongated nebula, swept-up ejecta, toroidal structure, and jets that characterize the Crab are now being found in a multitude of PWNe. Observations across the electromagnetic spectrum are helping to characterize the evolution of PWNe, and opening new questions regarding the nature of the wind and its interactions with its environs. These, in turn, have stimulated new theoretical work on relativistic flows, shocks, and jets with application to a broad range of topics in high energy astrophysics.    \\begin{theacknowledgments} I would like to thank the conference organizers for the opportunity to present this work, and also Bryan Gaensler, Jack Hughes, David Helfand, Steve Reynolds, and my many other colleagues with whom I have worked on studies of pulsar winds. \\end{theacknowledgments}"
    },
    "0711/0711.1372_arXiv.txt": {
        "abstract": "%  Are the organic molecules crucial for life on Earth abundant in early-epoch galaxies?  To address this, we searched for organic molecules in extragalactic sources via their absorption features, known as diffuse interstellar bands (DIBs).  There is strong evidence that DIBs are associated with polycyclic aromatic hydrocarbons (PAHs) and carbon chains.  Galaxies with a preponderance of DIBs may be the most likely places in which to expect life.  We use the method of quasar absorption lines to probe intervening early-epoch galaxies for the DIBs.  We present the equivalent width measurements of DIBs in one neutral hydrogen (\\HI) abundant galaxy and limits for five DIB bands in six other {\\HI}-rich galaxies (damped Lyman-$\\alpha$ systems--DLAs).  Our results reveal that {\\HI}-rich galaxies are dust poor and have significantly lower reddening than known DIB-rich Milky Way environments.  We find that DIBs in {\\HI}-rich galaxies do not show the same correlation with hydrogen abundance as observed in the Milky Way; the extragalactic DIBs are underabundant by as much as 10 times.  The lower limit gas-to-dust ratios of four of the {\\HI}-rich early epoch galaxies are much higher than the gas-to-dust ratios found in the Milky Way.  Our results suggest that the organic molecules responsible for the DIBs are underabundant in {\\HI}-rich early epoch galaxies relative to the Milky Way. ",
        "introduction": "%  Since their discovery in 1921 \\citep{hege22}, the diffuse intersteller bands (DIBs) have remained the longest known interstellar absorption features without a positive identification.  There have been several hundred DIBs discovered to date \\citep{jenn94,tuai00,wese00}.  The DIBs span the visible spectrum between 4000 and 13000~{\\AA}.  Despite no positive identifications, several likely organic molecular candidates have emerged as the sources of the DIBs, including polycyclic aromatic hydrocarbons (PAHs), fullerenes, long carbon chains, and polycyclic aromatic nitrogen heterocycles (PANHs) \\citep{herb95, snow01, cox06a, hudg05}.  The organic-molecular origin of the DIBs may give them an importance to astrobiology; they are now considered an important early constituent to the inventory of organic compounds on Earth \\citep{bada02}.  There are several environmental factors that are known to enhance or inhibit DIB strengths.  Thus, measuring DIB strengths can give clues to the environments of galaxies like those with high neutral hydrogen (\\HI) content, known as damped Lyman-$\\alpha$ systems (DLAs). Measuring DIB strengths allow researchers to explore the quantities of gas, metallicity, dust, and radiation in galaxies and their influence on the strengths of the bands.  Furthermore, observing DIBs in galaxies at a higher redshift, $z$, (distance or lookback time related to the expansion of the Universe) allows us to test the evolution of these organics in cosmic time.  Two important environmental factors that are often probed in galaxies are the {\\HI} content and the reddening.  The {\\HI} content in DLAs is typically measured as a column density via their Lyman-$\\alpha$ line in absorption, as observed using a bright background source such as a quasar.  The column density is the number of atoms of a certain element as seen along a line-of-sight projected to a unit area at the observer (typically in units of atoms cm$^{-2}$ and denoted as $N$(\\HI) for neutral hydrogen).  The reddening is a measure of the dust content in the galactic environment.  Reddening is expressed as $E(B-V)$, or the magnitude of blue light minus visible light observed relative to what is expected from typical stars or galaxies.  Because dust preferentially obscures blue light, a high reddening is a signature for a high dust content.  $N$(\\HI) is a measure of the gas phase and $E(B-V)$ is a measure of the dust phase of the ISM in galaxies.  Due to their relatively weak absorption strengths, the DIBs have been difficult to detect in extragalactic sources.  Aside from the hundreds of detections within the Milky Way \\citep{jenn94,tuai00, wese00}, DIBs have been detected in the Magellanic Clouds \\citep{welt06,cox06b,cox07}, seven starburst galaxies \\citep{heck00}, the active galaxy Centaurus A via supernova 1986A \\citep{rich87}, spiral galaxy NGC 1448 via Supernovae 2001el and 2003hn \\citep{soll05}, one DLA galaxy at $z=0.524$ toward the quasar AO~0235+164 \\citep{junk04, lawt06}, and one galaxy selected by singly ionized calcium (\\CaII), J0013--0024, at $z=0.157$ from the Sloan Digital Sky Survey \\citep{elli07}.  We further the knowledge of DIB lines in extragalactic environments by cataloguing the strengths of the $\\lambda4428$, $\\lambda5780$, $\\lambda5797$, $\\lambda6284$, and $\\lambda6613$ DIBs relative to the $E(B-V)$ and $N$(\\HI) content of each of the seven DLAs in our sample. Observations were obtained, with seven facilities, of seven DLAs toward six QSO sightlines.  The facilities and instruments used for this project are the VLT/FORS2, VLT/UVES, APO/DIS, Keck/HIRES, WHT/ISIS, and Gemini/GMOS-S. ",
        "conclusions": ""
    },
    "0711/0711.3141_arXiv.txt": {
        "abstract": "We present new evolutionary synthesis models for Single Stellar Populations covering a wide range in age and metallicity. The most important difference with existing models is the use of NLTE atmosphere models for the hot stars (O, B, WR, post-AGB stars, and planetary nebulae) that have an important impact in the stellar cluster's ionizing spectra. ",
        "introduction": "We have used the synthesis code by \\cite{gmb98}, updated by \\cite{mgv00} and newly revised now.The basic grid is composed by Single Stellar Populations (SSP) for five different IMF's. The first one is a \\cite{sal55} power law with masses between 0.85 and 120 M$_{\\odot}$. The others IMFs are \\cite{sal55}, \\cite{fpp90}, \\cite{kro02} and \\cite{cha03} functions, with masses between 0.15 and 100 M$_{\\odot}$.  The isochrones are those from \\cite{bgs98} for 6 different metallicities: Z $=$ 0.0004, 0.001, 0.004, 0.008, 0.02 and 0.05. The age coverage is from $log{t}=$ 5.00 to 10.30 with a variable time resolution which is $\\Delta(log{t})=0.01$ in the youngest stellar ages. The WC and WN stars are identified in the isochrones according to their surface abundances.  The atmosphere models are from \\cite{lcb97} with an excellent coverage in effective temperature, gravity and metallicities, for stars with Teff$ \\leq 25000$ K.  For O, B and WR we have taken the NLTE blanketed models by \\cite{snc02} at Z$=$0.001, 0.004, 0.008, 0.02 and 0.04. There are 110 for O-B stars, with 25000 K $ < Teff \\leq 51500$ K and $2.95 \\leq \\log{g} \\leq 4.00$, and 120 for WR stars (60 WN $+$ 60 WC), with 30000 K $\\leq T^{*} \\leq 120000$ K and $1.3R_{\\odot}\\leq R^{*}\\leq 20.3 R_{\\odot}$ for WN, and with $ 40000 K \\leq T^{*} \\leq 140000$ K and $ 0.8R_{\\odot}\\leq R^{*}\\leq 9.3 R_{\\odot}$ for WC. T$^{*}$ and R$^{*}$ are the temperature and the radius at a Roseland optical depth of 10.  To assign a model to each WR star, we use the relationships among opacity, mass loss and wind velocity: $d\\tau=-\\kappa(r)\\rho(r)dr$, where $\\kappa(r)=-0.2(1+X_{\\rm S})$ and $X_{\\rm S}$, the H surface abundance, is taken as 0.2 for WN and 0 for WC. The mass loss is: $\\frac{dM}{dt}=4\\pi r^{2}\\rho(r)v(r)dr$ with $v(r)=v_{\\infty}(1-R_{S}/r)^{\\beta}$ taking $\\beta=1$.  Integrating these equations we find R$^{*}$ and then we select the closer atmosphere model.  For post-AGB and PN with Teff from 50000 to 220000 K we take the NLTE models by \\cite{rau03}. For higher temperatures we use black bodies. ",
        "conclusions": ""
    },
    "0711/0711.1144_arXiv.txt": {
        "abstract": "We study three quasar radio sources (B1257-326, B1519-273, and J1819+385) that show large amplitude intraday and annual scintillation variability produced by the Earth's motion relative to turbulent-scattering screens located within a few parsecs of the Sun.  We find that the lines of sight to these sources pass through the edges of partially ionized warm interstellar clouds where two or more clouds may interact. From the gas flow vectors of these clouds, we find that the relative radial and transverse velocities of these clouds are large and could generate the turbulence that is responsible for the observed scintillation.  For all three sight lines the flow velocities of nearby warm local interstellar clouds are consistent with the fits to the transverse flows of the radio scintillation signals. ",
        "introduction": "Interstellar scintillation (ISS) is the apparent variation in flux density of very compact radio sources due to propagation through the irregular refractive index of the ionized interstellar medium.  The phenomenon was discovered and investigated through its 100\\% variations in pulsar signals seen on time scales from minutes to days. Most extragalactic radio sources have angular diameters much too large to scintillate, in the same way that the angular size of planets suppresses their atmospheric twinkling.  However, the most compact parts of some active galactic nuclei do fluctuate at a level of 1--10\\% over times shorter than a few days in a phenomenon called intraday variation (IDV) \\citep[see,][]{heeschen87,quirrenbach92}. Although this was initially thought to be intrinsic variation of the sources interpreted as due to emission from regions with brightness temperatures very much higher than the inverse Compton limit, it was subsequently explained as ISS somewhere within the 1~kpc thick layer of the ``warm ionized medium'' (WIM) of our Galaxy, and implied maximum source brightness temperatures near the inverse Compton limit \\citep[e.g.,][]{rickett07}.   However, among intraday variables, a few sources stand out as having unusually fast high-amplitude variations at frequencies of 3--8~GHz. With time scales less than a few hours, they have been called ``intrahour variables'' (IHV), the best studied of which are B1257--326, B1519--273, and J1819+385.  Their short time scales can only be explained by scintillation in a local region whose electron density has much greater root-mean-squared variation than in the general WIM.  The distances to the scattering are estimated to be in the range 1--30~pc.  However, until now no specific site in the local ISM has been identified as responsible.  A common feature of the IHV observations is that there is a systematic annual variation in their time scales due to a changing velocity of the Earth relative to the interstellar-scattering medium.  Indeed several observers have estimated the characteristic transverse velocities of the medium relative to the Sun needed to fit the ISS observations \\citep{Bignall2006,Dennett2003,Jauncey2003}.   It is also interesting that only a very few sources varied on hourly time scales in the 5~GHz VLA survey of 525 compact flat-spectrum AGNs, even though about half of them varied at a low level on times of several days \\citep{Lovell2007}.  Thus the regions responsible for IHV cover a very small fraction of the sky.   The basic goal of this paper is to compare the velocities from the rapid ISS sources with the velocities of the warm local interstellar clouds identified by \\citet{Redfield2007} and so find the regions of the local ISM responsible for the IHV. We find that the turbulence responsible for the scintillation may be generated by cloud-cloud interactions that could also enhance the ionization at cloud interfaces. ",
        "conclusions": "We have fit the ISS time scale observations of the three quasars that show the largest amplitude intraday variations and have high quality data sets.  For each quasar, our search of a large range of possible transverse velocities of the scintillation screen yielded good agreement with the annular variation of the ISS data for a relatively narrow range of transverse velocities. For each quasar there are at least two warm interstellar clouds that have transverse velocities consistent with the acceptable range (reduced $\\chi^2 \\leq 1.2$), but we propose that the scattering screen is located in one of the clouds near its edge. The transverse velocity of the local standard of rest, the usual assumption in previous studies, is an unacceptable solution for all three quasars.  In all cases, the scintillation screen is asymmetrical. We compared the ellipse orientations $\\theta_A$ with the orientations of the nearest cloud edge to see whether the turbulence structure is elongated parallel to the cloud edge, perhaps by a magnetic field confining the outside of the cloud. We found no preferred alignments as the ellipse orientations differ typically by $25^{\\circ}$ to $50^{\\circ}$ from the orientation of the edge of the cloud that we believe is the location of the scattering screen. This apparent misalignment may indicate that the cloud edges have complex geometries that our coarse fitting to a modest number of lines of sight cannot reveal. Also \\citet{Rickett2007} and \\citet{Ramachandran2006} show that the turbulent regions producing the observed scintillation are likely very small on the basis of the arrival times of pulses from the millisecond pulsar B1937+21 and other arguments.   Previous authors have placed the locations of the scattering screens near to the Sun on the basis of the scintillation time scales and the time delays between the scintillation flux variations as seen by two widely separated radio telescopes. Our analysis confirms and places more stringent constraints on the distances to the scattering screens. For example, \\citet{Bignall2006} place the scattering screen toward B1257--326 within 10~pc of the Sun, and we argue that it is located within 6.7~pc. \\citet{Dennett2003} argue that the scattering screen for J1819+385 is located between 1 and 12~pc from the Sun, and \\citet{Macquart2007} present two models that fit their data with the scattering screen located at $3.8\\pm 0.3$~pc or $2.0\\pm 0.3$~pc. We argue that the screen toward this source is located at the edge of the Mic cloud within 5.1~pc of the Sun. If the Mic cloud is indeed compressed and heated by the LIC and G clouds, then the scattering screen in the Mic cloud is located much closer to the Sun than 5.1~pc. The scattering screen toward B1519--273 near the edge of the G cloud is likely closer than 3~pc from the Sun.   An important result is that for each source, the line of sight to the quasar passes through or very near to at least two warm clouds. Although we know only upper limits to the distances to most of these clouds, they are all nearby and could interact. For example, interaction of the LIC and G clouds is likely responsible for the shape and high temperature of the Mic cloud. Since typical velocity differences between clouds along the three lines of sight are 10--20~\\kms, the interactions could easily produce turbulence that would be responsible for the scintillation. Also, cloud edges are regions of high ionization because of the absence of shielding from the strong UV radiation from hot stars and white dwarfs. Figure~\\ref{fig:fig2} shows areas on the sky where the velocity differences between clouds along the line of sight are large and where cloud-cloud interactions could produce turbulence. We predict that these are regions where large amplitude IDV sources will be discovered. Two examples that deserve further study are the two IDV sources located in the southern hemisphere (PKS 0405--385 and PSR J0437--4715) that are plotted in Figure~\\ref{fig:fig2} but are not analyzed in this paper as high quality ISS data are not yet published."
    },
    "0711/0711.4885_arXiv.txt": {
        "abstract": "A Fourier analysis on galaxy number counts from redshift data of both the Sloan Digital Sky Survey and the 2dF Galaxy Redshift Survey indicates that galaxies have preferred periodic redshift spacings of $\\Delta z =$ 0.0102, 0.0246, and 0.0448 in the SDSS and strong agreement with the results from the 2dF GRS. The redshift spacings are confirmed by the mass density fluctuations, the power spectrum $P(z)$ and $N_{pairs}$ calculations. Application of the Hubble law results in galaxies preferentially located on co-moving concentric shells with periodic spacings. The combined results from both surveys indicate regular co-moving radial distance spacings of $31.7 \\pm 1.8$ $h^{-1} Mpc$, $73.4 \\pm 5.8$ $h^{-1} Mpc$ and $127 \\pm 21$ $h^{-1} Mpc$. The results are consistent with oscillations in the expansion rate of the universe over past epochs. ",
        "introduction": "Introduction} When modeling the large scale structure of the cosmos the cosmological principle is assumed, therefore what we see must be biased by our viewpoint. The assumption is that the universe has expanded over time and any observer at any place at the same epoch would see essentially the same picture of the large scale distribution of galaxies in the universe.  However evidence has emerged from both the 2dF Galaxy Redshift Survey \\citep{2dFGRS} (see fig. 17 of Colless et al \\citep{Colless}) and also the Sloan Digital Sky Survey \\citep{SDSS} (see fig. 2 of Tegmark et al \\citep{Tegmark2}) that seem to indicate that there is periodicity in the abundances of measured galaxy redshifts. This has emerged from the galaxy number counts ($N$) within a small redshift interval ($\\delta z$) as a function of redshift ($z$).  What has been found is the so-called \\textit{picket-fence} structure of the $N$-$z$ relation, first noticed in one dimension by Broadhurst et al \\citep{Broadhurst} from a pencil-beam survey of field galaxies. In this paper the galaxy surveys are analysed on 2D and 3D scales. The usual interpretation is that this is evidence of the Cosmic Web, including voids and long filaments of galaxies. An alternative interpretation is suggested in this paper; galaxies are preferentially found in redshift space at certain redshifts with higher number densities than at other redshifts. Of course, this does not rule out cosmic web structures in addition to this effect.  However the concept that this represents real space galaxy distances is incompatible with the cosmological principle, which assumes the uniformity of space at all epochs on sufficiently large enough scales.  And it has been demonstrated \\citep{Yoshida}, from a large number of numerical simulations using the Einstein-de Sitter and $\\Lambda$CDM models, that the probability of getting such a periodic spatial structure from clustering and Cosmic Web filaments is less than $10^{-3}$. This allows for an alternate explanation, namely that the redshift periodicity is evidence for past oscillations in the expansion rate of the universe.  The only information that has been sought from large scale galaxy surveys is the spatial power spectrum that describes the random yet uniform distribution of galaxies in the universe, where the only departure expected from the random distribution is some characteristic scale for the clustering of galaxies. Certainly very few expect to find evidence consistent with a periodic redshift distribution.  However Hirano et al \\citep{Hirano} have attempted to explain galaxy redshift abundance periodicity as a real effect resulting from the universe having a non-minimally coupled scalar field that manifests, in later cosmological times. The effect produced oscillations in the expansion rate as a function of time. They model not only the redshift space observations but also the effect their cosmology has on the CMB power spectrum and the type Ia supernova Hubble diagram.  In the usual analysis the spatial two-point or autocorrelation function is used to define the excess probability, compared to that expected for a random distribution, of finding a pair of galaxies at a given separation \\citep{Baugh}. As a result the power spectrum $P(k)$ is derived from the two-point correlation function \\citep{Peebles}.  The power spectrum is predicted by theories for the formation of large scale structure in the universe and compared with that measured, or more precisely calculated from the available data.  The power spectrum for the SDSS has been calculated using a set of basis functions defined in \\citep{Hamilton2}. Three power spectra (galaxy-galaxy, galaxy-velocity, and velocity-velocity) were determined. Local motions of galaxies only contribute a radial component of a galaxy's total redshift, hence only affect the radial component of the power spectrum. Because the angular power spectrum is unaffected the assumption is also made that galaxy-galaxy spectrum  is equal to the redshift-space power spectrum in the transverse direction. In this paper we analyze the redshift hence radial component of the power spectrum $P(z)$, calculated in redshift space not $k$-space, and find periodicity consistent with that observed in $N(z)$.  Also an $N_{pairs}$ analysis is carried out on both survey data sets with confirming results.  The ``Bull's-eye'' effect has been studied \\citep{Praton, Melott, Thomas} via $N$-body simulations to result from large-scale infall plus small-scale virial motion of galaxies. It is believed that these two effects can bias measurements of large scale real space galaxy distributions. One is the so-called ``fingers-of-God'' (FOG) effect, which results from random peculiar motions of galaxies at the same distance from the observer. The second acts on much larger scales and where overdensities of galaxies occur, like at the ``Great Wall'' for instance. See Fig. 1.  Galaxies tend to have local motions toward the center of such structures. These motions are not random but coherent and add or subtract to the observer's line-of-sight redshift determination.  These effects preferentially distort the map in redshift space toward the observer due to the velocities of galaxies within clusters, i.e. non-cosmological redshift contributions. However the latter can only enhance in redshift space existing weak real space structures, it cannot create concentric structures centered on the Galaxy. And the FOG effect tends to smooth out finer detailed or higher Fourier frequency structures in  redshift space and is not very relevant to this discussion.  In the maps, fig. 18 of Colless et al \\citep{Colless}) and fig. 4 of Tegmark et al \\citep{Tegmark2}, there appears to the eye to be concentric structure in the arrangement of galaxies.  See the left hand side of the Fig. 1, the SDSS data with declinations between $0^{\\circ}$ and $6^{\\circ}$ mapped in Cartesian coordinates. Could it be that we inhabit a region of the Universe that is a bubble, part of a background of many such bubbles? But because redshift only gives us one dimensional information the conclusion could be drawn that if the galaxies in the Universe are evenly distributed throughout then redshift periodicity is indicative of oscillations in the expansion at past epochs. All observers would see the same shell structure in redshift space regardless of their location.  In this analysis we take a heuristic approach and look for periodicity in the redshift spacing between galaxies, in 2D in the case of slices and in 3D in case of the whole SDSS data set. In practice the discrete Fourier Transform ($\\nu_s$) is calculated from the $N$-$z$ relations, which are histograms determined by binning (counting) the observed redshifts of the survey galaxies between $z-\\delta z/2$ and $z+\\delta z/2$, as a function of redshift $z$. Within the $i$th bin there are $\\Delta N_i$ galaxies. Discrete Fourier amplitudes $\\nu_s$ are generated for each integer $s>1$ from  \\begin{equation} \\label{FFT} \\nu_s =\\frac{1}{\\sqrt{n}}\\sum_{i=1}^n u_s e^{2\\pi j(i-1)(s-1)/n}, \\end{equation} where $u_s = \\Delta N_i$, $j=\\sqrt{-1}$, $n$ is the total number of $\\Delta N_i$ bins.  Because $\\nu_s$ are complex their absolute value is taken in the analysis. The Fourier frequencies are calculated from \\begin{equation} \\label{Freq} \\mathcal{F} =\\frac{n \\delta z}{s} =\\Delta z. \\end{equation} where $\\Delta z$ is the periodic redshift interval.  The results suggest that there are redshifts, on the scale size discussed here, which is not primarily the result of clustering or the FOG and infall effects resulting in distortion of the redshift space maps. Both the SDSS and 2dF GRS redshift data where the redshifts are known with 95\\% confidence or greater are used. Fourier analysis can find periodic structure even where it may not be so apparent to the eye.  Data for 427,513 galaxies from the SDSS Fifth Data Release (DR5) were obtained where the data are primarily sampled from within about -10 to 70 degrees of the celestial equator. Also data for 229,193 galaxies were obtained from the 2dFGRS where the data are confined to within 2 degrees near the celestial equator and balanced between the Northern and Southern hemispheres. In the SDSS case the data are not so well balanced, yet there are nearly twice as many usable data. ",
        "conclusions": "A lot of work has been done in the last 20 years trying to show that the spatial structure determined from galaxy survey data is in agreement with what is predicted by the standard model. The standard model is assumed to be correct and when evidence questioning that conclusion is found (Broadhurst et al. \\cite{Broadhurst} and evidence for concentric shells) a special effort was immediately made to show how it could still be explained in the standard model \\cite{Kaiser,Praton}. But Kaiser and Peacock also concluded that their results ``did not prove that unconventional pictures for large-scale structures were ruled out.''  The appearance of concentric rings, like the \\textit{picket fence}, can be explained if appropriate assumptions are made (e.g. the Bull's Eye effect), but this too does not mean that that explanation is necessarily correct. Furthermore, surely one cannot attribute peaks in the power spectrum to a windowing function if the structures producing these peaks are already visible in the raw data, before the Fourier transform is taken.  This paper uses a simple Fourier analysis of number counts of galaxy redshifts in the SDSS and 2df GRS galaxy samples to search for periodic redshift spacings. Several spacings are found and are also confirmed using mass density fluctuations and the redshift separations between galaxies, though the latter are less sensitive methods. This analysis is similar in many ways to recent analyses carried out by other investigators, such as Tegmark et al. \\cite{Tegmark2} where the power spectra obtained from the SDSS and 2dF GRS samples are also found to have similar power peaks.  Acoustic oscillations in the plasma during the pre-recombination epoch are expected to have been imprinted on the distribution of matter in the Universe and to have survived until the present. In standard model the size, distribution, and clustering of galaxies were set in the very early Universe. In addition to the windowing function, non-linear evolution of structure and scale dependent bias, also play a role in our correct understanding of this. In this paper we have introduced another possible uncertainty into these studies that has not previously been considered. We suggest that the power peaks we find are also consistent with oscillations in the expansion rate at recent epochs. As unlikely as this may seem, it is not possible to rule it out completely since there is visible evidence in the raw data for an apparent concentric shell structure centered on the observer. It would be easier to rule it out if one could argue that such oscillations in the Hubble flow are impossible. However, astronomers are now arguing for a speed-up in the Hubble flow at the present epoch. Also a model was recently published (Hirano et al. \\cite{Hirano}) describing the origin of such oscillations.  One of the arguments against the power peaks being related to oscillations in the Hubble flow is the fact that the $P(k)$ distribution is a reasonable fit to what is predicted for the standard model. However, this does not rule oscillations out also. If there were no density fluctuations (no clustering  etc.), and the Universe was made up of a completely uniform distribution of galaxies, oscillations in the raw Hubble flow data would likely result in the appearance of concentric shells of apparent low and high density surrounding the observer because of the light travel time. To detect this there may be no need to use a windowing function. However, even if oscillations are present in the low-redshift Universe, any accurate model would also have to account for the density fluctuations (due to clusters, super clusters, etc.) that are also clearly visible.  And some have argued that the peaks found in this paper are just the scale sizes of clustering \\citep{Kaiser, Mo, Gonzalez}; this was considered improbable by \\citep{Yoshida} but a similar analysis needs to be performed on a larger data set in order to test this. Nevertheless the methods applied to the data here are consistent with the Universe undergoing oscillations in its expansion rate over past epochs. This means in redshift space there are preferred redshifts where galaxies are more dense due to a slower expansion rate at those epochs."
    },
    "0711/0711.4543_arXiv.txt": {
        "abstract": "It is argued that apart from the well known dependence of the Am phenomenon on the mass, age (effective temperature, gravity) and rotation there is also a complex dependence on the orbital parameters in binary systems. This is why the generally accepted scenario in which the Am star rotation plays a unique role needs to be revisited, the strong correlations between the rotation, orbital period and eccentricity need to be properly addressed and tidal effects taken into account. Recent observations of Am stars in binary systems are reviewed. ",
        "introduction": "Chemically peculiar (CP) stars is a well know group of anomalous stars on the upper main sequence (see e.g. Romanyuk 2007, Wahlgren 2004 for the most recent reviews). It is generally accepted that the slow rotation of these stars is the primordial cause of the Am and Ap peculiarity. Slow rotation translates into the low rotationally induced mixing and stable atmospheres void of disturbing turbulence. Then radiatively driven microscopic diffusion takes over and drives a vertical or horizontal stratification of chemical elements and observed abundance anomalies. At low effective temperatures CP phenomenon disappears due to the deep convection zones, at hotter temperatures it is because of the strong stellar winds. Large scale magnetic fields, if present, can strongly affect the diffusion, observed abundance anomalies, and differentiate between the magnetic Ap and nonmagnetic Am stars. Slow rotation is due to the magnetic breaking in single magnetic stars or tidal synchronization in binary stars.  With this picture in mind one can explain many of the fundamental observed properties of Am stars: \\begin{itemize} \\item The slow rotation with the cutoff of about 100 km/s; \\item abundance anomalies decreasing with the rotational velocity; \\item high fraction of close binaries in Am stars; \\item absence of Am binaries with very short orbital periods ($P_{orb}<1.2$ days, such systems rotate above the 100 km/s cutoff due to the spin-orbit synchronization and, thus, the abundance anomalies disappear). \\end{itemize} ",
        "conclusions": "A more complex and detailed study of Am stars in binary systems can shed more light on the fundamental questions of the origin and nature of the CP stars as well as on the interior hydrodynamics in binary systems."
    },
    "0711/0711.3822_arXiv.txt": {
        "abstract": "s{Einstein-\\ae ther theory is general relativity coupled to a dynamical unit timelike vector field. A brief review of current theoretical understanding and observational constraints on the four coupling parameters of the theory is given.} ",
        "introduction": "In general relativity (GR), spacetime structure is determined by a dynamical metric tensor field $g_{ab}$ and nothing else, and the theory is both diffeomorphism invariant and locally Lorentz invariant. Einstein-\\ae ther theory is the extension of GR that incorporates a dynamical unit timelike vector field $u^a$---the ``\\ae ther\"---which breaks the local Lorentz symmetry down to a 3d rotation subgroup. Direct coupling of matter to the \\ae ther would violate local Lorentz symmetry yet preserve diffeomorphism invariance. This paper presents a brief overview of the current theoretical and observational status of this theory, assuming that matter does not couple directly to the \\ae ther.   The action involving metric and \\ae ther is highly constrained. Besides the cosmological constant term, the only independent diffeomorphism invariant local  terms containing no more than two derivatives are \\beq S = -\\frac{1}{16\\pi G}\\int \\sqrt{-g}~ (R+K^{ab}_{mn} \\nabla_a u^m \\nabla_b u^n)~d^{4}x, \\label{action} \\eeq where $R$ is the Ricci scalar, $K^{ab}_{mn}$ is defined as \\beq K^{ab}_{mn} = c_1 g^{ab}g_{mn}+c_2\\d_m^a\\d_n^b +c_3\\d_n^a\\d_m^b +c_4u^au^bg_{mn} \\eeq with dimensionless coupling constants $c_i$, and the unit timelike constraint on the \\ae ther is implicit. (The metric signature is $({+}{-}{-}{-})$ and the speed of light defined by the metric $g_{ab}$ is unity.) Higher derivatives would be suppressed by powers of a (presumably) small length, e.g. the Planck length. It is assumed here that the \\ae ther is aligned at large scales with the rest frame of the microwave background radiation.  Einstein-\\ae ther theory---``\\ae -theory\" for short---is similar to the vector-tensor gravity theories studied by Will and Nordvedt,\\cite{willnord} but with the crucial difference that the vector field is constrained to have unit norm. This constraint eliminates a wrong-sign kinetic term for the length-stretching mode,\\cite{Elliott:2005va} hence gives the theory a chance to be viable. An equivalent theory using the tetrad formalism was first studied by Gasperini,\\cite{Gasperini} and in the above form it was introduced by Jacobson and Mattingly.\\cite{Jacobson:2000xp} ",
        "conclusions": "Einstein-\\ae ther theory is an intriguing theoretical laboratory in which gravitational effects of possible Lorentz violation can be meaningfully studied. There is a large (order unity) two-parameter space of Einstein-\\ae ther theories for which (i) the PPN parameters are identical to those of GR, (ii) the linear perturbations are stable and carry positive energy, (iii) there is no vacuum \\v{C}erenkov radiation, (iv) the dynamics of the cosmological scale factor and perturbations differ little from GR, (v) non-rotating neutron star and black hole solutions are close to those of GR, but might be distinguishable with future observations. Radiation damping from binaries, imposes an order $0.001$ constraint on one combination of the parameters. Strong self-field effects in neutron stars and black holes produce violations of the strong equivalence principle and higher order post-Newtonian effects which will constrain all the parameters $c_i$ to be less than around $0.01$, presuming that the sensitivity parameters for neutron stars (which have not yet been computed with the required precision) turn out to have the expected magnitude."
    },
    "0711/0711.0683_arXiv.txt": {
        "abstract": "{At the {\\sl New Horizons in Globular Cluster Astronomy} conference (Padova, June 2002), two members of the {\\sl VLT globular cluster team\\/} presented different views on the importance of heavy-element sedimentation in Population II stars: ``The lack of evidence for depletion of Fe and Li in the atmospheres of globular cluster subgiants led {\\em some people} to suspect that, for unknown reasons, Population II stars are not affected by this mechanism.'' \\citep{Castellani_2003} and ``There should be {\\em some mechanism} that prevents sedimentation.'' \\citep{Gratton_2003}.  In this review, I will argue that the scepticism behind both these statements is justified. We recently revisited the results on sedimentation in NGC 6397 stars presented by {\\citet{Gratton_etal_2001}} using higher-quality VLT/FLAMES-UVES data \\citep{Korn_etal_2006,Korn_etal_2007}. Element-specific abundance trends were identified which agree with atomic-diffusion predictions, {\\em if\\/} turbulent mixing below the convective envelope is accounted for in a parametrized way. Have we thus detected signatures of Atomic Diffusion in Old Stars (ADiOS)? Or are these trends mere artefacts of Conspiring Inaccuracies in Abundance Observations (CIAO)? ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.2686_arXiv.txt": {
        "abstract": "The cosmological constant $\\Lambda$, i.e., the energy density stored in the true vacuum state of all existing fields in the Universe, is the simplest and the most natural possibility to describe the current cosmic acceleration. However, despite its observational successes, such a possibility exacerbates the well known $\\Lambda$ problem, requiring a natural explanation for its small, but nonzero, value. In this paper we study cosmological consequences of a scenario driven by a varying cosmological term, in which the vacuum energy density decays linearly with the Hubble parameter, $\\Lambda \\propto H$. We test the viability of this scenario and study a possible way to distinguish it from the current standard cosmological model by using recent observations of type Ia supernova (Supernova Legacy Survey Collaboration), measurements of the baryonic acoustic oscillation from the Sloan Digital Sky Survey and the position of the first peak of the cosmic microwave background angular spectrum from the three-year Wilkinson Microwave Anisotropy Probe. ",
        "introduction": "The nature of the mechanism behind the current cosmic acceleration constitutes a major problem nowadays in Cosmology \\cite{rev}. Even though almost all observational data available so far are in good agreement with the simplest possibility, i.e., a vacuum energy plus cold dark matter ($\\Lambda$CDM) scenario, it is becoming rather consensual that in order to better understand the nature of the dark components of matter and energy, one must also consider more complex models as, for instance, scenarios with interaction in the dark sector \\cite{cq}.   In this regard, the simplest examples of interacting dark matter/dark energy  models are scenarios with vacuum decay [$\\Lambda$(t)CDM]. In reality,  $\\Lambda$(t)CDM cosmologies constitute the special case in  which the ratio of the dark energy pressure to its energy density, $\\omega$, is exactly $-1$ \\cite{Lambda(t),lambdat}. This kind of model may be based on the idea that dark energy is the manifestation of vacuum quantum fluctuations in the curved space-time, after a renormalization in which the divergent vacuum contribution in the flat space-time is subtracted. The resulting effective vacuum energy density will depend on the space-time curvature, decaying from high initial values to smaller ones as the universe expands \\cite{1}. As a result of conservation of total energy, implied by Bianchi identities, the variation of vacuum density leads either to particle production or to an increasing in the mass of dark matter particles, two general features of decaying vacuum or, more generally, of interacting dark energy models \\cite{jsa}.  Naturally, the precise law of vacuum density variation depends on a suitable  derivation of the vacuum contribution in the curved background, which is in general a difficult task. In this regard, however, a viable possibility is initially to consider a de Sitter space-time and estimate the renormalized vacuum contribution with help of thermodynamic reasonings as, e.g., those in line with the holographic conjecture \\cite{Holography}. The resulting ansatz is $\\Lambda \\approx (H + m)^4 - m^4$, where $H$ is the Hubble parameter and $m$ is a cutoff imposed to regularize the vacuum contribution in flat space-time (the next step is to consider a quasi-de Sitter background, with a slowly decreasing $H$. In this case, the above ansatz may be considered a good approximation, but with the vacuum density decaying with time.). In the early-time limit, with $H >> m$, we have $\\Lambda \\approx H^4$. By using this scaling law for the vacuum density and introducing a relativistic matter component, some of us have obtained from the Einstein equations an interesting solution with the following features \\cite{2}. Firstly, the Universe undergoes an empty, quasi-de Sitter phase, with $H \\approx 1$, which in the asymptotic limit of infinite past tends to de Sitter solution with $H = 1$ (in Planck units). However, at a given time,  the vacuum undertakes a fast phase transition, with $\\Lambda$ decreasing to nearly zero in a few Planck times, producing a considerable amount of radiation. In this sense, this can be understood as a non-singular inflationary scenario, with a semi-eternal quasi-de Sitter phase originating a radiation dominated universe.  The subsequent evolution of vacuum energy essentially depends on the masses of the produced particles.  If we apply the energy-time uncertainty relation to  the process of matter production, we will conclude that massive particles can be produced only at very early times, when $H$ is very high. Therefore, baryonic particles and massive dark matter (as supersymmetric particles and axions) stopped being produced before the time of electroweak phase transition.  On the other hand, the late-time production of photons and massless neutrinos must be forbidden by some selection rule, otherwise the Universe would be completely different from the one observed today. Thus, if no other particle is produced, the vacuum density stops decaying at very early times, and for late times we have a standard universe, with the presence of a genuine cosmological constant (some thermodynamic considerations, in line with the holographic principle, permit to infer the value of this constant, leading to $\\Lambda \\approx m^6$ \\cite{1,2}).  In order to have a decaying vacuum density at the present time, the produced dark particles should have a mass as small as $10^{-65}$g\\footnote{Some authors associate this mass to the graviton in a de Sitter background with $H$ and $\\Lambda$ of the orders currently observed \\cite{Graviton}.}. Here, we have considered the possibility of a late-time decaying $\\Lambda$, and compared the consequent cosmological scenario with the constraints imposed by current observations \\cite{3,4}. In such a limit,  $H << m$ and $\\Lambda \\approx m^3 H$, leading again to $\\Lambda \\approx m^6$ in the de Sitter limit. From a qualitative point of view, we have found no important difference between this $\\Lambda(t)$ scenario and the flat $\\Lambda$CDM model \\cite{3}. After the phase transition described above the Universe enters a radiation-dominated phase, followed by a matter-dominated era long enough for structure formation, which tends asymptotically to a de Sitter universe with vacuum dominating again. The only important novelty, related to matter production, is a late-time suppression of the density contrast of matter, which may constitute a potential solution to the cosmic coincidence problem \\cite{5}. Moreover, the analysis of the redshift-distance relation for supernovae of type Ia, particularly with the Supernova Legacy Survey (SNLS) data set, has shown a good fit, with present values of $H$ and the relative density of matter in accordance with other observations \\cite{4}.  In this paper, we go a little further in our investigation and study new observational consequences of the $\\Lambda$(t)CDM scenario described above. We use to this end distance measurements from type Ia supernovae (SNe Ia), measurements of the  baryonic acoustic oscillations (BAO) and the position of the first peak of the cosmic microwave background (CMB). We show that, besides the interesting cosmic history of this class of $\\Lambda$(t)CDM models, a conventional, spatially flat $\\Lambda$CDM model is only slightly favored over them by the current observational data. ",
        "conclusions": ""
    },
    "0711/0711.2365_arXiv.txt": {
        "abstract": "{ The $^{12}$C/$^{13}$C isotope ratio in the interstellar medium (ISM), and its evolution with time, is an important tracer of stellar yields. Spatial variations of this ratio can be used to study mixing in the ISM\\@.}  {We want to determine this ratio and its spatial variations in the local ISM from CH$^+$ absorption lines in the optical towards early-type stars. The aim is to determine the average value for the local ISM and study possible spatial variations. } {We observed a large number of early-type stars with Feros to extend the sample of suitable target stars for CH$^+$ isotope studies.  The best suited targets were observed with Uves with higher signal-to-noise ratio and spectral resolution to determine the isotope ratio from the interstellar CH$^+$ lines. This study significantly expands the number of $^{13}$CH$^+$ detections. }  {We find an average ratio of $\\langle R \\rangle = 76.27 \\pm 1.94$ or, for $f = 1/R$, $\\langle f \\rangle = (120.46 \\pm 3.02) \\times 10^{-4}$. The scatter in $f$ is $6.3\\times\\sigma(\\langle f \\rangle)$. This findings strengthens the case for chemical inhomogeneity in the local ISM, with important implications for the mixing in the ISM\\@. Given the large scatter, the present-day value in the ISM is not significantly larger than the solar value, which corresponds to the local value 4.5\\,Gyr ago.} {} ",
        "introduction": "The $^{12}$C/$^{13}$C ratio is a cornerstone of models of the nuclear history of our galactic interstellar matter (ISM), that is, the Galactic Chemical Evolution.  Low and intermediate mass stars increase (during the AGB phase) the $^{12}$C/$^{13}$C ratio to $\\sim300$, while stars more massive stars (undergoing CNO processing) increase the ratio to $^{12}$C/$^{13}$C$\\sim$3. A measurement of the $^{12}$C/$^{13}$C ratio in the ISM thus provides important information about the relative importance of stars of different mass to the chemical enrichment of the interstellar matter.  The $^{12}$C/$^{13}$C ratio in the solar system is 89 \\citep{1989GeCoA..53..197A}. Models of chemical evolution predict a decrease in the $^{12}$C/$^{13}$C ratio with time for a given galacto-centric distance, and a decrease with galacto-centric distance in the Galaxy \\citep{1996A&A...309..760P}.  Such models assume that mixing in the ISM is complete and restricted to material at a given galacto-centric distance (azimuthal mixing). A comparison of results between our Galaxy and other galaxies will give important data for the nuclear processing history \\citep[see, e.g.,][]{2000emw..conf..505T,2001coev.conf..233P}  The $^{12}$C/$^{13}$C ratio has been determined mostly from radio-astronomical measurements of CO and H$_2$CO \\citep[see e.g.,][]{1994ARA&A..32..191W} and millimeter observations of CN \\citep{2002ApJ...578..211S,2005ApJ...634.1126M}.  These results may be affected by interstellar chemistry in different ways, however: Chemical fractionation (which results from different zero-point energies of the isotopomers) enriches molecules in $^{13}$C \\citep{1976ApJ...205L.165W} and thus lowers the ratio, selective dissociation destroys the rarer species more and thus raises the ratio. The latter effect is due to self-shielding of the more abundant isotopemer, which therefore survives photo-destruction better than the less abundant counterpart.  The molecule CH$^{+}$ is thought to be produced only in hotter parts of photo-dissociation regions and is a rare species, so is not affected by either chemical fractionation or selective dissociation. In addition, the optical CH$^+$ lines have typically only small saturation effects. This molecule, which can be observed in optical absorption lines from the ground, is therefore considered to provide the most secure $^{12}$C/$^{13}$C ratios \\citep[see e.g.][]{1976ApJ...205L.165W}. We caution, however, that the formation process of CH$^+$ is probably not completely understood \\citep{1993A&A...269..477G}.  Two systems of CH$^{+}$ lines are strong enough to be easily observed from the ground. In addition, the absorption lines from $^{12}$CH$^{+}$ at 4232 and 3957 \\AA\\ are sufficiently shifted from their $^{13}$CH$^{+}$ component to allow unblended measurements, provided the velocity profiles in the intervening cloud are narrow enough. If both are measured simultaneously, we can make use of the fact that the $^{13}$CH$^{+}$ line is blue-shifted by 0.26 \\AA\\ relative to the $^{12}$CH$^{+}$ at 4232 \\AA\\ but red-shifted by 0.44 \\AA\\ relative to the $^{12}$CH$^{+}$ line at 3957 \\AA\\@. This way, a chance superposition with another cloud along the line-of-sight can be excluded. In addition, the two simultaneous measurements help to eliminate systematic effects.  This property has been used e.g.\\@ by \\citet{1992A&A...254..327S} to check that there is no accidental overlap of a weak velocity component of $^{12}$CH$^{+}$ at $-$18.8 km sec$^{-1}$ from the deepest $^{12}$CH$^{+}$ absorption line with the $^{13}$CH$^{+}$ line. Nearly all previous data were taken at 4232 \\AA\\ only.  These measurements in the optical require, however, observations with high spectral resolution and signal/noise ratio. Most of the published CH$^+$-measurements have therefore been obtained with very bright background stars at relatively small distances from the Sun.  The availability of efficient high-resolutions spectrograph at 8m-class telescopes in recent years has improved the situation considerably \\citep[][in the following paper\\,I]{2005A&A...441..181C}. The observations are hampered, however, by the scarcity of suitable published candidate stars with sufficiently strong and narrow CH$^+$ absorption lines.  In this paper we describe new observations which have been obtained at targets specifically selected for this purpose.  The outline of the paper is as follows: Section~\\ref{sec:obs} describes our observations. Section~\\ref{sec:model} gives details on our procedure to measure the isotopic ratio $R$. The results are discussed in section~\\ref{sec:results}, both the average properties of the sample and notes about individual targets. Section~\\ref{sec:conc} discusses the results and concludes the paper. ",
        "conclusions": "\\label{sec:conc}  Our systematic search for stars with clean line profiles in CH$^+$ and large equivalent widths, which are suitable for isotope studies, has resulted in a much increased sample of reliable $^{12}$C/$^{13}$C isotope ratios from optical observations.  The VLT-{\\sc Uves} spectra of CH$^+$ absorption in the lines of sight towards 14 stars have allowed us to measure the carbon isotopic ratio $R= ^{12}$C/$^{13}$C with unprecedented accuracy. The significant extension of the sample of \\citet{2005A&A...441..181C} allows us to strengthen their conclusion that there is truly a significant scatter in the local ISM\\@.  Averaging our measurements for the 14 lines of sight gives a value representative of the local ISM: $\\langle R \\rangle = 76.27 \\pm 1.94$.  This mean ratio is in excellent agreement with the local ratio determined from radio-astronomical measurements of CO and H$_2$CO \\citep[76$\\pm$7,][]{1994ARA&A..32..191W}, and also in good agreement with more recent determinations from CN \\citep[$\\approx$68,][]{2005ApJ...634.1126M}. This agreement of results obtained from different molecules suggests that chemical fractionation and selective dissociation do not play a major role for the CO, H$_2$CO and CN molecules. The results of \\citet{2007ApJS..168...58S} suggest, however, that CO may be affected by chemical fractionation  in cool and dense regions.  The weighted scatter of $f$ values is $6.3\\times\\sigma(\\langle f \\rangle)$ which directly shows that the $^{12}$C/$^{13}$C isotope ratio is variable on the relatively small spatial scales which we sample.  Considering this intrinsic scatter of the $^{12}$C/$^{13}$C ratio in the ISM, the solar carbon isotope ratio of $89$ is undistinguishable from the present-day value, even after 4.5\\,Gyr of galactic evolution. This statement is not affected by instrumental uncertainties, although a sampling bias could distort our ISM value. Of course, this conclusion is based on the assumption that the solar system ratio is representative for the galacto-centric distance R$_\\odot$ = 8.5 kpc, 4.5\\,Gyr ago. Since we have shown that the interstellar $^{12}$C/$^{13}$C isotope ratio is spatially variable now, it is not clear if the solar value is characteristic for the local $^{12}$C/$^{13}$C ratio, 4.5\\,Gyr ago. However, there is also no strong evidence that the solar value is untypical: Although it has been claimed that the sun is overabundant, especially in carbon and oxygen, compared to the local ISM, the solar carbon and oxygen abundances have been revised downwards.  Recent papers \\citep[e.g.][]{arXiv:astro-ph/0702429v1} find good agreement for the carbon and oxygen abundance in the sun, the local ISM and young stars. Comets show a $^{12}$C/$^{13}$C ratio consistent with the photospheric solar value \\citep[e.g.][]{2005A&A...432L...5M}. Therefore it seems also unlikely that the $^{12}$C/$^{13}$C ratio was modified from the value in the solar nebula by fractionation of carbon isotopes during the formation of the sun.  Therefore we conclude that there is no evidence that the solar $^{12}$C/$^{13}$C ratio is untypical for the solar neighborhood 4.5 Gyr ago. It is plausible however, that the average value for the solar neighborhood was somewhat different from the solar value by a similar amount as the present-day variations. A small decrease of the average $^{12}$C/$^{13}$C ratio in the solar neighborhood during the lifetime of the sun is therefore certainly possible.  The $^{12}$C/$^{13}$C ratio increases with galacto-centric distance \\citep{1994ARA&A..32..191W} by about 7.5$\\pm$1.9 per kpc. Therefore, a comparison of the present-day local $^{12}$C/$^{13}$C ratio with the solar photospheric value should also take into account that the sun was born at a different galacto-centric distance \\citep{1997A&A...326..139W}. Since the sun has, according to \\citet{1996A&A...314..438W}, migrated from its birth-place in the inner part of the Galaxy outwards by $1.9\\pm0.9$ kpc during its lifetime, this effect decreases the expected change in the $^{12}$C/$^{13}$C ratio significantly.  Early models for the $^{12}$C/$^{13}$C isotope ratio were able to model the significant decrease of the local ratio over the lifetime of the sun (about a factor two), which was thought at that time to be present. However, the models require some fine tuning and calibration, so that a smaller decrease is not in obvious conflict with the models. Nevertheless, some variation of the $^{12}$C/$^{13}$C isotope ratio is to be expected. New models of the chemical evolution of the Galaxy are required to take into account the new observational results.  A larger sample of stars will allow a better estimate of the average ISM value. In addition, these observations will allow to determine the scale of the variations of the isotope ratio and thus the efficiency of mixing on different scales."
    },
    "0711/0711.2962_arXiv.txt": {
        "abstract": "We have observed a  typical Galactic  plane field at $(l,b) = (\\timeform{28D46}, \\timeform{-0D20})$ with Suzaku\\/ for 100 ksec to carry out a precise spectral study of the Galactic Ridge X-ray Emission (GRXE). The field is known to be devoid of X-ray point sources brighter than $\\sim2 \\times  10^{-13}$ergs s$^{-1}$ cm$^{-2}$ (2--10 keV), and already deeply observed with Chandra. Thanks to the low and stable background and high spectral resolution of Suzaku, we were able to resolve, for the first time, three {\\em narrow}\\/ iron K-emission lines from low-ionized (6.41 keV), helium-like (6.67 keV), and hydrogenic ions (7.00 keV) in the GRXE spectrum. These line features constrain the GRXE emission mechanisms: The cosmic-ray ion charge exchange model or the non-equilibrium ionization  plasma model are unlikely, since they require either broad emission lines or lines at intermediate  ionization states. Collisional ionization equilibrium plasma is the likely origin for the 6.67 keV and 7.00 keV lines, while origin of the 6.41 keV line, which is due to fluorescence from cold material, is not elucidated. Low non-X-ray background and little stray-light contamination of Suzaku\\/ allowed us to  precisely measure the absolute X-ray surface brightness in the direction of the Galactic plane. Excluding the point sources brighter than $\\sim$2 $\\times 10^{-13}$ ergs s$^{-1}$ cm$^{-2}$ (2--10 keV), the total surface brightness  on the Galactic plane is $\\sim$6.1 $\\times 10^{-11}$ ergs s$^{-1}$ cm$^{-2}$ deg$^{-2}$ (2--10 keV), including the contribution of the cosmic X-ray background that is estimated to be $\\sim$1.3$\\times 10^{-11}$ ergs s$^{-1}$ cm$^{-2}$ deg$^{-2}$. ",
        "introduction": "The Galactic plane has been known to be a source of hard (2--10 keV) X-ray emission since the 1980s (Worrall et al.\\ 1982;  Warwick et al.\\ 1985; Koyama et al.\\ 1986). Although most of the X-ray satellites have spent significant amounts of observing time to study the Galactic Ridge X-ray Emission (GRXE), its origin has not been elucidated yet. Whether the GRXE is truly diffuse emission or composed of numerous dim point sources remains  inconclusive.  Ebisawa et al.\\ (2001, 2005) studied a blank Galactic plane area at $(l,b)\\approx(\\timeform{28.5D},\\timeform{0.0D})$ with Chandra\\/, consisting of two partially overlapping fields each exposed 100 ksec (sequence 900021 and 900125). They found that only 10 -- 15 \\%  of the observed 2--10 keV GRXE was resolved into point sources brighter than $\\sim3 \\times 10^{-15}$ ergs s$^{-1}$ cm$^{-2}$. They also found that the point source surface density in the 2 -- 10 keV band was not significantly higher than that at high Galactic latitudes, and therefore the majority of these hard X-ray sources is likely to be the background AGNs. On the other hand, numerous point sources were detected in the soft ($<2$ keV) band which are considered to be nearby Galactic sources, hardly contributing to the 2 -- 10 keV GRXE.  Ebisawa et al.\\ (2001, 2005) suggest that the number of Galactic hard X-ray sources starts to deplete ($\\log N-\\log S$ curve begins to saturate) around $\\sim 3 \\times 10^{-15}$ ergs s$^{-1}$ cm$^{-2}$ (2 -- 10 keV), and they conclude that most of the GRXE has the diffuse origin.  Meanwhile, Revnivtsev et al.\\ (2006) discovered a striking similarity between the global GRXE distribution and that of the near infrared (NIR) emission at 3.5 $\\mu$m. Since the NIR emission is considered to represent the stellar population, this spatial correlation is suggestive of the stellar origin of the GRXE. Assuming the low-luminosity X-ray source population based on the all-sky RXTE and ROSAT data of the sources in the solar neighborhood (Sazonov et al.\\ 2006), Revnivtsev et al.\\ (2006) proposed that the GRXE is ultimately resolved into point sources at $\\sim10^{-16}$--$10^{-16.5}$ ergs s$^{-1}$ cm$^{-2}$ (2 -- 10 keV). Revnivtsev and Sazonov\\ (2007) claim that, using the same Chandra\\/ data as those of Ebisawa et al.\\ (2005), $\\sim25 \\%$ of the GRXE in 1-- 7 keV is accounted for by point sources brighter than $1.2 \\times10^{-15}$ ergs s$^{-1}$ cm$^{-2}$ (interstellar absorption corrected flux), and that there is  no indication of saturation of the $\\log N- \\log S$ slope down to this flux limit. In order to test the point source origin, one would need more than an order of magnitude lower limiting flux than the current value. Currently only Chandra can make the required observations, which would have much deeper ($>$ 1 Msec) exposures than those already carried out.  Another important property of the GRXE is its energy spectrum. It has been known from previous observations that the spectrum is hard (e.g, Valinia et al.\\ 2000) and characterized by many intense emission lines from highly ionized ions of the abundant elements (e.g., Kaneda et al.\\ 1997). These emission lines, in particular the iron K-emission lines in 6--7 keV, can constrain the emission mechanism. Suzaku\\/ satellite is the best suited for this study, since it is equipped with the X-ray telescopes of large effective areas, and the focal plane detectors of superior spectral resolution. In addition, Suzaku\\/'s low and stable background and small stray-light contamination enable us to measure the total surface brightness on the Galactic plane accurately, from which we can estimate the absolute GRXE flux by subtracting the contribution of the cosmic X-ray background. In this paper, we report the results of a 100 ksec observation of the GRXE with Suzaku\\/ for the same Galactic plane field as that of the Chandra\\/ observation by Ebisawa et al (2001, 2005). ",
        "conclusions": "\\subsection{Emission Mechanism of the GRXE}  We have observed a typical Galactic plane region at $(l,b)\\approx(\\timeform{28.46D},\\timeform{-0.20D})$ (= the Chandra\\/ AO1 deep field) with Suzaku\\/. Thanks to the superior spectral resolution, the large effective area and the low background of Suzaku\\/ XIS, we obtained so far the best quality spectrum of the GRXE. For the first time, the iron K-line complex was clearly resolved into three individual peaks at 6.4 keV, 6.67 keV and 7.0 keV. Importantly, the line widths are also tightly constrained to be less than $\\sim$50 eV (1 $\\sigma$).  The central energy of the He-like K$\\alpha$ line, which   is in fact a blend of the resonance, inter-combination and forbidden lines,  is a good indicator of the plasma condition (see Koyama et al. 2007b). In the case of collisional ionization equilibrium (CIE), the centroid is 6.680--6.685 keV, depending on plasma densities and different plasma models (Koyama et al. 2007b).  The peak energy  is shifted downward if the  plasma is in a  non-equilibrium ionization  (NEI) state. Using ASCA SIS, Kaneda et al.\\ (1997) claimed to have detected a {\\em single}\\/ narrow emission line at $6.61\\pm0.04$ keV, that is significantly lower than what is expected in the CIE case, and suggested that the iron line is from a NEI plasma. Chandra observation by Ebisawa et al.\\ (2005)  also indicated a similar line center energy $6.52\\pm^{0.08}_{0.14}$ keV, if the line feature is fitted with a {\\em single}\\/ Gaussian. Meanwhile, thanks to Suzaku's superior sensitivity and energy resolution, we unambiguously detected the two separate lines at 6.41 keV and 6.670 keV, the flux of the former being one-third of that of the latter. The centroid of the two  lines will be at 6.58 keV, and the 7.0 keV line is too weak to be  detected with ASCA and Chandra. Hence, we consider that the lower He-like K$\\alpha$  line energies previously reported by  ASCA and Chandra were significantly affected by the presence of the 6.41 keV line,  which was unresolved by these instruments.  Koyama et al.\\ (2007b) determined the He-like K$\\alpha$ line energy from the Galactic Center (GC) plasma as $6.680\\pm0.001$ keV, which is consistent with the CIE plasma model. We obtained a very close value of  $6.670\\pm0.006$ keV, which does not show evidence of  the NEI conditions. Thanks to the brightness of the GC plasma illuminating  the entire CCD chip, Koyama et al.\\ (2007b) were able to perform the chip-segment dependent Charge Transfer Inefficiency (CTI) correction, which effectively increased the iron line energies by several electron-volts. We cannot apply the same correction,  since the Ridge emission is not so strong; however,  if we consider the possible upward shift of the line energies due to this effect, our line energy completely agrees with that from the GC and the CIE plasma. The observed  H-like/He-like line intensity ratio ($\\sim 0.2$) is also consistent with that expected for a CIE plasma of $kT$ = 5 -- 8 keV.  Earlier, Tanaka et al. (2000) reported that, analyzing the ASCA\\/ data of the extended emission in the  Galactic center region, the iron K-lines are significantly broadened with a width of $\\sim 70$ eV. They suggested a possibility of the charge exchange process for the origin of the lines, where low-energy cosmic-ray iron ions will undergo charge exchange as they slow down in the interstellar space. They interpreted the finite line width as due to the bulk motion of the cosmic-ray ions. Tanaka (2002) pointed out the close similarity of the shape of the GRXE spectrum (at the same direction as the present one) with that around the Galactic center, and considered that the iron lines in the GRXE may as well be broadened, while this could not be  examined with ASCA due to insufficient statistics.  The He-like K$\\alpha$ line energy expected from the charge exchange process is $6.666\\pm0.005$ keV (Wargelin et al.\\ 2005), which is not clearly excluded from the present result $6.670\\pm0.006$ keV unless the segment-dependent CTI correction is applied. However, our result supports that the He-like and H-like iron lines from the Galactic ridge are essentially narrow lines, which is against the line-broadening due to the cosmic-ray bulk motion and charge exchange. The Suzaku\\/ observation of the GC region also clearly refuted the broadening of the iron lines, as well as that the He-like K$\\alpha$ line energy does not agree with that expected from the charge exchange (Koyama et al. 2007b).  % Therefore, we conclude that the charge exchange origin is quite unlikely for GRXE,  as well as GC plasma.\\footnote{The reason why Tanaka et al. (2000) obtained the line width of $\\sim70$ eV is still unclear. As written in their paper, the observed spectra from eight CCD chips were stacked with reference to the 6.7-keV line peak position, such that the line width after summation was minimized. One possibility is that the tuning of the energy scales among each other was not sufficient.}    Thus, a CIE plasma is the most probable origin of the He-like and H-like iron K$\\alpha$ lines. By the same token, the lower-energy lines from highly ionized ions of other abundant elements are considered to be thermal emission. However, the thermal structure of the plasma is not simple. As shown in Table 2, the H-like/He-like line intensity ratio is similar among different elements from Mg through Ar, and all roughly 0.3 -- 0.5, and the same is true for Fe. This presumably requires a multi-temperature plasma.  There exist two alternative possibilities to explain the CIE plasma emission from the Galactic Ridge. The fundamental difference is whether such plasma is genuinely diffuse or gravitationally bound in the stellar sources. Consequently the ultimate questions are the following: (1) Does the superposition of all the point sources reproduce the observed GRXE spectrum, in particular, the complex emission line features? (2) Is the number of such Galactic point sources sufficient to explain the total GRXE  flux?  As regards the origin of the 6.4 keV line, which is not emitted  by the hot plasma but due to fluorescence from cold material, another explanation is required. Based on the diffuse model, Valinia et al.\\ (2000)  proposed that interaction between the interstellar medium  and cosmic supra/non-thermal electrons is responsible for the 6.4 keV line, as well as for the hard-tail observed above $\\sim10 $ keV (Yamasaki et al.\\ 1997; Valinia et al.\\ 2000). On the other hand, superposition of numerous, different kinds of point sources may create such complex iron line spectra (Revnivtsev et al.\\ 2006). In fact, quiescent cataclysmic variables are known to emit three iron lines, where the 6.4 keV line is due to reflection from the white-dwarf surfaces (e.g., Ezuka \\& Ishida 1999). Also, cataclysmic variables often reveal multi-temperature spectra with various emission lines as seen in GRXE.    \\subsection{Absolute flux of the GRXE}  Excluding the point sources brighter than $2 \\times 10^{-13}$ ergs s$^{-1}$cm$^{-2}$ (2-- 10 keV), we have measured the total flux in  the direction of the Galactic plane (including cosmic X-ray background) as  $6.1 \\times 10^{-11}$ ergs s$^{-1}$ cm$^{-2}$ deg$^{-2}$ (2--10 keV), which is consistent with the Chandra\\/ measurement on the same region within 20 \\%. Currently, only $\\sim$15 \\% of the 2 -- 10 keV of the total surface brightness is explained by the  point sources (either Galactic or extragalactic) brighter than  $\\sim3 \\times 10^{-15}$ ergs s$^{-1}$cm$^{-2}$  (Ebisawa et al.\\ 2005). In 1 -- 7 keV, where Galactic point sources are more dominant, $\\sim25$ \\% of the total flux is resolved into point sources with interstellar absorption corrected fluxes higher than $ 1.2 \\times 10^{-15}$ ergs s$^{-1}$cm$^{-2}$ (Revnivtsev \\& Sazonov 2007).  For comparison, we estimate  the cosmic X-ray background (CXB) flux seen through the Galactic plane. The 2-- 10 keV CXB flux determined from the high galactic region is $\\sim2 \\times 10^{-11}$ ergs s$^{-1}$ cm$^{-2}$ deg$^{-2}$ (Moretti et al.\\ 2003), where point sources dimmer than $8.0 \\times 10^{-12}$ ergs s$^{-1}$ cm$^{-2}$ are included. In order to compare with our results where  point sources brighter than $2.0 \\times 10^{-13}$ ergs s$^{-1}$ cm$^{-2}$ are excluded, we subtract the cumulative flux of the sources between $2.0 \\times 10^{-13}$ ergs s$^{-1}$ cm$^{-2}$ and $8.0 \\times 10^{-12}$ ergs s$^{-1}$ cm$^{-2}$ using the model extragalactic $\\log N - \\log S$ curve by Moretti et al.\\ (2003); thus we obtained the unabsorbed CXB flux $\\sim1.8 \\times 10^{-11}$ ergs s$^{-1}$ cm$^{-2}$ deg$^{-2}$.  The hydrogen column density through the Galactic plane is estimated to be $\\sim 6 \\times 10^{22}$  cm$^{-2}$ in this direction (Ebisawa et al.\\ 2001), which attenuates the background CXB down to $\\sim70$ \\% in 2 -- 10 keV, assuming a  photon-index of 1.7. Consequently,  $\\sim1.3 \\times 10^{-11}$ ergs s$^{-1}$ cm$^{-2}$ deg$^{-2}$ is expected from the background CXB on the Galactic plane, which is ultimately resolved into individual AGNs. Chandra\\/ already resolved point-sources down to $3 \\times 10^{-15}$ ergs s$^{-1}$cm$^{-2}$, whose cumulative flux is $\\sim9 \\times 10^{-12}$ ergs s$^{-1}$ cm$^{-2}$ deg$^{-2}$. Therefore, we see that a significant part of the dim hard Chandra\\/ sources is extragalactic.  In fact, only $\\sim$20 \\% of the Chandra\\/ hard X-ray sources had near-infrared counterparts down to $K\\approx16$ (hence considered to be Galactic), whereas almost all the soft X-ray sources had counterparts (Ebisawa et al.\\ 2005).   In conclusion, we measure the flux of the the Galactic Ridge X-ray emission as $\\sim4.8\\times 10^{-11}$ ergs s$^{-1}$ cm$^{-2}$ deg$^{-2}$ (2 -- 10 keV), excluding the Galactic and extragalactic sources brighter than $2 \\times 10^{-13}$ ergs s$^{-1}$cm$^{-2}$ and the estimated contribution of  dimmer extragalactic sources. How much fraction of this emission is ultimately resolved into Galactic point sources has been a long standing question (e.g., Worrall \\& Marshall 1983; Ottmann \\& Schmitt 1992; Mukai \\& Shiokawa 1992; Ebisawa et al.\\ 2005; Revnivtsev et al.\\ 2006; Revnivtsev \\& Sazonov 2007).  A future ultra-deep Chandra  observation on the Galactic Ridge is expected to give the answer.   \\subsection{Plasma Density Diagnostics of the GRXE}  We obtained the GRXE flux (excluding the extragalactic sources) as  $\\sim4.8\\times 10^{-11}$ ergs s$^{-1}$ cm$^{-2}$ deg$^{-2}$ (2 -- 10 keV) in the direction of $(l,b)\\approx(\\timeform{28.46D},\\timeform{-0.20D})$. Assuming that distance to the Galactic center is  8.5 kpc and radius of the Galactic disk 20 kpc, distance to the edge of the Galactic disk in this direction  is $\\sim25$ kpc. Hence, we estimate the GRXE emissivity as $\\sim2.6\\times10^{-29}$ erg s$^{-1}$ cm$^{-3}$ or $\\sim7.7 \\times 10^{26}$ erg s$^{-1}$ pc$^{-3}$. Thermal bremsstrahlung emissivity is given as $\\sim2.4 \\times 10^{-27}\\: T^{1/2}\\:  N_e^2 $ erg cm$^{-3}$ s$^{-1}$ (Zombeck 2007), thus if we assume uniformly distributed hot plasma with $T \\approx 6 \\times 10^7 $K,  the electron density is $N_e \\sim 1.2 \\times 10^{-3}$ cm$^{-3}$, which is extremely tenuous.  The Galactic volume corresponding to 1 deg$^2$ in the direction of $(l,b)\\approx(\\timeform{28.5D},\\timeform{0.0D})$ is $\\sim1.6 \\times 10^9$ pc$^3$. If the GRXE is ultimately resolved into point sources such as cataclysmic variables or active binary stars, $\\sim10^5$ sources are expected per deg$^2$ (Revnivtsev et al.\\ 2006).  Hence the spatial density of such sources is $\\sim6.3 \\times 10^{-5}$ pc$^{-3}$, and the average luminosity is $\\sim1.1 \\times 10^{31}$ erg s$^{-1}$ per source. We estimate volume of the plasma associated with each point source as $(f  R_\\odot)^3$, where $f$ is the scaling factor, $\\approx 1$ for active binaries and $\\lesssim 0.01$ for cataclysmic variables (white dwarfs). In the case of magnetic cataclysmic variables, which is supposed to be responsible for hard ($\\gtrsim$10 keV) part of the GRXE (Revnivtsev et al.\\ 2006), $f$ is expected to be particularly small, since accretion flow is concentrated on the small magnetic pole regions. Again assuming $T \\approx 6 \\times 10^7 $K, we obtain $N_e \\sim 4 \\times 10^{10} \\: f^{-3/2}$ cm$^{-3}$.  Thus, we see that the GRXE plasma densities expected for the diffuse case and the point source case are  many orders of magnitude different, which is distinguishable from plasma diagnostics. In near future, if we could resolve the He-like K$\\alpha$ complex into individual lines, we may constrain density of the plasma, hence origin of the plasma emission. As a matter of fact, the resonance line (w) at 6699 eV is always the strongest, whereas emissivity of the forbidden line (z) at 6634 eV and those of the inter-combination lines (x,y) at 6680 eV and 6665 eV are dependent on the plasma density. In the high density plasma, the forbidden line is not  supposed to be observed, while in the low-density coronal limit, the forbidden line is stronger than the inter-combination lines. Such precise spectral observations with a few eV  resolution  would have been already possible with Suzaku XRS (Kelley et al.\\ 2007), and will be certainly realized by future calorimeter missions such as NeXT (Inoue and Kunieda 2004)."
    },
    "0711/0711.2153_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:1}  The FIGGS (Faint Irregular Galaxy GMRT Survey) is aimed at creating a multi-wavelength observational data base for a volume limited sample of the faintest gas rich galaxies. As described in more detail in the contribution by Begum et al. in these proceedings, the galaxies form  an HI flux and optical diameter limited subsample of the Karachentsev et al.(2004) catalog of galaxies within 10 Mpc.  The sample consists of  65 galaxies with M$_B \\widetilde{>}-14.5$ with median ${\\rm{M_B}} \\sim -13$ and a median HI mass $\\sim 3 \\times 10^7$~M$_\\odot$. HI aperture synthesis data (from the Giant Meterwave Radio Telescope - GMRT) has been obtained for all galaxies in the sample. Because the GMRT has a hybrid configuration (see Swarup et al. (1991)) images at a variety of spatial resolutions (ranging from $\\sim 40^{``}$ to $\\sim 3^{``}$) can be made from a single GMRT observation run. Galaxies in the FIGGS survey have substantially lower M$_{\\rm HI}$ and L$_{B}$ that typical of galaxies in earlier aperture synthesis surveys.  The GMRT observations also used a velocity resolution ($\\sim$1.6 kms$^{-1}$), that is $\\sim$ 4 times better than most earlier interferometric studies of such faint dwarf galaxies. This high velocity resolution is crucial to detect large scale velocity gradients, which cannot be clearly distinguished in lower velocity resolution observations (see e.g.  Begum et al. 2003a, 2003b, 2004a, 2004b, and for contrast Lo et al. 1993). In this paper we discuss two very gas rich galaxies that were observed as part of the FIGGS survey, viz. NGC~3741 and AndIV. ",
        "conclusions": ""
    },
    "0711/0711.4964_arXiv.txt": {
        "abstract": "In various astronomical projects it is crucial to have coordinates indexed tables. All sky optical and IR catalogues have up to 1 billion objects that will increase with forthcoming projects. Also partial sky surveys at various wavelengths can collect information (not just source lists) which can be saved in coordinate ordered tables. Selecting a sub-set of these entries or cross-matching them could be un-feasible if no indexing is performed. Sky tessellation with various mapping functions have been proposed. It is a matter of fact that the astronomical community is accepting the HTM and HEALPix schema as the default for object catalogues and for maps visualization and analysis, respectively. Within the MCS library project, we have now made available as MySQL-callable functions various HTM and HEALPix facilities. This is made possible thanks to the capability offered by MySQL 5.1 to add external plug-ins. The DIF (Dynamic Indexing Facilities) package distributed within the MCS library, creates and manages a combination of Views, Triggers, DB-engine and plug-ins allowing the user to deal with database tables indexed using one or both these pixelisation schema in a completely transparent way. ",
        "introduction": "Spatial indexes has always been an important issue for multi dimensional data sets in relational databases (DBs), in particular for those dealing with spherical coordinates, e.g. latitude/longitude for Earth locations or RA/Dec for celestial objects. Some DB servers offer built-in capabilities to create indexes on these (coordinate) columns which consequently speed up the execution of queries involving them. However 1. the use of these facilities could be not easy, 2. they typically use a syntax quite different from the astronomical one, 3. their performance is inadequate for the astronomical use.  Within the MCS library project (Calderone \\& Nicastro 2007; Nicastro \\& Calderone 2006, 2007; \\htmladdURL{ross.iasfbo.inaf.it/MCS/}) we have implemented the DIF package, a tool which performs and manages in a fully automatic way the sky pixelisation with both the HTM (Kunszt et al.\\ 2001) and HEALPix (G\\'orski et al.\\ 2005) schema. Using a simple tool, any DB table with sky coordinates columns can be easily indexed. This is achieved by using the facilities offered by the MySQL DB server (which is the only server MCS supports at the moment), i.e. triggers, views and plugins. Having a table with sky coordinates, the user can make it fully indexed in order to perform quick queries on rectangular and circular regions (cone) or to create an HEALPix map file. An SQL query to select objects in a cone will look like this: \\verb|SELECT * FROM MyCatalogue WHERE| \\verb|EntriesInCone(20, 30, 5)|, where (20,30) are the coordinates of the center in degrees and 5 is the radius in arcmin. The important thing to note is that the DB manager needs to supply only a few parameters in the configuration phase, whereas the generic user does not need to know anything about the sky pixelisation either for \\texttt{SELECT} or \\texttt{INSERT} or \\texttt{UPDATE} queries. It also demonstrates that there is no need to extend standard SQL for astronomical queries (see ADQL), at least if MySQL is used as DB server. ",
        "conclusions": ""
    },
    "0711/0711.3858_arXiv.txt": {
        "abstract": "The $^6$Li abundance observed in metal poor halo stars appears to exhibit a plateau as a function of metallicity similar to that for $^7$Li,  suggesting a big bang  origin. However,  the inferred primordial abundance of  $^6$Li is $\\sim$1000 times larger than that predicted by  standard big bang nucleosynthesis for the baryon-to-photon ratio  inferred from the WMAP data. Also, the inferred $^7$Li primordial abundance is 3 times smaller than the big bang prediction.    We here describe in detail a possible simultaneous solution to both the problems of underproduction of $^6$Li and overproduction of $^7$Li in big bang nucleosynthesis.  This solution involves a hypothetical massive, negatively-charged leptonic particle that would bind to the light nuclei produced in big bang nucleosynthesis, but would decay long before it could be detected.   We consider only the $X$-nuclear reactions and assume that the effect of decay products is negligible, as would be the case if lifetime were large or the mass difference between the charged particle and its daughter were small. An interesting  feature of this paradigm is that, because the particle remains bound to the existing nuclei after the cessation of the usual big bang nuclear reactions, a second longer epoch of nucleosynthesis can occur among $X$-nuclei which have reduced Coulomb barriers.  The existence of the $X^-$ particles thus extends big bang nucleosynthesis from the first few minutes to roughly five hours. We confirm that reactions in which the hypothetical particle is transferred can occur that greatly enhance the production of $^6$Li while depleting $^7$Li. We also identify a new reaction that destroys large amounts of $^7$Be, and hence reduces the ultimate $^7$Li abundance.  Thus, big-bang  nucleosynthesis in the presence of these hypothetical particles, together with or without an event of stellar processing, can simultaneously solve the two Li abundance problems. ",
        "introduction": "\\label{sec1}  A long-standing effort in astrophysics has involved searches for signatures of unstable particles that might have existed in the early Universe, but have long since become extinct.  The precision that exists in the abundances of the nuclides produced in  big bang nucleosynthesis (BBN) suggests that primordial nucleosynthesis  is a good place to look for such signatures.  In that context, it is of considerable interest that recent observations~\\citep[e.g.][]{asp06} of metal poor halo stars (MPHSs) indicate that the primordial abundances of both $^6$Li and $^7$Li  may  not be in agreement with the predictions of standard BBN. Specifically, the $^6$Li appears to  have an abundance plateau similar to that for $^7$Li in very low metallicity stars.  This suggests a primordial abundance.  However the abundance value is roughly a factor of 1000 larger than predicted in standard BBN.  A similar, though much less severe, problem exists for $^7$Li~\\citep{rya00,mel04,asp06}; the BBN value that would be consistent with the baryon-to-photon ratio inferred from analysis of the Wilkinson Microwave Anisotropy Probe (WMAP)-power spectrum suggests an abundance that is roughly a factor of three higher than the observed value.  In response to this intriguing  $^6$Li result, a number of solutions have been suggested.  Some relate the Li anomalies to the possible existence of unstable heavy particles in the early Universe~\\citep{dim88a,dim88b,dim89,jed00,cyb03,jed04,kaw05,kus06} which lead to $^6$Li production via non-thermal reactions induced by the particle decay.  Some have suggested a possible epoch of enhanced cosmic ray nucleosynthesis via the $\\alpha + \\alpha$ reactions~\\citep{rol06,pra06,kus07a}.  Of interest to the present work, however,  is the suggestion in several recent works~\\citep{cyb06,kap06,pos07,koh07,ham07,kus07b,jedamzik08a,jedamzik08b} that heavy negatively charged unstable particles could modify BBN and lead to $^6$Li production.  However, this mechanism would operate in a  rather different way from the previous suggestions.  In this paradigm, the heavy particles (here denoted as $X^-$) would bind to the nuclei synthesized during BBN to produce exotic nuclei (hereafter denoted as $X$-nuclei).  These massive $X^-$ particles would be bound in orbits that would have radii comparable to those of the nuclei to which they were attached.  Their presence, however,  would reduce the Coulomb barrier seen by incident charged particles, thereby enhancing the thermonuclear reaction rates and allowing a longer duration of BBN in an extended epoch of nucleosynthesis.  In the present work we expand upon our previous study~\\citep{kus07b} and describe details of a revised analysis of $X^-$ particle effects on BBN.  The focus of the present work is to better quantify the effects of $X^-$ particles on BBN.  It is possible, however, that both $X^-$ effects and the $X$-particle decay contribute to the abundances.  In a future work we will consider this, but for now we only address the effects of $X^-$ particles and show that this alone can explain the observed lithium abundances.  In standard BBN,  $^6$Li  production is suppressed.  It is synthesized primarily via the $^4$He($d$,$\\gamma$)$^6$Li reaction, which has a very small cross section.  However,~\\citet{pos07}  suggested that a large enhancement of the $^6$Li abundance could result from an $X^-$ bound to $^4$He (denoted $^4$He$_X$).  This allows for the $X^-$ transfer reaction of $^4$He$_X$($d$,$X^-$)$^6$Li, which can enhance the  production of $^6$Li by as much as seven orders of magnitude.  Subsequently,~\\citet{ham07} carried out a theoretical calculation of the cross section for this $X^-$ transfer reaction in a quantum three-body model.  Their value was about an order of magnitude smaller than that of~\\citet{pos07}.  This difference can be traced  to their exact treatment of the quantum tunneling in the fusion process and  the use of a better nuclear potential.  In other work,~\\citet{cyb06} identified the $X^-$ as the supersymmetric counterpart of a tau, i.e., a stau., thereby providing some plausibility  for the existence of such particles.  They also considered the $X^-$ transfer reactions for $^6$Li, $^7$Li, and $^7$Be production as Pospelov suggested.  \\citet{kap06} observed that the decay of an $X^-$ when bound to $^4$He  occasionally knocks out a proton or neutron to produce $^3$He or $^3$H, thereby enhancing their abundances in BBN.  Those nuclides could then interact with other primordial $^4$He nuclei at higher energies than those normally associated with the BBN production of  $^6$Li.  \\citet{koh07} and~\\citet{kus07b} studied the recombination of nuclei with $X^-$ particles, and suggested the possibility that BBN with charged massive particles also destroys $^7$Li and so solves the $^7$Li problem as well.  In addition, a resonant reaction $^7$Be$_X$($p$,$\\gamma$)$^8$B$_X$ has recently been proposed (Bird et al. 2007) that further destroys $^7$Be$_X$ through an atomic excited state of $^8$B$_X$.  In the present work we have estimated new reaction rates for the $X^-$ transfer reactions suggested by~\\citet{pos07} as well as others that could occur.  We have reanalyzed the $X^-$ transfer reaction rates estimated in~\\citet{cyb06} and  found them to be negligible based upon the reaction dynamics discussed below.  We have adopted a simple model to estimate the binding energy due to the bound $X^-$ particles for each nucleus involved in BBN.  The reaction rates are corrected for the modified nuclear charges and the effective mass resulting from the binding of the $X^-$ particle.  The reaction rates, however,   are not particularly sensitive to the mass of the $X^-$ particle.   Therefore, we have simply assumed that they are  much heavier than the nucleon mass.  In this paper we study the potential consequences of $X^-$ particles on BBN, including recombination of the $X^-$ particles by the existing normal nuclei.  Several new reaction processes have been included that can produce or destroy $^6$Li and $^7$Li in the early universe.  We utilize a fully dynamical description which is solved numerically, and which includes all relevant details of kinetic and chemical equilibrium.  The purpose of this present work is to better quantify whether there are  observable consequences of the resulting reaction network.  The ultimate goal is to better explain  the observed overproduction (1000 times) and underproduction (3 times) of  the abundances of $^6$Li and $^7$Li.  In Section~\\ref{sec2} we outline some of the details of the calculations we have performed to estimate the thermonuclear reaction rates, and to describe the evolution and abundance  of the different nuclei, $X^-$ particles, and $X$-nuclei during the BBN.  In Section~\\ref{sec3}, we show the results of BBN with the reactions involving the nuclides that could be formed with the embedded $X^-$ particles.     As was done in~\\citet{jedamzik08a}, we identify the parameter space for solving both the the $^6$Li and $^7$Li BBN abundance problems.  The present work, however,  differs from that of \\citet{jedamzik08a} in several important ways.   Besides differences in some of the relevant reaction rates, there are two major differences.  One is the inclusion of $X$-transfer reactions.  \\citet{jedamzik08a} found that $X$-transfer reactions involving the charge $Z = 1$ $X$-nuclei could possibly change the light-element abundances.  That result, however,  relied on  reaction rates which were calculated within the framework of the Born approximation.  Recently, however, a detailed study of one of the relevant $X$-transfer reactions by~\\citet{ham07} shows that a more realistic calculation gives a much smaller (factor of 10)  $X$-transfer reaction rate.  Based upon that study, we conclude that $X$-transfer reactions can be neglected in this work.    Another difference  is the assumption that hadronic and electromagnetic (EM)  decays are unimportant.  That is, we only consider the $X$-nuclear reactions and assume that the effect the of decay products is negligible.  This is possible if the decay lifetimes are long and/or  the mass difference between the $X^-$  and daughter particle is small. In supesymmetric  scenarios for example~\\citep[e.g.][]{feng03a,feng03b}, nearly equal masses  would also imply a larger lifetime. Very large lifetimes are ruled out but lifetimes of order a year  are very interesting for structure formation \\citep{Sigurdson04}.  In Section~\\ref{sec4}, the effects  $X^-$ decay are discussed.  Our conclusions are summarized in Section~\\ref{sec5}. ",
        "conclusions": "\\label{sec5}  We have investigated light-element nucleosynthesis during the big bang in the presence  of massive, negatively-charged  $X^-$ particles.  Such particles would bind to light nuclei in the early universe. As suggested by many authors, they would facilitate BBN by enhancing the nuclear reaction rates  both by reducing the charge of the bound $X$-nuclei, and by enabling transfer reactions involving the $X^-$ particles.  We considered the recombination processes of $X^-$ particles and normal nuclei. Our conclusions are as follows:  First, as suggested in previous studies, the $X^-$ particles greatly enhance the production of $^6$Li.  The main production process of $^6$Li is the sequence of $^4$He$_X$ production through $X^-$ capture on $^4$He followed by  the $X^-$ transfer reaction $^4$He$_X$($d$,$X^-$)$^6$Li.  The resultant $^7$Li abundance, however is almost the same as normal BBN value unless there is a large $X^-$ abundance $Y_X \\gtrsim 0.1$.  Secondly, when the lifetime of the $X^-$ particle is much longer than the period of normal BBN, the $^6$Li abundance monotonically increases with the $X^-$ particle abundance, except for very small $X^-$ abundance ($Y_X \\lesssim 10^{-10}$).  On the other hand, the $^7$Li abundance is nearly independent of  the $X^-$ particle abundance unless the $X^-$ particle abundance is larger than $\\sim 0.1$ times the total abundance of baryons.  In this case the $^7$Li abundance decreases with the $X^-$ particle abundance due to the resonance reaction of $^7$Be$_X$($p$,$\\gamma$)$^8$B$_X$ which reduces the $^7$Li abundance.  Thirdly, the $^6$Li/H and $^7$Li/H observed in MPHSs can constrain the lifetime and abundance of an $X^-$ particle.  These observational constraints  require the lifetime and abundance to be in the ranges of $\\tau_X \\approx (1.0-1.8) \\times 10^3$~s and $Y_X \\gtrsim 0.9$.  When the reaction $^7$Be$_X \\rightarrow ^7$Li+$X^0$ is also taken into account, these ranges change to $\\tau_X \\approx (1.4-2.6) \\times 10^3$~s and $Y_X \\approx 0.04-0.2$.  Therefore, introducing $X^-$ particles with an adequate lifetime and abundance can be a solution for both of the factor of $\\sim 1000$ underproduction of $^6$Li and the factor of 3-4 overproduction of $^7$Li in standard BBN.  Fourthly, a constraint on the $X^-$ particle mass can be made from the dark-matter content deduced from the WMAP analysis of the CMB. If the abundance of $X^-$ particles is $Y_X\\gtrsim 0.1-1$ as summarized above, the mass of dark matter particles $Y^0$, which are produced from the decay of $X^-$ particles, turns out to be $m_Y\\lesssim 10-100$~GeV, thus leading to the constraint $m_X\\lesssim O(100~{\\rm GeV})$ when $m_X \\sim m_Y$.  In summary, although several possible solutions have been proposed to solve the underproduction problem of $^6$Li, they do not necessarily resolve the overproduction problem of $^7$Li simultaneously.  \\citet{pos07} and~\\citet{cyb06} proposed the $X^-$ transfer reactions to produce $^6$Li and $^7$Li.  Regarding $^6$Li production, however,~\\citet{ham07} have shown that the assumed reaction cross section is not as large as the value deduced by Pospelov.  \\citet{bir07} proposed new destruction processes of $^7$Be$_X$ through atomic excitations, and we here investigated yet another destruction process through a $^8$B$^\\ast_X$ nuclear excited resonance in order to resolve the $^7$Li overproduction problem.  The destruction efficiency of these newly proposed processes depends on the excitation energies above the $p$+$^7$Be$_X$ separation threshold.  We then found that a better estimation of the binding energies leads to negligible effect from the nuclear excited resonance. It would be useful, however, in future work to predict more precisely the binding energies and excited states of exotic $X$-nuclei and their reaction cross sections utilizing  a more realistic quantum mechanical treatment."
    },
    "0711/0711.4363_arXiv.txt": {
        "abstract": "We study the correlation between the radio, mid-infrared and far-infrared properties for a sample of 28 blue compact dwarf (BCD) and low metallicity star-forming galaxies observed by {\\em Spitzer}. We find that these sources extend the same far-infrared to radio correlation typical of star forming late type galaxies to lower luminosities. In BCDs, the 24$\\mu$m (or 22$\\mu$m) mid-infrared to radio correlation is similar to starburst galaxies, though there is somewhat larger dispersion in their q$_{24}$ parameter compared to their q$_{\\rm FIR}$. No strong correlations between the q parameter and galaxy metallicity or effective dust temperature have been detected, though there is a hint of decreasing q$_{24}$ at low metallicities. The two lowest metallicity dwarfs in our sample, IZw18 and SBS0335-052E, despite their similar chemical abundance, deviate from the average q$_{24}$ ratio in opposite manners, displaying an apparent radio excess and dust excess respectively. ",
        "introduction": "More than 35 years ago, a correlation between the 10$\\mu$m mid-infrared (mid-IR) and 1.4\\,GHz luminosities in galactic nuclei was first noticed by \\citet{van71}. In 1983, the launch of the {\\em Infrared Astronomical Satellite (IRAS)} brought a new era in collecting large samples of data in the infrared, and sparked numerous studies investigating the above mentioned correlation in star forming galaxies. Both the radio and infrared emission are closely related to star formation activities. The radio continuum is known to originate from two processes: thermal free-free emission from ionized gas in H{\\sc ii} regions and non-thermal synchrotron radiation from relativistic electrons that are accelerated in the supernovae remnants (SNRs) \\citep[see][ for a review]{Condon92}. The thermal emission usually has a rather flat spectrum with f$_{\\nu}$\\,$\\propto$\\,$\\nu^{-0.1}$ while the non-thermal component often has a much steeper spectral slope with f$_{\\nu}$\\,$\\propto$\\,$\\nu^{-0.8}$.  In normal galaxies the relative contribution of the two components varies with frequency and at 1.4\\,GHz the radio continuum is dominated (at nearly $\\sim$90\\%) by the non-thermal component. The infrared emission is due to the thermal re-radiation of starlight from dust surrounding H{\\sc ii} regions \\citep{Harwit75}.  Based mostly on {\\em IRAS} data, it has been shown that the far-infrared (FIR, 40-120\\,$\\mu$m) to radio correlation holds well over a remarkably wide range of star forming galaxies, spanning several orders of magnitude in luminosities \\citep{Helou85, Dejong85, Condon91,Yun01}.  The availability of deep observations of distant galaxies by the {\\em Infrared Space Observatory (ISO)} and recently by {\\em Spitzer} \\citep{Werner04} showed that the correlation between the IR and radio emission is not limited to the local universe but also extends in galaxies at higher redshifts \\citep{Garrett02,Appleton04}. It was also revealed that in addition to the FIR, the mid-IR emission at $\\sim$24$\\mu$m correlates with the 20cm radio continuum, though with more scatter \\citep{Elbaz02,Gruppioni03,Appleton04,Wu05}. A number of very deep {\\em Spitzer} mid-IR and FIR surveys can now probe a population of galaxies with low infrared luminosities for which ancillary data, including deep radio imaging, are becoming available \\citep[e.g.][]{Jannuzi99,Sanders07,Rosenberg06}. This is particularly interesting since in the near future $\\sim20\\mu$m is the longest wavelength which will likely be probed by the James Web Space Telescope.  It is thus instructive to examine the mid-IR to radio correlation in more detail in these low luminosity nearby systems, for which little is known to date.  As ubiquitous as the FIR/radio correlation appears to be, there are a few significant deviations from it.  Radio excess exists in some instances, such as galaxies hosting an active galactic nucleus (AGN). External magnetic field compression due to the interaction with nearby galaxies could also produce extra emission in the radio continuum \\citep{Miller01}. Conversely, synchrotron deficiency has been found in some nascent starburst galaxies studied by \\citet{Roussel03,Roussel06} which was attributed to the lack of time for massive young stars to evolve into supernovae (SN) since these galaxies are just at the onset of a starburst episode. On the other hand, in the infrared, dust emission can be damped in an optically-thin environment, because the ultraviolet (UV) and optical light may not be fully reprocessed by the dust, as seen in low luminosity dwarf galaxies.   Blue compact dwarf galaxies are a group of extra-galactic objects that are characterized by their blue optical colors, small sizes ($\\le$1\\,kpc) and low luminosities (M$_B>$-18). These galaxies do not display any AGN signature\\footnote{A possible exception among our sources is CG0752 in which the high ionization lines [NeV]$\\lambda$14.3/24.3$\\mu$m are detected \\citep[see][]{Hao07}.} and have recent bursts of star formation in a relatively unevolved chemical environment. As such they have been proposed as nearby analogs of star formation in young galaxies in the early universe.  In a metal poor environment, star forming regions are usually optically thin and emit less in the infrared.  However, \\citet{Devereux89} suggested that in a low luminosity galaxy, the radio emission also decreases, and probably faster than the infrared.  The deficiency in both the non-thermal radio and FIR emission may counterbalance each other and result in a similar FIR/radio ratio to the one observed in normal spiral galaxies \\citep{Klein91}.  This was examined by a study of star formation rates (SFRs) in BCDs performed by \\citet{Hopkins02}, in which the authors found an excellent agreement between the SFRs estimated from 1.4\\,GHz and 60\\,$\\mu$m luminosities.  As \\citet{Bell03} has pointed out though, the FIR/radio correlation is almost linear, not because the IR and radio emission reflect the SFRs correctly, but because in low luminosity galaxies they are both underestimated by similar factors.  This is in agreement with \\citet{Helou93}, who found from their modelling work in disk galaxies that the transparency of the disk was about the same to both re-emission processes.  \\citet{Hunt05a} in their analysis of the spectral energy distributions (SEDs) of low metallicity BCDs, noticed that these systems do not follow several of the usual correlations between the mid-IR, FIR and radio emission and display a scatter of a factor of $\\sim$10.  The {\\em Spitzer} Space telescope has enabled us to study the infrared properties of a large sample of BCDs, probing the lower end of the luminosity and metallicity range. In this paper, the fifth in a series \\citep{Houck04b,Wu06,Wu07a,Wu07b}, we examine their mid-IR and FIR to radio correlation extending the work of \\citet{Hopkins02}. We describe the sample selection and the observational data in \\S 2. A detailed study of mid-IR and FIR/radio correlation, as well as its dependence with other parameters, such as metallicity and dust temperature are presented in \\S 3.  We also discuss two extreme cases, IZw18 and SBS0335-052 in \\S 3. We summarize our conclusions in \\S 4. ",
        "conclusions": "We have studied the mid-IR and FIR to radio correlation in a sample of dwarf star-forming galaxies spanning a metallicity range from 7.2$<$12+log(O/H)$<$8.9, using {\\em Spitzer} IRS/MIPS, as well as radio 1.4\\,GHz data obtained from the literature.  The BCD sample appears to follow the same FIR/radio correlation as normal star forming galaxies. The analysis based on mid-IR and radio data reveals a similar correlation, though the scatter is larger, probably due to an intrinsically higher variation in the 15-30$\\mu$m SEDs of dwarf galaxies.  When comparing the q$_{24}$ ratios with metallicity or effective dust temperature, we find no strong correlation, though there is a general trend of lower q ratios at lower metallicity for galaxies with 12+log(O/H)$<$8.0 and the correlation flattens out toward higher metallicity. In general the SFRs estimated from the radio 1.4\\,GHz continuum and the mid-IR data are in good agreement. Two extremely metal poor BCDs, IZw18 and SBS0335-052E appear to deviate from the average q$_{24}$ by $\\sim$2$\\sigma$, with one displaying a radio excess and the other an infrared excess."
    },
    "0711/0711.4822_arXiv.txt": {
        "abstract": "{The Lockman Hole field represents the sky area of lowest Galactic line-of-sight column density N$_{\\rm H}$=5.7$\\times$10$^{19}$ cm$^{-2}$. It was observed by the XMM-{\\em Newton} X-ray observatory in 18 pointings for a total of 1.16 Msec (raw EPIC pn observing time) constituting the deepest XMM-{\\em Newton} exposure so far. After the removal of time intervals with high particle background, the final effective exposure was 637 ks.} {We present a catalogue of the  X-ray sources detected in the central 0.196 deg$^{2}$ of the field and discuss the derived number counts and X-ray colours.} {The analysis was performed using  the XMM-SAS data analysis package version 6.0. The detection sensitivity and fraction of spurious detections was calibrated using Monte Carlo simulations.} {In the 0.5--2.0~keV band, a sensitivity limit (defined as the faintest detectable source) of 1.9$\\times10^{-16}$ erg cm$^{-2}$ s$^{-1}$ was reached. The 2.0--10.0~keV band and 5.0--10.0~keV band sensitivity limits were 9$\\times10^{-16}$ erg cm$^{-2}$ s$^{-1}$ and 1.8$\\times10^{-15}$ erg cm$^{-2}$ s$^{-1}$, respectively. A total of 409 sources above a detection likelihood of 10 (3.9 sigma) were found within a radius of 15$^\\prime$ off the field centre, of which 340, 266, and 98 sources were detected in the soft, hard, and very hard bands, respectively. The number counts in each energy band are in close agreement with results from previous surveys and with the synthesis models of the  X-ray background. A 6\\%  of Compton-thick source candidates have been selected from the  X-ray colour-colour diagram. This fraction is consistent with the most recent predictions of  X-ray background population synthesis models at our flux limits. We also estimated, for the first time, the logN-logS relation for Compton-thick AGN. } {} ",
        "introduction": "Due to its extremely low line of sight column density of N$_{\\rm H}$=5.7$\\times$10$^{19}$ cm$^{-2}$, the Lockman Hole field is one of the main extragalactic deep-survey target areas of recent X-ray observatories. The choice of a low N$_{\\rm H}$  field on the one hand optimizes the soft-band sensitivity compared to fields with higher N$_{\\rm H}$. An even more important aspect is that for many sources we can determine spectra to the lowest possible energies and thus get a better handle on the continuum and intrinsic absorption.  The field was observed for a total of 1.1 Msec, by the ROSAT HRI detector, reaching a sensitivity limit of 1.2~$\\cdot 10^{-15}$ erg cm$^{-2}$ s$^{-1}$ in the 0.5--2.0~keV band, for the first time resolving the majority of the soft X-ray background into discrete sources (Hasinger et al. \\cite{hasinger98}). The Lockman Hole was also the target of the deepest observations conducted by the XMM-{\\em Newton} observatory, considerably improving the sensitivity limit in the soft band and extending the survey into the hard X-ray band up to an energy of 10 keV. Although the sensitivity limit of the XMM-{\\em Newton} Lockman Hole survey does not reach the  sensitivity of the deepest surveys performed by the {\\em Chandra} observatory (i.e. the {\\em  Chandra} deep field North and South, Bauer et al. \\cite{bauer}, Giacconi et al. \\cite{giac}), the larger collecting area of XMM-{\\em Newton} offers the possibility to perform a spectral analysis of a large fraction of the detected sources down to the confusion limit, opening up a powerful window into the cosmological evolution of the X-ray universe.  The Lockman Hole field dataset allowed us for example to study the Fe K line emission of one of the most distant X-ray selected cluster of galaxies, RX J1053.7+5735l. Hashimoto et al. (\\cite{hashimoto}) were able to determine the X-ray redshift of the cluster and its metallicity. The XMM-{\\em Newton} Lockman Hole survey is also well suited to investigate the X-ray spectra and cosmological evolution of both absorbed and unabsorbed AGN.  With its  large  collecting area,  XMM-{\\em Newton} is able to obtain detailed spectral information even for the faintest sources. Mainieri et al. (\\cite{mainieri}) and Mateos et al. (\\cite{mateos}) performed a detailed spectral analysis and they discovered several candidate Compton thick AGNs. According to the recent model of Gilli et al. (2007),  40\\% of the flux  of the  X-ray background is expected to arise in  highly  absorbed AGNs (N$_{\\rm H}>$10$^{22}$ cm$^{-2}$). These sources are mostly faint and detectable with deep X-ray surveys  such as the XMM-{\\em Newton} Lockman Hole  survey. Worsley et al. (\\cite{worsley04}) estimated that  the XMM-{\\em Newton} Lockman Hole observations resolve $\\sim$100\\% of the soft  X-ray background into discrete sources while at  energies above 5 keV this fraction drops to $\\sim$50\\%. The missing part therefore could be due to very faint, highly  absorbed AGNs (N$_{\\rm H}>$10$^{22}$ cm$^{-2}$) to which XMM-{\\em Newton} is not sensitive. ",
        "conclusions": "We analyzed a set of 18 XMM-{\\em Newton} pointings centred on the Lockman Hole which constitute the deepest exposure performed by XMM-{\\em Newton}, reaching a sensitivity limit of 1.9$\\times10^{-16}$ erg cm$^{-2}$ s$^{-1}$,  9$\\times 10^{-16}$ erg cm$^{-2}$ s$^{-1}$ and 1.8$\\times10^{-15}$ erg cm$^{-2}$ s$^{-1}$, in the 0.5--2.0 keV, 2.0--10.0 keV and 5.0--10.0~keV energy bands, respectively. 409 sources were detected within the survey area of 0.20 sq. degrees. 340, 266, and 98 objects were detected in the soft, hard, and very hard band, respectively. A catalogue including the main X-ray characteristics of the sources in the Lockman hole survey  is presented. The number counts derived from the XMM-{\\em Newton} data are in close agreements with previous surveys and the most recent synthesis models of the  X-ray background. Differences in the  normalization of our  logN-logS with those of other surveys have been  discussed and explained in terms of low counting statistics typical of deep pencil beam surveys. The high throughput of the XMM-{\\em Newton} telescope allowed us to compute and constrain the parameters of the logN-logS in the hard and very hard energy bands. This region of the X-ray spectrum is fundamental for understanding absorbed AGNs which are thought to be the most important  contributors to the total flux of the  X-ray background. We also  present the colour-colour diagram for our X-ray sources, which confirms, together with the spectral analysis on this field performed by Mateos et al. (\\cite{mateos}) and Mainieri et al. (\\cite{mainieri}), the presence of a conspicuous number of highly absorbed AGNs together with evidence of the increasing fraction of this kind of objects at low X-ray fluxes.   \\begin{sidewaystable*} \\begin{minipage}[t][180mm]{\\textwidth} \\caption[]{Catalogue of X-ray sources with likelihood of detection in 0.5--10~keV band greater than 10 (3.9 $\\sigma$). Description of columns: (1) XMM-Newton source number as displayed in Fig. 3, (2) ROSAT source number (Lehmann et al. \\cite{lehmann}), (3) IAU source name, (4--6) J2000 coordinates and errors, (7--18) likelihood of detection (rounded to nearest integer), fluxes and flux errors [$10^{-16}$ erg cm$^{-2}$ s$^{-1}$] in 0.5--10, 0.5--2, 2--10, and 5--10~keV band, (19--22) hardness ratios and errors (see text for band definitions). All errors are 1 $\\sigma$ errors.} \\label{souCat} \\centering {\\scriptsize \\begin{tabular}{rrcllcrrrrrrrrrrrrrrrr} \\hline \\multicolumn{2}{c|}{SRC \\#}&\\multicolumn{1}{c|}{}&\\multicolumn{2}{c}{RA~~~(J2000)~~~DEC}&err&\\multicolumn{3}{|c|}{0.5--10.0~keV}&\\multicolumn{3}{|c|}{0.5--2.0~keV}&\\multicolumn{3}{|c|}{2.0--10.0~keV}&\\multicolumn{3}{|c|}{5.0--10.0~keV}&&&& \\\\ XMM&\\multicolumn{1}{r|}{R}&\\multicolumn{1}{c|}{IAU name}&hh mm ss&$^{\\circ}$~~~~$^\\prime$~~~$^{\\prime\\prime}$&$^{\\prime\\prime}$&\\multicolumn{1}{|r}{$\\mathcal{L}$}&flux&err&\\multicolumn{1}{|r}{$\\mathcal{L}$}&flux&err& \\multicolumn{1}{|r}{$\\mathcal{L}$}&flux&err&\\multicolumn{1}{|r}{$\\mathcal{L}$}&flux&\\multicolumn{1}{r|}{err}&HR1&err&HR2&err \\\\ \\hline 352 &     &XMMU J105053.0+573028 & 10 50 53.05 & 57 30 28.5 & 2.4 &     13 &    21.0 &  6.1 &     12 &     5.6 &  1.2 &      3 &    15.4 &  6.0 &      0 &     2.6 &  7.6 & -0.39 &  0.19 & -0.79 &  0.56 \\\\ 358 &  17 &XMMU J105104.0+573053 & 10 51  4.04 & 57 30 53.8 & 0.3 &   2242 &   233.4 &  8.7 &   1896 &    87.1 &  2.3 &    311 &   146.3 &  8.4 &     37 &   111.7 & 15.0 & -0.58 &  0.02 & -0.30 &  0.07 \\\\ 458 &     &XMMU J105106.4+573435 & 10 51  6.46 & 57 34 35.6 & 0.4 &    744 &   238.9 & 10.7 &    134 &    18.6 &  1.5 &    492 &   220.3 & 10.6 &    118 &   239.6 & 20.0 &  0.30 &  0.04 & -0.13 &  0.05 \\\\ 430 &  11 &XMMU J105108.3+573345 & 10 51  8.38 & 57 33 45.4 & 0.2 &   3406 &   363.1 & 10.4 &   2522 &   102.0 &  2.4 &    792 &   261.1 & 10.2 &     92 &   186.7 & 16.9 & -0.42 &  0.02 & -0.33 &  0.04 \\\\ 1522 &     &XMMU J105108.4+572833 & 10 51  8.46 & 57 28 33.6 & 1.7 &     15 &    18.9 &  4.3 &      0 &     0.0 &  0.5 &     11 &    18.9 &  4.3 &      5 &    33.9 & 10.1 &  1.00 &  0.30 &  0.12 &  0.19 \\\\ 500 &     &XMMU J105108.9+573553 & 10 51  8.97 & 57 35 53.0 & 1.0 &     78 &    62.0 &  8.1 &     40 &    11.1 &  1.4 &     28 &    51.0 &  8.0 &     10 &    79.1 & 18.5 & -0.16 &  0.10 &  0.05 &  0.14 \\\\ 217 &     &XMMU J105111.8+572635 & 10 51 11.87 & 57 26 35.4 & 0.2 &   1816 &   160.1 &  6.5 &   1531 &    57.1 &  1.6 &    270 &   103.0 &  6.3 &     17 &    56.9 & 10.6 & -0.56 &  0.02 & -0.44 &  0.08 \\\\ 357 &     &XMMU J105112.8+573035 & 10 51 12.87 & 57 30 35.7 & 1.4 &     20 &    18.8 &  4.1 &     14 &     4.1 &  0.8 &      7 &    14.7 &  4.0 &      0 &     9.1 &  9.4 & -0.28 &  0.16 & -0.39 &  0.45 \\\\ 1361 &     &XMMU J105114.8+573551 & 10 51 14.86 & 57 35 51.4 & 1.2 &     31 &    37.1 &  6.1 &      0 &     0.6 &  0.8 &     23 &    36.5 &  6.1 &     10 &    67.1 & 15.5 &  0.81 &  0.22 &  0.13 &  0.14 \\\\ 1489 &     &XMMU J105116.0+572015 & 10 51 16.03 & 57 20 15.2 & 1.7 &     10 &    25.1 &  7.0 &      4 &     3.7 &  1.2 &      5 &    21.4 &  6.8 &      1 &    30.6 & 15.7 & -0.05 &  0.23 &  0.01 &  0.30 \\\\ 1417 &     &XMMU J105116.8+572109 & 10 51 16.82 & 57 21  9.8 & 0.9 &     32 &    35.2 &  5.9 &     15 &     5.6 &  1.1 &     16 &    29.5 &  5.8 &      1 &    25.2 & 13.5 & -0.09 &  0.14 & -0.25 &  0.27 \\\\ 41 &  46 &XMMU J105120.0+571848 & 10 51 20.10 & 57 18 48.9 & 0.4 &    415 &   117.5 & 10.2 &    370 &    45.7 &  2.4 &     38 &    71.8 &  9.9 &      7 &    78.5 & 20.9 & -0.60 &  0.05 & -0.13 &  0.15 \\\\ 202 &     &XMMU J105120.4+572603 & 10 51 20.49 & 57 26  3.7 & 0.7 &    131 &    36.2 &  4.3 &    106 &    10.9 &  0.9 &     25 &    25.3 &  4.1 &      2 &    19.5 &  8.7 & -0.46 &  0.07 & -0.29 &  0.22 \\\\ 226 &     &XMMU J105120.6+572701 & 10 51 20.66 & 57 27  1.7 & 0.3 &    973 &    96.7 &  5.3 &    834 &    35.4 &  1.3 &    127 &    61.3 &  5.1 &     14 &    44.8 &  9.1 & -0.57 &  0.03 & -0.32 &  0.10 \\\\ 2134 &     &XMMU J105120.7+572515 & 10 51 20.74 & 57 25 15.2 & 1.0 &     23 &    26.5 &  4.4 &     12 &     5.5 &  0.9 &     10 &    21.0 &  4.3 &      1 &    22.5 &  8.9 & -0.24 &  0.13 & -0.14 &  0.22 \\\\ 174 &     &XMMU J105120.8+572458 & 10 51 20.82 & 57 24 58.5 & 0.5 &     23 &    60.9 &  5.3 &     12 &    21.9 &  1.3 &     10 &    39.0 &  5.2 &      1 &    21.0 & 10.3 & -0.56 &  0.05 & -0.45 &  0.20 \\\\ 1395 &     &XMMU J105121.2+571938 & 10 51 21.21 & 57 19 38.3 & 2.9 &     13 &    22.7 &  6.8 &     12 &     6.3 &  1.4 &      2 &    16.4 &  6.7 &      0 &    19.8 & 16.3 & -0.42 &  0.19 & -0.08 &  0.46 \\\\ 313 &     &XMMU J105121.2+572932 & 10 51 21.25 & 57 29 32.5 & 1.0 &     41 &    31.6 &  4.4 &      1 &     1.1 &  0.6 &     38 &    30.5 &  4.4 &      3 &    23.4 &  8.8 &  0.62 &  0.17 & -0.30 &  0.18 \\\\ 111 &     &XMMU J105122.3+572214 & 10 51 22.30 & 57 22 14.6 & 0.6 &    122 &    35.0 &  4.9 &    111 &    12.8 &  1.1 &     13 &    22.2 &  4.7 &      0 &    12.3 & 10.7 & -0.57 &  0.08 & -0.44 &  0.36 \\\\ 530 &     &XMMU J105122.6+573734 & 10 51 22.63 & 57 37 34.7 & 1.3 &     21 &    28.5 &  6.1 &     14 &     6.0 &  1.2 &      8 &    22.5 &  6.0 &      0 &    11.6 & 13.2 & -0.25 &  0.16 & -0.47 &  0.46 \\\\ 89 &     &XMMU J105123.3+572048 & 10 51 23.34 & 57 20 48.1 & 0.7 &     85 &    55.9 &  6.4 &      0 &     0.4 &  0.6 &     57 &    55.5 &  6.4 &     30 &   108.0 & 15.9 &  0.92 &  0.13 &  0.16 &  0.09 \\\\ 1458 &     &XMMU J105125.2+572903 & 10 51 25.28 & 57 29  3.8 & 2.5 &    100 &    14.1 &  3.5 &      7 &     0.6 &  0.5 &      1 &    13.5 &  3.5 &      0 &    28.5 &  7.6 &  0.56 &  0.32 &  0.20 &  0.18 \\\\ 354 &  75 &XMMU J105125.4+573050 & 10 51 25.40 & 57 30 50.7 & 0.2 &   2206 &   154.3 &  6.0 &   1873 &    58.4 &  1.6 &    294 &    95.9 &  5.8 &     41 &    74.5 &  9.8 & -0.59 &  0.02 & -0.29 &  0.07 \\\\ 490 &     &XMMU J105125.4+573542 & 10 51 25.45 & 57 35 43.0 & 1.1 &     63 &    17.9 &  4.5 &     65 &    11.4 &  1.2 &      0 &     6.5 &  4.3 &      0 &    11.5 & 11.6 & -0.83 &  0.10 &  0.11 &  0.59 \\\\ 109 &     &XMMU J105126.2+572208 & 10 51 26.24 & 57 22  8.1 & 0.8 &     90 &    20.9 &  4.1 &     91 &    11.0 &  1.0 &      2 &     9.9 &  4.0 &      0 &     0.9 &  5.6 & -0.75 &  0.09 & -0.87 &  0.70 \\\\ 474 &     &XMMU J105128.2+573502 & 10 51 28.20 & 57 35  2.5 & 0.4 &    816 &    88.1 &  5.6 &    772 &    46.9 &  1.8 &     45 &    41.2 &  5.4 &      2 &    28.6 & 11.3 & -0.76 &  0.03 & -0.34 &  0.18 \\\\ 243 &     &XMMU J105128.2+572740 & 10 51 28.23 & 57 27 40.1 & 0.2 &   3042 &   308.8 &  8.2 &   1126 &    40.4 &  1.3 &   1618 &   268.4 &  8.1 &    299 &   220.0 & 12.5 &  0.03 &  0.02 & -0.27 &  0.03 \\\\ 2342 &     &XMMU J105128.7+572856 & 10 51 28.78 & 57 28 56.1 & 1.4 &     11 &    15.3 &  3.3 &      0 &     0.0 &  0.2 &     12 &    15.3 &  3.3 &      1 &    15.5 &  7.3 &  1.00 &  0.16 & -0.17 &  0.25 \\\\ 280 &     &XMMU J105128.9+572823 & 10 51 28.96 & 57 28 23.0 & 1.5 &     11 &     7.6 &  3.0 &     12 &     3.0 &  0.6 &      0 &     4.6 &  2.9 &      0 &     2.7 &  4.8 & -0.61 &  0.21 & -0.41 &  0.78 \\\\ 523 &     &XMMU J105129.9+573736 & 10 51 29.93 & 57 37 36.9 & 0.2 &   1550 &   188.3 &  8.4 &   1346 &    73.4 &  2.2 &    189 &   114.9 &  8.1 &     16 &    85.1 & 16.2 & -0.60 &  0.02 & -0.31 &  0.09 \\\\ 112 &     &XMMU J105130.5+572220 & 10 51 30.50 & 57 22 20.3 & 0.5 &    162 &    14.4 &  3.4 &    167 &    14.4 &  1.0 &      0 &     0.0 &  3.2 &      0 &     1.3 &  4.9 & -1.00 &  0.07 &  1.00 &  7.18 \\\\ 235 &     &XMMU J105130.8+572705 & 10 51 30.89 & 57 27  5.7 & 1.4 &     27 &    22.0 &  3.7 &      5 &     2.1 &  0.6 &     20 &    20.0 &  3.6 &      2 &    17.9 &  7.1 &  0.21 &  0.17 & -0.22 &  0.21 \\\\ 444 &   8 &XMMU J105131.0+573440 & 10 51 31.01 & 57 34 40.0 & 0.1 &  82940 &  1457.9 & 12.5 &  80900 &  1071.2 &  6.4 &   1987 &   386.6 & 10.7 &     56 &   114.1 & 12.8 & -0.89 &  0.00 & -0.66 &  0.03 \\\\ 1405 &     &XMMU J105131.4+571928 & 10 51 31.45 & 57 19 28.9 & 1.3 &     24 &    25.3 &  5.6 &     19 &     6.2 &  1.1 &      6 &    19.1 &  5.5 &      0 &    15.6 & 14.0 & -0.34 &  0.15 & -0.27 &  0.44 \\\\ 213 &     &XMMU J105132.0+572618 & 10 51 32.02 & 57 26 18.7 & 0.6 &    183 &    37.9 &  3.9 &    140 &    11.3 &  0.9 &     35 &    26.7 &  3.8 &      8 &    32.9 &  8.3 & -0.45 &  0.06 & -0.07 &  0.14 \\\\ 160 &     &XMMU J105132.3+572414 & 10 51 32.35 & 57 24 14.1 & 1.5 &     31 &    18.5 &  3.6 &     24 &     4.4 &  0.7 &      9 &    14.1 &  3.5 &      0 &     7.7 &  7.7 & -0.32 &  0.13 & -0.44 &  0.41 \\\\ 2506 &     &XMMU J105133.7+573525 & 10 51 33.75 & 57 35 25.9 & 2.0 &    100 &    13.3 &  4.0 &      7 &     3.1 &  1.0 &      5 &    10.1 &  3.9 &      1 &    13.1 &  9.2 & -0.32 &  0.22 & -0.05 &  0.40 \\\\ 578 &     &XMMU J105134.9+574019 & 10 51 34.95 & 57 40 19.8 & 1.7 &     16 &    47.2 &  9.1 &      0 &     0.4 &  0.9 &     16 &    46.8 &  9.1 &      2 &    46.4 & 21.0 &  0.90 &  0.21 & -0.18 &  0.24 \\\\ 248 &     &XMMU J105135.7+572740 & 10 51 35.72 & 57 27 40.4 & 0.9 &     99 &    37.8 &  3.9 &     24 &     4.0 &  0.6 &     62 &    33.8 &  3.8 &     14 &    34.2 &  7.0 &  0.15 &  0.10 & -0.17 &  0.11 \\\\ 2507 &     &XMMU J105135.7+573521 & 10 51 35.75 & 57 35 21.9 & 2.0 &    100 &     2.4 &  2.0 &      3 &     1.9 &  0.9 &      0 &     0.5 &  1.8 &      0 &     0.0 &  1.8 & -0.92 &  0.28 & -1.00 &  5.08 \\\\ 99 &     &XMMU J105136.0+572135 & 10 51 36.04 & 57 21 35.6 & 0.6 &    156 &    52.6 &  5.0 &     87 &    10.2 &  1.0 &     59 &    42.3 &  4.9 &     11 &    48.1 & 10.8 & -0.21 &  0.07 & -0.11 &  0.12 \\\\ 415 &     &XMMU J105136.3+573303 & 10 51 36.33 & 57 33  3.6 & 0.7 &    129 &    29.1 &  3.6 &    109 &     9.8 &  0.8 &     20 &    19.3 &  3.5 &      1 &    14.0 &  7.3 & -0.52 &  0.07 & -0.32 &  0.25 \\\\ 437 &     &XMMU J105136.4+573333 & 10 51 36.50 & 57 33 33.4 & 1.4 &     18 &    19.5 &  3.7 &      3 &     1.8 &  0.7 &     15 &    17.7 &  3.6 &      1 &    13.4 &  7.8 &  0.23 &  0.20 & -0.31 &  0.28 \\\\ 316 &     &XMMU J105136.9+572940 & 10 51 37.00 & 57 29 40.8 & 0.9 &     57 &    20.7 &  3.2 &     43 &     5.4 &  0.7 &     15 &    15.3 &  3.1 &      0 &     5.5 &  5.8 & -0.38 &  0.10 & -0.59 &  0.35 \\\\ 295 & 837 &XMMU J105137.1+572857 & 10 51 37.18 & 57 28 57.8 & 0.6 &    114 &    19.4 &  3.2 &    102 &     8.7 &  0.8 &      6 &    10.8 &  3.1 &      7 &    26.5 &  6.9 & -0.67 &  0.08 &  0.27 &  0.18 \\\\ 529 &     &XMMU J105137.3+573755 & 10 51 37.38 & 57 37 56.0 & 0.2 &   3616 &   357.3 & 10.7 &   2910 &   118.3 &  2.7 &    634 &   238.9 & 10.4 &     74 &   184.7 & 18.6 & -0.52 &  0.02 & -0.29 &  0.05 \\\\ 353 &  19 &XMMU J105137.3+573043 & 10 51 37.38 & 57 30 43.6 & 0.3 &   1063 &    93.4 &  4.6 &    853 &    29.8 &  1.1 &    202 &    63.7 &  4.5 &     10 &    30.9 &  7.3 & -0.49 &  0.03 & -0.49 &  0.09 \\\\ 212 &     &XMMU J105139.2+572615 & 10 51 39.21 & 57 26 15.5 & 2.2 &     11 &     6.9 &  2.9 &     11 &     2.6 &  0.6 &      0 &     4.2 &  2.8 &      1 &    11.5 &  6.5 & -0.60 &  0.23 &  0.31 &  0.39 \\\\ 234 &     &XMMU J105139.4+572707 & 10 51 39.47 & 57 27  7.1 & 0.7 &     87 &    26.8 &  3.4 &     55 &     6.0 &  0.7 &     28 &    20.8 &  3.3 &      4 &    20.0 &  6.6 & -0.29 &  0.09 & -0.19 &  0.18 \\\\ 491 &     &XMMU J105142.0+573555 & 10 51 42.07 & 57 35 55.2 & 0.4 &    517 &    78.6 &  5.2 &    424 &    25.5 &  1.2 &     94 &    53.1 &  5.0 &      1 &    18.2 &  9.0 & -0.50 &  0.04 & -0.61 &  0.16 \\\\ 467 &     &XMMU J105142.1+573449 & 10 51 42.10 & 57 34 49.3 & 1.0 &     56 &    14.8 &  3.3 &     53 &     6.8 &  0.8 &      2 &     8.0 &  3.2 &      2 &    18.5 &  7.9 & -0.69 &  0.11 &  0.24 &  0.28 \\\\ 1414 &     &XMMU J105143.7+573856 & 10 51 43.75 & 57 38 56.9 & 1.4 &     14 &    25.7 &  5.4 &      0 &     1.3 &  0.9 &     12 &    24.4 &  5.3 &      2 &    32.6 & 13.8 &  0.50 &  0.28 & -0.03 &  0.24 \\\\ 314 &     &XMMU J105143.8+572937 & 10 51 43.83 & 57 29 38.0 & 0.8 &     79 &    13.7 &  2.7 &     73 &     6.7 &  0.7 &      3 &     6.9 &  2.6 &      4 &    18.9 &  6.1 & -0.72 &  0.09 &  0.32 &  0.22 \\\\ 595 &     &XMMU J105143.8+574122 & 10 51 43.87 & 57 41 22.6 & 0.9 &     63 &    53.9 & 11.9 &     63 &    23.1 &  2.6 &      3 &    30.8 & 11.6 &      0 &     4.1 & 14.6 & -0.65 &  0.11 & -0.83 &  0.56 \\\\ 229 &     &XMMU J105144.3+572651 & 10 51 44.37 & 57 26 51.2 & 1.1 &     25 &     9.5 &  2.5 &     24 &     3.6 &  0.6 &      2 &     5.8 &  2.4 &      0 &     3.4 &  4.5 & -0.59 &  0.14 & -0.41 &  0.57 \\\\ 261 &  80 &XMMU J105144.9+572808 & 10 51 44.91 & 57 28  8.0 & 0.2 &   1731 &   116.4 &  4.7 &   1374 &    36.5 &  1.1 &    338 &    79.9 &  4.6 &     20 &    37.1 &  6.6 & -0.49 &  0.02 & -0.51 &  0.07 \\\\ \\hline \\end{tabular}} \\vfill \\end{minipage} \\end{sidewaystable*} \\begin{sidewaystable*} \\begin{minipage}[t][180mm]{\\textwidth} \\bigskip \\centering {\\scriptsize \\begin{tabular}{rrcllcrrrrrrrrrrrrrrrr} \\hline \\multicolumn{2}{c|}{SRC \\#}&\\multicolumn{1}{c|}{}&\\multicolumn{2}{c}{RA~~~(J2000)~~~DEC}&err&\\multicolumn{3}{|c|}{0.5--10.0~keV}&\\multicolumn{3}{|c|}{0.5--2.0~keV}&\\multicolumn{3}{|c|}{2.0--10.0~keV}&\\multicolumn{3}{|c|}{5.0--10.0~keV}&&&& \\\\ XMM&\\multicolumn{1}{r|}{R}&\\multicolumn{1}{c|}{IAU name}&hh mm ss&$^{\\circ}$~~~~$^\\prime$~~~$^{\\prime\\prime}$&$^{\\prime\\prime}$&\\multicolumn{1}{|r}{$\\mathcal{L}$}&flux&err&\\multicolumn{1}{|r}{$\\mathcal{L}$}&flux&err& \\multicolumn{1}{|r}{$\\mathcal{L}$}&flux&err&\\multicolumn{1}{|r}{$\\mathcal{L}$}&flux&\\multicolumn{1}{r|}{err}&HR1&err&HR2&err \\\\ \\hline 2366 &     &XMMU J105145.3+573215 & 10 51 45.36 & 57 32 15.7 & 1.5 &     13 &    12.9 &  3.0 &      0 &     0.1 &  0.3 &     10 &    12.8 &  3.0 &      4 &    20.4 &  6.8 &  0.93 &  0.26 &  0.06 &  0.20 \\\\ 426 &     &XMMU J105145.3+573307 & 10 51 45.36 & 57 33  7.2 & 1.3 &     22 &     8.9 &  2.8 &     23 &     4.0 &  0.7 &      1 &     4.8 &  2.8 &      0 &     7.9 &  6.6 & -0.68 &  0.16 &  0.07 &  0.50 \\\\ 372 &     &XMMU J105146.2+573110 & 10 51 46.25 & 57 31 11.0 & 0.9 &     35 &    20.3 &  3.1 &      0 &     0.6 &  0.5 &     29 &    19.7 &  3.1 &      7 &    21.8 &  6.0 &  0.69 &  0.22 & -0.12 &  0.16 \\\\ 351 &     &XMMU J105146.5+573036 & 10 51 46.57 & 57 30 36.6 & 0.4 &    462 &    38.4 &  3.2 &    430 &    18.0 &  0.9 &     32 &    20.4 &  3.1 &      2 &    13.3 &  5.5 & -0.70 &  0.04 & -0.37 &  0.19 \\\\ 16 &     &XMMU J105146.9+571715 & 10 51 46.97 & 57 17 15.0 & 0.5 &    118 &    50.8 &  6.9 &     98 &    16.2 &  1.5 &     17 &    34.6 &  6.7 &      4 &    51.9 & 16.4 & -0.50 &  0.08 &  0.03 &  0.19 \\\\ 1572 &     &XMMU J105147.4+571840 & 10 51 47.50 & 57 18 40.6 & 1.3 &     14 &    24.8 &  5.2 &      2 &     2.3 &  0.9 &     11 &    22.6 &  5.1 &      1 &    26.1 & 12.8 &  0.22 &  0.22 & -0.10 &  0.27 \\\\ 565 &     &XMMU J105148.6+573932 & 10 51 48.63 & 57 39 32.9 & 0.5 &    152 &    59.2 &  6.3 &    112 &    16.4 &  1.4 &     34 &    42.8 &  6.1 &      6 &    58.4 & 16.2 & -0.42 &  0.07 & -0.02 &  0.16 \\\\ 407 &  12 &XMMU J105148.8+573248 & 10 51 48.85 & 57 32 48.4 & 0.1 &   5104 &   375.6 &  8.3 &   1845 &    48.9 &  1.3 &   2634 &   326.7 &  8.2 &    626 &   282.3 & 12.2 &  0.03 &  0.02 & -0.24 &  0.02 \\\\ 391 &     &XMMU J105149.1+573153 & 10 51 49.13 & 57 31 53.7 & 2.1 &     10 &     3.3 &  1.6 &     11 &     2.5 &  0.6 &      0 &     0.8 &  1.5 &      0 &     7.2 &  5.4 & -0.90 &  0.18 &  0.72 &  0.48 \\\\ 1396 &     &XMMU J105149.1+572858 & 10 51 49.16 & 57 28 58.9 & 1.0 &     38 &     8.9 &  2.4 &     39 &     4.3 &  0.6 &      1 &     4.6 &  2.4 &      0 &     0.0 &  4.5 & -0.71 &  0.13 & -1.00 &  1.37 \\\\ 186 &     &XMMU J105150.0+572521 & 10 51 50.10 & 57 25 21.6 & 0.2 &   2842 &   232.1 &  6.7 &   1462 &    40.1 &  1.2 &   1208 &   192.0 &  6.6 &    173 &   123.5 &  9.1 & -0.14 &  0.02 & -0.38 &  0.03 \\\\ 395 &     &XMMU J105151.4+573219 & 10 51 51.48 & 57 32 19.9 & 1.4 &     59 &    12.8 &  2.7 &     58 &     6.6 &  0.7 &      2 &     6.3 &  2.6 &      0 &     5.5 &  4.7 & -0.74 &  0.10 & -0.24 &  0.45 \\\\ 555 &     &XMMU J105152.1+573906 & 10 51 52.16 & 57 39  6.2 & 0.3 &    813 &   127.4 &  7.2 &    680 &    43.3 &  1.8 &    131 &    84.0 &  7.0 &      5 &    43.0 & 13.5 & -0.53 &  0.03 & -0.47 &  0.13 \\\\ 1566 &     &XMMU J105152.2+574104 & 10 51 52.28 & 57 41  4.5 & 1.6 &     20 &    17.5 &  5.8 &     17 &     7.0 &  1.3 &      1 &    10.6 &  5.7 &      3 &    46.8 & 17.8 & -0.61 &  0.18 &  0.52 &  0.24 \\\\ 492 &     &XMMU J105153.7+573555 & 10 51 53.73 & 57 35 55.6 & 0.6 &    157 &    59.4 &  5.2 &     11 &     3.2 &  0.7 &    105 &    56.3 &  5.1 &     41 &    82.8 & 10.7 &  0.48 &  0.09 &  0.02 &  0.08 \\\\ 57 &     &XMMU J105154.1+571924 & 10 51 54.16 & 57 19 24.7 & 1.0 &     53 &    22.2 &  4.2 &     43 &     7.1 &  0.9 &      7 &    15.1 &  4.1 &      4 &    32.4 & 10.5 & -0.50 &  0.11 &  0.20 &  0.20 \\\\ 456 &   9 &XMMU J105154.5+573437 & 10 51 54.50 & 57 34 37.6 & 0.1 &   5850 &   254.9 &  6.7 &   5102 &   104.5 &  1.9 &    701 &   150.4 &  6.4 &     48 &    72.3 &  8.9 & -0.63 &  0.01 & -0.49 &  0.05 \\\\ 255 &     &XMMU J105154.5+572751 & 10 51 54.54 & 57 27 51.7 & 0.9 &     40 &    16.9 &  2.7 &     23 &     3.2 &  0.5 &     15 &    13.7 &  2.7 &      2 &     7.5 &  4.6 & -0.20 &  0.12 & -0.44 &  0.26 \\\\ 156 &     &XMMU J105154.7+572407 & 10 51 54.71 & 57 24  7.4 & 0.4 &    584 &    59.4 &  3.9 &    487 &    19.7 &  0.9 &     94 &    39.7 &  3.8 &      5 &    21.3 &  6.4 & -0.52 &  0.04 & -0.45 &  0.13 \\\\ 312 &     &XMMU J105155.2+572932 & 10 51 55.27 & 57 29 32.6 & 0.5 &    178 &    25.9 &  3.0 &    157 &    10.1 &  0.8 &     20 &    15.8 &  2.9 &      2 &    13.3 &  5.7 & -0.60 &  0.06 & -0.26 &  0.22 \\\\ 180 &     &XMMU J105155.3+572458 & 10 51 55.38 & 57 24 58.2 & 0.4 &    222 &    31.2 &  3.3 &    196 &    11.5 &  0.8 &     27 &    19.7 &  3.2 &      0 &     8.7 &  5.8 & -0.57 &  0.06 & -0.52 &  0.25 \\\\ 359 &     &XMMU J105155.3+573043 & 10 51 55.39 & 57 30 43.4 & 0.7 &    198 &    31.5 &  3.0 &    154 &     9.3 &  0.7 &     43 &    22.2 &  3.0 &      3 &    14.1 &  5.3 & -0.45 &  0.06 & -0.38 &  0.17 \\\\ 274 &     &XMMU J105155.8+572812 & 10 51 55.82 & 57 28 12.0 & 1.5 &     20 &    11.5 &  2.5 &      7 &     1.9 &  0.5 &      9 &     9.6 &  2.4 &      5 &    16.6 &  5.0 & -0.10 &  0.18 &  0.10 &  0.19 \\\\ 436 &     &XMMU J105156.4+573331 & 10 51 56.43 & 57 33 31.3 & 1.6 &     21 &    17.2 &  3.1 &      0 &     0.8 &  0.5 &     19 &    16.4 &  3.0 &      2 &    14.8 &  5.9 &  0.53 &  0.23 & -0.22 &  0.21 \\\\ 541 &     &XMMU J105156.4+573802 & 10 51 56.44 & 57 38  2.2 & 1.3 &     15 &    13.9 &  3.9 &     12 &     3.7 &  0.8 &      3 &    10.2 &  3.8 &      0 &    16.0 &  9.8 & -0.39 &  0.18 &  0.05 &  0.36 \\\\ 61 &     &XMMU J105157.6+571940 & 10 51 57.66 & 57 19 40.7 & 0.5 &    112 &    53.9 &  5.2 &     25 &     5.1 &  0.8 &     75 &    48.8 &  5.1 &     13 &    50.6 & 10.6 &  0.21 &  0.09 & -0.16 &  0.11 \\\\ 256 &     &XMMU J105158.7+572758 & 10 51 58.79 & 57 27 59.0 & 1.3 &     28 &    16.9 &  2.7 &      6 &     1.7 &  0.5 &     21 &    15.2 &  2.7 &      2 &    10.7 &  4.8 &  0.16 &  0.16 & -0.34 &  0.22 \\\\ 405 &     &XMMU J105158.8+573229 & 10 51 58.80 & 57 32 30.0 & 0.7 &    151 &    22.2 &  3.0 &    139 &     9.8 &  0.8 &     11 &    12.4 &  2.9 &      2 &    14.5 &  5.7 & -0.67 &  0.07 & -0.10 &  0.23 \\\\ 163 &  33 &XMMU J105200.1+572423 & 10 52  0.11 & 57 24 23.5 & 0.4 &    523 &    60.1 &  3.9 &    397 &    17.5 &  0.9 &    113 &    42.6 &  3.8 &     13 &    30.6 &  6.5 & -0.44 &  0.04 & -0.33 &  0.10 \\\\ 31 &     &XMMU J105200.4+571806 & 10 52  0.41 & 57 18  6.2 & 0.5 &    258 &   103.2 &  6.9 &     71 &    10.7 &  1.1 &    173 &    92.4 &  6.9 &     14 &    70.5 & 14.0 &  0.15 &  0.06 & -0.30 &  0.10 \\\\ 439 &     &XMMU J105201.5+573337 & 10 52  1.58 & 57 33 37.4 & 1.1 &     15 &    14.1 &  3.0 &      3 &     1.5 &  0.5 &     11 &    12.6 &  3.0 &      1 &    12.1 &  6.3 &  0.16 &  0.21 & -0.19 &  0.27 \\\\ 77 &     &XMMU J105201.6+572017 & 10 52  1.68 & 57 20 17.6 & 0.6 &     96 &    26.0 &  3.8 &     85 &     8.9 &  0.9 &     13 &    17.1 &  3.7 &      0 &     5.8 &  6.9 & -0.53 &  0.08 & -0.61 &  0.38 \\\\ 512 &     &XMMU J105202.2+573647 & 10 52  2.22 & 57 36 47.8 & 1.3 &     41 &    10.3 &  3.1 &     43 &     6.1 &  0.8 &      0 &     4.3 &  3.0 &      0 &     8.7 &  7.5 & -0.80 &  0.13 &  0.18 &  0.54 \\\\ 2278 &     &XMMU J105203.8+572708 & 10 52  3.86 & 57 27  8.7 & 1.2 &     19 &    12.6 &  2.5 &      0 &     0.3 &  0.4 &     15 &    12.3 &  2.5 &      5 &    17.4 &  4.9 &  0.73 &  0.28 & -0.00 &  0.17 \\\\ 471 &     &XMMU J105203.8+573500 & 10 52  3.87 & 57 35  0.1 & 1.2 &     22 &    12.8 &  3.1 &     11 &     2.7 &  0.6 &      6 &    10.1 &  3.0 &      5 &    23.3 &  7.0 & -0.26 &  0.17 &  0.24 &  0.20 \\\\ 142 &     &XMMU J105203.9+572338 & 10 52  4.00 & 57 23 38.9 & 0.5 &    277 &    37.4 &  3.4 &    234 &    12.4 &  0.8 &     45 &    25.1 &  3.3 &      0 &     6.9 &  5.6 & -0.51 &  0.05 & -0.67 &  0.23 \\\\ 2045 &     &XMMU J105204.3+574108 & 10 52  4.35 & 57 41  8.7 & 0.3 &    206 &   210.2 & 11.2 &    190 &    81.6 &  3.1 &     16 &   128.6 & 10.7 &      1 &    51.0 & 18.3 & -0.60 &  0.03 & -0.56 &  0.13 \\\\ 30 &     &XMMU J105204.8+571806 & 10 52  4.87 & 57 18  6.7 & 0.7 &     92 &    38.2 &  5.2 &     64 &     9.7 &  1.0 &     21 &    28.5 &  5.0 &      7 &    48.5 & 12.7 & -0.36 &  0.09 &  0.09 &  0.16 \\\\ 568 &     &XMMU J105205.2+573946 & 10 52  5.23 & 57 39 46.9 & 1.0 &     51 &    40.9 &  5.5 &     20 &     5.4 &  1.0 &     30 &    35.4 &  5.4 &      1 &    25.3 & 12.1 &  0.01 &  0.12 & -0.33 &  0.22 \\\\ 576 &     &XMMU J105205.4+574014 & 10 52  5.42 & 57 40 14.5 & 1.0 &     79 &    33.1 &  5.6 &     69 &    11.9 &  1.2 &      8 &    21.3 &  5.4 &      3 &    36.4 & 13.4 & -0.56 &  0.09 &  0.09 &  0.22 \\\\ 233 &     &XMMU J105205.4+572659 & 10 52  5.46 & 57 26 59.2 & 1.3 &     16 &    11.7 &  2.5 &      0 &     0.8 &  0.4 &     11 &    10.9 &  2.5 &      5 &    18.2 &  5.3 &  0.34 &  0.25 &  0.08 &  0.18 \\\\ 584 &     &XMMU J105206.1+574124 & 10 52  6.19 & 57 41 24.2 & 0.2 &    434 &   259.4 & 11.3 &    379 &   112.8 &  3.3 &     55 &   146.6 & 10.8 &      3 &    87.8 & 19.5 & -0.66 &  0.02 & -0.41 &  0.10 \\\\ 411 &  53 &XMMU J105206.2+571525 & 10 52  6.24 & 57 15 25.9 & 0.5 &    958 &   500.4 & 23.7 &    496 &    86.2 &  4.3 &    411 &   414.2 & 23.3 &     51 &   350.2 & 43.4 & -0.14 &  0.04 & -0.25 &  0.06 \\\\ 126 &     &XMMU J105206.4+572301 & 10 52  6.46 & 57 23  1.2 & 1.0 &     28 &    17.7 &  3.0 &      3 &     1.5 &  0.5 &     20 &    16.3 &  2.9 &      5 &    19.6 &  6.2 &  0.27 &  0.19 & -0.08 &  0.18 \\\\ 2046 &     &XMMU J105206.5+574109 & 10 52  6.51 & 57 41  9.5 & 0.5 &    206 &   106.5 &  9.0 &    190 &    42.6 &  2.4 &     17 &    63.9 &  8.7 &      1 &    47.0 & 17.5 & -0.62 &  0.05 & -0.32 &  0.18 \\\\ 110 & 825 &XMMU J105206.5+572213 & 10 52  6.51 & 57 22 14.0 & 0.4 &    275 &    42.4 &  3.7 &    224 &    12.9 &  0.8 &     53 &    29.5 &  3.6 &      0 &    10.5 &  6.4 & -0.47 &  0.05 & -0.60 &  0.20 \\\\ 306 &     &XMMU J105206.8+572924 & 10 52  6.81 & 57 29 24.3 & 0.3 &    835 &    89.2 &  4.0 &    327 &    12.2 &  0.7 &    451 &    76.9 &  3.9 &     58 &    47.8 &  5.5 & -0.00 &  0.04 & -0.39 &  0.05 \\\\ 522 &     &XMMU J105207.0+573731 & 10 52  7.06 & 57 37 31.7 & 1.0 &     35 &    16.7 &  3.6 &     29 &     5.3 &  0.8 &      5 &    11.4 &  3.5 &      1 &    17.5 &  8.3 & -0.49 &  0.13 &  0.04 &  0.28 \\\\ 548 & 832 &XMMU J105207.4+573838 & 10 52  7.44 & 57 38 38.3 & 0.4 &    374 &    57.3 &  5.0 &    339 &    22.2 &  1.2 &     37 &    35.1 &  4.9 &      0 &     8.1 &  8.8 & -0.60 &  0.05 & -0.72 &  0.27 \\\\ 51 &     &XMMU J105207.6+571905 & 10 52  7.70 & 57 19  5.9 & 1.4 &     15 &    16.7 &  4.1 &      0 &     1.0 &  0.7 &      8 &    15.7 &  4.0 &      7 &    39.6 & 10.7 &  0.45 &  0.31 &  0.28 &  0.17 \\\\ 557 &     &XMMU J105207.7+573904 & 10 52  7.79 & 57 39  4.6 & 0.3 &    485 &    76.6 &  5.4 &    406 &    25.7 &  1.3 &     78 &    50.8 &  5.3 &      2 &    26.0 & 10.3 & -0.52 &  0.04 & -0.47 &  0.16 \\\\ 451 &     &XMMU J105207.9+573414 & 10 52  7.92 & 57 34 14.5 & 0.3 &    702 &    77.3 &  4.4 &    517 &    20.7 &  1.0 &    171 &    56.6 &  4.3 &     15 &    34.3 &  6.8 & -0.40 &  0.04 & -0.40 &  0.09 \\\\ 149 &     &XMMU J105208.2+572348 & 10 52  8.24 & 57 23 48.2 & 0.9 &     56 &    16.2 &  2.8 &     46 &     4.7 &  0.6 &     10 &    11.5 &  2.7 &      0 &     8.9 &  5.4 & -0.44 &  0.11 & -0.29 &  0.30 \\\\ 2081 &     &XMMU J105208.3+572044 & 10 52  8.35 & 57 20 44.4 & 0.8 &     52 &    42.5 &  3.9 &     18 &     3.8 &  0.7 &     29 &    24.5 &  3.8 &      4 &    25.3 &  8.3 &  0.01 &  0.12 & -0.16 &  0.18 \\\\ 336 &     &XMMU J105208.3+573014 & 10 52  8.39 & 57 30 14.2 & 0.4 &    247 &    40.1 &  3.2 &      2 &     1.0 &  0.4 &    144 &    39.2 &  3.2 &    103 &    66.2 &  6.0 &  0.73 &  0.10 &  0.09 &  0.06 \\\\ 88 &  39 &XMMU J105209.8+572104 & 10 52  9.81 & 57 21  4.4 & 0.3 &    884 &    83.6 &  4.8 &    755 &    30.3 &  1.2 &    124 &    53.3 &  4.6 &      6 &    28.9 &  8.2 & -0.56 &  0.03 & -0.45 &  0.12 \\\\ 534 &     &XMMU J105209.9+573743 & 10 52  9.99 & 57 37 43.3 & 0.8 &     67 &    28.9 &  4.3 &     51 &     7.4 &  0.9 &     17 &    21.5 &  4.2 &      0 &    11.6 &  8.4 & -0.37 &  0.10 & -0.45 &  0.30 \\\\ 393 &     &XMMU J105210.9+573202 & 10 52 10.96 & 57 32  2.6 & 0.4 &    322 &    60.9 &  3.8 &     91 &     6.6 &  0.6 &    210 &    54.3 &  3.8 &     22 &    31.1 &  5.4 &  0.13 &  0.06 & -0.42 &  0.08 \\\\ 136 &     &XMMU J105212.2+572322 & 10 52 12.23 & 57 23 22.3 & 1.3 &     12 &     7.0 &  2.4 &     11 &     2.2 &  0.5 &      1 &     4.7 &  2.4 &      0 &     5.6 &  5.1 & -0.50 &  0.21 & -0.09 &  0.51 \\\\ \\hline \\end{tabular}} \\vfill \\end{minipage} \\end{sidewaystable*} \\begin{sidewaystable*} \\begin{minipage}[t][180mm]{\\textwidth} \\bigskip \\centering {\\scriptsize \\begin{tabular}{rrcllcrrrrrrrrrrrrrrrr} \\hline \\multicolumn{2}{c|}{SRC \\#}&\\multicolumn{1}{c|}{}&\\multicolumn{2}{c}{RA~~~(J2000)~~~DEC}&err&\\multicolumn{3}{|c|}{0.5--10.0~keV}&\\multicolumn{3}{|c|}{0.5--2.0~keV}&\\multicolumn{3}{|c|}{2.0--10.0~keV}&\\multicolumn{3}{|c|}{5.0--10.0~keV}&&&& \\\\ XMM&\\multicolumn{1}{r|}{R}&\\multicolumn{1}{c|}{IAU name}&hh mm ss&$^{\\circ}$~~~~$^\\prime$~~~$^{\\prime\\prime}$&$^{\\prime\\prime}$&\\multicolumn{1}{|r}{$\\mathcal{L}$}&flux&err&\\multicolumn{1}{|r}{$\\mathcal{L}$}&flux&err& \\multicolumn{1}{|r}{$\\mathcal{L}$}&flux&err&\\multicolumn{1}{|r}{$\\mathcal{L}$}&flux&\\multicolumn{1}{r|}{err}&HR1&err&HR2&err \\\\ \\hline 1591 &     &XMMU J105212.6+572452 & 10 52 12.70 & 57 24 52.6 & 1.3 &     15 &     4.3 &  1.8 &     17 &     2.7 &  0.5 &      0 &     1.6 &  1.8 &      0 &     2.8 &  3.7 & -0.82 &  0.18 &  0.10 &  0.84 \\\\ 307 &     &XMMU J105212.9+572919 & 10 52 12.99 & 57 29 19.7 & 0.9 &     64 &    21.5 &  2.6 &      0 &     0.2 &  0.3 &     54 &    21.4 &  2.6 &     12 &    20.3 &  4.5 &  0.91 &  0.14 & -0.20 &  0.12 \\\\ 400 &  13 &XMMU J105213.3+573221 & 10 52 13.35 & 57 32 21.2 & 0.2 &   3273 &   134.4 &  4.5 &   2818 &    52.4 &  1.2 &    423 &    82.0 &  4.3 &     34 &    39.7 &  5.7 & -0.60 &  0.02 & -0.49 &  0.06 \\\\ 201 &     &XMMU J105213.6+572559 & 10 52 13.63 & 57 25 59.1 & 0.6 &    141 &    24.7 &  2.6 &    103 &     6.4 &  0.6 &     39 &    18.3 &  2.5 &      0 &     6.5 &  4.2 & -0.37 &  0.07 & -0.60 &  0.21 \\\\ 386 &     &XMMU J105214.0+573140 & 10 52 14.06 & 57 31 40.6 & 0.9 &     46 &    19.0 &  2.7 &      8 &     1.7 &  0.4 &     32 &    17.3 &  2.6 &      6 &    16.1 &  4.3 &  0.23 &  0.14 & -0.21 &  0.15 \\\\ 588 &     &XMMU J105216.2+574108 & 10 52 16.27 & 57 41  8.6 & 0.4 &    263 &    79.6 &  7.3 &    220 &    25.3 &  1.7 &     44 &    54.3 &  7.1 &      1 &    27.4 & 14.7 & -0.49 &  0.06 & -0.47 &  0.21 \\\\ 566 &     &XMMU J105216.5+573946 & 10 52 16.50 & 57 39 46.5 & 1.3 &     27 &    24.3 &  4.8 &     15 &     4.7 &  0.9 &     10 &    19.6 &  4.7 &      2 &    29.5 & 11.4 & -0.20 &  0.15 &  0.03 &  0.23 \\\\ 526 &     &XMMU J105216.6+573721 & 10 52 16.68 & 57 37 21.3 & 1.5 &     11 &     6.2 &  2.8 &     12 &     3.1 &  0.7 &      0 &     3.0 &  2.7 &      0 &     1.7 &  4.3 & -0.73 &  0.21 & -0.43 &  1.10 \\\\ 72 &  84 &XMMU J105216.9+572018 & 10 52 16.95 & 57 20 18.0 & 0.2 &   4957 &   407.6 &  9.3 &   2383 &    65.9 &  1.6 &   2215 &   341.8 &  9.2 &    360 &   263.9 & 14.2 & -0.10 &  0.02 & -0.29 &  0.03 \\\\ 549 &     &XMMU J105216.9+573843 & 10 52 16.99 & 57 38 43.6 & 1.1 &     42 &    35.7 &  4.9 &      0 &     0.7 &  0.7 &     39 &    35.0 &  4.8 &      4 &    31.4 & 10.1 &  0.77 &  0.20 & -0.22 &  0.17 \\\\ 1595 &     &XMMU J105217.4+573637 & 10 52 17.42 & 57 36 37.8 & 1.6 &     10 &     8.7 &  2.9 &      0 &     0.8 &  0.5 &      3 &     7.9 &  2.9 &      6 &    26.2 &  7.4 &  0.21 &  0.35 &  0.40 &  0.19 \\\\ 97 &     &XMMU J105217.7+572127 & 10 52 17.72 & 57 21 27.1 & 1.3 &     13 &    11.7 &  3.1 &      2 &     1.4 &  0.6 &      6 &    10.4 &  3.0 &      5 &    22.6 &  7.0 &  0.09 &  0.25 &  0.21 &  0.20 \\\\ 196 &     &XMMU J105218.8+572539 & 10 52 18.80 & 57 25 39.7 & 0.8 &     55 &    20.5 &  2.6 &      9 &     1.9 &  0.4 &     39 &    18.7 &  2.6 &      7 &    17.2 &  4.6 &  0.23 &  0.13 & -0.21 &  0.14 \\\\ 91 &     &XMMU J105218.8+572056 & 10 52 18.83 & 57 20 56.4 & 1.4 &     21 &    18.8 &  3.5 &      4 &     1.7 &  0.6 &     16 &    17.1 &  3.4 &      2 &    16.7 &  7.0 &  0.23 &  0.19 & -0.18 &  0.22 \\\\ 413 &     &XMMU J105219.4+573253 & 10 52 19.47 & 57 32 53.9 & 0.8 &     76 &    12.7 &  2.3 &     73 &     5.7 &  0.6 &      4 &     7.0 &  2.3 &      0 &     6.8 &  4.4 & -0.67 &  0.09 & -0.19 &  0.35 \\\\ 554 &     &XMMU J105219.4+573903 & 10 52 19.47 & 57 39  3.4 & 0.7 &    198 &    45.2 &  4.7 &    170 &    15.7 &  1.1 &     30 &    29.5 &  4.5 &      0 &    10.9 &  9.5 & -0.54 &  0.06 & -0.59 &  0.29 \\\\ 107 &     &XMMU J105219.6+572213 & 10 52 19.70 & 57 22 13.5 & 0.3 &    698 &    69.8 &  4.2 &    539 &    22.2 &  1.0 &    130 &    47.6 &  4.1 &     30 &    51.1 &  7.7 & -0.49 &  0.04 & -0.14 &  0.08 \\\\ 125 & 607 &XMMU J105220.1+572306 & 10 52 20.17 & 57 23  6.8 & 0.4 &    483 &    50.1 &  3.5 &    406 &    17.2 &  0.9 &     77 &    32.9 &  3.4 &      2 &    12.6 &  5.6 & -0.54 &  0.04 & -0.57 &  0.15 \\\\ 1499 &     &XMMU J105220.5+572942 & 10 52 20.53 & 57 29 42.5 & 1.2 &     18 &     8.7 &  2.0 &     14 &     2.1 &  0.4 &      5 &     6.5 &  2.0 &      0 &     0.0 &  1.1 & -0.34 &  0.16 & -1.00 &  0.23 \\\\ 525 & 802 &XMMU J105222.4+573735 & 10 52 22.48 & 57 37 35.0 & 0.3 &   1016 &    47.0 &  3.5 &   1019 &    39.1 &  1.4 &      2 &     8.0 &  3.2 &      0 &     0.9 &  4.1 & -0.94 &  0.02 & -0.85 &  0.62 \\\\ 1526 &     &XMMU J105222.7+573303 & 10 52 22.71 & 57 33  3.7 & 1.6 &     11 &     8.7 &  2.3 &      0 &     0.2 &  0.3 &      8 &     8.4 &  2.3 &      5 &    15.4 &  4.8 &  0.71 &  0.34 &  0.13 &  0.20 \\\\ 1233 &     &XMMU J105223.2+572946 & 10 52 23.24 & 57 29 46.8 & 1.1 &     27 &     8.2 &  2.0 &     24 &     3.0 &  0.5 &      3 &     5.2 &  1.9 &      0 &     6.3 &  3.7 & -0.57 &  0.14 & -0.08 &  0.35 \\\\ 591 &     &XMMU J105223.3+574121 & 10 52 23.36 & 57 41 21.6 & 0.4 &    847 &   112.2 &  7.3 &    778 &    50.1 &  2.0 &     64 &    62.1 &  7.0 &      7 &    56.3 & 15.0 & -0.67 &  0.03 & -0.22 &  0.14 \\\\ 188 &     &XMMU J105223.9+572531 & 10 52 23.93 & 57 25 31.4 & 0.3 &    394 &    34.3 &  2.8 &    351 &    13.5 &  0.7 &     46 &    20.8 &  2.7 &      0 &     4.3 &  4.1 & -0.61 &  0.04 & -0.74 &  0.21 \\\\ 24 &     &XMMU J105224.3+571739 & 10 52 24.36 & 57 17 39.8 & 0.9 &     62 &    28.5 &  4.8 &     54 &     8.8 &  1.0 &     10 &    19.7 &  4.7 &      0 &     8.2 &  9.4 & -0.48 &  0.10 & -0.55 &  0.41 \\\\ 321 &  23 &XMMU J105224.6+573008 & 10 52 24.69 & 57 30  9.0 & 0.2 &   2075 &    95.4 &  3.7 &   1696 &    33.2 &  0.9 &    334 &    62.1 &  3.6 &     46 &    40.3 &  5.1 & -0.54 &  0.02 & -0.37 &  0.06 \\\\ 105 &     &XMMU J105225.0+572155 & 10 52 25.01 & 57 21 55.5 & 1.2 &     18 &    12.6 &  3.0 &     13 &     2.8 &  0.6 &      6 &     9.8 &  2.9 &      0 &     7.1 &  5.8 & -0.29 &  0.17 & -0.32 &  0.39 \\\\ 120 &     &XMMU J105225.2+572305 & 10 52 25.22 & 57 23  5.3 & 0.2 &    344 &   123.0 &  4.6 &    287 &    51.5 &  1.3 &     50 &    71.5 &  4.4 &      8 &    45.6 &  6.8 & -0.64 &  0.02 & -0.38 &  0.07 \\\\ 999 & 870 &XMMU J105225.3+572247 & 10 52 25.34 & 57 22 47.3 & 0.4 &    359 &    44.0 &  3.6 &    305 &    17.1 &  0.9 &     47 &    26.9 &  3.5 &      8 &    23.5 &  6.0 & -0.60 &  0.04 & -0.24 &  0.13 \\\\ 199 &     &XMMU J105225.5+572550 & 10 52 25.50 & 57 25 50.7 & 0.2 &   1010 &   102.4 &  4.2 &    353 &    13.1 &  0.7 &    573 &    89.3 &  4.1 &     84 &    61.0 &  6.0 &  0.04 &  0.04 & -0.35 &  0.05 \\\\ 499 &     &XMMU J105225.5+573605 & 10 52 25.59 & 57 36  5.9 & 1.0 &     50 &    16.7 &  3.0 &     40 &     4.9 &  0.7 &      9 &    11.8 &  2.9 &      1 &    13.0 &  6.3 & -0.45 &  0.11 & -0.13 &  0.27 \\\\ 50 & 817 &XMMU J105225.8+571906 & 10 52 25.88 & 57 19  6.7 & 0.6 &    252 &    46.9 &  4.5 &    217 &    16.2 &  1.1 &     38 &    30.8 &  4.3 &      0 &    10.1 &  8.9 & -0.54 &  0.06 & -0.63 &  0.27 \\\\ 2141 &     &XMMU J105227.0+573320 & 10 52 27.00 & 57 33 20.4 & 1.4 &     13 &    13.0 &  2.6 &      4 &     1.6 &  0.5 &      9 &    11.4 &  2.6 &      0 &     9.2 &  4.7 &  0.04 &  0.19 & -0.27 &  0.26 \\\\ 427 &     &XMMU J105227.9+573330 & 10 52 28.00 & 57 33 30.7 & 0.4 &     13 &    31.4 &  3.0 &      4 &     9.8 &  0.7 &      9 &    21.6 &  2.9 &      0 &    16.2 &  4.7 & -0.48 &  0.06 & -0.31 &  0.14 \\\\ 364 &  18 &XMMU J105228.3+573104 & 10 52 28.34 & 57 31  4.3 & 0.3 &     69 &    44.7 &  2.8 &     67 &    17.8 &  0.8 &      4 &    26.8 &  2.7 &      0 &    24.4 &  4.5 & -0.62 &  0.03 & -0.22 &  0.10 \\\\ 577 &     &XMMU J105228.7+574014 & 10 52 28.72 & 57 40 14.8 & 1.4 &     12 &    18.2 &  4.4 &      4 &     2.4 &  0.8 &      7 &    15.8 &  4.3 &      1 &    20.6 & 11.0 &  0.01 &  0.21 & -0.04 &  0.30 \\\\ 2254 &     &XMMU J105229.5+574332 & 10 52 29.51 & 57 43 32.6 & 1.0 &     60 &   106.9 & 25.6 &     59 &    46.9 &  6.1 &      3 &    60.0 & 24.8 &      0 &    58.3 & 54.7 & -0.66 &  0.12 & -0.19 &  0.49 \\\\ 482 &     &XMMU J105229.7+573529 & 10 52 29.73 & 57 35 29.9 & 1.2 &     30 &    16.2 &  3.0 &     19 &     3.1 &  0.6 &     11 &    13.1 &  3.0 &      0 &     6.7 &  5.6 & -0.20 &  0.14 & -0.47 &  0.34 \\\\ 317 &     &XMMU J105229.8+572947 & 10 52 29.88 & 57 29 47.2 & 0.9 &     40 &    13.8 &  2.3 &     28 &     3.2 &  0.5 &     14 &    10.7 &  2.3 &      0 &     0.0 &  2.8 & -0.31 &  0.12 & -1.00 &  0.36 \\\\ 2086 &     &XMMU J105230.0+573103 & 10 52 30.00 & 57 31  3.4 & 0.7 &     69 &    14.2 &  2.3 &     67 &     6.4 &  0.6 &      4 &     7.9 &  2.2 &      0 &     0.0 &  2.7 & -0.67 &  0.08 & -1.00 &  0.49 \\\\ 553 &   2 &XMMU J105230.0+573914 & 10 52 30.10 & 57 39 14.1 & 0.1 &  11794 &   487.8 &  9.6 &  10324 &   211.3 &  2.9 &   1312 &   276.6 &  9.2 &    159 &   187.3 & 14.1 & -0.66 &  0.01 & -0.35 &  0.04 \\\\ 340 &     &XMMU J105231.1+573016 & 10 52 31.12 & 57 30 16.2 & 0.7 &    108 &    25.9 &  2.8 &     13 &     2.2 &  0.5 &     63 &    23.7 &  2.7 &     33 &    34.3 &  5.0 &  0.26 &  0.11 &  0.01 &  0.09 \\\\ 2393 &     &XMMU J105231.2+571547 & 10 52 31.29 & 57 15 47.2 & 2.0 &     11 &    22.4 &  5.7 &      6 &     3.7 &  1.0 &      6 &    18.7 &  5.6 &      0 &     9.5 & 11.8 & -0.11 &  0.20 & -0.47 &  0.50 \\\\ 435 &     &XMMU J105231.6+573341 & 10 52 31.62 & 57 33 41.5 & 0.8 &     82 &    15.3 &  2.5 &     72 &     5.5 &  0.6 &      9 &     9.8 &  2.4 &      2 &    10.8 &  4.8 & -0.56 &  0.09 & -0.12 &  0.25 \\\\ 179 &     &XMMU J105231.8+572500 & 10 52 31.85 & 57 25  0.6 & 0.5 &    286 &    50.7 &  3.4 &     58 &     5.0 &  0.6 &    176 &    45.7 &  3.4 &     52 &    48.5 &  5.8 &  0.18 &  0.06 & -0.14 &  0.07 \\\\ 166 &     &XMMU J105232.1+572428 & 10 52 32.18 & 57 24 28.2 & 0.4 &    660 &    90.9 &  4.3 &     87 &     6.4 &  0.6 &    456 &    84.5 &  4.3 &    117 &    79.3 &  6.8 &  0.35 &  0.05 & -0.20 &  0.05 \\\\ 486 &     &XMMU J105232.8+573543 & 10 52 32.81 & 57 35 43.9 & 1.1 &     54 &    24.5 &  3.3 &     19 &     3.4 &  0.6 &     31 &    21.1 &  3.2 &      4 &    19.1 &  6.2 & -0.00 &  0.12 & -0.22 &  0.17 \\\\ 192 &     &XMMU J105232.9+572537 & 10 52 32.93 & 57 25 37.9 & 2.0 &     13 &     9.7 &  2.3 &      5 &     1.4 &  0.4 &      7 &     8.3 &  2.2 &      2 &     9.0 &  4.0 & -0.02 &  0.20 & -0.14 &  0.26 \\\\ 26 &     &XMMU J105233.0+571749 & 10 52 33.07 & 57 17 49.3 & 0.7 &    200 &    51.9 &  5.4 &    160 &    15.5 &  1.2 &     35 &    36.3 &  5.2 &      5 &    39.8 & 11.8 & -0.46 &  0.06 & -0.13 &  0.16 \\\\ 260 &     &XMMU J105233.4+572800 & 10 52 33.49 & 57 28  0.2 & 0.8 &     74 &    11.6 &  2.0 &     69 &     4.8 &  0.5 &      6 &     6.8 &  2.0 &      0 &     3.6 &  3.7 & -0.64 &  0.09 & -0.46 &  0.42 \\\\ 1443 &     &XMMU J105233.8+571644 & 10 52 33.87 & 57 16 44.7 & 1.3 &     18 &    10.9 &  4.3 &     20 &     5.6 &  1.0 &      0 &     5.3 &  4.2 &      0 &     5.1 &  8.7 & -0.74 &  0.19 & -0.19 &  0.90 \\\\ 291 &     &XMMU J105234.1+572852 & 10 52 34.10 & 57 28 52.2 & 1.1 &     39 &     9.2 &  2.0 &     36 &     3.3 &  0.5 &      4 &     5.9 &  1.9 &      0 &     0.0 &  0.8 & -0.56 &  0.12 & -1.00 &  0.20 \\\\ 104 &     &XMMU J105234.7+572144 & 10 52 34.75 & 57 21 44.3 & 1.1 &     25 &    12.7 &  2.9 &     17 &     3.1 &  0.6 &      6 &     9.6 &  2.9 &      2 &    14.3 &  5.9 & -0.34 &  0.16 &  0.03 &  0.26 \\\\ 242 &     &XMMU J105235.3+572725 & 10 52 35.38 & 57 27 25.5 & 0.9 &     43 &    13.3 &  2.2 &     29 &     3.0 &  0.5 &     14 &    10.3 &  2.1 &      0 &     5.0 &  3.6 & -0.30 &  0.12 & -0.49 &  0.28 \\\\ 2097 &     &XMMU J105235.4+573347 & 10 52 35.43 & 57 33 47.4 & 1.4 &     16 &     8.6 &  2.4 &     15 &     2.8 &  0.6 &      2 &     5.7 &  2.3 &      0 &     3.6 &  4.1 & -0.52 &  0.17 & -0.39 &  0.52 \\\\ 227 &     &XMMU J105235.5+572652 & 10 52 35.55 & 57 26 52.3 & 0.7 &    107 &    21.3 &  2.5 &     63 &     4.6 &  0.5 &     35 &    16.7 &  2.4 &     10 &    17.8 &  4.2 & -0.26 &  0.08 & -0.14 &  0.14 \\\\ 528 &     &XMMU J105236.2+573742 & 10 52 36.23 & 57 37 42.2 & 0.5 &    358 &    45.1 &  4.0 &    318 &    18.4 &  1.0 &     36 &    26.7 &  3.9 &      6 &    26.4 &  7.8 & -0.63 &  0.05 & -0.18 &  0.16 \\\\ 63 &     &XMMU J105236.3+571602 & 10 52 36.39 & 57 16  3.0 & 0.4 &    671 &   118.7 &  7.8 &    580 &    44.6 &  2.0 &     81 &    74.1 &  7.6 &     11 &    73.7 & 16.4 & -0.58 &  0.04 & -0.18 &  0.12 \\\\ 598 &     &XMMU J105236.5+574146 & 10 52 36.57 & 57 41 47.0 & 1.7 &     29 &    19.0 &  5.5 &     30 &     8.9 &  1.3 &      1 &    10.1 &  5.3 &      0 &     2.9 &  9.1 & -0.69 &  0.14 & -0.66 &  0.89 \\\\ \\hline \\end{tabular}} \\vfill \\end{minipage} \\end{sidewaystable*} \\begin{sidewaystable*} \\begin{minipage}[t][180mm]{\\textwidth} \\bigskip \\centering {\\scriptsize \\begin{tabular}{rrcllcrrrrrrrrrrrrrrrr} \\hline \\multicolumn{2}{c|}{SRC \\#}&\\multicolumn{1}{c|}{}&\\multicolumn{2}{c}{RA~~~(J2000)~~~DEC}&err&\\multicolumn{3}{|c|}{0.5--10.0~keV}&\\multicolumn{3}{|c|}{0.5--2.0~keV}&\\multicolumn{3}{|c|}{2.0--10.0~keV}&\\multicolumn{3}{|c|}{5.0--10.0~keV}&&&& \\\\ XMM&\\multicolumn{1}{r|}{R}&\\multicolumn{1}{c|}{IAU name}&hh mm ss&$^{\\circ}$~~~~$^\\prime$~~~$^{\\prime\\prime}$&$^{\\prime\\prime}$&\\multicolumn{1}{|r}{$\\mathcal{L}$}&flux&err&\\multicolumn{1}{|r}{$\\mathcal{L}$}&flux&err& \\multicolumn{1}{|r}{$\\mathcal{L}$}&flux&err&\\multicolumn{1}{|r}{$\\mathcal{L}$}&flux&\\multicolumn{1}{r|}{err}&HR1&err&HR2&err \\\\ \\hline 443 &     &XMMU J105237.1+573400 & 10 52 37.18 & 57 34  0.4 & 0.3 &     16 &    85.6 &  4.3 &     15 &    15.0 &  0.8 &      2 &    70.6 &  4.2 &      0 &    56.3 &  6.6 & -0.15 &  0.04 & -0.28 &  0.06 \\\\ 355 &     &XMMU J105237.3+573103 & 10 52 37.35 & 57 31  3.7 & 0.2 &   1561 &    88.2 &  3.7 &   1252 &    30.0 &  0.9 &    287 &    58.2 &  3.6 &     24 &    28.7 &  4.8 & -0.53 &  0.02 & -0.48 &  0.07 \\\\ 424 &     &XMMU J105238.1+573321 & 10 52 38.12 & 57 33 21.9 & 0.2 &   1859 &   171.2 &  5.5 &    415 &    16.4 &  0.8 &   1153 &   154.8 &  5.4 &    291 &   136.3 &  8.1 &  0.20 &  0.03 & -0.23 &  0.03 \\\\ 2297 &     &XMMU J105238.5+571644 & 10 52 38.51 & 57 16 44.4 & 1.5 &     22 &    28.9 &  5.5 &     10 &     4.2 &  1.0 &     12 &    24.7 &  5.4 &      1 &    22.4 & 12.6 & -0.04 &  0.16 & -0.22 &  0.29 \\\\ 1495 &     &XMMU J105239.4+572842 & 10 52 39.49 & 57 28 42.2 & 1.7 &     16 &     7.6 &  1.9 &     12 &     1.9 &  0.4 &      4 &     5.7 &  1.8 &      0 &     5.9 &  3.5 & -0.36 &  0.17 & -0.15 &  0.33 \\\\ 148 &  32 &XMMU J105239.7+572431 & 10 52 39.76 & 57 24 31.5 & 0.1 & 102046 &  1413.9 & 11.4 &  87917 &   633.0 &  3.6 &  12131 &   780.9 & 10.8 &   1999 &   434.1 & 12.0 & -0.67 &  0.00 & -0.44 &  0.01 \\\\ 610 &     &XMMU J105239.8+574310 & 10 52 39.89 & 57 43 10.4 & 1.0 &     89 &    70.9 & 12.7 &     88 &    34.1 &  3.5 &      4 &    36.8 & 12.2 &      0 &     0.0 & 11.8 & -0.71 &  0.09 & -1.00 &  0.45 \\\\ 468 &     &XMMU J105241.0+573448 & 10 52 41.08 & 57 34 48.2 & 1.0 &     43 &    21.7 &  3.1 &     13 &     2.6 &  0.5 &     29 &    19.1 &  3.0 &      2 &    11.7 &  5.1 &  0.08 &  0.13 & -0.40 &  0.19 \\\\ 2378 &     &XMMU J105241.4+573229 & 10 52 41.44 & 57 32 29.7 & 1.5 &     39 &     7.2 &  2.0 &     42 &     4.6 &  0.6 &      0 &     2.6 &  1.9 &      0 &     1.4 &  2.6 & -0.84 &  0.11 & -0.45 &  0.82 \\\\ 131 &     &XMMU J105241.6+572320 & 10 52 41.61 & 57 23 20.9 & 0.8 &     41 &    10.5 &  2.3 &     40 &     4.3 &  0.6 &      3 &     6.2 &  2.2 &      0 &     0.0 &  4.3 & -0.63 &  0.12 & -1.00 &  0.98 \\\\ 505 & 104 &XMMU J105241.6+573650 & 10 52 41.65 & 57 36 50.2 & 0.1 &   6575 &   392.6 &  8.9 &   4511 &   101.5 &  2.0 &   1802 &   291.1 &  8.7 &    263 &   194.6 & 12.1 & -0.38 &  0.02 & -0.36 &  0.03 \\\\ 350 &     &XMMU J105241.7+573038 & 10 52 41.78 & 57 30 38.1 & 0.4 &    456 &    44.9 &  3.0 &    344 &    12.3 &  0.7 &    109 &    32.6 &  2.9 &      4 &    12.1 &  3.9 & -0.41 &  0.04 & -0.59 &  0.11 \\\\ 521 &     &XMMU J105242.3+573717 & 10 52 42.33 & 57 37 17.9 & 0.7 &     73 &    39.4 &  4.6 &     22 &     4.7 &  0.8 &     46 &    34.7 &  4.6 &      5 &    29.0 &  8.7 &  0.08 &  0.11 & -0.26 &  0.15 \\\\ 385 &  14 &XMMU J105242.3+573158 & 10 52 42.38 & 57 31 58.3 & 0.1 &   5765 &   232.1 &  5.4 &   3967 &    60.8 &  1.2 &   1580 &   171.3 &  5.3 &    219 &    97.2 &  6.5 & -0.38 &  0.02 & -0.43 &  0.03 \\\\ 54 &     &XMMU J105242.4+571913 & 10 52 42.46 & 57 19 14.0 & 1.5 &     16 &     4.2 &  2.9 &     18 &     4.2 &  0.8 &      0 &     0.0 &  2.8 &      0 &     8.2 &  8.5 & -1.00 &  0.21 &  1.00 &  0.97 \\\\ 289 &     &XMMU J105242.7+572910 & 10 52 42.72 & 57 29 10.1 & 0.6 &    163 &    25.5 &  2.4 &    114 &     6.2 &  0.5 &     49 &    19.2 &  2.4 &      1 &     6.2 &  3.4 & -0.34 &  0.07 & -0.63 &  0.17 \\\\ 1461 &     &XMMU J105242.7+572847 & 10 52 42.76 & 57 28 47.6 & 1.5 &     11 &     2.1 &  1.5 &     14 &     2.1 &  0.4 &      0 &     0.0 &  1.4 &      0 &     0.0 &  0.8 & -1.00 &  0.21 &  -$^a$ &  -$^a$ \\\\ 5 &  52 &XMMU J105243.2+571545 & 10 52 43.23 & 57 15 45.3 & 0.1 &     96 &   278.5 &  9.9 &     36 &   121.1 &  2.8 &     60 &   157.4 &  9.5 &      1 &    93.7 & 17.9 & -0.66 &  0.02 & -0.41 &  0.08 \\\\ 2187 &     &XMMU J105243.2+571722 & 10 52 43.28 & 57 17 22.2 & 0.9 &     39 &    31.0 &  5.3 &     27 &     7.6 &  1.1 &     11 &    23.5 &  5.2 &      1 &    25.5 & 11.9 & -0.34 &  0.12 & -0.13 &  0.25 \\\\ 254 & 486 &XMMU J105243.5+572759 & 10 52 43.51 & 57 27 59.5 & 0.2 &   1541 &    98.1 &  3.9 &   1091 &    26.3 &  0.9 &    399 &    71.8 &  3.8 &     52 &    43.1 &  5.2 & -0.40 &  0.03 & -0.40 &  0.06 \\\\ 2379 &     &XMMU J105243.5+572504 & 10 52 43.55 & 57 25  4.7 & 1.0 &     52 &    19.1 &  2.6 &     26 &     3.6 &  0.6 &     25 &    15.6 &  2.6 &      1 &     9.2 &  4.2 & -0.19 &  0.11 & -0.41 &  0.20 \\\\ 1235 &     &XMMU J105243.5+571607 & 10 52 43.57 & 57 16  7.2 & 0.8 &     97 &    70.0 &  7.2 &     36 &     8.8 &  1.2 &     60 &    61.2 &  7.1 &      1 &    29.8 & 14.5 &  0.05 &  0.09 & -0.49 &  0.19 \\\\ 489 &     &XMMU J105243.7+573546 & 10 52 43.77 & 57 35 46.8 & 0.7 &    116 &    21.8 &  3.0 &     99 &     7.9 &  0.7 &     14 &    13.9 &  3.0 &      4 &    17.5 &  6.0 & -0.56 &  0.08 & -0.06 &  0.20 \\\\ 322 &     &XMMU J105243.7+572947 & 10 52 43.79 & 57 29 48.0 & 1.1 &     27 &    12.0 &  2.2 &     18 &     2.6 &  0.5 &     11 &     9.4 &  2.2 &      0 &     0.0 &  3.4 & -0.27 &  0.14 & -1.00 &  0.51 \\\\ 465 &     &XMMU J105244.2+573442 & 10 52 44.23 & 57 34 42.2 & 0.8 &     61 &    16.5 &  2.7 &     45 &     4.5 &  0.6 &     13 &    11.9 &  2.6 &      3 &    14.3 &  5.0 & -0.41 &  0.10 & -0.09 &  0.21 \\\\ 15 &     &XMMU J105244.2+571710 & 10 52 44.28 & 57 17 10.5 & 0.8 &     59 &    37.8 &  5.2 &     39 &     8.5 &  1.0 &     16 &    29.4 &  5.1 &      4 &    31.6 & 11.8 & -0.29 &  0.10 & -0.14 &  0.20 \\\\ 502 &     &XMMU J105244.7+573613 & 10 52 44.75 & 57 36 13.6 & 1.6 &     10 &     2.7 &  1.4 &     12 &     2.7 &  0.6 &      0 &     0.0 &  1.2 &      0 &     0.0 &  5.2 & -1.00 &  0.15 &  -$^a$ &  -$^a$ \\\\ 23 &  47 &XMMU J105244.9+571733 & 10 52 44.92 & 57 17 33.6 & 0.9 &     40 &    31.3 &  4.7 &     28 &    13.3 &  1.1 &     12 &    18.0 &  4.5 &      1 &     9.3 &  9.9 & -0.65 &  0.08 & -0.46 &  0.43 \\\\ 96 & 814 &XMMU J105244.9+572122 & 10 52 44.92 & 57 21 22.2 & 0.2 &    298 &   108.1 &  4.9 &    253 &    44.6 &  1.3 &     41 &    63.5 &  4.7 &      5 &    47.7 &  7.9 & -0.63 &  0.02 & -0.31 &  0.08 \\\\ 2514 &     &XMMU J105245.5+574042 & 10 52 45.54 & 57 40 42.0 & 2.0 &     23 &    11.8 &  4.1 &     26 &     8.1 &  1.3 &      1 &     3.6 &  3.9 &      1 &    10.6 & 10.8 & -0.87 &  0.13 &  0.34 &  0.65 \\\\ 532 & 801 &XMMU J105245.6+573745 & 10 52 45.63 & 57 37 45.8 & 0.6 &    284 &    34.6 &  3.7 &    273 &    18.3 &  1.1 &     14 &    16.3 &  3.5 &      0 &     0.0 &  7.1 & -0.75 &  0.05 & -1.00 &  0.61 \\\\ 139 &     &XMMU J105245.7+572329 & 10 52 45.78 & 57 23 29.7 & 0.6 &     12 &     3.6 &  1.3 &     11 &     2.8 &  0.5 &      0 &     0.8 &  1.2 &      2 &    12.6 &  5.1 & -0.91 &  0.13 &  0.84 &  0.24 \\\\ 2513 &     &XMMU J105245.8+574033 & 10 52 45.84 & 57 40 33.0 & 2.0 &     80 &    26.4 &  4.8 &     81 &    14.5 &  1.4 &      5 &    11.9 &  4.6 &      1 &    11.8 & 10.5 & -0.77 &  0.08 & -0.18 &  0.47 \\\\ 334 &     &XMMU J105246.2+573009 & 10 52 46.22 & 57 30  9.5 & 0.8 &     74 &    13.4 &  2.2 &     67 &     4.8 &  0.5 &      9 &     8.6 &  2.1 &      0 &     0.6 &  2.0 & -0.56 &  0.09 & -0.91 &  0.31 \\\\ 279 &     &XMMU J105246.3+572820 & 10 52 46.37 & 57 28 20.5 & 1.2 &     35 &    17.5 &  2.5 &      1 &     0.8 &  0.4 &     34 &    16.6 &  2.4 &      1 &     7.6 &  3.9 &  0.53 &  0.18 & -0.51 &  0.20 \\\\ 2029 &     &XMMU J105246.6+572105 & 10 52 46.62 & 57 21  5.7 & 0.5 &    102 &    34.2 &  4.2 &     85 &    12.3 &  1.1 &     18 &    21.9 &  4.0 &      1 &    17.3 &  7.3 & -0.56 &  0.07 & -0.28 &  0.21 \\\\ 2512 & 840 &XMMU J105247.0+574044 & 10 52 47.04 & 57 40 44.0 & 2.0 &    110 &    23.4 &  4.5 &    114 &    16.9 &  1.4 &      2 &     6.5 &  4.3 &      0 &     0.0 &  6.9 & -0.89 &  0.07 & -1.00 &  1.49 \\\\ 1289 &     &XMMU J105247.0+574100 & 10 52 47.05 & 57 41  0.7 & 1.8 &     41 &    23.5 &  5.3 &     40 &    10.1 &  1.3 &      3 &    13.4 &  5.2 &      0 &     5.7 &  9.0 & -0.65 &  0.12 & -0.54 &  0.57 \\\\ 390 &     &XMMU J105247.1+573200 & 10 52 47.11 & 57 32  0.4 & 1.0 &     45 &     7.7 &  2.0 &     46 &     4.3 &  0.5 &      1 &     3.4 &  2.0 &      0 &     0.0 &  3.7 & -0.78 &  0.12 & -1.00 &  1.51 \\\\ 503 &     &XMMU J105247.9+573620 & 10 52 47.93 & 57 36 20.9 & 0.6 &    123 &    25.9 &  3.3 &     29 &     4.2 &  0.6 &     33 &    21.7 &  3.2 &     61 &    71.6 &  8.0 & -0.10 &  0.10 &  0.40 &  0.08 \\\\ 481 &     &XMMU J105248.0+573531 & 10 52 48.01 & 57 35 31.2 & 1.7 &     11 &     4.5 &  2.3 &     13 &     2.4 &  0.5 &      0 &     2.1 &  2.2 &      0 &     1.1 &  3.2 & -0.76 &  0.23 & -0.47 &  1.22 \\\\ 448 &     &XMMU J105248.0+573407 & 10 52 48.01 & 57 34  7.4 & 0.4 &    307 &    33.2 &  3.1 &    265 &    12.0 &  0.7 &     37 &    21.3 &  3.0 &      6 &    17.8 &  5.1 & -0.56 &  0.05 & -0.26 &  0.15 \\\\ 90 &  37 &XMMU J105248.2+572116 & 10 52 48.21 & 57 21 16.7 & 0.2 &    108 &   207.5 &  6.4 &     89 &    90.8 &  1.9 &     20 &   116.7 &  6.1 &      0 &    70.3 &  9.0 & -0.66 &  0.02 & -0.40 &  0.06 \\\\ 1312 &     &XMMU J105248.2+572505 & 10 52 48.29 & 57 25  5.1 & 1.3 &     22 &     3.0 &  1.5 &     25 &     3.0 &  0.5 &      0 &     0.0 &  1.4 &      0 &     0.0 &  3.9 & -1.00 &  0.14 &  -$^a$ &  -$^a$ \\\\ 594 &     &XMMU J105248.4+574128 & 10 52 48.42 & 57 41 28.8 & 0.4 &    483 &   104.5 &  7.9 &    425 &    41.1 &  2.1 &     59 &    63.4 &  7.6 &      2 &    33.9 & 15.4 & -0.61 &  0.04 & -0.45 &  0.19 \\\\ 2221 &     &XMMU J105249.5+573209 & 10 52 49.59 & 57 32  9.3 & 1.2 &     13 &     6.4 &  2.1 &     10 &     2.0 &  0.5 &      1 &     4.4 &  2.0 &      2 &     9.5 &  4.0 & -0.49 &  0.20 &  0.21 &  0.30 \\\\ 74 & 905 &XMMU J105251.1+572013 & 10 52 51.16 & 57 20 13.3 & 0.6 &    201 &    29.0 &  3.6 &    191 &    13.8 &  1.0 &     11 &    15.2 &  3.5 &      0 &     9.9 &  7.2 & -0.70 &  0.06 & -0.37 &  0.33 \\\\ 394 &     &XMMU J105251.4+573200 & 10 52 51.48 & 57 32  0.2 & 0.4 &    185 &    41.7 &  3.0 &    127 &    14.2 &  0.7 &     57 &    27.6 &  2.9 &      2 &    13.3 &  4.3 & -0.53 &  0.04 & -0.49 &  0.13 \\\\ 498 &     &XMMU J105251.5+573606 & 10 52 51.58 & 57 36  6.3 & 0.7 &     49 &    15.0 &  2.8 &     42 &     4.6 &  0.6 &      8 &    10.4 &  2.7 &      0 &     5.6 &  5.6 & -0.47 &  0.11 & -0.45 &  0.42 \\\\ 560 &     &XMMU J105251.7+573907 & 10 52 51.78 & 57 39  7.3 & 1.2 &     24 &    14.2 &  3.6 &     20 &     4.3 &  0.8 &      4 &    10.0 &  3.5 &      0 &    13.4 &  8.6 & -0.46 &  0.15 & -0.03 &  0.37 \\\\ 377 &     &XMMU J105252.1+573134 & 10 52 52.11 & 57 31 34.7 & 0.3 &    978 &    92.4 &  4.0 &    434 &    15.0 &  0.7 &    462 &    77.4 &  3.9 &     83 &    54.4 &  5.4 & -0.10 &  0.03 & -0.34 &  0.05 \\\\ 606 &     &XMMU J105252.5+574217 & 10 52 52.52 & 57 42 17.4 & 2.0 &     15 &     9.5 &  3.1 &     18 &     9.5 &  1.8 &      0 &     0.0 &  2.5 &      0 &     0.0 &  5.4 & -1.00 &  0.08 &  -$^a$ &  -$^a$ \\\\ 404 &     &XMMU J105252.7+573219 & 10 52 52.73 & 57 32 19.1 & 0.7 &     82 &    22.6 &  2.6 &     44 &     4.2 &  0.5 &     38 &    18.4 &  2.6 &      0 &     6.8 &  4.0 & -0.18 &  0.09 & -0.58 &  0.20 \\\\ 290 & 901 &XMMU J105252.9+572859 & 10 52 52.91 & 57 28 59.9 & 0.2 &    833 &   155.5 &  2.9 &    290 &    11.1 &  0.6 &    123 &    32.9 &  2.8 &    420 &   136.1 &  6.9 & -0.36 &  0.05 &  0.49 &  0.04 \\\\ 507 &     &XMMU J105253.0+573629 & 10 52 53.02 & 57 36 29.3 & 1.6 &     18 &    17.1 &  3.2 &      4 &     1.6 &  0.6 &     15 &    15.5 &  3.2 &      0 &     7.2 &  5.9 &  0.21 &  0.19 & -0.51 &  0.31 \\\\ 135 & 513 &XMMU J105254.4+572342 & 10 52 54.45 & 57 23 42.4 & 0.1 &  10554 &   321.8 &  6.5 &   8736 &   122.7 &  1.8 &   1595 &   199.2 &  6.2 &    224 &   124.3 &  8.2 & -0.59 &  0.01 & -0.39 &  0.03 \\\\ 388 &     &XMMU J105255.1+573146 & 10 52 55.18 & 57 31 46.5 & 0.5 &    237 &    40.0 &  3.1 &    135 &     7.7 &  0.6 &     97 &    32.3 &  3.0 &      7 &    15.8 &  4.3 & -0.20 &  0.06 & -0.49 &  0.11 \\\\ 65 &     &XMMU J105255.4+571950 & 10 52 55.43 & 57 19 50.7 & 0.4 &    468 &    75.9 &  5.1 &    328 &    18.9 &  1.1 &    126 &    57.1 &  4.9 &     15 &    43.5 &  8.8 & -0.35 &  0.05 & -0.30 &  0.10 \\\\ 2084 &     &XMMU J105256.7+572359 & 10 52 56.77 & 57 23 59.6 & 0.6 &    104 &    24.9 &  3.1 &     90 &     9.8 &  0.8 &     14 &    15.0 &  2.9 &      1 &    12.1 &  4.9 & -0.61 &  0.07 & -0.28 &  0.21 \\\\ \\hline \\end{tabular}} \\vfill \\end{minipage} \\end{sidewaystable*} \\begin{sidewaystable*} \\begin{minipage}[t][180mm]{\\textwidth} \\bigskip \\centering {\\scriptsize \\begin{tabular}{rrcllcrrrrrrrrrrrrrrrr} \\hline \\multicolumn{2}{c|}{SRC \\#}&\\multicolumn{1}{c|}{}&\\multicolumn{2}{c}{RA~~~(J2000)~~~DEC}&err&\\multicolumn{3}{|c|}{0.5--10.0~keV}&\\multicolumn{3}{|c|}{0.5--2.0~keV}&\\multicolumn{3}{|c|}{2.0--10.0~keV}&\\multicolumn{3}{|c|}{5.0--10.0~keV}&&&& \\\\ XMM&\\multicolumn{1}{r|}{R}&\\multicolumn{1}{c|}{IAU name}&hh mm ss&$^{\\circ}$~~~~$^\\prime$~~~$^{\\prime\\prime}$&$^{\\prime\\prime}$&\\multicolumn{1}{|r}{$\\mathcal{L}$}&flux&err&\\multicolumn{1}{|r}{$\\mathcal{L}$}&flux&err& \\multicolumn{1}{|r}{$\\mathcal{L}$}&flux&err&\\multicolumn{1}{|r}{$\\mathcal{L}$}&flux&\\multicolumn{1}{r|}{err}&HR1&err&HR2&err \\\\ \\hline 403 &     &XMMU J105257.1+573235 & 10 52 57.13 & 57 32 35.6 & 0.5 &    198 &    21.8 &  2.5 &    181 &     8.8 &  0.6 &     19 &    13.0 &  2.4 &      0 &     2.1 &  3.0 & -0.62 &  0.06 & -0.79 &  0.26 \\\\ 176 &  30 &XMMU J105257.3+572507 & 10 52 57.32 & 57 25  7.3 & 0.1 &   5813 &   180.8 &  4.8 &   5179 &    82.2 &  1.5 &    593 &    98.6 &  4.6 &     43 &    44.6 &  5.9 & -0.68 &  0.01 & -0.52 &  0.05 \\\\ 19 &     &XMMU J105257.3+571729 & 10 52 57.39 & 57 17 29.7 & 1.2 &     30 &    27.8 &  5.0 &     17 &     4.9 &  0.9 &     13 &    22.9 &  4.9 &      0 &    17.6 & 11.1 & -0.15 &  0.14 & -0.30 &  0.30 \\\\ 2375 &     &XMMU J105257.3+574152 & 10 52 57.40 & 57 41 52.8 & 0.6 &    186 &    83.7 &  8.1 &    161 &    39.0 &  2.8 &     24 &    44.7 &  7.6 &      2 &    46.6 & 18.8 & -0.69 &  0.05 & -0.15 &  0.21 \\\\ 569 &     &XMMU J105257.5+573947 & 10 52 57.56 & 57 39 47.8 & 1.2 &     21 &    24.5 &  4.5 &      0 &     1.1 &  0.7 &     18 &    23.5 &  4.5 &      3 &    28.3 & 10.0 &  0.55 &  0.24 & -0.08 &  0.20 \\\\ 121 & 434 &XMMU J105258.3+572251 & 10 52 58.32 & 57 22 51.0 & 0.3 &    842 &    69.8 &  4.0 &    723 &    25.4 &  1.0 &    116 &    44.4 &  3.9 &      4 &    18.1 &  5.7 & -0.57 &  0.03 & -0.55 &  0.11 \\\\ 1291 &     &XMMU J105258.4+573102 & 10 52 58.42 & 57 31  2.2 & 1.0 &     36 &    11.2 &  2.3 &      0 &     0.0 &  0.2 &     15 &    11.1 &  2.3 &     22 &    27.1 &  4.6 &  0.97 &  0.22 &  0.26 &  0.12 \\\\ 153 &     &XMMU J105258.5+572355 & 10 52 58.52 & 57 23 56.0 & 0.4 &    104 &    51.5 &  3.5 &     90 &    16.7 &  0.8 &     14 &    34.8 &  3.4 &      1 &    16.9 &  5.0 & -0.50 &  0.04 & -0.49 &  0.12 \\\\ 433 &     &XMMU J105258.8+573334 & 10 52 58.89 & 57 33 34.8 & 0.6 &    145 &    23.2 &  2.7 &    119 &     7.4 &  0.6 &     25 &    15.7 &  2.6 &      2 &     9.7 &  4.3 & -0.50 &  0.07 & -0.39 &  0.20 \\\\ 447 &     &XMMU J105258.9+573358 & 10 52 58.94 & 57 33 58.9 & 1.2 &     27 &     3.7 &  1.9 &     30 &     3.7 &  0.6 &      0 &     0.0 &  1.8 &      0 &     0.0 &  4.3 & -1.00 &  0.15 &  -$^a$ &  -$^a$ \\\\ 332 &  77 &XMMU J105259.1+573030 & 10 52 59.19 & 57 30 30.2 & 0.1 &   1526 &   254.1 &  5.3 &   1072 &    85.9 &  1.4 &    361 &   168.2 &  5.1 &     94 &   118.8 &  6.8 & -0.53 &  0.01 & -0.33 &  0.03 \\\\ 278 &     &XMMU J105259.3+572831 & 10 52 59.32 & 57 28 31.0 & 0.7 &     96 &     7.8 &  1.9 &     99 &     6.0 &  0.5 &      0 &     1.8 &  1.9 &      0 &     0.0 &  0.7 & -0.91 &  0.09 & -1.00 &  0.51 \\\\ 387 &  15 &XMMU J105259.8+573155 & 10 52 59.89 & 57 31 55.1 & 0.3 &   1020 &    57.5 &  3.1 &    910 &    23.7 &  0.9 &    111 &    33.8 &  3.0 &      2 &     8.8 &  3.9 & -0.63 &  0.03 & -0.69 &  0.12 \\\\ 48 &     &XMMU J105259.9+571900 & 10 52 59.94 & 57 19  0.3 & 1.0 &     25 &    18.6 &  3.6 &     17 &     4.1 &  0.8 &     10 &    14.5 &  3.5 &      0 &     3.0 &  5.9 & -0.28 &  0.14 & -0.75 &  0.44 \\\\ 34 &     &XMMU J105300.8+571812 & 10 53  0.88 & 57 18 12.4 & 0.9 &     67 &    31.6 &  4.7 &     52 &     8.0 &  1.0 &     17 &    23.6 &  4.6 &      0 &     8.8 &  9.3 & -0.36 &  0.10 & -0.58 &  0.35 \\\\ 597 &     &XMMU J105300.9+574208 & 10 53  0.99 & 57 42  8.5 & 0.1 &  35165 &  1283.9 & 14.5 &  35120 &  1204.2 & 10.5 &     53 &    79.7 & 10.0 &      0 &     3.8 & 12.0 & -0.98 &  0.00 & -0.93 &  0.20 \\\\ 607 &     &XMMU J105301.6+571459 & 10 53  1.64 & 57 14 60.0 & 0.5 &    576 &   188.6 & 12.9 &    482 &    61.0 &  2.9 &     92 &   127.7 & 12.6 &      4 &    72.6 & 24.4 & -0.50 &  0.04 & -0.43 &  0.14 \\\\ 1527 &     &XMMU J105302.0+574012 & 10 53  2.09 & 57 40 12.1 & 1.5 &     14 &    22.4 &  4.8 &      0 &     1.1 &  0.8 &     13 &    21.4 &  4.7 &      2 &    25.4 & 11.1 &  0.52 &  0.27 & -0.09 &  0.24 \\\\ 2015 &     &XMMU J105302.1+573032 & 10 53  2.17 & 57 30 32.1 & 0.8 &    105 &    21.1 &  2.6 &     78 &     5.7 &  0.6 &     26 &    15.4 &  2.5 &      1 &     7.3 &  3.9 & -0.40 &  0.08 & -0.50 &  0.21 \\\\ 527 &   5 &XMMU J105302.5+573756 & 10 53  2.52 & 57 37 56.7 & 0.2 &   4274 &   223.8 &  6.8 &   3695 &    89.7 &  1.9 &    514 &   134.1 &  6.5 &     67 &    93.2 & 10.2 & -0.62 &  0.02 & -0.34 &  0.05 \\\\ 2838 &     &XMMU J105302.8+573055 & 10 53  2.90 & 57 30 55.9 & 0.5 &     15 &     5.3 &  2.2 &     12 &     2.4 &  0.5 &      3 &     2.8 &  2.1 &      0 &     0.0 &  2.3 & -0.69 &  0.20 & -1.00 &  1.17 \\\\ 60 &     &XMMU J105303.5+571936 & 10 53  3.52 & 57 19 36.3 & 1.4 &     12 &    13.5 &  3.4 &      7 &     2.7 &  0.7 &      5 &    10.8 &  3.4 &      0 &    10.0 &  7.9 & -0.22 &  0.19 & -0.21 &  0.40 \\\\ 476 & 827 &XMMU J105303.5+573529 & 10 53  3.53 & 57 35 29.5 & 0.4 &    384 &    48.9 &  3.9 &    316 &    16.8 &  0.9 &     57 &    32.1 &  3.8 &     12 &    29.6 &  6.6 & -0.54 &  0.05 & -0.21 &  0.12 \\\\ 300 & 426 &XMMU J105303.8+572924 & 10 53  3.88 & 57 29 24.7 & 0.2 &   3169 &   136.6 &  4.4 &   2597 &    47.8 &  1.1 &    535 &    88.7 &  4.2 &     38 &    36.8 &  5.1 & -0.55 &  0.02 & -0.55 &  0.05 \\\\ 1415 &     &XMMU J105304.0+573715 & 10 53  4.08 & 57 37 15.8 & 1.5 &     14 &     9.9 &  3.1 &     10 &     2.6 &  0.6 &      2 &     7.2 &  3.0 &      2 &    16.7 &  7.1 & -0.40 &  0.20 &  0.24 &  0.28 \\\\ 145 &     &XMMU J105304.6+572345 & 10 53  4.65 & 57 23 45.7 & 1.0 &     24 &     7.9 &  2.3 &     23 &     3.2 &  0.5 &      1 &     4.7 &  2.2 &      0 &     6.1 &  4.4 & -0.63 &  0.15 & -0.05 &  0.44 \\\\ 362 &     &XMMU J105304.7+573100 & 10 53  4.71 & 57 31  0.4 & 2.2 &     14 &     5.0 &  2.1 &     16 &     2.7 &  0.5 &      0 &     2.3 &  2.1 &      0 &     0.0 &  2.0 & -0.76 &  0.19 & -1.00 &  1.19 \\\\ 2502 &     &XMMU J105304.8+574105 & 10 53  4.87 & 57 41  5.6 & 2.0 &    100 &     1.9 &  3.6 &      2 &     1.9 &  1.1 &      0 &     0.0 &  3.5 &      2 &    25.9 & 16.1 & -1.00 &  0.59 &  1.00 &  0.38 \\\\ 259 &     &XMMU J105305.6+572810 & 10 53  5.66 & 57 28 10.1 & 0.2 &   1659 &   143.3 &  5.0 &     29 &     3.4 &  0.5 &   1065 &   139.8 &  5.0 &    566 &   180.1 &  8.1 &  0.73 &  0.04 & -0.05 &  0.03 \\\\ 431 &     &XMMU J105305.8+573330 & 10 53  5.80 & 57 33 30.4 & 1.0 &     31 &    17.3 &  2.9 &     11 &     2.2 &  0.5 &     20 &    15.1 &  2.8 &      1 &     9.6 &  4.8 &  0.04 &  0.15 & -0.38 &  0.23 \\\\ 93 &     &XMMU J105306.0+572112 & 10 53  6.04 & 57 21 12.6 & 1.9 &     11 &    11.7 &  3.0 &      7 &     2.3 &  0.6 &      5 &     9.4 &  2.9 &      0 &     0.0 &  4.4 & -0.21 &  0.20 & -1.00 &  0.65 \\\\ 450 & 477 &XMMU J105306.0+573424 & 10 53  6.09 & 57 34 24.5 & 0.2 &   1423 &    80.7 &  4.0 &   1264 &    34.6 &  1.1 &    135 &    46.1 &  3.8 &     26 &    38.6 &  6.3 & -0.65 &  0.03 & -0.26 &  0.09 \\\\ 599 &  54 &XMMU J105307.3+571505 & 10 53  7.37 & 57 15  5.1 & 0.4 &    938 &   216.0 & 13.2 &    823 &    82.5 &  3.3 &    102 &   133.5 & 12.8 &     15 &   119.0 & 24.4 & -0.59 &  0.03 & -0.23 &  0.11 \\\\ 477 &     &XMMU J105308.0+573519 & 10 53  8.07 & 57 35 19.6 & 2.0 &     12 &     9.3 &  2.6 &      7 &     2.2 &  0.6 &      3 &     7.1 &  2.5 &      2 &    13.1 &  5.9 & -0.32 &  0.20 &  0.13 &  0.28 \\\\ 2257 &     &XMMU J105308.6+573437 & 10 53  8.65 & 57 34 37.1 & 1.1 &     16 &    10.5 &  2.7 &     11 &     2.5 &  0.6 &      4 &     8.0 &  2.6 &      1 &     9.7 &  5.5 & -0.33 &  0.18 & -0.08 &  0.32 \\\\ 270 & 120 &XMMU J105309.3+572820 & 10 53  9.39 & 57 28 21.0 & 0.1 &   4573 &   135.0 &  4.0 &   4101 &    61.6 &  1.2 &    434 &    73.4 &  3.8 &     40 &    37.4 &  5.1 & -0.68 &  0.01 & -0.47 &  0.06 \\\\ 138 &     &XMMU J105309.4+572330 & 10 53  9.40 & 57 23 30.2 & 1.0 &     44 &    12.6 &  2.6 &     42 &     4.6 &  0.6 &      5 &     8.0 &  2.5 &      0 &     0.0 &  3.1 & -0.57 &  0.12 & -1.00 &  0.54 \\\\ 1392 &     &XMMU J105309.4+573723 & 10 53  9.49 & 57 37 23.4 & 1.3 &     17 &    17.6 &  3.6 &      0 &     0.0 &  0.5 &     16 &    17.6 &  3.6 &      2 &    19.7 &  7.8 &  1.00 &  0.36 & -0.12 &  0.22 \\\\ 220 &     &XMMU J105309.6+572634 & 10 53  9.63 & 57 26 34.2 & 0.9 &     51 &     8.8 &  2.1 &     52 &     4.4 &  0.5 &      2 &     4.5 &  2.0 &      0 &     0.0 &  3.2 & -0.72 &  0.11 & -1.00 &  1.00 \\\\ 86 &     &XMMU J105309.7+572057 & 10 53  9.78 & 57 20 57.5 & 0.4 &    508 &    51.5 &  4.1 &    461 &    22.3 &  1.1 &     44 &    29.2 &  4.0 &      4 &    23.7 &  7.5 & -0.66 &  0.04 & -0.27 &  0.16 \\\\ 2515 &     &XMMU J105309.8+574147 & 10 53  9.84 & 57 41 47.1 & 2.0 &    100 &     3.9 &  2.9 &      5 &     3.9 &  1.4 &      0 &     0.0 &  2.5 &      0 &     0.0 &  5.9 & -1.00 &  0.20 &  -$^a$ &  -$^a$ \\\\ 497 &     &XMMU J105311.1+573600 & 10 53 11.15 & 57 36  0.1 & 0.7 &    112 &    26.3 &  3.4 &     97 &     9.9 &  0.9 &     15 &    16.4 &  3.3 &      1 &    12.2 &  6.5 & -0.58 &  0.07 & -0.31 &  0.26 \\\\ 2147 &     &XMMU J105311.3+573829 & 10 53 11.38 & 57 38 29.4 & 0.7 &     72 &    25.6 &  3.8 &     63 &    10.0 &  1.0 &     11 &    15.6 &  3.7 &      0 &     0.0 &  4.5 & -0.61 &  0.08 & -1.00 &  0.41 \\\\ 604 &     &XMMU J105311.7+574216 & 10 53 11.70 & 57 42 16.1 & 0.6 &    686 &   170.8 & 11.1 &    582 &    66.7 &  3.1 &     88 &   104.1 & 10.7 &     17 &   127.1 & 24.3 & -0.60 &  0.04 & -0.07 &  0.11 \\\\ 124 & 634 &XMMU J105311.8+572307 & 10 53 11.80 & 57 23  7.0 & 0.3 &    976 &    74.4 &  4.0 &    809 &    26.2 &  1.0 &    153 &    48.2 &  3.9 &     16 &    33.1 &  6.4 & -0.55 &  0.03 & -0.35 &  0.09 \\\\ 183 &  82 &XMMU J105312.3+572506 & 10 53 12.37 & 57 25  6.4 & 0.2 &   2456 &   106.8 &  4.1 &   2156 &    45.6 &  1.2 &    266 &    61.3 &  3.9 &     35 &    41.3 &  5.9 & -0.65 &  0.02 & -0.36 &  0.07 \\\\ 453 & 804 &XMMU J105312.4+573424 & 10 53 12.50 & 57 34 24.7 & 0.3 &    694 &    68.0 &  4.2 &    587 &    24.3 &  1.1 &    104 &    43.7 &  4.1 &      5 &    20.6 &  6.3 & -0.56 &  0.04 & -0.50 &  0.12 \\\\ 543 &     &XMMU J105312.8+573818 & 10 53 12.86 & 57 38 18.3 & 0.5 &     72 &    52.0 &  5.0 &     59 &    11.8 &  1.0 &     13 &    40.2 &  4.9 &      1 &    36.9 &  9.6 & -0.30 &  0.07 & -0.21 &  0.14 \\\\ 339 &     &XMMU J105313.0+573016 & 10 53 13.08 & 57 30 16.4 & 0.5 &    153 &    38.0 &  3.2 &     32 &     3.4 &  0.5 &    106 &    34.6 &  3.2 &     16 &    24.5 &  4.8 &  0.24 &  0.08 & -0.33 &  0.10 \\\\ 197 &     &XMMU J105313.1+572549 & 10 53 13.17 & 57 25 49.6 & 1.5 &     14 &     5.1 &  2.0 &     15 &     2.5 &  0.5 &      0 &     2.6 &  2.0 &      0 &     0.6 &  2.3 & -0.71 &  0.19 & -0.72 &  0.93 \\\\ 101 &     &XMMU J105313.2+572139 & 10 53 13.24 & 57 21 39.9 & 1.7 &     11 &     3.8 &  1.9 &     11 &     2.6 &  0.6 &      0 &     1.3 &  1.9 &      1 &    11.8 &  6.4 & -0.86 &  0.20 &  0.74 &  0.36 \\\\ 2339 &     &XMMU J105313.4+572024 & 10 53 13.47 & 57 20 24.8 & 1.3 &     30 &     8.0 &  2.7 &     33 &     5.8 &  0.8 &      0 &     2.3 &  2.6 &      0 &     0.0 &  1.7 & -0.88 &  0.13 & -1.00 &  1.08 \\\\ 283 &     &XMMU J105313.6+572834 & 10 53 13.69 & 57 28 34.6 & 1.3 &     14 &    10.5 &  2.2 &      5 &     1.5 &  0.4 &     10 &     9.1 &  2.2 &      0 &     0.0 &  2.1 & -0.01 &  0.19 & -1.00 &  0.32 \\\\ 214 &     &XMMU J105315.2+572629 & 10 53 15.23 & 57 26 29.4 & 0.3 &    847 &    79.5 &  4.0 &    586 &    19.7 &  0.9 &    248 &    59.8 &  3.9 &     14 &    25.6 &  5.4 & -0.35 &  0.03 & -0.54 &  0.08 \\\\ 162 &     &XMMU J105315.2+572418 & 10 53 15.26 & 57 24 18.1 & 1.1 &     15 &    10.6 &  2.5 &      7 &     2.0 &  0.5 &      7 &     8.6 &  2.5 &      1 &    10.9 &  5.3 & -0.18 &  0.19 & -0.06 &  0.28 \\\\ 1608 &     &XMMU J105315.6+571751 & 10 53 15.65 & 57 17 51.4 & 1.5 &     13 &    10.7 &  4.5 &     14 &     4.7 &  1.0 &      0 &     6.0 &  4.4 &      0 &     0.0 &  7.7 & -0.66 &  0.21 & -1.00 &  1.81 \\\\ 172 &     &XMMU J105315.7+572448 & 10 53 15.76 & 57 24 48.6 & 0.4 &    371 &    69.0 &  4.2 &     95 &     7.2 &  0.7 &    233 &    61.8 &  4.1 &     43 &    50.9 &  6.6 &  0.16 &  0.06 & -0.26 &  0.07 \\\\ 52 &     &XMMU J105316.0+571909 & 10 53 16.07 & 57 19  9.7 & 1.1 &     24 &    16.5 &  3.9 &     21 &     5.2 &  0.9 &      4 &    11.2 &  3.8 &      0 &     9.7 &  9.2 & -0.49 &  0.14 & -0.24 &  0.48 \\\\ 1520 &     &XMMU J105316.2+574123 & 10 53 16.28 & 57 41 23.2 & 2.3 &     10 &    31.3 &  7.1 &      0 &     1.3 &  1.3 &     11 &    30.1 &  7.0 &      0 &    16.0 & 16.8 &  0.58 &  0.34 & -0.45 &  0.43 \\\\ \\hline \\end{tabular}} \\vfill \\end{minipage} \\end{sidewaystable*} \\begin{sidewaystable*} \\begin{minipage}[t][180mm]{\\textwidth} \\bigskip \\centering {\\scriptsize \\begin{tabular}{rrcllcrrrrrrrrrrrrrrrr} \\hline \\multicolumn{2}{c|}{SRC \\#}&\\multicolumn{1}{c|}{}&\\multicolumn{2}{c}{RA~~~(J2000)~~~DEC}&err&\\multicolumn{3}{|c|}{0.5--10.0~keV}&\\multicolumn{3}{|c|}{0.5--2.0~keV}&\\multicolumn{3}{|c|}{2.0--10.0~keV}&\\multicolumn{3}{|c|}{5.0--10.0~keV}&&&& \\\\ XMM&\\multicolumn{1}{r|}{R}&\\multicolumn{1}{c|}{IAU name}&hh mm ss&$^{\\circ}$~~~~$^\\prime$~~~$^{\\prime\\prime}$&$^{\\prime\\prime}$&\\multicolumn{1}{|r}{$\\mathcal{L}$}&flux&err&\\multicolumn{1}{|r}{$\\mathcal{L}$}&flux&err& \\multicolumn{1}{|r}{$\\mathcal{L}$}&flux&err&\\multicolumn{1}{|r}{$\\mathcal{L}$}&flux&\\multicolumn{1}{r|}{err}&HR1&err&HR2&err \\\\ \\hline 475 &   6 &XMMU J105316.8+573550 & 10 53 16.86 & 57 35 50.9 & 0.1 &  91979 &  2154.7 & 17.9 &  72626 &   777.2 &  4.9 &  16338 &  1377.5 & 17.3 &   3016 &   914.3 & 21.6 & -0.56 &  0.00 & -0.36 &  0.01 \\\\ 66 &     &XMMU J105317.3+572010 & 10 53 17.38 & 57 20 11.0 & 0.8 &    101 &    21.9 &  3.6 &     98 &    11.1 &  1.0 &      5 &    10.7 &  3.4 &      0 &    11.8 &  7.8 & -0.74 &  0.08 & -0.13 &  0.36 \\\\ 103 &     &XMMU J105317.8+572150 & 10 53 17.89 & 57 21 50.1 & 0.7 &    160 &    29.6 &  3.5 &    140 &    10.9 &  0.9 &     21 &    18.7 &  3.4 &      0 &    10.8 &  6.5 & -0.57 &  0.07 & -0.42 &  0.26 \\\\ 2353 &     &XMMU J105318.6+572719 & 10 53 18.68 & 57 27 19.6 & 2.3 &     11 &     2.5 &  1.6 &     13 &     2.5 &  0.5 &      0 &     0.0 &  1.5 &      0 &     3.4 &  4.1 & -1.00 &  0.19 &  1.00 &  1.22 \\\\ 69 &  41 &XMMU J105318.9+572044 & 10 53 18.97 & 57 20 44.7 & 0.3 &    967 &    72.5 &  4.7 &    945 &    45.7 &  1.6 &     25 &    26.9 &  4.4 &      0 &     2.7 &  5.2 & -0.83 &  0.03 & -0.87 &  0.24 \\\\ 39 &  45 &XMMU J105319.0+571852 & 10 53 19.03 & 57 18 52.9 & 0.2 &   2189 &   181.7 &  7.3 &   1930 &    77.1 &  2.1 &    232 &   104.5 &  7.0 &     29 &    89.3 & 13.8 & -0.65 &  0.02 & -0.25 &  0.08 \\\\ 1537 &     &XMMU J105319.3+571520 & 10 53 19.37 & 57 15 20.1 & 1.7 &     19 &    56.8 & 12.3 &     10 &     8.6 &  2.0 &      9 &    48.2 & 12.2 &      0 &    22.6 & 23.7 & -0.06 &  0.17 & -0.50 &  0.40 \\\\ 540 &     &XMMU J105320.3+573806 & 10 53 20.31 & 57 38  6.9 & 1.0 &     32 &    14.4 &  3.9 &     32 &     6.4 &  0.9 &      1 &     8.0 &  3.8 &      0 &     0.0 &  5.8 & -0.67 &  0.14 & -1.00 &  1.01 \\\\ 384 &     &XMMU J105321.6+573147 & 10 53 21.69 & 57 31 47.2 & 0.3 &    714 &    63.6 &  3.7 &    555 &    19.1 &  0.9 &    149 &    44.5 &  3.6 &     11 &    24.0 &  5.4 & -0.46 &  0.04 & -0.45 &  0.10 \\\\ 2509 &     &XMMU J105321.9+573524 & 10 53 21.95 & 57 35 24.1 & 2.0 &     89 &    15.9 &  3.0 &     91 &     8.5 &  0.8 &      5 &     7.4 &  2.9 &      0 &     1.0 &  3.7 & -0.76 &  0.09 & -0.82 &  0.59 \\\\ 287 & 821 &XMMU J105322.1+572852 & 10 53 22.18 & 57 28 52.2 & 0.3 &    677 &    51.8 &  3.2 &    568 &    18.5 &  0.8 &    103 &    33.3 &  3.1 &      7 &    17.7 &  4.8 & -0.56 &  0.04 & -0.46 &  0.11 \\\\ 572 &     &XMMU J105322.2+574008 & 10 53 22.23 & 57 40  8.9 & 0.9 &     80 &    37.1 &  6.0 &     73 &    13.8 &  1.4 &      9 &    23.3 &  5.8 &      0 &     1.3 &  7.4 & -0.58 &  0.09 & -0.93 &  0.41 \\\\ 470 &     &XMMU J105322.9+573501 & 10 53 22.97 & 57 35  1.1 & 0.7 &     68 &    16.3 &  2.9 &     60 &     6.3 &  0.7 &      6 &    10.0 &  2.8 &      2 &    14.8 &  6.3 & -0.60 &  0.10 &  0.02 &  0.26 \\\\ 25 &     &XMMU J105323.5+571736 & 10 53 23.58 & 57 17 36.3 & 0.9 &     36 &    20.0 &  5.1 &     35 &     7.8 &  1.1 &      2 &    12.2 &  5.0 &      0 &     0.0 &  6.9 & -0.60 &  0.14 & -1.00 &  0.80 \\\\ 545 &     &XMMU J105324.2+573828 & 10 53 24.26 & 57 38 28.5 & 0.8 &     64 &    21.3 &  4.0 &     60 &     8.4 &  0.9 &      6 &    12.9 &  3.9 &      0 &    12.6 &  8.8 & -0.61 &  0.10 & -0.18 &  0.37 \\\\ 267 & 428 &XMMU J105324.6+572818 & 10 53 24.67 & 57 28 18.9 & 0.2 &   2324 &   109.2 &  4.2 &   1981 &    41.9 &  1.1 &    305 &    67.3 &  4.0 &     39 &    42.7 &  5.8 & -0.59 &  0.02 & -0.38 &  0.06 \\\\ 146 &     &XMMU J105324.7+572348 & 10 53 24.78 & 57 23 48.2 & 1.4 &     10 &     7.3 &  2.6 &      9 &     2.2 &  0.5 &      1 &     5.1 &  2.5 &      0 &     8.1 &  5.6 & -0.45 &  0.22 &  0.06 &  0.42 \\\\ 1492 &     &XMMU J105325.1+573202 & 10 53 25.11 & 57 32  2.3 & 1.3 &     11 &     2.5 &  1.2 &     14 &     2.5 &  0.5 &      0 &     0.0 &  1.1 &      0 &     0.0 &  3.7 & -1.00 &  0.14 &  -$^a$ &  -$^a$ \\\\ 305 &     &XMMU J105325.5+572910 & 10 53 25.56 & 57 29 10.3 & 1.5 &     26 &     7.1 &  2.2 &     27 &     3.4 &  0.5 &      1 &     3.6 &  2.1 &      0 &     0.0 &  1.9 & -0.71 &  0.15 & -1.00 &  0.74 \\\\ 247 &     &XMMU J105325.9+572740 & 10 53 25.91 & 57 27 40.9 & 1.8 &     17 &    11.1 &  2.4 &     10 &     2.1 &  0.5 &      8 &     8.9 &  2.4 &      0 &     0.0 &  3.5 & -0.21 &  0.17 & -1.00 &  0.55 \\\\ 1496 &     &XMMU J105326.0+572035 & 10 53 26.07 & 57 20 35.8 & 0.8 &     10 &     4.8 &  2.5 &     12 &     3.5 &  0.8 &      0 &     1.3 &  2.4 &      0 &     2.6 &  8.2 & -0.89 &  0.19 &  0.16 &  1.81 \\\\ 1476 &     &XMMU J105326.8+573904 & 10 53 26.86 & 57 39  4.7 & 1.4 &     17 &    23.1 &  4.7 &      7 &     3.1 &  0.8 &     11 &    20.0 &  4.6 &      0 &     0.0 &  7.3 &  0.02 &  0.18 & -1.00 &  0.51 \\\\ 102 &     &XMMU J105326.8+572145 & 10 53 26.89 & 57 21 45.2 & 0.7 &     87 &    40.9 &  4.3 &     21 &     4.1 &  0.7 &     62 &    36.8 &  4.3 &      6 &    27.4 &  8.0 &  0.17 &  0.10 & -0.31 &  0.14 \\\\ 189 &     &XMMU J105326.8+572521 & 10 53 26.90 & 57 25 21.7 & 0.5 &    129 &    38.2 &  3.5 &     13 &     2.5 &  0.5 &     92 &    35.7 &  3.5 &     24 &    38.5 &  6.3 &  0.38 &  0.10 & -0.14 &  0.09 \\\\ 459 &     &XMMU J105327.2+573429 & 10 53 27.23 & 57 34 29.0 & 0.5 &    184 &    49.4 &  4.0 &     46 &     5.5 &  0.7 &    112 &    44.0 &  4.0 &     26 &    45.0 &  7.3 &  0.12 &  0.08 & -0.16 &  0.09 \\\\ 137 &     &XMMU J105327.4+572335 & 10 53 27.41 & 57 23 35.9 & 0.5 &    216 &    33.5 &  3.6 &    189 &    12.2 &  0.9 &     27 &    21.3 &  3.5 &      2 &    14.5 &  6.2 & -0.56 &  0.06 & -0.35 &  0.20 \\\\ 488 & 815 &XMMU J105328.6+573540 & 10 53 28.67 & 57 35 40.7 & 1.2 &     40 &     6.3 & 1.7 &     43 &     6.2 &  0.8 &      0 &     0.1 &  1.5 &      0 &     0.0 &  2.9 & -0.99 &  0.07 & -1.00 & 37.59 \\\\ 195 &     &XMMU J105328.7+572537 & 10 53 28.74 & 57 25 37.8 & 1.0 &     39 &    18.0 &  2.9 &     25 &     3.8 &  0.6 &     15 &    14.2 &  2.8 &      0 &     4.9 &  5.0 & -0.25 &  0.12 & -0.61 &  0.33 \\\\ 2053 &     &XMMU J105329.2+572046 & 10 53 29.27 & 57 20 46.7 & 1.0 &     85 &    21.8 &  3.9 &     84 &    10.8 &  1.0 &      2 &    10.9 &  3.7 &      0 &     8.5 &  8.1 & -0.72 &  0.08 & -0.29 &  0.46 \\\\ 85 &  38 &XMMU J105329.5+572104 & 10 53 29.52 & 57 21  4.2 & 0.2 &     70 &   143.1 &  5.9 &     72 &    62.9 &  1.7 &      0 &    80.3 &  5.6 &      0 &    36.3 &  9.0 & -0.66 &  0.02 & -0.52 &  0.09 \\\\ 20 &     &XMMU J105330.4+571731 & 10 53 30.50 & 57 17 31.1 & 0.6 &    217 &    90.3 &  7.3 &    119 &    16.0 &  1.3 &     93 &    74.4 &  7.2 &      5 &    50.0 & 15.2 & -0.15 &  0.06 & -0.36 &  0.14 \\\\ 563 &     &XMMU J105330.9+573922 & 10 53 30.95 & 57 39 22.9 & 0.4 &    579 &   175.5 &  9.1 &     60 &    10.8 &  1.2 &    382 &   164.7 &  9.0 &    136 &   237.7 & 19.2 &  0.42 &  0.05 &  0.01 &  0.05 \\\\ 590 &     &XMMU J105331.4+574113 & 10 53 31.42 & 57 41 13.9 & 0.9 &     69 &    47.9 &  7.7 &     57 &    16.0 &  1.9 &     11 &    31.8 &  7.5 &      1 &    42.8 & 21.1 & -0.52 &  0.10 & -0.03 &  0.27 \\\\ 84 &     &XMMU J105331.6+572033 & 10 53 31.64 & 57 20 33.4 & 0.6 &    179 &    60.3 &  5.1 &     82 &     9.4 &  0.9 &     90 &    50.9 &  5.0 &      7 &    38.7 & 10.1 & -0.08 &  0.07 & -0.30 &  0.13 \\\\ 168 &  31 &XMMU J105331.8+572454 & 10 53 31.84 & 57 24 54.3 & 0.1 &  22507 &   643.0 &  9.4 &  18931 &   256.9 &  2.8 &   3158 &   386.1 &  9.0 &    420 &   222.7 & 11.4 & -0.62 &  0.01 & -0.42 &  0.02 \\\\ 315 &     &XMMU J105332.7+572942 & 10 53 32.75 & 57 29 42.9 & 1.0 &     28 &     9.0 &  2.3 &     25 &     3.3 &  0.5 &      2 &     5.7 &  2.3 &      2 &    10.8 &  4.8 & -0.57 &  0.15 &  0.15 &  0.29 \\\\ 487 &     &XMMU J105333.4+573536 & 10 53 33.41 & 57 35 37.0 & 2.4 &     18 &     7.6 &  3.1 &     20 &     4.1 &  0.7 &      0 &     3.4 &  3.0 &      0 &     2.8 &  4.8 & -0.77 &  0.18 & -0.27 &  0.89 \\\\ 258 &     &XMMU J105333.6+572812 & 10 53 33.67 & 57 28 12.4 & 2.1 &     10 &     6.4 &  2.4 &     10 &     2.2 &  0.5 &      1 &     4.2 &  2.3 &      0 &     1.4 &  3.2 & -0.53 &  0.22 & -0.63 &  0.72 \\\\ 116 &     &XMMU J105334.7+572227 & 10 53 34.73 & 57 22 27.3 & 0.6 &    189 &    66.3 &  5.0 &     28 &     4.7 &  0.7 &    151 &    61.6 &  5.0 &     10 &    34.6 &  8.2 &  0.35 &  0.08 & -0.43 &  0.10 \\\\ 147 &     &XMMU J105334.9+572344 & 10 53 34.97 & 57 23 44.0 & 0.9 &     51 &    21.9 &  3.3 &     35 &     5.0 &  0.7 &     18 &    16.9 &  3.2 &      0 &     2.3 &  4.2 & -0.30 &  0.11 & -0.82 &  0.29 \\\\ 191 &  29 &XMMU J105335.1+572542 & 10 53 35.12 & 57 25 42.4 & 0.1 &  37756 &   913.1 & 11.0 &  31005 &   348.5 &  3.1 &   5824 &   564.5 & 10.6 &    928 &   350.8 & 13.2 & -0.59 &  0.01 & -0.39 &  0.02 \\\\ 1455 &     &XMMU J105335.4+573334 & 10 53 35.46 & 57 33 34.6 & 1.5 &     11 &     9.7 &  2.8 &      9 &     2.4 &  0.6 &      3 &     7.3 &  2.7 &      0 &     3.2 &  4.5 & -0.34 &  0.20 & -0.53 &  0.53 \\\\ 518 &     &XMMU J105336.3+573726 & 10 53 36.33 & 57 37 26.4 & 0.4 &    352 &    86.1 &  5.8 &    215 &    19.5 &  1.3 &    129 &    66.6 &  5.7 &      9 &    43.2 & 10.4 & -0.30 &  0.05 & -0.37 &  0.11 \\\\ 535 & 232 &XMMU J105336.3+573801 & 10 53 36.33 & 57 38  1.0 & 0.1 &  22734 &   452.4 &  6.4 &  22713 &   422.0 &  4.2 &     28 &    30.4 &  4.8 &      0 &     0.9 &  5.5 & -0.98 &  0.00 & -0.96 &  0.24 \\\\ 28 &     &XMMU J105336.3+571756 & 10 53 36.34 & 57 17 56.6 & 0.7 &     80 &    40.9 &  5.9 &     63 &    11.5 &  1.3 &     17 &    29.4 &  5.7 &      0 &    23.1 & 13.9 & -0.42 &  0.09 & -0.29 &  0.29 \\\\ 44 &     &XMMU J105336.5+571846 & 10 53 36.52 & 57 18 46.2 & 0.8 &     81 &    22.9 &  4.6 &     75 &    11.6 &  1.2 &      3 &    11.3 &  4.5 &      4 &    41.0 & 12.7 & -0.73 &  0.09 &  0.44 &  0.20 \\\\ 409 &     &XMMU J105336.8+573259 & 10 53 36.90 & 57 32 59.0 & 1.4 &     87 &    11.6 &  2.7 &     89 &     8.7 &  0.8 &      0 &     2.9 &  2.6 &      0 &     8.9 &  5.9 & -0.90 &  0.09 &  0.37 &  0.48 \\\\ 1501 &     &XMMU J105337.1+574153 & 10 53 37.13 & 57 41 53.7 & 1.6 &     15 &    74.2 & 17.2 &      6 &     8.9 &  2.8 &     10 &    65.3 & 17.0 &      0 &     0.2 & 18.3 &  0.07 &  0.20 & -1.00 &  0.39 \\\\ 122 &     &XMMU J105337.5+572247 & 10 53 37.50 & 57 22 47.4 & 0.8 &     63 &    24.3 &  3.6 &     47 &     6.2 &  0.8 &     17 &    18.1 &  3.5 &      0 &    10.1 &  7.3 & -0.37 &  0.10 & -0.44 &  0.31 \\\\ 562 &     &XMMU J105338.1+573919 & 10 53 38.12 & 57 39 19.0 & 1.5 &     29 &    31.0 &  5.8 &     18 &     6.5 &  1.2 &     11 &    24.6 &  5.7 &      0 &    20.7 & 13.3 & -0.25 &  0.14 & -0.25 &  0.32 \\\\ 271 &     &XMMU J105338.3+572811 & 10 53 38.32 & 57 28 12.0 & 1.3 &     16 &    11.2 &  2.7 &     12 &     2.6 &  0.6 &      5 &     8.7 &  2.7 &      0 &     6.5 &  5.7 & -0.31 &  0.17 & -0.31 &  0.42 \\\\ 511 &     &XMMU J105338.3+573654 & 10 53 38.33 & 57 36 54.6 & 0.4 &    351 &   102.0 &  6.3 &    131 &    13.4 &  1.1 &    211 &    88.6 &  6.2 &     10 &    42.0 & 10.0 &  0.02 &  0.05 & -0.50 &  0.09 \\\\ 472 & 228 &XMMU J105339.7+573520 & 10 53 39.78 & 57 35 20.2 & 0.4 &   1581 &   133.9 &  6.1 &   1411 &    58.3 &  1.8 &    162 &    75.6 &  5.8 &      9 &    37.7 &  9.1 & -0.66 &  0.02 & -0.48 &  0.10 \\\\ 342 &  16 &XMMU J105339.8+573104 & 10 53 39.82 & 57 31  4.7 & 0.1 &  44056 &   937.6 & 10.8 &  38819 &   437.3 &  3.6 &   4646 &   500.3 & 10.3 &    594 &   276.9 & 12.5 & -0.69 &  0.01 & -0.44 &  0.02 \\\\ 144 & 131 &XMMU J105340.1+572352 & 10 53 40.17 & 57 23 52.7 & 0.5 &    502 &    28.1 &  2.9 &    507 &    25.7 &  1.2 &      0 &     2.4 &  2.7 &      0 &     3.1 &  5.0 & -0.97 &  0.03 & -0.05 &  0.98 \\\\ 117 &     &XMMU J105340.5+572230 & 10 53 40.51 & 57 22 30.8 & 1.0 &     50 &    28.8 &  4.0 &     21 &     4.2 &  0.7 &     28 &    24.6 &  3.9 &      2 &    19.8 &  8.3 & -0.04 &  0.12 & -0.28 &  0.21 \\\\ 2380 &     &XMMU J105341.4+573024 & 10 53 41.42 & 57 30 24.7 & 1.0 &     37 &    20.6 &  3.2 &     18 &     3.8 &  0.7 &     19 &    16.8 &  3.2 &      0 &     8.3 &  5.6 & -0.17 &  0.13 & -0.48 &  0.27 \\\\ 58 &     &XMMU J105341.9+571935 & 10 53 41.93 & 57 19 35.7 & 0.6 &    220 &    50.8 &  5.8 &    200 &    20.4 &  1.4 &     19 &    30.4 &  5.7 &      2 &    31.0 & 12.1 & -0.62 &  0.06 & -0.16 &  0.21 \\\\ 397 &     &XMMU J105342.0+573220 & 10 53 42.06 & 57 32 20.1 & 0.7 &    123 &    39.7 &  4.1 &     48 &     6.2 &  0.8 &     58 &    33.5 &  4.0 &     17 &    41.4 &  8.0 & -0.08 &  0.09 & -0.07 &  0.11 \\\\ 190 &     &XMMU J105344.3+572532 & 10 53 44.35 & 57 25 32.1 & 0.8 &     59 &    26.1 &  3.9 &     16 &     3.7 &  0.7 &     24 &    22.4 &  3.8 &     19 &    44.7 &  8.2 & -0.02 &  0.13 &  0.17 &  0.12 \\\\ 277 &  25 &XMMU J105344.9+572840 & 10 53 44.97 & 57 28 40.5 & 0.1 &   7891 &   301.4 &  6.9 &   6554 &   111.2 &  1.9 &   1187 &   190.2 &  6.7 &    151 &   117.8 &  9.1 & -0.57 &  0.01 & -0.39 &  0.04 \\\\ \\hline \\end{tabular}} \\vfill \\end{minipage} \\end{sidewaystable*} \\begin{sidewaystable*} \\begin{minipage}[t][180mm]{\\textwidth} \\bigskip \\centering {\\scriptsize \\begin{tabular}{rrcllcrrrrrrrrrrrrrrrr} \\hline \\multicolumn{2}{c|}{SRC \\#}&\\multicolumn{1}{c|}{}&\\multicolumn{2}{c}{RA~~~(J2000)~~~DEC}&err&\\multicolumn{3}{|c|}{0.5--10.0~keV}&\\multicolumn{3}{|c|}{0.5--2.0~keV}&\\multicolumn{3}{|c|}{2.0--10.0~keV}&\\multicolumn{3}{|c|}{5.0--10.0~keV}&&&& \\\\ XMM&\\multicolumn{1}{r|}{R}&\\multicolumn{1}{c|}{IAU name}&hh mm ss&$^{\\circ}$~~~~$^\\prime$~~~$^{\\prime\\prime}$&$^{\\prime\\prime}$&\\multicolumn{1}{|r}{$\\mathcal{L}$}&flux&err&\\multicolumn{1}{|r}{$\\mathcal{L}$}&flux&err& \\multicolumn{1}{|r}{$\\mathcal{L}$}&flux&err&\\multicolumn{1}{|r}{$\\mathcal{L}$}&flux&\\multicolumn{1}{r|}{err}&HR1&err&HR2&err \\\\ \\hline 165 &     &XMMU J105345.0+572419 & 10 53 45.01 & 57 24 19.5 & 1.7 &     13 &    12.1 &  3.2 &      9 &     2.6 &  0.6 &      4 &     9.5 &  3.2 &      0 &     5.8 &  6.3 & -0.27 &  0.19 & -0.40 &  0.48 \\\\ 370 &     &XMMU J105346.1+573104 & 10 53 46.18 & 57 31  4.4 & 1.1 &     37 &     9.6 &  2.8 &     39 &     5.3 &  0.7 &      0 &     4.3 &  2.7 &      0 &     3.3 &  4.6 & -0.77 &  0.13 & -0.30 &  0.69 \\\\ 1544 &     &XMMU J105346.2+573608 & 10 53 46.23 & 57 36  8.5 & 1.9 &     12 &     4.4 &  2.1 &     13 &     4.0 &  0.8 &      0 &     0.4 &  1.9 &      1 &    16.3 &  9.5 & -0.97 &  0.14 &  0.93 &  0.31 \\\\ 461 & 229 &XMMU J105346.4+573510 & 10 53 46.47 & 57 35 10.8 & 0.3 &    932 &    94.3 &  5.9 &    888 &    50.5 &  1.9 &     48 &    43.8 &  5.6 &      0 &     6.0 &  7.3 & -0.76 &  0.03 & -0.83 &  0.20 \\\\ 200 &     &XMMU J105346.9+572606 & 10 53 46.90 & 57 26  6.2 & 0.3 &    856 &    76.0 &  4.5 &    741 &    29.4 &  1.2 &    105 &    46.6 &  4.3 &     11 &    35.3 &  7.9 & -0.60 &  0.03 & -0.30 &  0.11 \\\\ 442 & 805 &XMMU J105347.3+573350 & 10 53 47.37 & 57 33 50.0 & 0.4 &    668 &    84.1 &  5.1 &    543 &    27.7 &  1.3 &    122 &    56.4 &  5.0 &      4 &    25.7 &  8.3 & -0.51 &  0.04 & -0.51 &  0.12 \\\\ 268 &     &XMMU J105348.1+572816 & 10 53 48.15 & 57 28 16.3 & 0.3 &    761 &   109.9 &  5.4 &    382 &    19.3 &  1.0 &    336 &    90.6 &  5.3 &     43 &    61.6 &  8.1 & -0.15 &  0.04 & -0.35 &  0.06 \\\\ 326 & 117 &XMMU J105348.8+573034 & 10 53 48.84 & 57 30 34.2 & 0.1 &   8944 &   488.6 &  9.5 &   5360 &   106.2 &  1.9 &   2989 &   382.4 &  9.4 &    595 &   296.9 & 13.4 & -0.27 &  0.01 & -0.29 &  0.02 \\\\ 383 &     &XMMU J105348.9+573133 & 10 53 48.91 & 57 31 33.2 & 0.7 &    101 &     9.2 &  1.1 &    105 &     9.2 &  0.8 &      0 &     0.0 &  0.8 &      0 &     3.0 &  4.6 & -1.00 &  0.03 &  1.00 &  0.73 \\\\ 2020 &  27 &XMMU J105350.2+572709 & 10 53 50.28 & 57 27  9.7 & 0.2 &   1607 &   171.5 &  6.5 &   1313 &    72.1 &  1.9 &    259 &    99.4 &  6.2 &     36 &    50.6 &  8.2 & -0.64 &  0.02 & -0.47 &  0.07 \\\\ 2374 &     &XMMU J105350.3+573000 & 10 53 50.33 & 57 30  0.4 & 1.3 &     19 &     5.6 &  2.2 &     19 &     4.0 &  0.7 &      0 &     1.6 &  2.1 &      1 &    13.5 &  6.4 & -0.88 &  0.15 &  0.72 &  0.34 \\\\ 452 &     &XMMU J105350.4+573406 & 10 53 50.42 & 57 34  6.2 & 1.4 &     30 &    25.6 &  4.2 &      8 &     2.7 &  0.7 &     21 &    22.9 &  4.1 &      2 &    18.1 &  8.2 &  0.14 &  0.16 & -0.28 &  0.22 \\\\ 1551 &     &XMMU J105350.7+572733 & 10 53 50.73 & 57 27 34.0 & 0.8 &     85 &    20.2 &  3.4 &     81 &     9.6 &  0.9 &      5 &    10.7 &  3.3 &      1 &    12.0 &  6.9 & -0.70 &  0.08 & -0.11 &  0.32 \\\\ 182 & 128 &XMMU J105350.8+572513 & 10 53 50.81 & 57 25 13.6 & 0.3 &   1316 &   114.1 &  5.5 &   1102 &    40.6 &  1.4 &    194 &    73.5 &  5.3 &     22 &    49.6 &  8.7 & -0.55 &  0.03 & -0.35 &  0.08 \\\\ 508 &     &XMMU J105351.5+573637 & 10 53 51.56 & 57 36 37.7 & 0.4 &    286 &    56.5 &  5.3 &    247 &    21.2 &  1.3 &     36 &    35.3 &  5.2 &      5 &    37.6 & 11.8 & -0.58 &  0.05 & -0.14 &  0.17 \\\\ 222 &     &XMMU J105351.7+572702 & 10 53 51.77 & 57 27  2.6 & 0.2 &   1627 &   268.3 &  7.8 &   1281 &    83.5 &  2.0 &    300 &   184.7 &  7.6 &     47 &    94.6 &  9.7 & -0.48 &  0.02 & -0.47 &  0.04 \\\\ 438 &     &XMMU J105352.1+573344 & 10 53 52.12 & 57 33 44.2 & 0.7 &     68 &    33.5 &  4.2 &     33 &     5.4 &  0.8 &     32 &    28.1 &  4.1 &      3 &    23.7 &  8.2 & -0.10 &  0.10 & -0.25 &  0.18 \\\\ 35 &     &XMMU J105352.9+571816 & 10 53 52.94 & 57 18 16.1 & 0.7 &    116 &    52.6 &  8.5 &    108 &    23.5 &  2.2 &      7 &    29.1 &  8.2 &      3 &    58.9 & 23.1 & -0.67 &  0.08 &  0.18 &  0.23 \\\\ 266 &     &XMMU J105354.5+572804 & 10 53 54.59 & 57 28  4.3 & 1.0 &     39 &    26.0 &  3.8 &      6 &     2.2 &  0.6 &     30 &    23.8 &  3.7 &      3 &    18.7 &  7.2 &  0.27 &  0.15 & -0.29 &  0.19 \\\\ 547 &     &XMMU J105355.6+573842 & 10 53 55.69 & 57 38 42.4 & 0.2 &   1491 &   412.1 & 16.8 &    964 &    98.0 &  3.7 &    489 &   314.1 & 16.3 &     39 &   203.2 & 28.0 & -0.33 &  0.03 & -0.37 &  0.06 \\\\ 206 &     &XMMU J105355.9+572611 & 10 53 55.95 & 57 26 11.5 & 0.8 &     78 &    39.7 &  4.4 &     12 &     3.1 &  0.7 &     59 &    36.6 &  4.3 &      7 &    30.5 &  8.2 &  0.30 &  0.11 & -0.26 &  0.14 \\\\ 151 &     &XMMU J105356.0+572356 & 10 53 56.03 & 57 23 56.4 & 0.6 &    108 &    46.2 &  4.9 &      2 &     1.7 &  0.7 &     70 &    44.5 &  4.8 &     37 &    79.7 & 11.1 &  0.62 &  0.13 &  0.12 &  0.09 \\\\ 406 & 828 &XMMU J105357.2+573241 & 10 53 57.20 & 57 32 41.8 & 0.2 &   1323 &   103.4 &  5.4 &   1186 &    45.3 &  1.5 &    116 &    58.1 &  5.2 &     23 &    56.2 &  9.6 & -0.66 &  0.03 & -0.19 &  0.09 \\\\ 1487 &     &XMMU J105357.5+572229 & 10 53 57.59 & 57 22 29.4 & 1.8 &     15 &    19.6 &  4.8 &      0 &     1.4 &  0.9 &      9 &    18.2 &  4.7 &      7 &    44.3 & 12.3 &  0.35 &  0.29 &  0.26 &  0.18 \\\\ 304 & 861 &XMMU J105358.3+572925 & 10 53 58.40 & 57 29 25.8 & 0.4 &    605 &    83.8 &  5.0 &    428 &    22.2 &  1.1 &    158 &    61.6 &  4.9 &     20 &    47.5 &  8.6 & -0.39 &  0.04 & -0.30 &  0.09 \\\\ 299 &     &XMMU J105358.8+572856 & 10 53 58.85 & 57 28 56.8 & 1.7 &     16 &     8.6 &  3.0 &     16 &     3.6 &  0.7 &      1 &     5.0 &  3.0 &      0 &    11.9 &  7.3 & -0.64 &  0.18 &  0.26 &  0.40 \\\\ 106 &     &XMMU J105359.3+572157 & 10 53 59.32 & 57 21 57.7 & 1.4 &     29 &    37.7 &  6.1 &      8 &     4.2 &  1.1 &     22 &    33.6 &  6.0 &      0 &    10.2 & 11.8 &  0.12 &  0.16 & -0.65 &  0.34 \\\\ 203 &     &XMMU J105359.9+572601 & 10 53 59.94 & 57 26  1.1 & 0.6 &    211 &    51.4 &  4.8 &    158 &    13.1 &  1.0 &     54 &    38.3 &  4.7 &      1 &    13.4 &  7.8 & -0.37 &  0.06 & -0.60 &  0.19 \\\\ 454 &     &XMMU J105400.2+573418 & 10 54  0.23 & 57 34 18.2 & 0.9 &     44 &    17.5 &  3.9 &     39 &     6.4 &  0.9 &      4 &    11.1 &  3.8 &      2 &    23.3 &  9.0 & -0.57 &  0.12 &  0.19 &  0.25 \\\\ 419 &     &XMMU J105400.3+573322 & 10 54  0.39 & 57 33 22.8 & 0.3 &   1003 &   143.7 &  6.7 &    624 &    31.5 &  1.3 &    330 &   112.2 &  6.6 &     50 &    94.9 & 11.7 & -0.28 &  0.03 & -0.25 &  0.06 \\\\ 1427 &     &XMMU J105400.7+572850 & 10 54  0.73 & 57 28 50.0 & 1.7 &     13 &    17.2 &  3.6 &      1 &     1.0 &  0.6 &     12 &    16.1 &  3.5 &      0 &    10.8 &  7.1 &  0.42 &  0.24 & -0.36 &  0.30 \\\\ 1574 &     &XMMU J105401.5+572815 & 10 54  1.57 & 57 28 16.0 & 2.3 &     10 &    14.6 &  3.6 &      0 &     0.9 &  0.6 &      8 &    13.7 &  3.5 &      2 &    18.0 &  7.9 &  0.41 &  0.29 & -0.04 &  0.25 \\\\ 119 &     &XMMU J105402.5+572235 & 10 54  2.55 & 57 22 36.0 & 0.5 &    315 &    70.2 &  6.5 &    286 &    31.4 &  1.9 &     29 &    38.8 &  6.2 &      2 &    33.8 & 15.2 & -0.67 &  0.05 & -0.24 &  0.22 \\\\ 2392 &     &XMMU J105402.5+573650 & 10 54  2.58 & 57 36 50.4 & 1.5 &     11 &    12.1 &  5.1 &     11 &     5.0 &  1.1 &      0 &     7.1 &  5.0 &      0 &    16.2 & 13.8 & -0.64 &  0.22 &  0.23 &  0.52 \\\\ 1438 &     &XMMU J105403.2+573024 & 10 54  3.25 & 57 30 24.4 & 2.0 &     13 &     7.7 &  3.1 &     14 &     3.5 &  0.7 &      0 &     4.1 &  3.0 &      0 &     0.0 &  4.5 & -0.69 &  0.20 & -1.00 &  1.54 \\\\ 1304 &     &XMMU J105403.9+572450 & 10 54  3.95 & 57 24 50.8 & 2.1 &     16 &    20.9 &  4.9 &      0 &     0.8 &  0.7 &     10 &    20.1 &  4.9 &      7 &    41.8 & 11.3 &  0.60 &  0.31 &  0.19 &  0.17 \\\\ 78 &     &XMMU J105404.3+572037 & 10 54  4.32 & 57 20 37.3 & 0.2 &   1413 &   163.8 &  8.9 &   1365 &    99.0 &  3.2 &     51 &    64.8 &  8.3 &      0 &    23.0 & 17.0 & -0.81 &  0.02 & -0.60 &  0.24 \\\\ 1518 &     &XMMU J105406.1+572413 & 10 54  6.17 & 57 24 13.4 & 1.3 &     12 &    23.4 &  5.1 &      0 &     0.7 &  0.8 &     13 &    22.7 &  5.0 &      0 &    17.0 & 11.6 &  0.68 &  0.33 & -0.31 &  0.32 \\\\ 421 &     &XMMU J105406.8+573310 & 10 54  6.85 & 57 33 10.1 & 1.0 &     43 &    28.4 &  4.7 &     19 &     4.4 &  0.8 &     18 &    24.0 &  4.6 &      7 &    38.2 & 10.0 & -0.07 &  0.13 &  0.06 &  0.16 \\\\ 469 &     &XMMU J105407.1+573525 & 10 54  7.16 & 57 35 25.3 & 0.1 &  10997 &   638.7 & 13.1 &   9676 &   287.0 &  4.2 &   1179 &   351.7 & 12.5 &    144 &   247.8 & 19.8 & -0.67 &  0.01 & -0.34 &  0.04 \\\\ 331 &     &XMMU J105407.5+573000 & 10 54  7.57 & 57 30  0.8 & 1.1 &     35 &     8.4 &  2.9 &     37 &     6.2 &  0.8 &      0 &     2.3 &  2.7 &      0 &    12.1 &  8.1 & -0.89 &  0.13 &  0.58 &  0.46 \\\\ 140 &     &XMMU J105408.4+572337 & 10 54  8.44 & 57 23 37.6 & 1.3 &     31 &    35.7 &  6.0 &     11 &     5.1 &  1.2 &     19 &    30.6 &  5.9 &      2 &    32.5 & 14.2 & -0.02 &  0.15 & -0.14 &  0.23 \\\\ 338 &  20 &XMMU J105410.3+573038 & 10 54 10.31 & 57 30 38.5 & 0.1 &   1005 &   172.4 &  5.8 &    779 &   125.9 &  2.4 &    211 &    46.5 &  5.3 &     16 &     2.8 &  5.9 & -0.89 &  0.01 & -0.92 &  0.16 \\\\ 2024 &     &XMMU J105410.9+573057 & 10 54 11.00 & 57 30 57.4 & 0.2 &   1115 &   154.3 &  6.8 &    875 &    56.9 &  1.9 &    224 &    97.3 &  6.6 &     17 &    58.2 & 10.6 & -0.57 &  0.03 & -0.41 &  0.08 \\\\ 2508 &     &XMMU J105412.1+573420 & 10 54 12.15 & 57 34 20.1 & 2.0 &    100 &     0.5 &  2.1 &      0 &     0.5 &  0.7 &      0 &     0.0 &  2.0 &      0 &     0.0 & 11.2 & -1.00 &  1.38 &  -$^a$ &  -$^a$ \\\\ 1379 &     &XMMU J105414.1+572744 & 10 54 14.14 & 57 27 44.9 & 1.6 &     29 &    36.5 &  5.8 &      4 &     2.8 &  1.0 &     26 &    33.7 &  5.7 &      1 &    21.7 & 11.5 &  0.31 &  0.17 & -0.38 &  0.24 \\\\ 429 &     &XMMU J105414.2+573330 & 10 54 14.23 & 57 33 30.2 & 0.6 &    163 &    54.4 &  6.0 &    137 &    17.6 &  1.4 &     28 &    36.8 &  5.8 &      0 &    13.0 & 11.6 & -0.50 &  0.07 & -0.60 &  0.29 \\\\ 218 &     &XMMU J105414.3+572635 & 10 54 14.33 & 57 26 35.4 & 1.4 &     23 &    29.6 &  5.5 &     15 &     6.3 &  1.2 &     10 &    23.2 &  5.4 &      0 &     1.5 &  6.7 & -0.27 &  0.14 & -0.91 &  0.38 \\\\ 185 &     &XMMU J105415.2+572502 & 10 54 15.22 & 57 25  2.2 & 1.0 &     87 &    51.9 &  7.1 &     72 &    16.4 &  1.7 &     17 &    35.5 &  6.9 &      0 &     8.3 & 11.0 & -0.49 &  0.08 & -0.72 &  0.33 \\\\ 230 &     &XMMU J105415.9+572656 & 10 54 15.97 & 57 26 56.3 & 0.8 &     63 &    35.6 &  5.9 &     52 &    11.3 &  1.3 &     10 &    24.2 &  5.7 &      2 &    29.7 & 12.5 & -0.49 &  0.10 & -0.07 &  0.24 \\\\ 167 &     &XMMU J105417.8+572423 & 10 54 17.81 & 57 24 23.9 & 1.1 &     31 &    44.2 &  7.1 &      2 &     3.5 &  1.4 &     23 &    40.7 &  7.0 &      6 &    57.3 & 17.0 &  0.29 &  0.20 & -0.00 &  0.17 \\\\ 171 &  28 &XMMU J105421.1+572544 & 10 54 21.14 & 57 25 44.8 & 0.1 & 215082 &   10757 & 62.8 & 152426 &    3234 & 15.4 &  52892 &    7524 & 60.9 &   9765 &    5139 & 76.9 & -0.46 &  0.00 & -0.35 &  0.01 \\\\ 1403 &     &XMMU J105423.0+572725 & 10 54 23.06 & 57 27 25.1 & 2.0 &     22 &    31.4 &  6.6 &     13 &     5.4 &  1.3 &     13 &    26.0 &  6.5 &      2 &    21.8 & 13.3 & -0.13 &  0.17 & -0.26 &  0.31 \\\\ 154 &     &XMMU J105426.1+572405 & 10 54 26.18 & 57 24  5.8 & 0.9 &    114 &    63.3 & 10.4 &    112 &    30.3 &  2.8 &      5 &    33.0 & 10.0 &      0 &     0.0 & 16.6 & -0.71 &  0.08 & -1.00 &  0.71 \\\\ 1367 &     &XMMU J105427.4+572712 & 10 54 27.49 & 57 27 12.3 & 1.7 &     25 &    33.4 &  7.2 &     18 &     8.3 &  1.6 &      7 &    25.1 &  7.0 &      0 &    24.5 & 16.0 & -0.35 &  0.15 & -0.18 &  0.34 \\\\ 281 &     &XMMU J105432.5+572841 & 10 54 32.52 & 57 28 41.4 & 0.7 &    224 &   101.3 & 11.4 &    201 &    39.3 &  2.9 &     25 &    62.1 & 11.0 &      0 &    20.7 & 17.3 & -0.60 &  0.06 & -0.62 &  0.26 \\\\ \\hline \\end{tabular}} \\begin{list}{}{} \\item[$^a$] hardness ratio undefined: fluxes in both constituting bands are zero \\end{list} \\vfill \\end{minipage} \\end{sidewaystable*}  \\begin{table*} \\caption{Catalogue of extended X-ray sources with likelihood of extent greater than 15.0 (5$\\sigma$). Description of columns: (1) XMM-Newton source number as displayed in Fig. 3, (2) ROSAT source number (Lehmann et al. \\cite{lehmann}), (3) IAU source name, (4--6) J2000 coordinates and errors, (7--8) likelihood of detection and likelihood of extent (rounded to nearest integer), (9) source extent (core radius of King profile), (10--11) fluxes and flux errors [$10^{-16}$ erg cm$^{-2}$ s$^{-1}$] in 0.5--10~keV band, (12--15) hardness ratios and errors. All errors are 1$\\sigma$ errors.} \\label{souCatExt} \\centering {\\scriptsize \\begin{tabular}{rrcllcrrrrrrrrr} \\hline \\multicolumn{2}{c|}{SRC \\#}&\\multicolumn{1}{c|}{}&\\multicolumn{2}{c}{RA~~~(J2000)~~~DEC}&err&\\multicolumn{2}{|c|}{$\\mathcal{L}$}&\\multicolumn{1}{c}{extent}&\\multicolumn{2}{|c|}{0.5--10.0~keV}& \\\\ XMM$^a$&R&\\multicolumn{1}{|c|}{IAU name}&hh mm ss&$^{\\circ}$~~~~$^\\prime$~~~$^{\\prime\\prime}$&$^{\\prime\\prime}$&\\multicolumn{1}{|r}{detection}&\\multicolumn{1}{r|}{extent}&\\multicolumn{1}{c|}{``}&flux&err&\\multicolumn{1}{|r}{HR1}&err&HR2&err \\\\ \\hline 395   &&XMMU J105151.6+573225& 10 51 51.62 & 57 32 25.0 & 2.3 &  134 &  62 & 12.3 & 127.7 &  40.6 & -0.75 &  0.15 & -0.55 &  0.76 \\\\ 355   &&XMMU J105237.3+573104& 10 52 37.30 & 57 31 04.5 & 0.3 & 1761 &  97 &  1.9 & 122.3 &  46.7 & -0.53 &  0.02 & -0.42 &  0.06 \\\\ 2378  &&XMMU J105242.1+573237& 10 52 42.15 & 57 32 37.0 & 2.2 &   76 &  33 &  6.3 &  25.8 &   6.2 & -0.94 &  0.08 & -1.00 &  2.36 \\\\ 2512- &840&XMMU J105246.1+574044& 10 52 46.18 & 57 40 44.8 & 1.5 &  376 & 134 &  9.7 & 217.4 &  22.8 & -0.80 &  0.05 & -0.56 &  0.46 \\\\ 2515  \\\\ 476   &827&XMMU J105303.4+573529& 10 53  3.48 & 57 35 29.9 & 0.6 &  431 &  27 &  1.5 &  66.4 &   4.9 & -0.50 &  0.04 & -0.24 &  0.11 \\\\ 69   & 41&XMMU J105318.8+572045& 10 53 18.87 & 57 20 45.1 & 0.7 & 1560 & 479 &  7.0 & 283.6 &  14.0 & -0.78 &  0.02 & -0.67 &  0.25 \\\\ 2394  &&XMMU J105319.9+573536& 10 53 19.96 & 57 35 36.2 & 5.7 &   38 &  27 & 19.3 & 164.8 &  71.6 & -0.91 &  0.16 & -0.08 &  1.44 \\\\ 409   &&XMMU J105336.8+573257& 10 53 36.89 & 57 32 57.4 & 1.5 &  125 &  41 &  4.1 &  26.2 &   6.0 & -0.91 &  0.08 &  0.55 &  0.38 \\\\ 472   &228&XMMU J105339.9+573522& 10 53 39.94 & 57 35 22.6 & 0.5 & 2408 & 565 &  6.3 & 414.9 &  14.8 & -0.64 &  0.02 & -0.51 &  0.09 \\\\ 144   &131&XMMU J105340.1+572352& 10 53 40.10 & 57 23 52.2 & 0.6 &  565 &  39 &  1.8 &  41.5 &   4.4 & -0.94 &  0.04 &  0.16 &  0.49 \\\\ 58   &&XMMU J105341.9+575525& 10 53 41.98 & 57 19 36.8 & 0.9 &  235 &  18 &  1.9 &  83.3 &   7.9 & -0.56 &  0.06 & -0.29 &  0.20 \\\\ 461   &229&XMMU J105346.5+573509& 10 53 46.54 & 57 35 09.2 & 0.7 & 1463 & 534 & 12.2 & 759.0 &  34.7 & -0.70 &  0.02 & -0.88 &  0.16 \\\\ 185   &&XMMU J105416.2+572458& 10 54 16.24 & 57 24 58.4 & 1.8 &  131 &  31 &  6.0 & 180.9 &  19.3 & -0.50 &  0.06 & -0.81 &  0.28 \\\\ \\hline \\end{tabular}} \\begin{list}{}{} \\item[$^a$]source number of corresponding object from point source detection list (Table~\\ref{souCat}), except for object 2394 not previously detected as a point source. \\end{list} \\end{table*}"
    },
    "0711/0711.2290_arXiv.txt": {
        "abstract": "We report on optical and infrared observations of the anomalous X-ray pulsar (AXP) \\axp, made during its ongoing X-ray flare which started in 2007 March. We detected the source in the optical $I$ and near-infrared $K_s$ bands in two ground-based observations and obtained deep flux upper limits from four observations, including one with the \\textit{Spitzer Space Telescope} at 4.5 and 8.0 $\\mu$m. The detections indicate that the source was approximately 1.3--1.6 magnitudes brighter than in 2003--2006, when it was at the tail of a previous similar X-ray flare. Similar related flux variations have been seen in two other AXPs during their X-ray outbursts, suggesting common behavior for large X-ray flux variation events in AXPs. The \\spz\\ flux limits are sufficiently deep that we can exclude mid-infrared emission similar to that from the AXP \\axpu, which has been interpreted as arising from a dust disk around the AXP. The optical/near-infrared emission from \\axp\\ probably has a magnetospheric origin.  The similarity in the flux spectra of \\axp\\ and \\axpu\\ challenges the dust disk model proposed for the latter. ",
        "introduction": "Supported by extensive observational studies over the past few years, it is generally believed that anomalous X-ray pulsars (AXPs) are magnetars---young neutron stars possessing ultra-high $\\sim$10$^{14}$~G magnetic fields.  While AXPs are primarily known as X-ray sources, exhibiting a variety of variability behavior related to their magnetar nature \\citep{wt06,kas07}, we now know that 5 of 10 identified AXPs---4U~0142+61, \\axp, 1RXS~J170849$-$400910, XTE~J1810$-$197, and 1E~2259+586---are bright at near-infrared (NIR) wavelengths (\\citealt{wt06}; McGill AXP online catalog\\footnote{\\url{www.physics.mcgill.ca/pulsar/magnetar/main.html}}).  Among them, the AXP 4U 0142+61 is detected from optical to mid-infrared (MIR) wavelengths, indicating a spectral energy distribution (SED) that can be described by a two-component model: one a power-law spectrum over optical $V\\/RI$ and NIR $J$ bands, presumably arising from the magnetosphere, and one thermal blackbody-like over the 2.2--8 $\\mu$m range \\citep{wck06}, arising from a debris disk. The discovery of such an optical/IR SED from the AXP was unexpected, and the two-component model remains controversial (e.g., \\citealt{dv06d}). Ideally, in order to understand the optical/IR emission mechanism for AXPs and in particular determine whether the IR emission from 4U 0142+61 is unusual or generic among AXPs, SEDs of other AXPs must be observed. However, most known AXPs are highly extincted in the optical range and often extremely faint in the IR.  Among the known AXPs, \\axp\\ is peculiar in that it has exhibited two long-term X-ray flares \\citep{gk04, tam+08} which have not been seen from other AXPs thus far. Monitored with the \\textit{Rossi X-ray Timing Explorer (RXTE)}, the first flare was found to start in 2002 April and lasted approximately two years. At the beginning of the flare, the NIR counterpart was discovered \\citep{wc02}, and in observations more than a year later, the counterpart was found to be approximately 2 magnitudes dimmer in $JK_s$ \\citep{dv05}, suggesting large flux variations associated with the flare. Such flaring events thus provide an opportunity to study optical/IR emission from an AXP other than 4U 0142+61. In addition, during their X-ray outbursts, the AXPs 1E~2259+586 and XTE J1810$-$197 both exhibited related NIR brightening. In the first source, its NIR and X-ray flux changes were correlated \\citep{tam+04} while in the latter, its NIR flux did not exactly follow the decline of the X-ray flux (\\citealt{rea+04}; \\citealt{cam+07}). Optical/IR observations of \\axp\\ during another X-ray flare would allow us to verify if the source exhibits similar emission behavior.  For these reasons, after a second flare started on 2007 March 24 \\citep{dib+07}, we made multi-wavelength observations of \\axp\\  with large ground-based telescopes and the \\textit{Spitzer Space Telescope}. In this paper, we report on our optical/IR observations during this on-going X-ray flare. ",
        "conclusions": "\\label{sec:disc}  The optical/NIR counterpart of \\axp\\  was observed a few times during 2003--2006 after the discovery of the counterpart, and the source's brightness had been low (e.g., $K_s\\simeq 21.0$--21.5; see \\citealt{tam+08} for details). Compared to the flux measurements, particularly those obtained in 2003 April and June when the counterpart was detected in both the optical $I$ and NIR $JK_s$ bands (\\citealt{dv05}; also see Table~\\ref{tab:mea}), our results clearly indicate an optical/NIR brightening during the 2007 X-ray flare.  Indeed, assuming $A_V= 5.4$ mag, the unabsorbed optical/NIR--to--X-ray flux ratios for the 2003 and our measurements are nearly identical (see Figure~\\ref{fig:comp}). In this comparison, the unabsorbed 2--10 keV X-ray fluxes at two epochs given in Table~\\ref{tab:mea} are used, and $A_V$ is estimated from $N_{\\rm H}=9.7\\times 10^{21}$ cm$^{-2}$ (\\citealt{tam+08}) by using the relation $A_V= N_{\\rm H}/1.79\\times 10^{21}$ cm$^{-2}$ \\citep{ps95}.  Also we note that even though our second $K_s$ measurement was an upper limit ($K_s>20.1$), it was only 0.1~mag smaller than that obtained from a detection two days later by \\citet{isr+07}. Thus in these two sets of observations, the optical/NIR and X-ray fluxes appeared to be correlated.  \\begin{center} \\includegraphics[scale=0.46]{f1.eps} \\figcaption{Unabsorbed optical/IR flux and flux upper limits (diamonds) of \\axp\\ measured in 2007 April--August, normalized by the unabsorbed 2--10 keV X-ray flux obtained on 2007 April 28 (see Table~\\ref{tab:mea}). The dotted line indicates our two nearly simultaneous observations in the $I$ and $K_s$ bands, and the dashed line indicates the optical and NIR measurements obtained in 2003 April and June \\citep{dv05}. For comparison, the unabsorbed optical/IR--to--X-ray flux ratios for \\axpu\\ are also shown (squares; \\citealt{dv06d,wk08}), with the vertical and horizontal lengths of the bars indicating the flux variation range and effective bandwidth at each band, respectively. $A_V=3.5$ \\citep{dv06a} is used for dereddening. The unabsorbed 2--10 keV X-ray flux used for \\axpu\\ was 6.6$\\times 10^{-11}$ ergs s$^{-1}$ cm$^{-2}$ \\citep{gon+07}. \\label{fig:comp} } \\end{center}  However, as shown by \\citet{tam+08}, the observed NIR brightening at the beginning of the 2002 X-ray flare preceded the pulsed X-ray flux peak by $\\sim$2 months. If we assume that the relation between the pulsed flux $F_{\\rm p}$ and total flux $F_{\\rm t}$, $F_{\\rm p} \\propto F_{\\rm t}^{0.54}$ \\citep{tam+08}, is generally true for this AXP, the 2--10 keV unabsorbed total flux derived at the NIR observation time would have been $\\simeq$$10^{-11}$ ergs s$^{-1}$ cm$^{-2}$, 3.6 times lower than the 2007 peak flux.  As a result, the 2002 NIR--to--X-ray flux ratios would be approximately 6 times higher than that shown in Figure~\\ref{fig:comp} (note $K_s=19.4$ in the 2002 detection). This would indicate that the NIR and X-ray fluxes were not correlated all the time. We note that since the $F_{\\rm p}$ and $F_{\\rm t}$ relation is mainly based on the total flux measurements obtained during the 2007 flare and there were not such measurements for the 2002 flare, it is possible that the total flux was much larger at the start of the 2002 flare, and not reflected in the \\textit{RXTE} pulsed flux measurements. In any case, we confirm that in this AXP, the large optical/NIR flux variations were related to the X-ray flares. Similar related NIR brightening has been seen in 1E~2259$+$586 and XTE~J1810$-$197 during their X-ray outbursts \\citep{tam+04,cam+07}, suggesting common behavior for emission from AXPs. However for \\axpu, while its X-ray flux has been relatively stable \\citep{gon+07}, its optical and NIR flux was found to have large, rapid variations \\citep{dv06d}. This limits the related NIR/X-ray behavior to only large X-ray flux variation events.  It is instructive to compare \\axp\\ to \\axpu, since the latter has well measured fluxes from the optical to MIR. In Figure~\\ref{fig:comp}, we compare the optical/IR fluxes and flux upper limits of \\axp\\ to those of \\axpu, with each set of data points normalized by the approximately contemporaneous unabsorbed 2--10 keV X-ray flux of each source. The optical $V\\/RI$ and NIR $JHK_s$ fluxes for \\axpu\\ are from Durant \\& van Kerkwijk (2006b; because of relatively large flux variations from the source, we show the range of the flux found in each band in the figure), the two MIR data points at 4.5 and 8.0 $\\mu$m are the 2005 flux measurements in Wang \\& Kaspi (2008), and the unabsorbed 2--10 keV X-ray flux used is 6.6$\\times 10^{-11}$ ergs s$^{-1}$ cm$^{-2}$, obtained on 2004 July 24 \\citep{gon+07}. As can be seen from the figure, while the flux ratios of \\axp\\ are roughly two times larger than those of \\axpu, the deep limit at 4.5 $\\mu$m from our \\spz\\ IRAC observation is approximately 3 times lower than the detection of \\axpu. This strongly suggests that our target does not have similar MIR emission. This difference in the MIR is not very sensitive to uncertainties on the reddening to the two sources, since MIR emission is only weakly extincted by the interstellar medium. The X-ray flux from \\axpu\\ has been stable \\citep{gon+07}. A possible explanation could be the non-simultaneous X-ray flux of \\axp\\ we use for the comparison. However, \\textit{RXTE} X-ray monitoring observations of the source have indicated that the pulsed flux, which is found to trace the total flux \\citep{tam+08}, has remained stable and high during our observations (Dib et al. 2008, in preparation).  In order to raise the 4.5 $\\mu$m upper limit to the flux ratio of \\axpu, the X-ray flux would have had to have been as low as 11$\\times 10^{-12}$ ergs s$^{-1}$ cm$^{-2}$, only 40\\% larger than the average quiescent flux (7.7$\\times 10^{-12}$ ergs s$^{-1}$ cm$^{-2}$) of the source and approximately one third of those obtained in 2007 April \\citep{tam+08}.  The lack of MIR emission similar to that of \\axpu\\  from \\axp\\ seems robust. In the dust disk model proposed for \\axpu\\ \\citep{wck06}, the MIR flux, arising from X-ray irradiation of the disk, is proportional to the X-ray flux of the pulsar. Therefore, even though our deep 4.5 $\\mu$m limit was obtained during the X-ray flare, we can exclude the existence of a similar disk around \\axp. However, since a disk could be further away from the central source, deep observations at longer wavelengths are needed in order to exclude the existence of a disk more conclusively. The MIR non-detection may suggest that debris disks are not commonly found around AXPs and thus the putative disk in \\axpu\\ is unique. Two other AXPs have also been observed by \\spz\\  with no detection; however, the derived upper limits are far above the MID--to--X-ray flux ratio of \\axpu\\ \\citep{wkh07}.  The optical and NIR emission from \\axp\\ probably has a magnetospheric origin, given that the flux spectrum is similar to that of \\axpu\\ (\\citealt{wc02,dv05}). For the latter source, its optical emission is known to be pulsed at the spin period and has a pulsed fraction much higher than that in X-rays, excluding a disk origin \\citep{km02,dhi+05}. Details of how optical and NIR emission is produced in the magnetosphere of a magnetar are not known, though possible radiation mechanisms have been suggested (e.g., \\citealt{bt07}). We note that in the disk model for \\axpu\\ \\citep{wck06}, the $K$-band flux primarily arises from the disk, not from the magnetosphere. Therefore, the similarity in the $K$-band--to--X-ray flux ratios for the AXPs \\citep{dv05}, including \\axp\\ as confirmed by our measurements, challenges the disk model. The $K$-band emission thus should originate from the magnetosphere, which may be supported by the fact that in the absence of X-ray variability, strong $K$-band variability is seen in \\axpu\\ \\citep{dv06d}."
    },
    "0711/0711.0295_arXiv.txt": {
        "abstract": "{} {We investigate the reconstruction of pulsar magnetospheres close to the light cylinder surface to study the curvature drift instability (CDI) responsible for the twisting of magnetic field lines in the mentioned zone. The influence of plasma dynamics on the saturation process of the CDI is studied.} {On the basis of the Euler, continuity, and induction equations, we derive the increment of the CDI and analyze parametrically excited drift modes. The dynamics of the reconstruction of the pulsar magnetosphere is studied analytically.} {We show that there is a possibility of a parametrically excited rotational-energy pumping-process in the drift modes. It is indicated by the generation of a toroidal component of the magnetic field that transforms the field lines into such a configuration, in which plasma particles do not experience any forces. At this stage, the instability process saturates and the further amplification of the toroidal component to the magnetic field lines is suspended.} {} ",
        "introduction": "Exploring the magnetic field of the Crab nebula, \\cite{pid} was the first to discover the presence of a central object in the nebula, with frozen magnetic field inside. It was assumed that the rotation of a central object provokes the generation of the toroidal component of magnetic field. Investigations have also shown that this kind of magnetic field characterizes magnetized star outflows (\\cite{weber}). Pulsars are one of the most distinctive examples of rotationally powered magnetized stars producing the relativistic outflows commonly known as pulsar winds.  One of the major problems for pulsar winds concerns the transition of the magnetized plasma flows by means of the so-called \"light cylinder\" surface, a hypothetical zone, where the linear velocity of rotation equals the speed of light. According to the model of pulsar magnetosphere (\\cite{gj}), the relativistic plasma flow, emanating from the pulsar surface streams along very strong magnetic field lines. The typical values of the magnetic field are of the order of $10^{12}$ G (\\cite{manch}). Such a strong magnetic field forces the particles, streaming toward the light cylinder zone, to co-rotate with a pulsar. On the other hand, once the relativistic effect of the mass increment is taken into account, the radial acceleration of the particles appears to be limited. This means that the particles will never cross the light cylinder surface if the rigid rotation is preserved (\\cite{mr}). Therefore, the co-rotation of the magnetized plasma cannot be maintained close to the mentioned zone and a further exploration of this phenomenon is needed.  The simplest means of explaining the transition in the plasma outflow at the light cylinder surface is to constrain the motion of the flow to become asymptotically close to a regime, in which particles do not experience any forces (see e.g., \\cite{mmsh2000,r03}). For this purpose, if the magnetic field is still strong, one needs to twist the field lines in an appropriate way. In other words, it is necessary to generate the toroidal component of the magnetic field. The pulsars maintain the dipolar magnetic field with a certain curvature of the field lines. The particles moving along the curved field lines drift perpendicularly to the curvature plane. The existence of the drift motion creates the necessary conditions for developing a CDI (see e.g., \\cite{kmm89,kmm91a,kmm91b,shmmkh03}). The mechanism excites transverse drift modes providing the magnetic field with the toroidal component. The CDI becomes effective in the vicinity of the light cylinder surface, where the pulsar's dipolar field is comparable to the toroidal component of the excited mode. The kinetic energy of the primary beam particles, of Goldreich-Julian (GJ) density, can indeed supply the generation of the drift waves. Mostly, it is assumed that the toroidal component of the magnetic field is created by the GJ current. However, the kinetic energy density, ${W}_b$, of the primary beam is insufficient to change the configuration of the dipolar magnetic field, $B_0$, significantly, since ${W}_b<<B_0^2/(4\\pi)$. Therefore, the process of generating the toroidal component, $B_r$ (of the order of $B_0$), requires an additional energy supply, which can be extracted from the pulsar rotational energy. We have already shown that the rotational energy pumping into the drift modes can be implemented by the \"parametric\" instability (\\cite{mnras}). This mechanism is called parametric, because the effect is caused by the relativistic centrifugal force, which, as a parameter, changes in time and induces the instability. The aim of the present paper is to study the saturation process of this instability close to the light cylinder zone.  The problem analysis is completed well if one considers the mechanical analog of the motion of plasma flow along the pulsar magnetic field lines and studies the dynamics of a single particle moving inside the rotating channels. In the present paper, we apply the method developed in (\\cite{r03}) and study the plasma process that converts the configuration of the magnetic field lines to the Archimedes' spiral: it appears that in order to suppress the reaction force, the field lines, rather being straight, should deviate back and lag behind the rotation; consequently, the curvature will eventually increase, inevitably leading to the decrement of the reaction force; this process will last until the magnetic field lines form the shape of the Archimedes' spiral, which removes the reaction force and saturates the instability.  The paper is arranged as follows. In section 2, we examine the parametric curvature drift instability and derive an expression for its corresponding growth rate. In section 3, we consider the kinematics of a particle moving along the rotating magnetic field line. In section 4, we apply the method for 1-second pulsars and the Crab pulsar, and in section 5 we summarize our results.    \\begin{figure} \\par\\noindent {\\begin{minipage}{0.35\\linewidth} \\includegraphics[width=\\textwidth] {12113fg3.ps} \\end{minipage} } \\hfill {\\begin{minipage}{0.35\\linewidth} \\includegraphics[width=\\textwidth] {12113fg4.ps} \\end{minipage} } \\caption[ ] {\\footnotesize $a$) The geometry in which the set of main Eqs. (\\ref{eul}-\\ref{ind}) is considered; (${\\bf e}_{\\theta},{\\bf e}_r,{\\bf e}_x$) denotes orthonormal basis of unit vectors; ${\\bf e}_x$ is directed perpendicularly to the plane of the figure; $C$ is the center of the curvature; and ${\\bf e}_{\\theta}$ and ${\\bf e}_r$ are the tangential and perpendicular (with respect to the field line) unit vectors respectively. $b$) The geometry for deriving Eq. (\\ref{time}): the curved line denotes the twisted magnetic field, ${\\bf B}$, generated due to the raising of magnetic perturbation, ${\\bf B}_r$. Note, that ${\\bf B}_r$ and ${\\bf B}_{\\theta}$ are oriented with respect to the initial `quasi straight' magnetic field line. }\\label{fig3} \\end{figure} ",
        "conclusions": "\\label{sec:discus}  We consider Eq. (\\ref{time}) and plot the timescale of the instability versus the wavelength, $\\lambda_x$, for several values of initial toroidal perturbations $B^0_r$. In Fig. \\ref{fig4}, we display the behavior of $T(\\lambda_x)$, when $\\tan\\chi=1$, for several values of initial perturbation $B^0_r/B_{\\theta}\\equiv\\delta\\in \\{10^{-1}; 10^{-3}; 10^{-5}; 10^{-7}\\}$. Two major applications are examined: (a) {\\it 1-second} pulsar and (b) the Crab pulsar. From Eq. (\\ref{time}), we can infer that the timescale is a continuously decreasing function of the initial perturbation, $B^0_r$: as becomes smaller the $\\delta$ parameter, so does the initial perturbation and, consequently, the magnetic field lines need more time to achieve the required structure. As observations show, the ratio $P/\\dot{P}$ ranges from $10^{11}s$ (PSR 0531 - Crab pulsar) to $10^{18}s$ (PSR 1952+29). However, the greatest values of the twisting timescales, shown in Fig. \\ref{fig4}, vary between $10^4s$ and $10^2s$, which are shorter by many orders of magnitude than $P/\\dot{P}$, illustrating the high efficiency of the instability.  The reconstruction of the magnetic field requires a certain amount of energy, therefore it is essential to estimate pulsar's slowdown luminosity ($L_p$) and compare it to the, so-called, \"magnetic luminosity\" ($L_m\\equiv \\triangle E_m/T$, where $\\triangle E_m$ is the variation in the magnetic field energy due to the instability). For $L_p$, one has \\begin{figure} \\par\\noindent {\\begin{minipage}{0.48\\linewidth} \\includegraphics[width=\\textwidth] {12113fg5.ps} \\end{minipage} } \\hfill {\\begin{minipage}{0.48\\linewidth} \\includegraphics[width=\\textwidth] {12113fg6.ps} \\end{minipage} } \\caption[ ] {\\footnotesize Here we show the behaviour of the timescale of sweepback of the magnetic field lines versus the wave length, $\\lambda\\approx\\lambda_x$, for the {\\it 1-second} pulsar [see (a)] and for the Crab pulsar, $P\\approx 0.033s$ [see (b)], respectively. The set of parameters is: $\\delta = \\{10^{-1}; 10^{-3}; 10^{-5}; 10^{-7}\\}$, $\\gamma_b\\sim 10^6$ and $\\lambda_{\\theta} = 1000R_{c}$.}\\label{fig4} \\end{figure}  \\begin{equation} \\label{rot1} L_{p} = I\\Omega\\dot{\\Omega} = I\\frac{4\\pi^2}{P^2}\\frac{\\dot{P}}{P}, \\end{equation} where $I\\sim MR_p^2$ is the moment of inertia of the pulsar, and $M\\sim M_{\\odot}\\approx 2\\times 10^{33}g$ and $R_p\\sim 10^6cm$ are the pulsar's mass and radius, respectively. For typical pulsars, $\\dot{P}/P\\sim 10^{-15}s^{-1}$, the slowdown luminosity is estimated to be \\begin{equation} \\label{rot2} L_{p}\\approx 7.9\\times 10^{31}erg/s. \\end{equation} The \"magnetic luminosity\", $L_m$, can be estimated straightforwardly \\begin{equation} \\label{mag1} L_{m} = \\frac{B_{r}^2}{4\\pi T}\\triangle V, \\end{equation} where $\\triangle V\\sim R_c^2\\triangle r=R_c^3 \\varepsilon$ ($\\varepsilon\\equiv\\triangle r/R_c<<1$) is the volume, where the process of twisting takes place. Since we are studying the instability in the nearby zone of the LC, the magnetic field components must be given as $B^0_r\\sim \\delta\\cdot B_{\\theta}=\\delta\\cdot B_{p}(R_p/R_c)^3$, where $B_{p}\\sim 10^{12}G$ is the magnetic field near the pulsar's surface. Substituting all quantities into Eq. (\\ref{mag1}) corresponding to the energy gain for $\\lambda_x\\sim 2.5\\times 10^{11}cm$, and $T\\approx 10^3s$, (see Fig. \\ref{fig4}a), and bearing in mind Eqs. (\\ref{br},\\ref{time}), one can show that for $\\varepsilon\\sim 0.1$ and $\\delta\\sim 0.1$ the value of the \"magnetic luminosity\" is approximately equal to  \\begin{equation} \\label{mag2} L_{m} \\approx 7.3\\times 10^{25}erg/s. \\end{equation} A direct comparison between Eqs. (\\ref{rot2}) and (\\ref{mag2}) infers that $L_{m}<<L_{p}$, meaning that approximately only $0.0001\\%$ of the total energy budget goes to the reconstruction of the magnetic field lines.  In Fig. \\ref{fig4}b, we show the behavior in $T(\\lambda_x)$ for the Crab pulsar in the different cases of initial magnetic perturbations. The dependence does not change qualitatively but the twisting process changes quantitatively as the corresponding timescale is now of the order of $\\sim 10\\d-10^2\\,s$. The \"magnetic luminosity\" of the Crab pulsar for $\\delta = 0.1$, $\\lambda_x = 7\\times 10^6cm$, $T\\approx 4\\,s$ (see Fig. \\ref{fig4}b), and $\\varepsilon = 0.1$ approximately equals [see Eq. (\\ref{mag1})] \\begin{equation} \\label{cmag2} L_{m}^{Crab} \\approx 4.7\\times 10^{32}erg/s. \\end{equation} On the other hand, a direct calculation of the Crab pulsar ($P\\approx 0.033\\,s$) luminosity, estimated by Eq. (\\ref{rot1}), yields \\begin{equation} \\label{crot2} L_{p}^{Crab}\\approx 9.3\\times 10^{38}erg/s. \\end{equation} Therefore, as in the case of the Crab pulsar, the energy required to reconstruct the magnetosphere averages $\\sim 0.00005\\%$ of the pulsar's energy and the twisting process becomes feasible in this case as well.  The sweepback mechanism described in this paper, therefore, appears to be extremely efficient for pulsar magnetospheres. The curvature drift instability may lead to the reconstruction of the magnetic field lines in such a way that the dynamics of magnetosphere becomes force-free, which in its turn, completely disrupts the instability, thus saturating the CDI.     We have analytically examined the non-stationary pattern of the pulsar magnetosphere near the LC zone. The GJ current can be estimated by $J_{GJ} = en_b c$, giving the value of the corresponding magnetic field, $B_r\\approx 4\\pi J_{GJ}R_n/c$, where $R_n\\approx B(R)/B'(R) = R/3$ (where $B'\\equiv dB/dR$) is the length scale of the spatial inhomogeneity of the magnetic field. If we assume a dipolar configuration, then, taking the value of the GJ density, $n_b\\approx \\Omega B/(2\\pi ec)$, into account, one can show that the toroidal magnetic field equals $\\frac{2R}{3R_c}B$. Inside the light cylinder ($R< R_c$), this value is less than the required one-$B$ and reaches its maximum value, $\\frac{2}{3}B$, on the light surface. Therefore, the GJ current cannot significantly change the configuration of the magnetic field. This current evidently exceeds the curvature drift current, because $J_{GJ}/J_{cur}\\approx c/u_b\\gg 1$. However, the drift current, $J_{cur}$, is not the source of the toroidal component, $B_r$, of the magnetic field, but it is a trigger mechanism for generation of the perturbed current, $J_1= e(n_b^0v^1_{b_x}+n_b^1u_b)$  responsible for the creation of $B_r$ (see Eq. (\\ref{indp})). The source of the instability of the current and the resulting magnetic field is the pulsar's rotational energy. We have found that the instability is achieved by the parametric mechanism, which effectively pumps the pulsar rotational energy directly into the generated mode. The process lasts until the plasma dynamics reaches the force-free regime of motion and the overall magnetosphere relaxes to the steady state configuration. \\cite{buc} considered the dynamics of rotating pulsar winds by performing the numerical solution of relativistic magnetohydrodynamic (RMHD) equations in the Schwarzschild metric. The system was allowed to relax to a steady state configuration rapidly approaching the force-free regime. We suppose that the full set of RMHD equations comprises the terms applicable to CDI. Therefore, we highlight the numerical results obtained by \\cite{buc} as confirmation of our theory. On the other hand, the aim of the present work was to explain the saturation of the CDI in terms of the generation of currents, which makes the physics of the process more transparent.  The main aspects of the present work can be summarized as follows:  \\begin{enumerate} \\item Examining the pulsar magnetospheric relativistic plasma, we have studied the role of the parametrically excited CDI in the process of sweepback of magnetic field lines and the saturation process of the instability. The present parametric instability is based on the method developed by Silin (1973), but differs from this in principal, since we study an alternating centrifugal force instead of an alternating electric force.  \\item The linear analysis of the Euler, continuity, and induction equations yields the dispersion relation governing the CDI.  \\item Considering the resonance frequencies of the sweepback process, an expression of the instability increment has been obtained.  \\item On the basis of the expression of the instability growth rate, we have derived the formula of the transition timescale of quasi-linear configuration of field lines into the Archimedes' spiral. The particles' motion is force-free along these magnetic field lines, leading to the saturation of the instability.   \\item The transition timescale has been compared to the wavelength for {\\it 1-second} pulsars and the Crab pulsar. For both cases it was shown that the corresponding timescale is shorter than pulsar's spin down rates by many orders of magnitude illustrating the high efficiency of the discussed process.    \\end{enumerate}"
    },
    "0711/0711.4130_arXiv.txt": {
        "abstract": "We present the results\\footnote{to appear in The Astrophysical Journal, July 10, 2008, v. 681} of the analysis of the first 9 months of data of the {\\it Swift} BAT survey of AGN in the $14-195$ keV band. Using archival X-ray data or follow-up {\\it Swift} XRT observations, we  have identified 129  (103 AGN) of 130 objects  detected at $|b|> 15\\degr$ and with significance $>4.8\\sigma$. One source remains  unidentified. These same X-ray data have allowed measurement of the X-ray properties of the objects. We fit a power law to the $\\log N - \\log S$ distribution, and find the slope to be $1.42\\pm0.14$. Characterizing  the differential luminosity function data as   a broken power  law,  we find a break luminosity $\\log L_*$(erg s$^{-1}$)$=43.85\\pm0.26         $, a low luminosity power law slope $a=0.84^{+0.16}_{-0.22}$, and a high luminosity power law slope  $b=2.55^{+0.43}_{-0.30}$, similar to the values  that have been reported  based on {\\it INTEGRAL} data. We obtain a mean photon index $1.98$ in the $14-195$ keV band, with an {\\it rms} spread of  0.27. Integration of our luminosity function gives a local volume density of AGN above $10^{41}$ erg s$^{-1}$ of $2.4\\times 10^{-3}$ Mpc$^{-3}$, which is about 10\\% of the total luminous local galaxy density above $M_*=-19.75$. We have obtained X-ray spectra from the literature and from {\\it Swift}  XRT follow-up observations. These show that the distribution of $\\log\\: n_H$   is essentially flat from $n_H=10^{20}$ cm$^{-2}$ to $10^{24}$ cm$^{-2}$, with 50\\% of the objects having column densities of less than $10^{22}$ cm$^{-2}$. BAT Seyfert galaxies have a median redshift of 0.03, a maximum $\\log$ luminosity of 45.1, and approximately half have $\\log n_H > 22$. ",
        "introduction": "It is now realized that most of the AGN in the Universe have  high column densities of absorbing material in our line of sight, which significantly changes their  apparent  properties across much of the electromagnetic spectrum. In many well studied objects this material significantly reduces the soft X-ray, optical, and UV signatures of an active nucleus essentially ``hiding'' the object. While it is commonly believed that extinction-corrected [OIII] can be used as an ``unbiased'' tracer of AGN activity (Risaliti et al. 1999), there is a large scatter between [OIII] and $2-10$ keV X-ray flux (Heckman et al. 2005) and between [OIII] and BAT flux (Mel\\'endez et al 2008).  We acknowledge that some Compton thick AGN are detected in [OIII] that cannot be detected in hard X-rays, but Compton thick AGN are outside the scope of this paper. Therefore,  surveys of AGN which rely primarily on rest frame optical and UV studies are very incomplete and have led to misleading results concerning the number, luminosity function, and evolution of active galaxies  (e.g.,                       Barger et al. 2005).   While the distribution of column densities is under intensive investigation, it is clear from both X-ray (Tozzi et al. 2006, Cappi et al. 2006) and IR data (Alonso-Herrero, et al. 2006) that  a large fraction of AGN  have column densities greater than $3\\times 10^{22}$ cm$^{-2}$ in the line of sight. Using the galactic reddening  law (Predehl \\& Schmitt 1995), this is equivalent to $A_V>13$, making the nuclei essentially invisible in the optical and UV bands. This effect seems to dominate the population seen in  deep X-ray surveys (e.g., Barger et al. 2005, Brandt \\& Hasinger 2005) where a large fraction of the X-ray selected objects do not have optical counterparts with classical AGN signatures.   There are only two spectral bands in which the nuclear emission is strong and  where, provided the column densities are less than $1.5\\times10^{24}$ cm$^{-2}$ (Compton-thin objects),  this obscuring material is relatively optically thin. These bands, the hard X-ray ($E>20$ keV) and the IR ($5-50\\mu$m),  are optimal for unbiased searches for AGN (Treister et al. 2005). While recent results from {\\it Spitzer} are finding many AGN via their IR emission, IR selection is  hampered by several effects (Barmby et al. 2006, Weedman et al. 2006, Franceschini et al. 2006): (1) the strong emission from star formation, (2) the lack of a unique ``IR color'' to  distinguish AGN from other luminous objects (Stern et al. 2005), and (3) the wide range in IR spectral parameters (Weedman et al. 2006). Thus,   while an IR  survey   yields  many objects,  it is very difficult to quantify its completeness and  how much of the IR luminosity of a particular galaxy is due to an  active nucleus. These  complications are not present in a hard X-ray survey  since at $E>20$ keV virtually all the radiation comes from the nucleus and selection effects are absent for Compton thin sources. Even for moderately Compton thick sources ($\\Lambda < 2.3$ is absorption $< 90$\\%), a hard X-ray survey has significant sensitivity, but without an absorption correction the luminosity will be underestimated. Essentially every object more luminous that $10^{42}$ erg s$^{-1}$ is an AGN.  A hard X-ray survey is thus unique in its ability to find all Compton thin AGN in a uniform, well-defined fashion,  and to determine their intrinsic luminosity. However, due to  the relative rarity of bright AGN    (even the {\\it ROSAT}  all sky survey has only $\\sim$1 src deg$^{-2}$ at its threshold $-$ Voges et al. [1999]), one needs  a very large solid angle survey to find the bright, easily studied objects.   With the recent  {\\it Chandra} and  {\\it XMM} data (e.g., Alexander et al. 2003, Giacconi et al. 2002, Yang et al. 2004; Mainieri et al. 2002, Szokoly et al. 2004, Zheng et al. 2004, Mainieri et al. 2005, Barger et al. 2001, 2003) there has been great progress in understanding the origin of the X-ray background and the evolution of AGN. It is now clear that much of the background at $E>8$ keV is not produced by the sources detected in the $2-8$ keV band (Worsley et al. 2005), and is likely to  come from a largely unobserved population of AGN with high column density and  low redshift $z<1$.   Thus the source of the bulk of the surface brightness of the X-ray background, which peaks at $E\\sim$30 keV (Gruber et al. 1999) is  uncertain. The measurement of the space density and evolution of this putative population of highly absorbed AGN and the derivation of the distribution of their column densities as a function of  luminosity and of redshift is crucial for modeling the X-ray background and the evolution of active galaxies. Progress in this area requires both a hard X-ray survey of sufficient sensitivity, angular resolution and solid angle coverage to find and identify large numbers of sources, {\\it  and} follow-up observations with softer X-ray measurements to obtain precise positions and detailed X-ray spectral properties.    Due to a lack of instrumentation  with sufficient angular resolution to permit  identification of unique counterparts in other wavelength bands and with sufficient solid angle and sensitivity (Krivonos et al. 2005)  to produce a large sample, there has been little progress in hard X-ray surveys  for over 25 years (e.g. Sazonov et al. 2005, 2007). This situation has been radically changed by the {\\it Swift} BAT survey (Markwardt et al. 2005)  and recent {\\it INTEGRAL} results (Beckmann et al. 2006b, Sazonov et al. 2007, Krivonos et al. 2005,  Bird et al. 2007) which have detected more than 100 hard X-ray selected AGN, thus providing the first unbiased sample of Compton thin AGN in the local Universe.       \\begin{figure} \\centering \\epsscale{1.15} \\plotone{f1.ps} \\vskip -2.cm \\figcaption{ Percentage of the sky covered as a function of limiting flux in erg cm$^{-2}$s$^{-1}$ (14$-$195 keV) and of effective exposure (upper scale). As only the sky $|b|>15\\degr$ is considered here, the maximum value is 74\\%. The corresponding curves as a function of limiting flux for the analyses of {\\it INTEGRAL} data by Beckmann et al. (2006b) and by Sazonov et al. (2007) are shown for comparison, the flux having been converted assuming a power law spectrum with index $-2$. \\label{sky_cov_plot}} \\smallskip \\end{figure}            In this paper we describe results from the first 9 months of the hard X-ray survey using the BAT instrument (Barthelmy et al. 2005) on the {\\it Swift} mission (Gehrels et al. 2005), concentrating on sources with $|b|> 15\\degr$.  Above this latitude limit, we have identified all but  one of the sources detected at  $>4.8\\sigma$  with optical counterparts using {\\it Swift} XRT and archival X-ray data. With these same data we have also obtained   X-ray spectra. With a median positional uncertainty of 1.7' and a sensitivity limit of a few times $10^{-11}$ erg cm$^{-2}$ s$^{-1}$ in the $14-195$ keV band, the BAT data are about 10 times more sensitive than the previous all-sky hard X-ray survey ({\\it HEAO 1} A-4: Levine et al. 1984) and the positions are accurate enough to  allow unique identifications of nearly all of the sources.  Spectra are characterized by a photon index $\\Gamma$, where $N(E)\\propto E^{-\\Gamma}$. Luminosities are calculated using $h_{70}=1, \\:\\Omega=0.3$. ",
        "conclusions": "\\subsection{Luminosity Function}  As shown above the low luminosity slope of the {\\gks luminosity function of} hard X-ray selected AGN is steeper than that of the $2-8$ keV function of Barger et al. 2005. We believe that this is  due to the high fraction of heavily absorbed objects at low BAT luminosities. Thus the contribution of low luminosity objects to the $10-100$ keV background is larger than originally calculated.   This is confirmed by the agreement of the slope of our luminosity function with the absorption corrected low luminosity slope of La Franca et al. (2005), which unlike Barger et al (2005) assumes an absorption that depends on luminosity. The break in the luminosity function is quite robust and thus is an intrinsic feature of the luminosity function and is not due to a spectral selection effect. Integration of our luminosity function gives a local volume density of $n(L_X>10^{41}$erg s$^{-1})= 2.4\\times 10^{-3}$ Mpc$^{-3}$, compared to a density of $0.02$ Mpc$^{-3}$  galaxies brighter than  $M_*=-19.75$ (Cross et al. 2001), and a local emissivity of 2.3$\\times 10^{39}$ erg  s$^{-1}$ Mpc$^{-3}$. The choice of  $M_*$ is that is the knee in the luminosity function and is the typical absolute magnitude for a galaxy. It is a simple way of estimating the galaxy density. The typical $J$ band absolute magnitude at the knee is $M*=-21.73$ (Cole et al 2001). The median BAT $J$ band absolute magnitude is $M=-23.8$ and only 3 BAT AGN have $M>-22$. Hence  $\\ga10$\\% of luminous galaxies in the local Universe are AGN with a hard X-ray luminosity $\\ga10^{41}$ erg s$^{-1}$. Because of the low median redshift of the sample, the BAT data are not sensitive to evolution in the luminosity function and  $V/V_{\\rm max} \\sim$0.5  is as expected.   \\subsection{Log N-Log S}  There have been numerous predictions of the hard X-ray $\\log N-\\log S$ (Treister et al. 2006, Gandhi \\& Fabian 2003) and our data allow a direct comparison of these models. We find that converting the observed BAT $\\log N-\\log S$ to the band predicted by these authors that we have good agreement with the predictions of Gandhi et al. (2004), but lie a factor of 2 lower than that predicted by Treister et al. (2006). Since each of these models makes different assumptions, our hard X-ray survey should be able to {\\gks determine which are valid.} %   \\subsection{The distribution of $n_H$}     In Figures \\ref{abs_frac} and \\ref{pl_index_distrib} the distribution of column densities over all objects is almost flat and appears to depend on hard X-ray luminosity. Similar results based on the {\\it RXTE} slew survey were obtained by Sazonov \\& Revnivtsev (2004). The standard unified model predicts that the ratio of absorbed to unabsorbed objects should be $4:1$, as opposed to our observed value of $1:1$. This difference is probably due to the neglect of the luminosity dependence of absorption in the simple unified model. The BAT results are roughly consistent with dependence of absorption on luminosity seen previously (Ueda et al. 2003, Steffen et al. 2003, Gilli et al. 2007). We note that the distribution of column densities in Tozzi et al. (2006) from the {\\it Chandra} deep fields is rather different from the BAT sample in that the Tozzi et al. sample seems to be missing the low $n_H$ half of the distribution. This has been confirmed by Wang et al. (2007) and by Gilli et al. (2007). Direct comparison of the $n_H$ distribution from the BAT sample and Tozzi et al shows apparent differences, especially at low $n_H$. Taken at face value, this would indicate an evolution of the $n_H$ distribution between the low median redshift of the BAT sample (0.03) and the redshift of the Tozzi sample ($\\sim0.7$). This is similar to the results reported by La Franca et al. (2005), however Hasinger et al (2008) find no such dependence."
    },
    "0711/0711.1379_arXiv.txt": {
        "abstract": "We present proper motions of water masers toward the Onsala 1 star forming region, observed with the Japanese VLBI network at three epochs spanning 290 days. We found that there are two water maser clusters (WMC1 and WMC2) separated from each other by \\timeform{1\".6} (2900 AU at a distance of 1.8 kpc). The proper motion measurement reveals that WMC1 is associated with a bipolar outflow elongated in the east-west direction with an expansion velocity of $69 \\pm 11$ km s$^{-1}$. WMC1 and WMC2 are associated with two 345 GHz continuum dust emission sources, and located \\timeform{2\"} (3600 AU) east from the core of an ultracompact H\\emissiontype{II} region traced by 8.4 GHz radio continuum emission. This indicates that star formation activity of Onsala 1 could move from the west side of ultracompact H\\emissiontype{II} region to the east side of two young stellar objects associated with the water masers. We also find that WMC1 and UC H\\emissiontype{II} region could be gravitationally bound. Their relative velocity along the line of sight is $\\sim 3$ km s$^{-1}$, and total mass is $\\sim 37 \\MO$. Onsala 1 seems to harbor a binary star at different evolutionary stage. ",
        "introduction": "Star forming regions are often associated with water maser emissions, and the Very Long Baseline Interferometry (VLBI) monitoring observations of the water maser provide a unique tool to study the structure and kinematics of star forming region. Analyses of spatial positions, Doppler velocities, and proper motions of water masers have revealed the 3-D gas kinematics often in the vicinity of young stellar objects (YSOs) (e.g., Orion-KL: \\cite{gen81}; Cepheus A: \\cite{tor01}; G192.16$-$3.84: \\cite{imai06}; G24.78+0.08: \\cite{mos07}).  Onsala 1 star forming region (hereafter ON1) has an ultracompact (UC) H\\emissiontype{II} region observed in radio continuum emissions at 8.4 GHz and 23.7 GHz (\\cite{zheng}; \\cite{argon}). Centimeter radio continuum luminosity of this UC H\\emissiontype{II} region is $10^{4.2} \\LO$, and indicates an exciting star of ZAMS type B0 (\\cite{macleod}). \\citet{kumar2004} reported the presence of multiple outflows from the UC H\\emissiontype{II} region. One outflow traced by $^{12}$CO $J$ = 2--1 line is elongated in the east-west direction with a velocity of 12 km s$^{-1}$ (respect to the systemic velocity of ON1 $V_{\\rm LSR} =$ 12 km s$^{-1}$), and other possible outflow traced by H$^{13}$CO$^{+}$ $J$ = 1--0 line is elongated in the northeast-southwest direction with a velocity of 4.5 km s$^{-1}$ (\\cite{kumar2004}). The H$^{13}$CO$^{+}$ outflow is, however on the other hand, interpreted as a rotating disk by the observations in the NH$_3$ and same H$^{13}$CO$^{+}$ lines (\\cite{zheng}; \\cite{lim02}). These molecular line observations suggest that there are several YSOs in or around the UC H\\emissiontype{II} region of ON1.  Submillimeter continuum emission at 345 GHz (0.85 mm) using the Submillimeter Array (SMA) have resolved two components separated in the northeast--southwest direction (\\cite{su}). The 345 GHz continuum emissions are elongated in the east side of the UC H\\emissiontype{II} region. It is known that YSOs at the youngest stage often show spectral energy distribution with peak flux at submillimeter wavelength, originating from the dust envelopes (\\cite{and93}). The strong submillimeter emissions from these two sources would, therefore, indicate that they are the youngest YSOs contained in ON1. Thus, YSOs at different evolutionary stages seem to be found in ON1.  VLBI observation of the water masers (\\cite{downes}) shows that ON1 has two clusters of water masers separated by $\\sim$\\timeform{2\"} (3600 AU at our assumed distance of 1.8 kpc) in the northeast-southwest direction. These two clusters appear to have offsets from the core of the UC H\\emissiontype{II} region traced by 23.7 GHz continuum emission (\\cite{zheng}). This fact suggests that the water masers are associated with the newly YSOs which formed around the UC H\\emissiontype{II} region. In order to confirm this, new multi-epoch VLBI observations to measure the proper motions of water masers are necessary.  The VLBI observation of \\citet{downes} detected water masers at $V_{\\rm LSR} = 5$--23 km s$^{-1}$. \\citet{kurtz} found that the water maser emission of ON1 has a larger velocity range ($V_{\\rm LSR} = -53$ to 63 km s$^{-1}$), and suggested that these high velocity components would give an important information about other star formation signposts in ON1.  Here, we present the distribution and proper motions of 22.2 GHz water masers towards ON1. Section 2 describes the observations with the Japanese VLBI network (JVN, e.g., \\cite{doi}; \\cite{imai06}) and data reduction. Section 3 shows the distribution and proper motion of the water masers. Section 4 discusses the driving sources of water masers and structure of ON1 through comparison with previous observations. The distance to ON1 is uncertain, however, a near kinematic distance of 1.8 kpc is favored by the most authors (e.g. \\cite{macleod}; \\cite{kumar2004}). We therefore adopt the distance of 1.8 kpc to this source. ",
        "conclusions": "The following conclusions are drawn from this study:  \\begin{enumerate}  \\item We carried out three epoch observations of the ON1 water masers with the JVN, and successfully measured the proper motions of ON1 water masers for the first time.  \\item ON1 has two major water maser clusters (WMC1 and WMC2) which are separated by 2900 AU. Both the WMCs are located at 3600 AU from the UC H\\emissiontype{II} region seen at 8.4 GHz continuum emission.  \\item Proper motion measurement reveals that WMC1 is associated with a bipolar outflow elongated in the east-west direction with high expansion velocity of $69 \\pm 11$ km s$^{-1}$. This outflow is coincident with a jet-like outflow found in CO 2--1 line.  \\item The WMC1 and WMC2 are associated with two 345 GHz continuum sources on the east side of the UC H\\emissiontype{II} region. These two YSOs are still surrounded by dusty envelopes.  \\item The star formation activity of ON1 appears to move from west side of the UC H\\emissiontype{II} region to east side of two WMCs and dense molecular cloud observed in CS lines.  \\item We suggests that the WMC1 and UC H\\emissiontype{II} region in ON1 form a binary system. The relative velocity and total mass of WMC1 and UC H\\emissiontype{II} region are estimated to be $\\Delta V_{\\rm LSR} \\sim 3$ km s$^{-1}$ and $M_{\\rm t} \\sim 37 \\MO$, respectively.  \\end{enumerate}    \\bigskip  The authors wish to thank to all staff members and students of VERA team for observing assistance and support. We hope VERA observe one thousand of water maser sources to perform the astrometry of galactic star forming regions and investigate the structure and dynamics of the Milky Way Galaxy. We thank an anonymous referee for very useful comments. T.O. and H.I. were supported by a Grant-in-Aid for Scientific Research from Japan Society for Promotion Science (17340055).   \\begin{table*}[h] \\begin{center} \\caption{Status of the telescopes, data reduction, and resulting performances in the individual epochs of the JVN observations.} \\label{tab:1} \\begin{tabular}{cccccccc} \\hline \\hline \\multicolumn{1}{c}{Epoch} & \\multicolumn{1}{c}{Date} & \\multicolumn{1}{c}{Duration} & \\multicolumn{1}{c}{Used} & \\multicolumn{1}{c}{Reference} & \\multicolumn{1}{c}{1-$\\sigma$ level} & \\multicolumn{1}{c}{Synthesized} & \\multicolumn{1}{c}{Number of} \\\\ & & & \\multicolumn{1}{c}{telescopes\\footnotemark[$*$]} & \\multicolumn{1}{c}{velocity\\footnotemark[$\\dagger$]} & \\multicolumn{1}{c}{noise} & \\multicolumn{1}{c}{beam\\footnotemark[$\\ddagger$]} & \\multicolumn{1}{c}{detected} \\\\ & & \\multicolumn{1}{c}{(hr)} & & \\multicolumn{1}{c}{(km s$^{-1}$)} & \\multicolumn{1}{c}{(Jy beam$^{-1}$)} & \\multicolumn{1}{c}{(mas)} & \\multicolumn{1}{c}{features} \\\\ \\hline 1 ... & 2005 Mar 24 & 10 & MZ, IR, OG, IS, KS\\footnotemark[$\\S$], NB & 16.5 & 0.045 & $1.9 \\times 0.7,~\\timeform{-38D}$ & 21 \\\\ 2 ... & 2005 Jun 1  & 10 & MZ, OG, IS, KS, NB     & 16.5 & 0.032 & $3.4 \\times 1.3,~\\timeform{-86D}$ & 18 \\\\ 3 ... & 2006 Jan 8  & 10 & MZ, IR, OG, IS, KS     & 16.5 & 0.039 & $1.7 \\times 0.9,~\\timeform{-51D}$ & 23 \\\\ \\hline \\multicolumn{8}{@{}l@{}} {\\hbox to 0pt{\\parbox{170mm} {\\footnotesize \\footnotemark[$*$] Telescopes that were effectively operated and whose recorded data were valid: MZ: the VERA 20-m telescope at Mizusawa, IR: the VERA 20-m telescope at Iriki, OG: the VERA 20-m telescope at Ogasawara Is, IS: the VERA 20-m telescope at Ishigakijima Is, KS: the NiCT 34-m telescope at Kashima, NB: the NRO 45-m telescope at Nobeyama. \\par\\noindent \\footnotemark[$\\dagger$] Local-standard-of-rest velocity of the spectral channel used for the phase reference in data reduction. \\par\\noindent \\footnotemark[$\\ddagger$] The synthesized beam made in natural weight; major and minor axis lengths and position angle. \\par\\noindent \\footnotemark[$\\S$] Ceasing operation for 8 hours due to strong winds. \\par\\noindent }\\hss}} \\end{tabular} \\end{center} \\end{table*}  \\begin{figure*}[h] \\begin{center} \\FigureFile(80mm,100mm){./fig1.eps} \\end{center} \\caption{The spectra of total flux of ON1 water masers, obtained in the three epochs with the Nobeyama 45-m and the Mizusawa 20-m telescopes The velocity resolution is equal to the channel width of 0.21 km s$^{-1}$.} \\label{fig:1} \\end{figure*}   \\begin{figure*}[h] \\begin{center} \\FigureFile(160mm,160mm){./fig2.eps} \\end{center} \\caption{(a): Observed water maser distributions (colored filled circle) superimposed on an 8.4 GHz radio continuum map (thick contour) of the UC H\\emissiontype{II} region (\\cite{argon}). The grey-scale map shows the 345 GHz radio continuum showing dust emission (\\cite{su}), where the lowest grey contour indicates a half of the peak brightness. Crosses show OH maser distribution (\\cite{fish}) and triangle indicates the 10$\\mu$m infrared source (\\cite{kumar2003}). The map origin with RA and DEC offsets = (0, 0) is at $\\alpha$(J2000)=\\timeform{20h10m09.201s} $\\pm$ \\timeform{0.004s} and $\\delta$(J2000)=\\timeform{31D31'36.02\"} $\\pm$ \\timeform{0.08\"}. 1000 mas corresponds to 1800 AU at a distance of 1.8 kpc. (b), (c): Close-up to the two water maser clusters (WMC1 and WMC2) with proper motion vectors. The colored arrows and number added for each feature denote LSR velocity. The arrows at the bottom-rihgt corner in (b) and (c) indicate a proper motion of 5.9 mas yr$^{-1}$ and 1.2 mas yr$^{-1}$ (50 km s$^{-1}$ and 10 km s$^{-1}$), respectively. The dashed line in (b) indicates the axis of bipolar outflow observed in $^{12}$CO $J$ = 2--1 line (\\cite{kumar2004})} \\label{fig:2} \\end{figure*}   \\begin{table*}[h] \\begin{center} \\caption{Parameters of the water maser features identified by proper motion toward ON1.} \\label{tab:2} \\begin{tabular}{rrrrrrrrrrrrr} \\hline \\hline \\multicolumn{1}{c}{ID\\footnotemark[$*$]} & \\multicolumn{4}{c}{Offset\\footnotemark[$\\dagger$]} & \\multicolumn{1}{c}{LSR velocity} & \\multicolumn{4}{c}{Proper motion\\footnotemark[$\\dagger$]} & \\multicolumn{3}{c}{Peak intensity at three epochs} \\\\ & \\multicolumn{4}{c}{(mas)}           & \\multicolumn{1}{c}{(km s$^{-1}$)}   & \\multicolumn{4}{c}{(mas yr$^{-1}$)} & \\multicolumn{3}{c}{(Jy beam$^{-1}$)} \\\\ & \\multicolumn{4}{c}{\\hrulefill \\ }   & \\multicolumn{1}{c}{\\hrulefill \\ }   & \\multicolumn{4}{c}{\\hrulefill \\ }   & \\multicolumn{3}{c}{\\hrulefill \\ }   \\\\ & \\multicolumn{1}{c}{$X$} & \\multicolumn{1}{c}{$\\sigma X$} & \\multicolumn{1}{c}{$Y$} & \\multicolumn{1}{c}{$\\sigma Y$} & \\multicolumn{1}{c}{$V_{\\rm LSR}$} & \\multicolumn{1}{c}{$\\mu_{x}$} & \\multicolumn{1}{c}{$\\sigma \\mu_{x}$} & \\multicolumn{1}{c}{$\\mu_{y}$} & \\multicolumn{1}{c}{$\\sigma \\mu_{y}$} & \\multicolumn{1}{c}{Epoch 1} & \\multicolumn{1}{c}{Epoch 2} & \\multicolumn{1}{c}{Epoch 3} \\\\ \\hline 1  &    129.04 & 0.05 &      48.58 & 0.05 &    15.69 &    1.49 & ...  & $-$0.44 & ...  &  0.40 &  ...  &  0.65 \\\\ 2  &     13.99 & 0.01 &       1.98 & 0.02 &    15.06 &    0.86 & 0.15 & $-$0.33 & 0.05 & 13.70 & 15.00 & 21.50 \\\\ 3  &     11.04 & 0.03 &       4.29 & 0.03 &    12.11 &    0.57 & 0.03 & $-$0.34 & 0.13 &  1.03 &  0.81 &  0.56 \\\\ 4  &      5.69 & 0.06 &      43.55 & 0.08 & $-$40.54 &    7.93 & ...  & $-$3.33 & ...  &  0.53 &  6.82 &  ...  \\\\ 5  &      0.00 & 0.06 &       0.00 & 0.04 &    16.53 &    0.00 & ...  &    0.00 & ...  &  6.55 &  6.75 &  5.28 \\\\ 6  & $-$189.05 & 0.02 &      42.17 & 0.02 &    58.45 & $-$2.48 & ...  &    0.15 & ...  &  3.09 &  5.50 &  ...  \\\\ 7  & $-$189.71 & 0.02 &      43.91 & 0.02 &    55.93 & $-$2.11 & ...  &    0.21 & ...  &  7.53 &  3.69 &  ...  \\\\ 8  & $-$190.73 & 0.01 &      52.90 & 0.01 &    53.40 & $-$4.07 & ...  &    0.42 & ...  & 10.60 & 17.90 &  ...  \\\\ 9  & $-$192.73 & 0.01 &      52.05 & 0.01 &    52.56 & $-$3.22 & ...  &    0.63 & ...  & 28.00 & 26.00 &  ...  \\\\ 10 & $-$818.83 & 0.13 & $-$1647.78 & 0.07 &    10.64 &    1.73 & ...  & $-$0.26 & ...  &  ...  &  1.63 &  1.90 \\\\ 11 & $-$823.19 & 0.03 & $-$1650.52 & 0.02 &    10.85 &    0.94 & 0.16 & $-$0.75 & 0.04 &  2.96 &  2.67 &  3.69 \\\\ 12 & $-$975.58 & 0.07 & $-$1268.37 & 0.06 &    14.85 &    0.96 & ...  &    0.31 & ...  &  1.31 &  ...  &  0.66 \\\\ 13 & $-$978.89 & 0.04 & $-$1268.50 & 0.04 &    11.90 &    0.66 & 0.20 & $-$0.01 & 0.03 &  0.89 &  1.16 &  0.71 \\\\ 14 & $-$981.02 & 0.04 & $-$1267.26 & 0.03 &     7.27 & $-$0.52 & 0.11 &    0.66 & 0.00 &  2.23 &  1.30 &  0.86 \\\\ \\hline \\multicolumn{13}{@{}l@{}} {\\hbox to 0pt{\\parbox{170mm}{\\footnotesize \\footnotemark[$*$] Feature ID number. \\par\\noindent \\footnotemark[$\\dagger$] Relative value with respect to the position-reference maser feature: ID 5. \\par\\noindent }\\hss}} \\end{tabular} \\end{center} \\end{table*}  \\begin{figure*}[h] \\begin{center} \\FigureFile(80mm,80mm){./fig3.eps} \\end{center} \\caption{Observed relative proper motions of water maser features in ON1. The proper motions of only the maser features detected at all three epochs are presented. A number added for each sub-panel shows the assigned one listed in Table \\ref{tab:2}. The dash line indicates a least-squares-fitted line assuming a constant velocity motion.} \\label{fig:3} \\end{figure*}  \\begin{table*}[h] \\begin{center} \\caption{Properties of B0 star in UC H\\emissiontype{II} region and YSOs in WMC1 and WMC2.} \\label{tab:3} \\begin{tabular}{cccccccc} \\hline \\hline \\multicolumn{1}{c}{No} & \\multicolumn{1}{c}{Name} & \\multicolumn{1}{c}{Emission} & \\multicolumn{1}{c}{Maser} & \\multicolumn{1}{c}{Velocity span\\footnotemark[$*$]} & \\multicolumn{1}{c}{Luminosity} & \\multicolumn{1}{c}{Age\\footnotemark[$\\dagger$]} & \\multicolumn{1}{c}{Mass\\footnotemark[$\\ddagger$]} \\\\ & & & & \\multicolumn{1}{c}{(km s$^{-1}$)} & \\multicolumn{1}{c}{(\\LO)} & \\multicolumn{1}{c}{(yr)}  & \\multicolumn{1}{c}{(\\MO)} \\\\ \\hline 1 ...      & B0 star    & Ionized gas\\footnotemark[$\\S$] + Dust\\footnotemark[$\\|$] & OH\\footnotemark[$\\#$] & $\\sim 15$\\footnotemark[$\\#$] & $10^{4.2}$\\footnotemark[$**$] & $\\sim 10^5$ & $\\sim 15$ \\\\ 2 ...      & WMC1 YSO   & Dust\\footnotemark[$\\S$] & Water      & $\\sim 100$ & $< 10^4$\\footnotemark[$\\dagger\\dagger$] & $\\sim 10^4$ & $\\sim$ 2--15 \\MO \\\\ 3 ...      & WMC2 YSO   & Dust\\footnotemark[$\\|$] & Water      & $\\sim 10$  & $< 10^4$\\footnotemark[$\\dagger\\dagger$] & ...        & $\\leq$ 2--15 \\MO \\\\ \\hline \\multicolumn{7}{@{}l@{}} {\\hbox to 0pt{\\parbox{170mm}{\\footnotesize \\footnotemark[$*$] LSR velocity span of masers. \\par\\noindent \\footnotemark[$\\dagger$] Estimated ages of B0 star and WMC1 YSO (see subsection 4.2 and 4.3). \\par\\noindent \\footnotemark[$\\ddagger$] Estimated masses (see subsection 4.3). \\par\\noindent \\footnotemark[$\\S$] Traced by 8.4 GHz continuum emission (\\cite{argon}). \\par\\noindent \\footnotemark[$\\|$] Traced by 345 GHz continuum emission (\\cite{su}). \\par\\noindent \\footnotemark[$\\#$] Results obtaind in OH maser observation (\\cite{fish}). \\par\\noindent \\footnotemark[$**$] Luminosity obtained in 1.3 to 20 cm radio continuum emissions (\\cite{macleod}). \\par\\noindent \\footnotemark[$\\dagger\\dagger$] Luminosity estimated from the far-infrared and radio continuum emissions (see subsection 4.1). \\par\\noindent }\\hss}} \\end{tabular} \\end{center} \\end{table*}  \\begin{figure*}[h] \\begin{center} \\FigureFile(100mm,80mm){./fig4.eps} \\end{center} \\caption{Schematic representation of the structure of ON1.} \\label{fig:4} \\end{figure*}"
    },
    "0711/0711.0180_arXiv.txt": {
        "abstract": "The VDFS comprises the system to pipeline process and archive the data from infrared observations taken by both the WFCAM instrument on UKIRT and the forthcoming VISTA telescope. These include the largest near-IR surveys to date, such as UKIDSS, which produce terabyte sized catalogues of over 10$^9$ rows. Such large data volumes present a performance challenge when the catalogue data, stored in a relational database, require many iterations of astrometric and photometric recalibration. Here we present the VDFS recalibration solution that will be employed in the WSA from the forthcoming UKIDSS Data Release 4 and VSA from its inception. ",
        "introduction": "The VISTA Data Flow System (VDFS) is designed to pipeline process and archive the data from infrared observations taken by both the UK Infrared Telescope's (UKIRT) Wide-Field Camera (WFCAM) and the forthcoming Visible and Infrared Survey Telescope (VISTA). These are currently the fastest near-IR survey instruments and will remain so for years to come, producing terabytes of catalogue data each year. These data are stored in the WFCAM and VISTA Science Archives (WSA \\footnote{http://surveys.roe.ac.uk/wsa/pre/index.html} and VSA respectively) designed and maintained by the Wide-Field Astronomy Unit (WFAU) in Edinburgh. The archives consist of relational databases of catalogue data and associated metadata, together with a data store of all the individual image and catalogue FITS files. The design of the existing WSA is described in detail by Hambly et al. (2008).  Calibration of WFCAM and VISTA data is performed at the pipeline processing stage by the Cambridge Astronomy Survey Unit (CASU, Hodgkin et al. 2008).  Improvements in the calibration have also been calculated by CASU in conjunction with a Calibration Working Group containing members from CASU, WFAU and UKIRT Infra-Red Deep Sky Survey (UKIDSS). Both the astrometric and the photometric calibration are calculated through comparison to the Two Micron All Sky Survey (2MASS, Skrutskie et al. 2006), which covers $99.998\\%$ of the sky resulting in a very uniform calibration accuracy.  However, it is a complicated task to convert from the raw photometric and astrometric quantities measured by the source extractor to calibrated quantities, removing all the instrumentational effects. Subtle effects are often only measured after significant amounts of data have been taken. The equations converting raw data to calibrated quantities will need to be adjusted to account for these effects thus reducing the error on these measurements.  In addition to improving the calibration of new data, existing data in the archive will require recalibration, and this process will happen many times throughout an archive's history. The WSA catalogues increase at a rate of $\\sim1$TB/year and the VSA catalogues will increase at a rate of $\\sim3$TB/year. After five years the VSA will contain 15TB of catalogue data, a factor of ten larger than the amount of data that was last recalibrated for UKIDSS Data Release 3.  \\vspace{-5mm} ",
        "conclusions": ""
    },
    "0711/0711.2716_arXiv.txt": {
        "abstract": "The fast rotating magnetized white dwarf, AE Aquarii, was observed with Suzaku, in October 2005 and October 2006 with exposures of 53.1 and 42.4 ks, respectively. In addition to clear spin modulation in the 0.5--10 keV band of the XIS data at the barycentric period of $33.0769\\pm0.0001$ s, the 10--30 keV HXD data in the second half of the 2005 observation also showed statistically significant periodic signals at a consistent period. On that occasion, the spin-folded HXD light curve exhibited two sharp spikes separated by $\\sim 0.2$ cycles in phase, in contrast to approximately sinusoidal profiles observed in energies below $\\sim$ 4 keV. The folded 4--10 keV XIS light curves are understood as a superposition of those two types of pulse profiles. The phase averaged 1.5--10~keV spectra can be reproduced by two thermal components with temperatures of $2.90_{-0.16}^{+0.20}$ keV and $0.53_{-0.13}^{+0.14}$ keV, but the 12-25 keV HXD data show a significant excess above the extrapolated model. This excess can be explained by either a power-law model with photon index of $1.12_{-0.62}^{+0.63}$ or a third thermal component with a temperature of $54_{-47}^{+26}$ keV. At a distance of 102 pc, the 4--30 keV luminosities of the thermal and the additional components become $1.7_{-0.6}^{+1.3}$ and $5.3_{-0.3}^{+15.3} \\times 10^{29}$ erg s$^{-1}$, respectively. The latter corresponds to 0.09\\% of the spin down energy of the object. Possible emission mechanisms of the hard pulsations are discussed, including in particular non-thermal ones. ",
        "introduction": "\\label{section:introduction} Since the discovery of cosmic-rays by Hess in 1912, the origin of such high-energy particles and the mechanism of their acceleration have remained a long-standing issue in astrophysics. Among various types of astrophysical particle acceleration sites, one important prototype is rotation-powered pulsars, namely fast rotating neutron stars (NSs) with strong magnetic field of $\\geq 10^{12}$ G. They are considered to accelerate particles using their rotation as the ultimate energy source, and strong electric fields as the acceleration tool. The electrostatic potential, induced by the strong surface field, $B$, and the fast rotation with a period, $P$, is expected to amount to \\begin{eqnarray} V &\\sim& \\left(\\frac{2 \\pi R}{P}\\right) B R \\nonumber \\\\ &\\sim& 6 \\times 10^{16} \\left(\\frac{P}{1 \\mbox{s}}\\right)^{-1} \\left(\\frac{B}{10^{12} \\mbox{G}}\\right) \\left(\\frac{R}{10^6 \\mbox{cm}}\\right)^2 \\mbox{V} \\label{eq:acc_volt} \\end{eqnarray} where $R$ is a typical radius from the NS center at which the electric acceleration takes place.  Magnetized white dwarfs (WDs) are similar systems to pulsars; rotating compact objects with strong magnetic fields. Since a typical magnetized WD has $P \\sim 5 \\times 10^3$ s, $B \\sim 10^6$ G, and $R \\geq 10^9$ cm, we expect $V \\sim 10^{13}$ V from equation (\\ref{eq:acc_volt}). Therefore, magnetized WDs should be a promising candidate of new particle acceleration sites, although efficiencies of the particle acceleration and the subsequent radiation could be different between pulsars and WDs.  Non-thermal incoherent radio emission has in fact been detected from seven WDs via systematic radio surveys \\citep{mcv_radio1,mcv_radio2,mcv_radio3,mcv_radio4,mcv_radio5}. Some objects like AE Aquarii \\citep{aeaqr_radio1, aeaqr_radio2} and AM Herculis \\citep{amher_radio1,amher_radio2} show signs of coherent radio emission. Then, higher-energy photons in the X-ray to gamma-ray bands may also be generated via such non-thermal processes as curvature radiation like in NSs, non-thermal bremsstrahlung, and inverse Compton scattering of soft photons \\citep{nonthermal_process1, nonthermal_process2,nonthermal_process3}, as well as electron-positron annihilation lines and $\\gamma$-ray emission via the $\\pi^0$ decay chain. In fact, some authors reported TeV gamma-ray emission from AE Aquarii and AM Herculis during optical flares \\citep{aeaqr_tev1, aeaqr_tev2,amher_tev}. Therefore, high energy electrons with energies reaching 1-100 MeV are thought to be present in these WDs \\citep{aeaqr_radio2,aeaqr_radio_theo}.  The object AE Aquarii (hereafter AE Aqr), located at a distance of $d = 102_{-23}^{+42}$ pc \\citep{aeaqr_distance}, is a magnetized cataclysmic variable (MCV) belonging to the DQ Herculis (or intermediate polar) class. It is spinning near break up \\citep{aeaqr_particle_acceleration} with a spin period of $P_{\\rm sp} = 33.0767$ s \\citep{aeaqr_optical}, which is the second fastest value among MCVs. This fast rotation, combined with the surface field intensity of $B \\sim 10^{5 \\mbox{--} 6}$ G which is a typical value among intermediate polars, yields via equation (\\ref{eq:acc_volt}) an induced potential of $V\\simeq 10^{14 \\mbox{--} 15}$ V, which is probably one of the highest found among MCVs. In addition, AE Aqr has been in  a spin-down state over the past 27 years, at a rather high rate of $1.8 \\times 10^{-6}$ s yr$^{-1}$ \\citep{aeaqr_spindown,aeaqr_spindown2}. The value implies a huge spin-down luminosity of $6 \\times 10^{33}$ erg s$^{-1}$ \\citep{aeaqr_particle_acceleration,aeaqr_spindown}. The object is one of the most asynchronous cataclysmic variables with an orbital period of $P_{\\rm orb} = 35567~{\\rm s} = 9.88$ hr \\citep{wd_period,aeaqr_spindown}. Flares reported in the radio band \\citep{aeaqr_radio1,aeaqr_radio2} are consistent with synchrotron emission \\citep{aeaqr_radio_theo}, and the pulsating TeV gamma-rays \\citep{aeaqr_tev1,aeaqr_tev2} are sometimes coincident with optical flares.  According to soft X-ray observations \\citep{aeaqr_asca1, aeaqr_newton}, the plasma temperature of thin thermal radiation from AE Aqr is significantly lower than the estimated gravitational potential of its WD, whose mass is about 0.8 -- 0.9 times the solar mass \\citep{wd_period,aeaqr_orbital_spin_param}. Furthermore, its UV to X-ray luminosity ($10^{31}$ erg s$^{-1}$) is three orders of magnitude lower than that in a simple accretion case \\citep{mcv_mdot} with a rate of $10^{17}$ g  s$^{-1}$. Therefore, the accreting matter may be inhibited from reaching the WD surface by the fast WD rotation \\citep{aeaqr_newton}. \\citet{aeaqr_propeller} argue that only 1 \\% of the accreting gas in AE Aqr reaches the WD surface, while the rest is expelled away due to propeller effects. These observed properties altogether suggest strongly that AE Aqr is a promising site of particle acceleration, as pointed out by \\citet{aeaqr_particle_acceleration}.  The hard X-ray band is most suited to the search for the expected non-thermal signals, because that particular band is expected to receive the largest number of non-thermal high-energy photons without being hampered by the soft X-rays arising from accretion. In the present paper, we report on the first successful detection of pulsating hard X-rays from the  most promising object, AE Aqr, achieved with Suzaku \\citep{suzaku07}, which carries two instruments; the X-ray CCD camera called the X-ray Imaging Spectrometer (XIS;\\cite{xis2007}) with the X-ray telescope (XRT; \\cite{xrt2007}) and the Hard X-ray Detector (HXD;\\cite{hxd2007a}). The emission is likely to be of non-thermal origin, thus reinforcing the view that MCVs are indeed particle accelerators. ",
        "conclusions": "\\label{section:discussion} \\subsection{Summary of the Observation} \\label{section:discussion:obs} Using Suzaku, we observed AE Aqr in 2005 and 2006. In the first observation, the HXD detected periodic signals synchronized with the rotation of the WD (section \\ref{section:timing:HXDperiod}) whose profile has a sharp peak (section \\ref{section:timing:pulse}) with a duration of $\\sim$0.1 in phase. The same sharp feature was also found in the spin-folded XIS light curves in energies above 4 keV. The phase-averaged spectra from the XIS above 1.5~keV can be explained with the two optically thin thermal plasma models with temperatures of $2.90_{-0.16}^{+0.20}$~keV and $0.53_{-0.13}^{+0.14}$~keV (section \\ref{section:spectrum}). The PIN spectrum up to 25~keV, however, shows an excess emission over the model. This excess can be fitted by a power-law model with a photon index of $1.1 \\pm 0.6$, or another thermal model with the temperature of $54_{-47}^{+26}$ keV. In the 13--30 keV range, the spectrally measured excess, $3.2^{+9.7}_{-2.5}$ c s$^{-1}$, is consistent within large errors with that contained in the sharp peaks in the spin-folded hard X-ray light curves.  The 4--30 keV flux of the additional hard component from the spectral analyses is obtained as $L_{\\rm HX} \\sim 1.7_{-0.1}^{+4.9}\\times 10^{-12}$ erg cm$^{-2}$ s$^{-1}$, which gives a luminosity of $5.3_{-0.3}^{+15.3} \\times 10^{29}$ erg s$^{-1}$ at the distance of 102~pc \\citep{aeaqr_distance}. This corresponds to $\\sim 0.01$\\% of the spin-down luminosity of $6 \\times 10^{33}$ erg s$^{-1}$ \\citep{aeaqr_particle_acceleration, aeaqr_spindown}. Note, however, that the hard X-ray pulsation has a duty ratio of only $\\sim$0.1. If the hard X-ray emission is highly anisotropic without a constant bias, and if the emission is beamed away from us in the other pulse phases, the true hard X-ray luminosity would be higher, amounting up to $5 \\times 10^{30}$ erg s$^{-1}$ which is $0.1$\\% of the spin-down luminosity. In any case, the ratio of the hard X-ray emission to the spin-down luminosity of AE Aqr comes in the range of 0.01--0.1\\%, which is the same as those measured in the 2--10 keV range from NS pulsars \\citep{ns_spindown1, ns_spindown2}.   \\subsection{The Nature of the Hard X-ray Component} \\label{section:discussion:hard_component} The observed hard X-ray emission pulsates at the WD rotation period, but not at the prograde or retrograde beat periods. In order to reinforce this result, we further folded the 4--10 keV XIS data at the prograde and retrograde beat periods, but the results were consistent with the absence of any periodic modulation, either in the 2005 or 2006 data sets. This implies that the hard X-ray emission region is fixed in the frame which is co-rotating with the WD, while excludes a possibility that the emission region is fixed to the binary frame.  Although approximately in phase with the soft X-ray spin modulation, the hard X-ray pulses exhibit a spiky profile with a duty ratio of only $\\sim 0.1$ in phase, unlike the sinusoidal profile in the XIS band below 4 keV. Furthermore, the spin phase of the major spike exhibits an energy dependence with a discontinuity at $\\sim$ 4 keV (figure~\\ref{fig:e_dependence_efold}). The soft and hard X-ray components are therefore likely to originate from distinct emission sites. If interpreted as hot thermal emission, the hard component, with its very high temperature, should be emitted from the bottom of the gravitational potential, i.e., the WD surface, while the softer component from some outer regions \\citep{aeaqr_newton}. However, it would be highly difficult to produce the observed sharp pulse profiles from the WD surface (e.g., magnetic poles) via any thermal process. Indeed, such X-ray pulsations, appearing only in the hard X-ray band with very sharp profiles, have never been observed from other accreting magnetic WDs \\citep{mcv_asca2004}, in which the thermal emission mechanism is considered to dominate. We thus consider that the hard component is more naturally interpreted as non-thermal emission, as evidenced by its successful power-law modeling. If so, a population of non-thermal particles must be accelerated in the AE Aqr system, at least occasionally, producing the hard X-ray pulses detected with Suzaku. In the following two subsections, we discuss possible particle acceleration mechanisms, and hard X-ray emission mechanisms from the accelerated electrons.  \\subsection{Possible Acceleration Mechanism} \\label{section:discussion:acc_mechanism} As mentioned in section \\ref{section:introduction}, the fast rotating WD in the AE Aqr system is expected to produce an induced electric potential reaching $V\\simeq 10^{14 \\mbox{--} 15}$ V. Therefore, it is most natural to assume that the non-thermal particles are acceleration by this electric potential.  Although the suggested acceleration mechanism is very similar to that in rotation-powered pulsars, the present WD system differs from pulsars in an important physical condition. While a rotation-powered pulsar is expected to reside in a low-density environment, AE Aqr is an accreting binary system. In particular, the radius of the light cylinder, $l_{\\rm lc} \\sim P_{\\rm sp} /2\\pi \\sim 5$~l-s, is comparable to the binary size of $\\sim 6$ l-s. In fact, \\citet{aeaqr_newton}, analyzing the XMM-Newton RGS data of AE Aqr, concluded that a plasma with the electron density $n_{\\rm e} \\sim 10^{11}$~cm$^{-3}$ is spread over a scale of $(2-3)\\times 10^{10}$~cm in this binary system. Since this plasma density exceeds the Godreich-Julian density \\citep{goldreich69} which is calculated as $\\sim 5\\times10^{4}$ cm$^{-3}$ in the present case \\citep{nonthermal_process3}, the induced potential would be short-circuited, and hence an efficient acceleration would not take place. However, both theoretical \\citep{aeaqr_propeller} and observational works suggest that the propeller effect produces a low-density cavity around the WD, by inhibiting the accretion plasma from getting closer to the WD than $\\sim 10^{10}$ cm \\citep{aeaqr_propeller}. The electric acceleration may be possible in such regions.  Even if the plasma density is sufficiently low, an efficient electric-field acceleration also requires evacuation of neutral matter from the acceleration region; otherwise, the electrons would suffer too heavy ionization losses to be accelerated efficiently. The mean-free-path of an electron with energy $E_{\\rm e}$, due to ionization loss, is given by \\begin{eqnarray} \\lambda_{\\rm e} &\\sim& \\frac{1}{n_{\\rm p}} \\frac{1}{\\ln \\Lambda} \\left( \\frac{\\pi e^2}{8 \\epsilon_0 E_{\\rm e}} \\right)^{-2} \\nonumber \\\\ &\\sim& 10^{7} \\left( \\frac{n_{\\rm p}}{10^{10} \\mbox{cm$^{-3}$}}\\right)^{-1} \\left( \\frac{E_{\\rm e}}{ 1 \\mbox{keV}}\\right)^2 \\mbox{cm}, \\end{eqnarray} where $n_{\\rm p}$ is the density of protons, $\\ln \\Lambda$ is a Coulomb logarithm (typically $\\sim$ 10 -- 20), $e$ is the elementary charge, and $\\epsilon_0$ is the vacuum dielectric constant (see e.g., \\cite{rybichi_lightman}). If an initial electron energy is $E_{\\rm e} \\sim$ 1 keV assuming a thermal plasma in the system, the energy gain of the electron accelerated along the length of $\\lambda_{\\rm e}$ can be roughly estimated to be $E_{\\rm acc} \\sim \\frac{V}{l_{\\rm lc}} \\cdot \\lambda_{\\rm e} \\sim 10^{10} \\left( \\frac{n_{\\rm p}}{10^{10} \\mbox{cm$^{-3}$}}\\right)^{-1}$ eV, with $V \\sim 10^{14}$ V and $l_{\\rm lc} \\sim 10^{11}$ cm. Thus, the energy gain $E_{\\rm acc}$ exceeds the initial energy $E_{\\rm e}$ when the surrounding density is $n_{\\rm p} < 10^{17}$ cm$^{-3}$. Therefore, in the AE Aqr system with $\\sim 10^{11}$ cm$^{-3}$ \\citep{aeaqr_newton}, the ionization loss does not affect the electron acceleration.  \\subsection{The Hard X-ray Emission Mechanism} \\label{section:discussion:emission_mechanism} In general, there are the following four elementary processes that can produce hard X-ray emission from accelerated electrons: non-thermal bremsstrahlung, inverse Compton scattering, synchrotron emission, and curvature radiation. Considering the observational results, we examine these processes.  If the hard X-rays are due to non-thermal bremsstrahlung in thick materials, electrons have to be accelerated only to an energy of $\\sim$ 100 keV to emit X-rays with an energy of several tens of keV. The emanating hard X-ray spectrum is expected to have a photon index of $\\sim$1, when the target is thick and hence the electrons have a flattened energy distribution due to Coulomb losses \\citep{loss_flattened_bremss}. This is in good agreement with the observed photon index, $1.12^{+0.63}_{-0.62}$. Candidates of the dense matter, co-rotating with the WD, includes the surface of the WD and the accretion stream toward the magnetic poles. However, to produce the sharp hard X-ray pulses, the electrons must be in the form of collimated beams (e.g., along the magnetic field lines), and furthermore, must be moderately relativistic ($\\geq$ 200 keV) so that the bremsstrahlung emissivity is sufficiently anisotropic \\citep{bremsstrahlung} to produce the spike in the hard X-ray light curve. It is not obvious whether these conditions can be fulfilled. Furthermore, a still larger difficulty is in energetics: in the 10 -- 100 keV range, an electron loses $\\sim 10^{4}$ times larger energies in ionization (or Coulomb) loss than in radiation. Then, the particle acceleration requires an energy input which is more than an order of magnitude higher than the spin-down luminosity. We therefore conclude that the bremsstrahlung interpretation is difficult.  In the inverse-Compton scenario, the seed photons can be either the UV photons from the WD or the optical photons from the companion star whose spectral type is K3--K5 \\citep{aeaqr_orbital_spin_param}. Then, the electrons are required to have a Lorentz factor of $\\gamma \\sim$10$^2$, or an energy of $\\sim 100$ MeV. The estimated 4--30 keV luminosity (after accounting for the duty ratio; section \\ref{section:discussion:obs}) translates to a hard-photon number flux of $2.6 \\times 10^{37}$ s$^{-1}$. This is a reasonably small fraction ($\\sim 10^{-7}$) of the optical photon number flux to be expected from a K4 star ($\\simeq 5 \\times 10^{44}$ s$^{-1}$ near the WD surface). % However, the inverse-Compton process is rather inefficient; an electron with a Lorentz factor $\\gamma \\simeq 10^{2}$ is expected to lose half its energy in $T_{\\rm 1/2} \\simeq \\frac{1}{\\gamma}\\frac{m_{\\rm e}c^2}{\\sigma_{\\rm T} c u_{\\rm s}} \\sim 4 \\times 10^{6}$ s, where $m_{\\rm e}$ is the electron mass, $\\sigma_{\\rm T}$ is the Thomson cross section, and $u_{\\rm s} \\simeq 0.1$ erg cm$^{-3}$ is the visible photon density at the position of the WD. Therefore, the electron is expected to leave the system after depositing only $\\sim 10^{-6}$ of its energy. The energy input needed to sustain the particle acceleration would be then $\\sim 10^{6} L_{\\rm HX}$, making the inverse Compton scenario highly unlikely.  As the third possibility, we consider synchrotron radiation by relativistic electrons. For a magnetic field strength of $B = B_5 \\times 10^{5}$ G, an electron is required to have a Lorentz factor $\\gamma \\sim 10^4 B_5^{-1}$ to emit $\\sim$ 30 keV photons. This value of $\\gamma$ is lower than the maximum possible energy expected from the induced voltage (section~\\ref{section:introduction}) by $\\sim$4 orders of magnitude. Furthermore, the synchrotron emission is highly efficient, with a very short lifetime of $\\sim 8 \\times 10^{-6} B_5^{-2} (\\frac{\\gamma}{10^4})^{-1}$ s (e.g, \\cite{synchrotron,synchrotron2}). Considering further that the synchrotron emission is highly anisotropic, we presume that this scenario is quite likely, if the electrons are accelerated along the field lines, and then at some stage, acquire cross-field velocity and start emitting synchrotron photons. The observed photon index $\\Gamma \\sim 1.1$ predicts a very flat electron spectrum with a spectral index of $\\sim$1.2. This does not matter, however, if direct electron field acceleration works. In fact, rotation-powered NS pulsars generally exhibit rather flat spectra in X-rays, $\\Gamma \\sim 1.4$ \\citep{Pulsar_index01}.  Finally, we consider the possibility of the curvature radiation process. According to \\citet{curvature_radiation}, the energy of the hard X-ray emission via curvature radiation can be written as \\begin{equation} E \\sim \\frac{3}{2} \\frac{\\hbar c \\gamma^3}{R_{\\rm crv}} \\sim 50 \\left(\\frac{\\gamma}{10^6}\\right)^3 \\left(\\frac{R_{\\rm crv}}{R_{\\rm WD}}\\right)^{-1} \\rm{keV}, \\label{eq:curvature} \\end{equation} where $\\hbar$ is the Plank constant, $R_{\\rm crv}$ is the curvature radius of magnetosphere field lines, and $R_{\\rm WD}$ is the WD radius of AE Aqr. On the right-hand side of equation~\\ref{eq:curvature}, $R_{\\rm crv}$ is normalized for convenience to $R_{\\rm WD}$. It is not easy to predict where the photon energy $E$ is maximized, because a smaller curvature radius can be realized closer to the WD, whereas the maximum $\\gamma$ is probably achieved at a larger radius. Generally, as we move away from the WD surface, $R_{\\rm crv}$ decreases gradually, while $\\gamma$ of the electrons being accelerated will increase rapidly. According to \\citet{nonthermal_process3}, the curvature-radiation photons will reach the maximum of $E \\sim 100$ keV at $R_{\\rm crv} \\sim l_{\\rm lc}$, together with $\\gamma \\sim 6 \\times 10^6$, which is feasible in the AE Aqr system. However, the efficiency of this process is so low that the luminosity is estimated to be $\\sim 3 \\times 10^{27}$ erg s$^{-1} \\sim 10^{-2} L_{\\rm HX}$. Therefore, it is difficult to explain the observation in terms of the curvature radiation.  We conclude that the hard X-ray pulse is probably caused by relativistic electrons accelerated by the fast rotation of the magnetized WD. The most promising emission mechanism is the synchrotron emission in the strong magnetic fields of the WD."
    },
    "0711/0711.1904_arXiv.txt": {
        "abstract": "We study the ionization and thermal evolution of the intergalactic medium during the epoch of \\heii reionization by means of radiation hydrodynamical cosmological simulations. We post-process baryonic density fields from a standard optically-thin IGM simulation with a homogeneous galaxy-dominated UV background (UVB) which reionizes \\hi and \\hei at z=6.5 but does not have any contribution to the ionization of \\heii. Therefore, we suppress the \\heii photoheating contribution to the gas temperature due to the homogeneous UVB. Quasars are introduced as point sources throughout the 100 Mpc simulation volume located at cold dark matter (CDM) density peaks consistent with the Pei luminosity function. We assume an intrinsic quasar spectrum $J(\\nu) \\propto \\nu^{-1.8}$ and a luminosity proportional to the halo mass. We evolve the spatial distribution of the \\heii ionizing radiation field at h$\\nu$ = 4, 8, and 16 Ryd using a time-implicit variable tensor Eddington factor radiative transfer scheme. Simultaneously, we also solve for the local ionization of \\heii to \\heii and the associated photoheating of the gas including opacity effects.  We find that the percolation of the \\heiii regions is essentially complete by z=2.5. When comparing to a self-consistent optically thin simulation at the same redshift, in which \\heii is also ionized by the uniform UVB, we find that inclusion of opacity effects results in higher IGM temperature by a factor of approximately 1.7 at the mean gas density level. We construct synthetic absorption line spectra from which we derive statistical parameters of the \\heii \\laf. We use 300 long ($\\Delta z = 0.2$) random lines of sight to compute at $\\bar{z} = 2.5 \\pm 0.1$ a mean \\heii \\lya line transmission of $\\bar{F} = 0.304 \\pm 0.002$. The error corresponds to a significant one standard deviation in the transmitted flux due to the sightline to sightline variance equal to $\\simeq 11\\%$ the mean value. The opacity effect on the gas temperature is shown by comparing the broadening width of the \\hi and \\heii \\lya lines to the results from the self-consistent optically thin simulation. We find a shift by approximately 1.25 km/s to higher b-parameter values for both \\hi and \\heii. Finally, we estimate the relative broadening width between the two forests and find that the \\heii median b-parameter is about 0.8 times the median \\hi broadening width. This implies that the \\heii absorbers are physically extended consistent with conclusions from observed lines of sight. ",
        "introduction": "The thermal evolution of the IGM after reionization is governed chiefly by photoheating (Efstathiou 1992; Miralda-Escud\\'{e} \\& Rees 1994). Models for the thermal evolution of the IGM (Miralda-Escud\\'{e} \\& Ostriker 1990; Giroux \\& Shapiro 1996; Abel \\& Haehnelt 1999)  generally assume that \\hi and \\hei are nearly fully ionized by $z \\sim 6$ by star forming galaxies (Fan, Carilli \\& Keating 2006), while \\heii ionizes somewhat later at z $\\sim$ 3 by quasars due to \\heii's higher ionization potential and recombination rate (Sokasian, Abel \\& Hernquist 2002; hereafter SAH). Observational support for late \\heii reionization is summarized in SAH. Abel \\& Haehnelt (1999) emphasized the importance of opacity effects during reionization in establishing the post-reionization temperature of the gas. They showed that models using the optically thin expression for the photoheating rate during \\heii reionization underestimate the IGM temperature at mean density by a factor of $\\sim$ 2.  Hydrodynamical cosmological simulations of the \\laf (Cen et al. 1994; Zhang, Anninos \\& Norman 1995; Hernquist et al. 1996; Miralda-Escud\\'{e} et al. 1996; Zhang et al. 1997, 1998; Theuns et al. 1998) generally adopt the optically thin expression for photoheating for simplicity and computational economy. This is a reasonable assumption for the \\hi and \\hei photoheating at $z\\sim3$ due to the low opacity of the IGM, but not for the \\heii photoheating as \\heii is in the process of being ionized by the percolation of \\heiii spheres (SAH). These simulations therefore underestimate the temperature of the IGM during the epoch of helium reionization. The standard approach taken in these simulations is to assume a homogeneous photoionizing background which evolves with redshift consistent with observed quasar and galaxy counts, such as that by Haardt \\& Madau (2001). The ionization and thermal state of the baryonic gas is then computed self-consistently with its dynamics by solving the equations of hydrodynamic cosmology including radiative heating and cooling (Cen 1992; Katz, Weinberg \\& Hernquist 1996; Anninos et al. 1997). It is found that the temperature of the low density IGM is determined almost entirely by photoheating balancing adiabatic cooling due to cosmic expansion. This results in a tight relationship between gas temperature and density, the so-called equation of state of the IGM (Hui \\& Gnedin 1997):  \\begin{equation} T=T_0\\Delta^\\beta \\end{equation} where $\\Delta=\\rho/\\bar{\\rho}$ is the gas overdensity and $\\beta$ is redshift dependent but in the range $0 \\leq \\beta(z) \\leq 0.6$. Abel \\& Haehnelt (1999) show that opacity effects during \\heii reionization raises $T_0$ and reduces $\\beta$ relative to an optically thin calculation. A technique often employed to include opacity effects within hydrodynamic simulations is to simply multiply the \\heii photoheating rate by a constant factor $X_{\\heii}\\geq 1$, with a value of 2-4 being sufficient to match observed temperatures (e.g., Bryan \\& Machacek 2000; Jena et al. 2005).  The combination of high spectral resolution quasar absorption line observations and hydrodynamical cosmological simulations provide a means for measuring the thermal evolution of the IGM.  The thermal state of the gas is generally deduced from \\hi \\lya linewidths (b-values) in high resolution spectra (Rauch et al. 1997; Schaye et al. 1999, 2000; Bryan \\& Machacek 2000; Theuns et al. 2000; Bolton et al. 2005; Jena et al. 2005) although the flux power spectrum has also been employed (Zaldarriaga, Hui \\& Tegmark 2001). If the temperature of the IGM alone determined the b-values, then it would be straightforward to measure the temperature from high resolution spectra. However, Hubble broadening is always of the same order as the thermal broadening in \\laf absorption lines (Schaye 1999), hence the need for comparison with simulations. Theuns, Schaye \\& Haehnelt (2000) used the b-parameter distribution to measure the temperature of the IGM at z=3.25, finding $T_0 \\geq 15,000$ K. Schaye et al. (2000) used the lower cutoff in the linewidth--column density scatter diagram to measure the temperature evolution of the IGM over the redshift range 2-4.5. They found evidence of late reheating at z$\\sim$3, which they ascribed to late \\heii reionization by quasars. Bryan \\& Machacek (2000) independently explored the same diagnostics and found that temperature estimates were sensitive to the assumed cosmology, in particular the amplitude of mass fluctuations on a few Mpc scales. Zaldarriaga, Hui \\& Tegmark (2001) used the falloff of the \\lya forest flux power spectrum at small scales to measure the IGM temperature, and found $T_0 \\sim 2 \\times 10^4$ K, dependent on their assumed $\\beta$.  A proper calculation of \\heii reionization would treat quasars as point sources within the simulation volume and solve the equation of radiative transfer for the spatial distribution of the ionizing background coupled self-consistently to the dynamical, thermal, and ionization evolution of the gas. This was done in an approximate way by SAH, who post-processed a series of density fields taken from a SPH hydrodynamical simulation of the IGM using the GADGET code. The hydrodynamical simulation used the optically thin prescription for ionizing and heating the gas, assuming all ionization states of hydrogen and helium were in ionization equilibrium with the UVB of Haardt \\& Madau (1996). In the post-processing step, the helium reionization calculation was recomputed for an evolving quasar source population treated as point sources within the volume. Quasars were assumed to have a constant lifetime of $10^7$ yr. Every $10^7$ yr, peaks within the density field of a suitably chosen data dump from the hydrodynamic simulation were populated with quasars consistent with an empirical luminosity function. Around each point source the static equation of radiative transfer for \\heii ionizing photons was solved using a photon-conserving ray-casting scheme. The ionization state of \\heii was then updated ignoring thermal feedback to the gas. SAH found that for reasonable parameter choices, \\heii reionization occurred in the range $3 \\leq z \\leq 4$ consistent with observations.  In this paper we present a simulation which is similar in spirit to SAH, but with several important differences. We also postprocess a series of snapshots from a hydrodynamic cosmological simulation with a radiative transfer code, however we keep track of the photoheating of the IGM as expanding \\heiii spheres percolate and eventually merge. Our hydrodynamic simulation was performed on an Eulerian grid of $512^3$ cells versus SAH's $224^3$ in the same volume, giving us $\\sim 12$ times the mass resolution and $\\sim 2.3$ times the spatial resolution in the low density IGM. We have also separated the effects of stellar and quasar populations of the ionization evolution of the IGM differently from SAH. Our hydrodynamic simulation computes the ionization of \\hi and \\hei due to stellar sources only using the homogeneous UVB of Haardt \\& Madau (2001) arising from galaxies \"GAL\". In the post-processing step, quasars are treated as point sources, which ionize \\heii to \\heiii. Although our radiative transfer scheme is based on completely different spatial and angular discretizations as SAH, the important difference for the purpose of this paper is that we solve the RT equation at three frequencies $h\\nu$=4, 8 and 16 Ryd in order to evaluate, albeit crudely, the local \\heii photoheating rate taking the processed QSO spectrum into account.  The outline of this paper is as follows. In \\S2 we describe the cosmological realization for the hydrodynamic simulation of homogeneous \\hi and \\hei reionization described in \\S3. In \\S2 we also introduce some physical concepts relating to late \\heii reionization and also describe our treatment of quasars in the simulation volume. In \\S4 we describe our method for simulating inhomogeneous \\heii reionization. In \\S4.1 our radiative transfer scheme is detailed, in \\S4.2 we present our single species ionization model for \\heii, and in \\S4.3 we present results in terms of globally integrated quantities.  Since we have split the calculation into two phases, the homogeneous reionization of \\hi and \\hei, followed by the inhomogeneous reionization of \\heii, our calculation is not fully self-consistent. Although we keep track of the late reheating, we do not modify the underlying density fields nor do we alter temperature-dependent recombination rates, which will affect the detailed ionization state of the gas. In \\S5 we present an analysis of our main result, which is the late reheating of the IGM due to inhomogeneous \\heii reionization, as well as the importance of neglecting these coupling effects. We show that these effects, while present, are small, and do not seriously undermine our estimate of the reheating. In \\S6 we present observational signatures of \\heii reionization based on synthetic \\hi and \\heii \\lya absorption line spectra derived from our simulation. In \\S7 we summarize our main results and conclude. In a series of appendices we derive the rate equation for species fraction in an expanding universe (Appendix A), provide more detail on the affect of neglecting \\heii photoheating on the ionization state of the gas (Appendix B), and document the reaction rate coefficients we use (Appendix C). ",
        "conclusions": "We have simulated the late inhomogeneous reionization of \\heii by quasars and the attendant photoheating of the IGM including opacity effects.  We post-process baryonic density fields from a standard optically thin IGM simulation with a homogeneous galaxy-dominated UV background which reionizes \\hi and \\hei at z=6.5. Quasars are introduced as point sources throughout the 100 Mpc simulation volume located at density peaks consistent with the Pei luminosity function. We assume an intrinsic quasar spectrum $J(\\nu) \\propto \\nu^{-1.8}$ and a luminosity proportional to the halo mass. We evolve the spatial distribution of the \\heii ionizing radiation field at h$\\nu$ = 4, 8, and 16 Ryd using a time-implicit variable tensor Eddington factor radiative transfer scheme which we describe. Simultaneously, we also solve for the local ionization of \\heii to \\heiii and the associated photoheating of the gas. This distinguishes our calculation from the work of Sokasian, Abel and Hernquist (2002) who simulated inhomogeneous \\heii reionization but did not study the thermal evolution of the gas.  We find that the cosmic evolution of the QSO population causes the individual \\heiii regions to overlap and subsequently ionize 90 \\% of the volume at \\heiii ionization fractions $\\psi_{3} = \\frac{n_{He_{III}}}{n_{He}} \\gtrsim 10^{-5}$ by $ z \\simeq 2.5$. In addition, to the \\heii and \\heiii number density calculation, we also update the local gas temperature due to the thermalization of the photoelectrons ejected by the HeII ionization process. As expected, the temperature is higher compared to the unprocessed simulation with the difference increasing rapidly at redshifts $z \\lesssim 3$. Relative to a self-consistent optically thin simulation where \\heii is also photoionized by a homogeneous UV background, we find that optical depth effects result in an increase in the temperature of the intergalactic medium at the cosmic mean density level by a factor of $\\approx 1.7$ at $z=2.5$. The results of our temperature calculation are consistent with analytic and numerical predictions of the HeII heating effect (Abel \\& Haehnelt 1999; Schaye et al. 2000; McDonald et al. 2000). Finally, we trace the redshift evolution of the \\heii Ly$\\alpha$ transmission using randomly casted synthetic spectra through the simulated volume. The analysis of the mean transmission allows for the derivation of the effective optical depth of the 304~\\AA~ line.  We have calculated that at $\\bar{z} = 2.5 \\pm 0.1$, the average of pixel flux among all lines of sight yields a value for the mean transmission of $\\bar{F} = 0.304 \\pm 0.002$ which corresponds to $\\tau^{eff}_{304} = 1.190 ~\\pm 0.007$. The length of each sightline ($\\Delta z=0.2$) is longer than the size of the simulation in order to minimize the transmission variance across the redshift path. We compute the error only due to the sightline to sightline variance and find that at z=2.5 the 1$\\sigma$ standard deviation is $11\\%$ the mean \\lya flux. When we compare to estimates from \\heii forest spectra observed with the FUSE (Far Ultraviolet Space Explorer) satellite and HST at the same redshift interval we find that the our value of the effective depth is comparable but slightly above the observed values. We attribute the disagreement to our approximate treatment of the inhomogeneous ionizing radiation that ignores the diffuse component and only focuses on the point sources' input.  In addition to the mean transmission estimate, we also compute the column density and b-parameter distribution of the identified \\heii Ly$\\alpha$ lines and compare them to the corresponding statistical properties of the HI Ly$\\alpha$ lines. We find that an optically thick calculation results in extra heating that shifts the b-parameter distributions to higher broadening widths. In addition and within the limitations of our coarse numerical resolution, which does not resolve the \\hi forest, our calculation shows that the median \\heii b-parameter at $\\bar{z} = 2.5$ is $82\\%$ the HI broadening width. The latter suggests that the Hubble differential flow may be the dominant line broadening mechanism which in turn indicates that \\heii lines primarily form in underdense to cosmic mean density regions.  ACKNOWLEDGEMENTS  This work was supported by NSF grants AST-9803137, AST-0307690 and AST-0507717. The simulations were carried out on the TITAN IA64 Linux cluster at the National Center for Supercomputing Applications and on the IBM Blue Horizon system at San Diego Supercomputing Center under LRAC allocation MCA098020.  \\appendix"
    },
    "0711/0711.1590_arXiv.txt": {
        "abstract": "The artificial viscosity is reconsidered in smoothed particle hydrodynamics to prevent inter-particle penetration, unwanted heating, and unphysical solutions. The coefficients in the Monaghan's standard artificial viscosity are considered as time variable, and a restriction on them is proposed such that avoiding the undesired effects in the subsonic regions. The shock formation in adiabatic and isothermal cases are used to study the ability of this modified artificial viscosity recipe. The computer experiments show that the proposal appears to work and the accuracy of this restriction is acceptable. ",
        "introduction": "Gas dynamical processes are believed to play an important role in the evolution of astrophysical systems on all length scales. A ubiquitous process in astrophysical fluid dynamics is the behavior of gases subjected by shock waves. These are almost the normal, rather than exceptional, type of astrophysical fluid flows that happen very often in cases of astrophysical interest. A number of examples are the onset of spiral structures and shock fronts in accretion disks around compact objects (Lanzafame et al.~2006), molecular cloud formation in shock-compressed layers (Koyama and Inutsuka~2000), the rapid and strong increase in pressure that can be produced by interaction of the fast wind and the slow envelope of a planetary nebulae (Gurzadyan~1997), and so on.  A powerful gridless particle method that was invented to solve complex fluid-dynamical problems in astrophysics is smoothed particle hydrodynamics~(SPH) (Lucy~1977, Gingold and Monaghan~1977). The SPH method has a number of attractive features such as its low numerical diffusion in comparison to grid based methods. Whilst the SPH originally developed for compressible flows, it has been extended to deal with free-surface incompressible fluids (e.g., Monaghan~1994, Ellero et al.~2007). A recent worthy review of the SPH methodology and its applications can be found in Monaghan~(2005).  An adequate scenario for SPH application is the unbound astrophysical problems, especially the behavior of gases subjected to compression (shock waves). The magnitude of the viscous term dose not affect the net shock-jump conditions, thus, many numerical schemes implicitly or explicitly incorporate the trick of artificial viscosity~(AV) for halting the ever-growing steepening tendency produced by nonlinear effects. The first use of viscosity in SPH equations was by Lucy~(1977) who introduced an AV to present a slow build-up of acoustic energy from integration errors in an SPH simulation. The AV is designed to allow shock phenomena to be simulated, or simply to stabilize a numerical algorithm. Indeed two forms of AV are applied in SPH equations, the bulk and the von Neumann-Richtmeyer viscosity, respectively. They prevent inter-particle penetration, allow shocks to form and damp the post-shock oscillations. Although many different functional forms for the AV have been proposed (see, e.g., Liu \\& Liu~2003), all contain problem sensitive parameters that are often set in a somewhat arbitrary manner. Here we try firstly to provide a guide to the relevant prescriptions for some AV that have been considered before.  A more effective viscosity which may conserves linear and angular momentum was suggested by Monaghan and Gingold (1983). They devised a viscosity by simple arguments about its form and its relation to gas viscosity. The viscous term between two particles, denoted by $\\Pi_{ab}$, is added to the pressure term of SPH equations. When two particles approach each other, the AV produces a repulsive force between them. When they recede from each other the force is attractive (Monaghan~1989).  An undesirable aspect of the Monaghan's standard AV is that it can introduce considerable shear viscosity into the flow. This problem may be reduced by a bulk viscosity in a general-purpose code of Hernquist and Katz (1989) for evolving three-dimensional, self-gravitating fluids in astrophysics. Clearly the artificial viscous dissipation increases the Reynolds number of a flow, artificially, with the result that, for example, the Kelvin-Helmholtz shear instabilities are heavily diffused. In this way a von Neumann stability analysis of the SPH equations along with a critical discussion of various parts of the algorithm was investigated by Balsara~(1995). He suggested reducing viscous dissipation by multiplying $\\Pi_{ab}$ by a symmetric factor.  Another problem of viscosity is that although AV is approximately successful for handling shocks but it can be too large in other parts of the flow. For this purpose, a switch to reduce the AV away from shocks was given by Morris and Monaghan (1997, hereafter MM97). They introduce the idea of time-varying switch which fits more naturally with a particle formulation. Each particle has a viscosity parameter which evolves according to a simple source and decay equation. The source causes the switch to grow when the particle enters a shock and the decay term causes it to decay to a small value beyond the shock.  In the present study we combine the Monaghan's standard AV with the time-varying coefficients like that of MM97's switch. we propose a modification to the time-dependent AV prescription designed by MM97 so that the aim is to maintain the shock-simulating capabilities of SPH. For this purpose we optimize the source term so that the AV is restricted to supersonic velocities in regions that are under compression. Section~2 is devoted to the SPH method and the suggested changes in the AV. These modifications have been tested in \\S3 for the one dimensional adiabatic and isothermal shock problems. Finally, a summary with conclusion is given in section~4. ",
        "conclusions": "In SPH method two forms of AV is necessary; the von Neumann-Richtmeyer viscosity and linear viscous term. In this work we revisited the relevant prescriptions for some AV, which was proposed to consider these two forms of viscosity in SPH methodology. As a good idea, MM97 introduced the time-varying switch which fits more naturally with the particle formulation. The shock formations in adiabatic and isothermal cases have been performed to test the effect of different AVs.  The simulation results in which the MM97's AV was used, show that the post-shock density has oscillations for both cases of the adiabatic and isothermal shocks. The oscillations in adiabatic shock grow further because the thermal energy in this case is restrained in the system. On the other hand, the post-shock density of adiabatic case is approximately equal to the analytical value, while the difference in isothermal case is more distinguishable.  To partly overcome some of these problems, we proposed a modification of the AV. Based on the method of MM97, we used time-variable $\\alpha$ and $\\beta$ in the standard AV. In this approach, the AV is used only in supersonic regions where the gas is not under compression, thus, the source term is restricted. Using this source multiplied by two regulating parameters ($z_\\alpha$ and $z_\\beta$) can virtually eliminate some problematic effects of using an artificial dissipation in SPH capability of shock simulation. Computer experiments have been performed to test the abilities of the proposal, and the results show that the accuracy of these restrictions on AV is acceptable. Although there are some satisfying results from shock formation tests, the proposed AV should be communicated to the astrophysical community, so that others can try it. There must be other tests to complete the idea."
    },
    "0711/0711.1842_arXiv.txt": {
        "abstract": "We present the first ever calculation of cosmic microwave background (\\textsc{CMB}) anisotropy power spectra from semilocal cosmic strings, obtained via simulations of a classical field theory. Semilocal strings are a type of non-topological defect arising in some models of inflation motivated by fundamental physics, and are thought to relax the constraints on the symmetry breaking scale as compared to models with (topological) cosmic strings.  We derive constraints on the model parameters, including the string tension parameter $\\mu$, from fits to cosmological data, and find that in this regard BPS semilocal strings resemble global textures more than topological strings.  The observed microwave anisotropy at $\\ell = 10$ is reproduced if $G\\mu = 5.3 \\times 10^{-6}$ ($G$ is Newton's constant). However as with other defects the spectral shape does not match observations, and in models with inflationary perturbations plus semilocal strings the $95\\%$ confidence level upper bound is $G\\mu<2.0\\times 10^{-6}$ when \\CMB data, Hubble Key Project and Big Bang Nucleosynthesis data are used (c.f.\\ $G\\mu<0.9\\times 10^{-6}$ for cosmic strings). We additionally carry out a Bayesian model comparison of several models with and without defects, showing models with defects are neither conclusively favoured nor disfavoured at present. ",
        "introduction": "Recent observational data are tightly constraining cosmological models, in many cases rendering them non-viable. One of the constraints comes from the fact that many potentially successful high-energy physics motivated models of inflation predict a cosmic string network \\cite{Vilenkin:1994book,Hindmarsh:1994re} in the post-inflationary era: this occurs in supersymmetric hybrid inflation \\cite{Lyth:1998xn}, in grand unified theories \\cite{Jeannerot:2003qv} and in some string theory inflation models \\cite{Majumdar:2002hy,Sarangi:2002yt,Copeland:2003bj,Dvali:2003zj}. These strings and other type of defects contribute to the cosmic microwave background (\\textsc{CMB}) anisotropies in addition to the primordial inflationary fluctuations, allowing \\CMB experiments to provide upper limits on the contribution from defects \\cite{Bevis:2004wk,Wyman:2005tu, Battye:2006pk, Fraisse:2006xc, Bevis:2007gh}. In the near future, observations will either reveal the presence of these defects or will significantly tighten the constraints upon the models from which they are predicted (see for example Refs.~\\cite{Seljak:2006hi,Bevis:2007qz,Battye:2007si,Fraisse:2007nu,Pogosian:2007gi}).  However, a simple and elegant method for evading the constraints coming from cosmic strings in these models was proposed in Refs.~\\cite{Urrestilla:2004eh,Dasgupta:2004dw}. A duplication of the matter fields responsible for the string formation, which arises quite naturally in several models \\cite{Dasgupta:2004dw,Achucarro:2005vz,Dasgupta:2007ds}, changes the nature of the strings from topological and therefore long-lived strings, to non-topological semilocal strings \\cite{Vachaspati:1991dz}.  All the desired inflationary properties of the model persist, and semilocal strings are believed to be cosmologically less harmful than their topological counterparts, in that their energy density may be lower for a given symmetry breaking scale. In this paper we confirm these expectations by performing the first \\CMB calculations for semilocal strings and determine the constraints upon semilocal strings imposed by current data.  Calculating the additional contribution that cosmic strings would make to the \\CMB temperature and polarization power spectra is not an easy task. They are highly non-linear entities, and the range of scales involved in the problem is enormous. However, as the width of the string is much smaller than the string length, one possibility is to neglect the finite width and perform Nambu--Goto type simulations \\cite{Allen:1996wi,Allen:1997ag,Contaldi:1998mx,Landriau:2003xf}. Unfortunately, these simulations ignore the string decay into gravitational waves or particles and the corresponding back-reaction onto the network. Their representation of small-scale loop production in the network is also under debate \\cite{Vincent:1996rb,Vanchurin:2005pa,Vanchurin:2005yb, Ringeval:2005kr,Martins:2005es}. Yet another level of simplification often used in \\CMB calculations for local strings is to simulate merely a stochastic ensemble of unconnected string segments \\cite{Albrecht:1997nt,Pogosian:1999np,Wyman:2005tu, Seljak:2006hi,Battye:2006pk,Battye:2007si}.  Lacking any dynamical equations, in this case the sub-horizon decay of strings must be taken into account by the random removal of segments, at a rate chosen to match that seen in network simulations, but qualitatively accurate results are obtained with much less computational cost. In the global defect case, the energy is not localized into the string cores and their resolution is not essential. This enables the calculation of \\CMB contributions from global defects using field theory simulations \\cite{Pen:1997ae,Durrer:1998rw}.  None of these three approaches is useful for the study of semilocal strings, since they are not well-modelled by Nambu--Goto strings while, unlike the global case, the string cores are potentially important. Fortunately, in a recent breakthrough we performed the first \\CMB calculations for local cosmic strings to derive from field-theoretical simulations \\cite{Bevis:2006mj}. We achieved this by the use of a parametrized approximation to reduce the range of scales required in the simulations, but which had no resolvable effect on the \\CMB predictions. That approach has opened the door for the possibility of simulating defects with no known approximating models.  We have therefore employed the machinery developed for the local string case to the simplest model containing semilocal strings, with the parameters chosen to give the Bogomol'nyi--Prasad--Sommerfield (BPS) limit (see Section \\ref{sec:model}). The BPS limit arises naturally in D-term SUSY inflation and $D3/D7$ brane inflation. We leave the study of non-BPS cases to a future article but it is the expectation that in the two opposite extremes between which the BPS case lies, the semilocal result will tend to that of global textures or to that of local strings \\cite{Achucarro:1997cx,Achucarro:1998ux,Achucarro:2005tu}. We perform a detailed comparison of our results with these two cases using our earlier results in Ref.~\\cite{\\longprd} for local strings. For the case of global textures, we  perform \\CMB calculations using the linear $\\sigma$ model.\\footnote{ See  for instance Ref.~\\cite{Durrer:1995sf} for a comparison of the dynamics of the linear and non-linear $\\sigma$ models.} We leave a discussion of our texture calculations to an Appendix.  We have also performed a multi-parameter comparison with current data, yielding results similar to those of Ref.~\\cite{Bevis:2007gh} where it was seen that local cosmic strings are mildly favoured by current \\CMB data. We find that the preference for semilocal strings is slightly greater and, as in Ref.~\\cite{Bevis:2007gh}, we explore this preference using Bayesian evidence calculations, as well as considering the implications of non-\\CMB data.  This paper is organized as follows: in the next Section we discuss the semilocal string model, before giving a brief overview of the simulation and \\CMB calculation method in Section \\ref{method}. We will then present our \\CMB results and consistency checks in Section~\\ref{CMBresults}, including the comparison with local strings and global textures. Section \\ref{MCMC} then presents the comparison with data. ",
        "conclusions": "\\label{discussion}  In this article we report on the first ever semilocal string simulations to predict \\CMB power spectra. This has been possible because of recent improvements in simulating field theoretical defects, as there is no Nambu--Goto type approximation for semilocal strings. The case of Abelian Higgs model strings is much better known, so we compared the semilocal string predictions to the Abelian Higgs string ones, trying to minimize the differences in the approaches. Since semilocal strings were expected to share properties with both an Abelian Higgs string prediction and a texture prediction, we also produced a set of simulations for texture-type defects to compare all three defect contributions.  Our first conclusion is that the amplitude of the spectra for the semilocal strings is indeed lower than in Abelian Higgs models (for a given energy scale); therefore, the normalization at $\\ell=10$ to match the \\WMAP value gives a higher value for semilocal strings $G\\mu_{10}=5.3\\times 10^{-6}$ than for Abelian Higgs strings $G\\mu_{10}=2.7\\times 10^{-6}$.  Moreover, the form of the power spectra for semilocal strings allows for their fractional contribution to the temperature power spectra at multipole $\\ell=10$ to be higher than the one for Abelian Higgs strings: the 95\\% confidence upper bounds are $\\fd<0.25$ ($\\fd<0.14$) for semilocal strings for \\CMB only (\\textsc{CMB}+\\textsc{BBN}+\\textsc{HKP}); whereas $\\fd<0.17$ ($\\fd<0.10$) for Abelian Higgs strings.  These two ingredients contribute to allow a higher value of $G\\mu$ for semilocal strings when \\CMB data (and also when \\HK and \\BBN data) are taking into account: $G\\mu<2.6\\times10^{-6}$ ($G\\mu<2.0\\times10^{-6}$) for semilocal strings for \\CMB only (\\textsc{CMB}+\\textsc{BBN}+\\textsc{HKP}); whereas $G\\mu<1.1\\times10^{-6}$ ($G\\mu<0.9\\times10^{-6}$) for Abelian Higgs strings. Thus, if a promising high-energy physics inflation model faces a problem of creating Abelian Higgs cosmic strings at a scale that it is too high for the \\textsc{CMB}, a possible resolution could be to extend the model to predict semilocal strings instead.  This weakens the $G\\mu$ constraint by a factor of nearly three, which may give the theory extra breathing space.  As in the case of textures \\cite{Bevis:2004wk} and local strings \\cite{Battye:2006pk,Bevis:2007gh}, the best fit for the cosmological parameters when taking into account semilocal strings gives a rather high spectral index $n_{\\rm s}$, comparing to the concordance model. For \\CMB only data, the preferred value of $n_{\\rm s}$ is higher than 1 ($n_{\\rm s}=1.02\\pm0.04$), whereas if \\BBN and \\HK data are included $n_{\\rm s}=0.98\\pm0.02$. This is in contrast to the Abelian Higgs data, where the \\CMB data alone predict $n_{\\rm s}=1$; but the result is very similar when \\BBN and \\HK data are included \\cite{Bevis:2007gh}.  This result also alleviates pressure on high--energy physics models (such as D-term and F-term hybrid inflation) which predict defects and an $n_s$ close to unity.  As $\\ns = 1$ is close to our best-fit values, we studied how well a combination of Harrison--Zel'dovich plus defects would fit the data, comparing it to the concordance model.  Using goodness of fit and Bayesian Evidence criteria, we find that models with defects are not ruled out (even when \\BBN and \\HK data are included), and should be considered as competitive models. In fact, taking into account only \\CMB data, there is significant evidence in favour of a Harrison--Zel'dovich plus semilocal strings model.  We have not included large-scale structure datasets into our analysis for two reasons.  Firstly, perturbations from defects are non-linear and non-Gaussian as soon as they are created \\cite{Bevis:2004wk}, and so it is not clear that the commonly-used model \\cite{Cole:2005sx} for non-linear corrections to the galaxy power spectrum applies. Secondly, recent analyses \\cite{Percival:2006gt,Sanchez:2007rc} show scale-dependent bias even on relatively large scales ($0.05 < k/h\\mathrm{Mpc}^{-1} < 0.15 $), and make clear that more work is needed before the galaxy power spectrum can be used for robust estimates of cosmological parameters.   The prediction for polarization power spectra from semilocal strings is similar to that of Abelian Higgs strings. As a matter of fact, the height of the \\BB polarization peak (normalized using the $95\\%$ confidence level upper bound on $\\fd$) is of the same order of magnitude for Abelian Higgs strings, semilocal strings and textures, though the position at which the maximum occurs vary from case to case. A detection of a \\BB signal on these scales could be an indicator of the existence of cosmic defects \\cite{Bevis:2007qz}.  In the present work we have studied only BPS semilocal strings, but we will in a future work extend the analysis to lower values of $\\beta$, which would arise in F-term supersymmetric models (or even in P-term models \\cite{Burrage:2007yu,Burrage:2007bv}). The behaviour of semilocal strings does depend strongly of $\\beta$ \\cite{Achucarro:2005tu}: the lower the value of $\\beta$ the more the semilocal strings behave as Abelian Higgs strings. Tracking how the \\CMB predictions change with respect to $\\beta$ will give us insight into how the different components in a string network contribute to the creation of anisotropies.  There is also a whole new class of strings coming from string inflation: $(p,q)$-strings \\cite{Copeland:2003bj,Sarangi:2002yt,Dvali:2003zj}. Even though there is some work about the networks of these defects \\cite{Saffin:2005cs,Hindmarsh:2006qn,Rajantie:2007hp,Urrestilla:07}, basic properties, such as scaling, are still not really understood. Provided some sound \\CMB predictions from these strings are calculated, it would be interesting to compare the results with those of Abelian Higgs or semilocal strings.  Ultimately, one would like to be able to discriminate between all the different defect types, be it Abelian Higgs, semilocal, textures of $(p,q)$-strings. As we have shown here, the predictions do look similar, but there are differences, e.g., the aforementioned position of the \\BB peak. The question is whether these differences are strong enough, or whether future data will be good enough, to pinpoint whether defects do appear and if so, which type.  In a future publication we will address this issue."
    },
    "0711/0711.0844_arXiv.txt": {
        "abstract": "The origin of the high energy emission (X-rays and ${\\gamma}$-rays) from black holes is still a matter of debate. We present new evidence that hard X-ray emission in the low/hard state may not be dominated by thermal Comptonization. We present an alternative scenario for the origin of the high energy emission that is well suited to explain the high energy emission from GRO~J1655$-$40. ",
        "introduction": "$ ${\\rm K}_{\\alpha}$ line}  Emission below 10\\,keV can be due to a combination of thermal emission from the disk, and non-thermal emission from a corona. It is difficult to study disks through their continuum spectra, so the most pragmatic way to study the disk is to use broad ${\\rm Fe}$\\,${\\rm K}$ emission lines. This line (as the reflection component that we will explain later) arises from the fluorescence of the inner disk from an irradiating source of high energy emission (George \\& Fabian, 1991). Depending on the ionization state of the disk this line is centered at 6.4-6.97\\,keV. From the study of the profiles of these lines, it can be inferred the spin of the black hole, which can be between zero for an Schwarzchild black hole (Fabian et~al., 1989) to a maximally rotating regime in the case of a Kerr black hole (Laor, 1991). In the hard X-ray emission (${\\rm E}{\\geq}10$\\,keV) a spectral hump usually appears, called the reflection bump, with the same origin. Thus, both reflection and broad ${\\rm Fe}$ ${\\rm K}_{\\alpha}$ line should appear together during the observations. Nevertheless, both effects are inhibited when a high ionization degree occurs, because the disk becomes a perfect reflector of the ${\\gamma}$-rays (Ross et~al.1999).  Current observations of relativistic lines suggest that a range of black hole spin parameters may be seen (for reviews, see Miller 2007 and Nandra et al. 2007). It is therefore important to take relativistic effects into account by fitting spectra with the appropriate line and reflection models. Inner disks can get very close to the black hole (${\\rm d}{\\leq}6$\\,${\\rm R}_{\\rm g}$) and this would imply a very concentrated source for the high energy emission, i.e., the high energy source has to be centrally compact. ",
        "conclusions": ""
    },
    "0711/0711.2064_arXiv.txt": {
        "abstract": "We compare the improved measurement of the \\Lya forest flux probability distribution at $1.7<z<3.2$ presented by Kim et al. (2007) to a large set of hydrodynamical simulations of the \\Lya forest with different cosmological parameters and thermal histories. The simulations are in  good agreement with the observational data if the temperature-density  relation for the low density intergalactic medium (IGM), $T=T_0 \\Delta^{\\gamma-1}$, is either close to isothermal or inverted ($\\gamma<1$). Our results suggest that the voids in the IGM may be significantly hotter and the thermal state of the low density IGM  may be substantially more complex than is usually assumed at these redshifts.   We discuss radiative transfer effects which alter  the spectral shape of ionising radiation during the epoch of \\HeII reionisation as a possible physical mechanism for achieving an inverted temperature-density relation at $z\\simeq 3$. ",
        "introduction": "Traditional analyses of the \\Lya forest observed in the spectra of high redshift quasi-stellar objects (QSOs) decompose the flux distribution into a series of discrete absorption profiles, generally characterised by a Voigt function (\\citealt{Hu95,Lu96,KirkmanTytler97,Kim97,Kim01}).  Voigt profiles provide an accurate  description for an absorption line if the absorber is a localised  gas cloud with a Gaussian velocity dispersion (\\citealt{Rauch98}). In the now widely established paradigm for the origin of the \\Lya forest most of the absorption is, however, caused by extended gas distributions broadened by the Hubble flow, removing much of the physical motivation for the decomposition of the flux distribution into Voigt profiles. Such a decomposition is nevertheless useful, but it is non-unique and the line fitting process is time consuming and somewhat subjective (see \\citealt{Kim07}, hereafter K07, for an exhaustive discussion).  Alternative characterisations of the flux distribution based on pixel statistics (\\citealt{Jenkins91, Rauch97,Gaztanaga99,McDonald00,Theuns00,Meiksin01,Viel04c,Lidz06b,Becker07}) and wavelet analysis (\\citealt{Meiksin00,TheunsZaroubi00,Zaldarriaga02}) have also been developed. These measures are simple and - at least in principle - easy to compare to the same quantities extracted from theoretical models of the \\Lya forest (\\citealt{Cen92,MiraldaEscude96,Bi97,Croft99, Dave99,Theuns98, Jena05}).  In particular, the power spectrum of the \\Lya forest flux distribution has recently been successfully developed into a quantitative tool for measuring the matter power spectrum on scales of $0.5 h^{-1}$ Mpc to $40h^{-1}$ Mpc (e.g. \\citealt{Croft02,Viel04b,McDonald05}).  The simplest pixel statistic is the \\Lya flux probability distribution function (PDF), which is sensitive to the density distribution and thermal state of the IGM ({\\it e.g.} \\citealt{Becker07}).  Several attempts have been made to obtain joint constraints on a variety of cosmological and astrophysical parameters using the flux PDF together with the flux power spectrum (\\citealt{Choudhury01, Desjacques05,Lidz06b,Desjacques07}).  However, the PDF is also sensitive to a range of systematic uncertainties, most notably the assumed continuum level, the noise properties of the spectra and contamination by metal absorption lines in the \\Lya forest.  In high resolution, high signal to noise data at  $z<3$, the first two uncertainties have a small impact on the shape of the PDF (K07). The $z<3$ metal absorption in the \\Lya forest, however, alters the shape and amplitude  of the PDF significantly, especially at $0.2<F<0.8$. Since many numerical simulations of the \\Lya forest do not include metals, and those that do must also correctly reproduce their uncertain distribution and abundance, this can lead to substantial systematic uncertainties when attempting to fit simulation data to the observed flux PDF of the \\Lya forest.  K07 have completed a detailed analysis of the impact of metal absorption lines on the \\Lya PDF at $1.7<z<3.2$.  They presented new measurements of the PDF in which all identifiable metal contamination has been removed from the \\Lya forest on a line-by-line basis.  The metal cleaned PDF had a significantly different shape when compared directly to the metal contaminated sample.  In this paper we shall investigate what can be inferred about the underlying matter distribution and the thermal state of the intergalactic medium from these new PDF measurements using state-of-the-art hydrodynamical simulations of the \\Lya forest.  Our approach is different to most recent studies of the flux PDF which have used either pseudo-hydrodynamical techniques, dark matter only simulations or semi-analytical models (\\citealt{Choudhury01,Desjacques05,Lidz06b,Becker07}) which model the effect of gas pressure on the underlying IGM density distribution in an approximate way.  The structure of this paper is as follows. We discuss the observational data and the hydrodynamical simulations in section 2. In section 3, we qualitatively compare the simulation data to the observations and determine the impact of various astrophysical and numerical effects on the simulated flux PDF.  A quantitative comparison is undertaken in section 4 and the best fitting models to the K07 data are presented along with a comparison to other recent studies in the literature.  We briefly discuss the implications of our results for measurements of the matter power spectrum from the \\Lya forest flux power spectrum in section 5 and conclude in section 6. ",
        "conclusions": "We have compared the improved measurements of the \\Lya flux PDF at $1.7<z<3.2$ by K07 obtained from a full Voigt profile analysis  of 18 high resolution, high signal to noise  UVES/VLT spectra to a large set of hydrodynamical simulations of the IGM.  Our main findings are as follows:  \\begin{itemize}  \\item {The shape of the \\Lya flux PDF depends only weakly on the fluctuation amplitude of the matter power spectrum characterised by $\\sigma_{8}$ and on the normalisation of  the temperature-density relation for the low density IGM, $T=T_0 \\Delta^{\\gamma-1}$ (\\citealt{HuiGnedin97,Valageas02}).  A moderately increased  scatter in the temperature-density relation, which may result from inhomogeneous radiative transfer effects during the epoch of \\HeII reionisation, also has little effect on the shape of the flux PDF. The shape of the \\Lya flux PDF is mainly determined by the  effective optical depth $\\tau_{\\rm eff}$ and by the slope of the temperature-density relation for the low density IGM.  The extreme ends of the PDF at flux levels of zero and unity are particularly sensitive to the slope of the temperature-density relation.\\\\}  \\item{A comparison of the full range of the \\Lya flux PDF to hydrodynamical simulations favours an inverted temperature-density relation.     The flux PDFs of our synthetic spectra with an inverted temperature-density relation, such that $\\gamma = [0.67,0.54,0.44]$ at $\\langle z \\rangle =[2.07,2.52,2.94]$,  provide an excellent fit to the observed metal-cleaned flux PDF measured by K07.   This suggests that the voids in the IGM  are significantly hotter than  previously assumed.\\\\}  \\item {If we restrict our analysis of the PDF to flux levels  $0.1 \\leq F \\leq 0.8$,  which are least affected by continuum fitting, noise, galactic winds  and the uncertainties of how to simulate the relevant physics in dense regions of the IGM, the comparison of observed  and synthetic spectra  still  favours a model where the temperature-density relation is either nearly isothermal or inverted.\\\\}  \\item{Raising or lowering the assumed continuum level of our synthetic spectra in order to mimic the possibility of a systematic over or underestimate  of the continuum level in the observed spectra rapidly worsens the quality of the fit to the K07 PDF for offsets larger than 1.5 per cent. This suggests that continuum fitting errors for the high quality data of K07 are under control at a level of 1-2 per cent, consistent with the estimates of the continuum uncertainties by K07.\\\\}  \\item{A comparison of our simulations with the full range of the  \\Lya flux PDF measured by M00 at $\\langle z \\rangle = 3$ from a smaller observational sample gives consistent results if the continuum on the synthetic spectra is lowered by $1.5$ per cent.  Using the currently preferred cosmological parameters, the M00 data also favour a nearly isothermal or inverted temperature-density relation.\\\\}  \\item{An inverted temperature-density relation has important implications when inferring the underlying matter power spectrum from the \\Lya forest flux distribution.  The effect of an inverted temperature-density relation on the flux power spectrum is similar to the effect of an larger normalisation, $\\sigma_8$, of the matter power spectrum.  This  could alleviate some of the existing strain between the measurements of $\\sigma_8$ from \\Lya forest data and the WMAP 3 year results  (\\citealt{VielHaehnelt06}).}  \\end{itemize}  The preference of the data for an inverted temperature-density relation would, if consolidated, be a rather profound result. It would suggest that the widely accepted paradigm of a tight IGM temperature-density relation at low densities, evolving from being almost isothermal following hydrogen reionisation to a relation where the temperature rises  with increasing density (\\citealt{HuiGnedin97,Theuns98,Valageas02}), is too simplistic.  Perhaps the most plausible explanation for an inverted temperature density relation is \\HeII reionisation. There is evidence to suggest this may end around $z\\sim 3$ (\\citealt{Songaila98,Schaye00,Theuns02d,Shull04,Bolton06,Faucher07}), and thus lies at the upper end of the redshift range considered in this work. Changes in the photoheating rate which occur as ionising photons propagate through optically thick regions can significantly raise the temperature of the IGM (\\citealt{AbelHaehnelt99}). First attempts at modeling inhomogeneous \\HeII reionisation and/or the associated radiative transfer effects also suggest a more complex and possibly multi-valued and inverted temperature-density relation is plausible (\\citealt{Bolton04,Gleser05,Tittley07, Furlanetto07}). Therefore, while a well-defined, inverted temperature-density relation may not faithfully represent the true thermal state of the IGM, a model with $\\gamma<1$ should mimic some of the expected  inhomogeneous radiative transfer induced effects on the IGM temperature following \\HeII reionisation.  Within this framework, this model provides a good fit to the K07 observational data and suggests that the thermal  state of the IGM may be significantly more complex at $z\\simeq 3$ than previously thought.    Further investigation of the thermal state of the IGM around $z\\simeq 3$ is therefore highly desirable. Note, however,  that  studies using wavelets have failed to  detect the expected fluctuations in the IGM  temperature at these  redshifts (\\citealt{Theuns02c,Zaldarriaga02}).   While radiative transfer effects and the inhomogeneity of \\HeII reionisation provide a plausible explanation for an increased temperature in underdense regions  of the IGM, it will  require full numerical radiative transfer simulations to confirm the validity of this picture. Helium reionisation is driven by hard photons emitted by rare, luminous and most likely short-lived sources (QSOs) under conditions where photon mean free paths are significantly longer and recombination time scales shorter (if compared to the Hubble time) than for hydrogen reionisation at  higher redshifts. Current 3D cosmological radiative transfer codes which incorporate multiple sources are designed primarily to investigate the topology of \\HI reionisation, and are less reliable when modeling the resulting thermal state of the IGM (\\citealt{Iliev06}).  Most detailed studies of radiative transfer effects on the IGM temperature therefore employ either reduced dimensionality (\\citealt{Bolton04,BoltonHaehnelt07}) or neglect the modeling of multiple sources (\\citealt{Maselli05,Tittley07}). Furthermore, the physical scales involved during \\HeII reionisation are much larger than those during \\HI reionisation at $z>6$, placing even greater demands on the required computational resources (although see \\citealt{Sokasian02,Paschos07}). Radiative transfer simulations which accurately  track the temperature evolution  are thus very challenging.   Developments in this area will be  key to interpreting the increasing amount of high quality \\Lya forest data.  Finally, we must also not lose sight of the possibility that other processes may contribute to the heating of the IGM (e.g. \\citealt{Nath99,MadauEfstathiou99,Inoue03,RicottiOstriker04,Samui05}). Our results may instead indicate that an alternative modification to the current models for the \\Lya forest and the thermal state of the IGM is required."
    },
    "0711/0711.2587_arXiv.txt": {
        "abstract": "In order to make asteroseismology a powerful tool to explore stellar interiors, different numerical codes should give the same oscillation frequencies for the same input phy\\-sics. Any differences found when comparing the numerical values of the eigenfrequencies will be an important piece of information regarding the numerical structure of the code.  The ESTA group was created to analyze the non-physical sources of these differences. The work presented in this report is a part of Task~2 of the ESTA group. Basically the work is devoted to test, compare and, if needed, optimize the seismic codes used to calculate the eigenfrequencies to be finally compared with observations. The first step in this comparison is presented here. The oscillation codes of nine research groups in the field have been used in this study. The same physics has been imposed for all the codes in order to isolate the non-physical dependence of any possible difference. Two equilibrium models with different grids, 2172 and 4042 mesh points, have been used, and the latter model includes an explicit modelling of semiconvection just outside the convective core. Comparing the results for these two models illustrates the effect of the number of mesh points and their distribution in particularly critical parts of the model, such as the steep composition gradient outside the convective core. A comprehensive study of the frequency differences found for the different codes is given as well. These differences are mainly due to the use of different numerical integration schemes. The number of mesh points and their distribution are crucial for interpreting the results. The use of a second-order integration scheme plus a Richardson extrapolation provides similar results to a fourth-order integration scheme. The proper numerical description of the Brunt-V\\\"ais\\\"al\\\"a frequency in the equilibrium model is also critical for some modes. This influence depends on the set of the eigenfunctions used for the solution of the differential equations. An unexpected result of this study is the high sensitivity of the frequency differences to the inconsistent use of values of the gravitational constant ($G$) in the oscillation codes, within the range of the experimentally determined ones, which differ from the value used to compute the equilibrium model. This effect can provide differences for a given equilibrium model substantially larger than those resulting from the use of different codes or numerical techniques; the actual differences between the values of $G$ used by the different codes account for much of the frequency differences found here. ",
        "introduction": "Asteroseismology is at present being developed as an efficient instrument in the study of stellar interiors and evolution. Pulsational frequencies are the most important asteroseismic observational inputs.  It is evident that a meaningful analysis of the observation, in terms of the basic physics of stellar interiors which is the ultimate target of the investigation, requires reliable computation of oscillation frequencies for specified physics. This is a two-step process, involving first the computation of stellar evolutionary models and secondly the computation of frequencies for the resulting models.  Lebreton et al. (this volume) provide an overview of the tests of stellar model calculations.  Here we consider the computation of the oscillation frequencies.  An evident goal is that the computed frequencies, for a given model, should have errors well below the observational error, which in the case of the CoRoT mission is expected to be below 0.1\\,$\\mu$Hz \\citep{baglin}.  For the Kepler mission \\citep[e.g.,][]{Christ2007}, with expected launch in early 2009, selected stars may be observed continuously for several years and errors as low as $10^{-3} \\, \\mu$Hz may be reachable, particularly for modes excited by the heat-engine mechanism.  Since errors resulting from numerical problems are typically highly systematic, they may affect the asteroseismic inferences even if they are substantially below the random errors in the observed frequencies.  This must be kept in mind in the assessment of the estimates of the numerical errors.  During the last decades a lot of codes obtaining numerical solutions of an adiabatic system of differential equations describing stellar oscillations have been developed. In order to ascertain whether any possible difference in the description of the same observational data by different numerical codes is due to physical descriptions or to different numerical integration schemes, the inter-comparison of these oscillation codes in a fixed and homogeneous framework is absolutely necessary. Some effort has been already done in the past but only regarding pairs of codes. Some codes have also developed a lot of internal precision tests. However, there is a lack of inter-comparison of a large enough set of codes. We aim in this study try to fix a set of minimum requirements for a code to be sure that any difference found is only due to a different physical assumption.  Ideally, for a given model there should be a set of `true' frequencies with which the results of the different codes could be compared.  This ideal situation could probably be approximated by considering polytropic models for which it is relatively straightforward to calculate the equilibrium structure with essentially arbitrary accuracy \\citep[see also][]{Christ1994}.  In practice, the situation for realistic stellar models is more complex.  Owing to the complexity of the stellar evolution calculation the models are often available on a numerical mesh which is not obviously adequate for the pulsation calculation.  The effect of this on the frequency computation depends on the detailed formulation of the equations in the pulsation codes. These formulations are equivalent only if the equilibrium model satisfies the `dynamical' equations of stellar structure, i.e., the mass equation and the equation of hydrostatic support, and this is obviously not exactly true for a model computed on a finite (possibly even relatively sparse) mesh.  One might define a consistent set of frequencies for a given model by interpolating it onto a very dense mesh and resetting it to ensure that the relevant equations of stellar structure are satisfied.  The model is fully characterized by the dependence on distance $r$ to the centre of density $\\rho$ and the adiabatic exponent $\\Gamma_1 = (\\partial \\ln P/ \\partial \\ln \\rho)_{\\rm ad}$, $P$ being pressure and the derivative being at constant specific entropy.  Thus one could interpolate $\\rho(r)$ and $\\Gamma_1(r)$ to a fine mesh, and recompute the mass distribution and pressure by integrating the mass equation and the equation of hydrostatic equilibrium.  Frequencies of this model should then be essentially independent of the formulation of the oscillation equation and would provide a suitable reference with which to compare other frequency calculations.  Such a test may be carried out at a later stage in the ESTA effort.  In Task~2, for now, we aim at testing, comparing and optimizing the seismic codes used to calculate the oscillations of existing models of different types of stars. In order to do so we consider steps in the comparison by addressing some of the most relevant items that must be compared regarding the seismic characterization of the models:  \\begin{itemize} \\item Step 1: comparison of the frequencies from different seismic codes for the same model. \\item Step 2: comparison of the frequencies from the same seismic code for different models of the same stellar case provided by different equilibrium codes. \\item Step 3: comparison of the frequencies for specific pulsators. \\end{itemize}  The work presented here is mostly focused on step~1. At this step three different equilibrium models have been used. Two of them have been computed using \\CESAM\\ \\citep{cesam}, % with 902 and 2172 mesh points, and a third one with 4042 mesh points provided by \\ASTEC\\ \\citep{astec}.% We present inter-comparisons using the two models with the larger numbers of mesh points. The same physics and physical constants (except $G$) are used for all the oscillation codes. Frequencies in the range of $[20,2500] \\,\\mu$Hz, belonging to spherical degrees $\\ell=0,1,2$ and 3 have been calculated, in order to recover most of the possible values we can find in the observational photometric data.  We present in Sect. 2 the equilibrium models used and we analyze their main features. Section 3 is devoted to the different oscillation codes used,   discussing   common  requirements   and   variations  in   the treatment.  In  Sect.  4  the  direct comparison  of  frequencies  for different  frequency  ranges and  spherical  degrees  is presented  in detail. Sections  5 and  6 analyze the  same inter-comparison  for the values of  the large and  small separations, respectively.   Section 7 discusses the  dominant effects that contribute to  the differences in frequencies determined  by the different codes. Conclusions are given in Sect. 8. ",
        "conclusions": "A complete inter-comparison of the frequencies calculated by using a set of oscillation codes has been presented.  Nine oscillation codes have been used: \\ADIPLS, \\FILOU, \\NOSC, \\LOSC, \\GRACO, \\OSCROX, \\PULSE, \\POSC\\ and \\ROMOSC. The only free parameter for each code has been the integration numerical scheme. In order to ensure that the same physics has been imposed to all the codes two equilibrium models of a $1.5M_\\odot$ star have been supplied, with 4042 (M4k) and 2172 (M2k) mesh points in their grids, respectively. No re-meshing has been allowed.  The model M2k pre\\-sents inadequate numerical resolution in the Brunt-V\\\"ais\\\"al\\\"a frequency close to the boundary of the convective core.  The model M4k does not suffer from these problems. In the present paper we have used these models only to show the sensitivities of the oscillation codes to possible numerical inaccuracies in the equilibrium models.  A complete study of the accuracies of the different evolutionary codes is provided by Lebreton et al. in this volume.  Several inter-comparisons have been performed: 1) Direct frequency inter-comparison, 2) large separation (LS), 3) small separation (SS), and 4) different experimental values of the gravitational constant $G$.  The main conclusions can be summarized as:  \\begin{itemize}  \\item When codes using the same numerical integration sche\\-mes are compared, the general precision obtained for model M4k ($0.02 \\,\\mu$Hz) is higher than the expected precision of the observational data obtained with the CoRoT mission. This precision is always one order of magnitude better or even more when compared with that obtained for model M2k ($0.5 \\, \\mu$Hz). The largest differences are obtained for the high-frequency region.  \\item The use of a second-order integration scheme plus Richardson extrapolation or a fourth-order integration scheme does not improve the agreement between the codes when compared with the use of a second-order scheme. In the cases here presented it was not possible to distinguish between the use of a fourth-order scheme or Richardson extrapolation, which always give very acceptable consistency for model M4k.  \\item However, the use of a fourth-order integration scheme (or a second-order plus Richardson extrapolation) introduces differences larger than $0.1 \\,\\mu$Hz ($1.5 \\,\\mu$Hz for M2k and $0.8 \\,\\mu$Hz for M4k) compared with a second-order scheme, mainly located in the regions of high order (high frequency for p-modes and low frequency for g-modes). This indicates that the second-order schemes have inadequate numerical precision on the mesh provided in M4k and, obviously, even more so on the sparser mesh in M2k.  \\item For ``large separations\" the agreement between different comparable codes using model M4k, with every possible numerical scheme is generally better than the expected precision of the data. However, the equation-of-state inconsistency in model M4k (see Sect.~\\ref{sec:dep}) introduces a signature that it potentially significant. The use of the Richardson extrapolation together with a second-order integration scheme, or a fourth-order scheme, gives differences for high-order p- and g-modes relative to the use of only a second-order scheme that are comparable to the expected observational accuracies, indicating that the accuracy of the latter scheme is inadequate.  \\item The ``small separations'' present differences for model M4k lower than the expected observational accuracies everywhere, although the inconsistency in M4k has a noticeable effect. Use of Richardson extrapolation has a potentially significant effect, indicating that the precision of second-order schemes may be inadequate.  \\item The use of $r/P$ as integration variable instead of $r$ has a small influence when a consistent value of $G$ is used. However, an inconsistency between the $G$ values used to obtain the equilibrium model and the oscillation frequencies can give substantial differences between the use of these integration variables, over the range of $G$ values in Table 2.  \\item For model M2k, the mixed modes near the avoided crossing present large differences, reaching even up to $4 \\,\\mu$Hz. These differences are not present when model M4k is used. The problems with $N^2$ in M2k are undoubtedly the main source of this difference. The same behaviour is found for the trapped g-modes. Therefore, the correct numerical description of $N^2$ is critical for the value of the frequencies of these trapped modes or in avoided crossing.  \\item The use of $\\delta P$ or $P^\\prime$ as eigenfunction, or the solution of the LAWE, may have a significant effect if the equilibrium model is based on thermodynamic quantities that are not internally consistent, as is the case for the OPAL tables used to calculate model M4k. Further effects arise when there are problems with the Brunt-V\\\"ais\\\"al\\\"a frequency (as for model M2k); here a different choice can change the frequencies of some modes since $N^2$ does not appear in the differential equations when $\\delta P$ or the LAWE are used and it does when using $P^\\prime$.  \\item The value of the gravitational constant $G$ in the oscillation calculations can introduce non-negligible differences as well, if it is not the same as the value used in the equilibrium model. When the two extreme values found in the literature are used, such inconsistency yields differences in the range $[-0.35,-0.08] \\,\\mu$Hz for model M4k, larger than those obtained when different numerical integration schemes are used. Differences between the values of $G$ actually used by the different codes, although less extreme, account for much of the difference between the computed frequencies.  \\end{itemize}  Therefore, for a proper pulsational study, we require that the number of mesh points and their distribution must be such as to yield an equilibrium model that satisfies the dynamical equations with sufficient accuracy also in the regions of the star where the physical quantities present rapid variations (e.g., the outer layers and $\\mu$-gradient zones). In addition, the mesh used in the pulsation calculation must properly resolve the eigenfunctions of the highest-order modes considered. In the present case model M4k, with 4042 points, appears to satisfy these conditions although the use of a fourth-order integration scheme, or a second-order scheme and Richardson extrapolation, is still needed in the oscillation calculation; these higher-order schemes give significant improvements compared with the use of a simple second-order scheme. The use of a second-order integration plus Richardson extrapolation scheme is not distinguishable, in accuracy and precision, from the use of a fourth-order integration scheme. A correct physical and numerical description of the Brunt-V\\\"ais\\\"al\\\"a frequency is essential when $P^\\prime$ is used as eigenfunction; in particular, inconsistencies in the equation of state can have serious effects on the frequencies. Inconsistency between the value of $G$ used in the oscillation calculation and the value used to compute the equilibrium model, within the range of the different values of $G$ found in the literature, may lead to substantial errors in the computed frequencies.  We note that the situation is somewhat different if consistent values of $G$ are used in the evolution and oscillation calculations. Then the effect on the frequencies is approximately given, according to homology arguments, as a scaling by $(GM)^{1/2}$. However, since the product $G M_\\odot$ is known extremely precisely from planetary motion in the solar system, any variation in $G$ should be reflected in a corresponding change in the assumed value of $M_\\odot$. If this is the case, and if the model is characterized by a given value of $M/M_\\odot$ (as is typically the case) the effect on the frequencies of changes in $G$ are very small \\citep[see also][]{Christ2005}.  In further tests more care is required to secure the full consistency of the models: a consistent equation of state should be used, and the value of $G$ should obviously be the same in the equilibrium-model and the pulsation calculations; indeed, this strongly argues for including the value of $G$ as one of the parameters in the model file. The main conclusion of this extensive investigation, however, it positive: with a properly resolved equilibrium model the broad range of oscillation codes likely to be involved in the asteroseismic analysis of data from CoRoT and other major upcoming projects generally give consistent results, well within the expected errors of the observations. Thus, although the remaining problems in the calculation evidently require attention, we can be reasonably confident in our ability to compute frequencies of given models and hence in the inferences concerning stellar structure drawn from comparing the computed frequencies with the observations."
    },
    "0711/0711.4474_arXiv.txt": {
        "abstract": "In the framework of the search of dark matter in galactic halos in form of massive compact halo object (MACHOs), we discuss the status of microlensing observations towards the Magellanic Clouds and the Andromeda galaxy, M31. The detection of a few microlensing events has been reported, but an unambiguous conclusion on the halo content in form on MACHOs has not been reached yet. A more detailed modelling of the expected signal and a larger statistics of observed events are mandatory in order to shed light on this important astrophysical issue. ",
        "introduction": "Gravitational microlensing, as first noted in \\cite{ref:pacz86}, is a very efficient tool for the detection and the characterisation of massive astrophysical halo compact objects (MACHOs), a possible component of dark matter halos. Following the first exciting detection of  microlensing events \\cite{ref:macho93,ref:eros93,ref:ogle93}, by now the detection of $\\sim 30$ events have been reported towards the Magellanic Clouds and our nearby galaxy, M31, and first interesting conclusions on this issue have been reported (Section~\\ref{sec:LMC} and Section~\\ref{sec:M31}). Soon enough, however, the Galactic bulge probed to be an almost an interesting target \\cite{ref:pacz91}, and indeed by now the number of observed microlensing events along this line of sight exceeds by two order of magnitudes that observed towards the Magellanic Clouds and M31. In that case the contribution from the dark matter halo is expected to be extremely small compared to that of either bulge or disc (faint) stars \\cite{ref:griest91}. Microlensing searches towards the Galactic bulge are therefore important as they allow to constrain the inner Galactic structure \\cite{ref:pacz94}. Recently, the MACHO \\cite{ref:popowski05}, OGLE \\cite{ref:sumi06} and EROS \\cite{ref:hamadache06} collaborations presented the results of their observational campaigns towards this target. A remarkable conclusion is the agreement among these different searches as for the observed value of the optical depth and the agreement with the theoretical expectations \\cite{ref:evans02,ref:hangould03}. For a more recent  discussion see also \\cite{ref:novati07}, where the issue of the bulge mass spectrum is treated. ",
        "conclusions": "Microlensing searches towards the Magellanic Clouds and the Andromeda galaxy, M31, have given, in recent years, first exciting though somewhat contradictory conclusions. The detection of microlensing events has been reported, however their interpretation with respect to the halo dark matter issue is still open to debate. Along both line of sights ``evidence'' for a MACHO signal as well as null results have been reported. (We may stress that, in the framework of galaxy formation theory, even a relatively ``small'' halo fraction contribution in form of MACHOs, say at the 10\\%-20\\% level, may turn out to be  relevant). New observational campaigns are currently under way to further address this interesting issue. The SuperMACHO collaboration \\cite{ref:rest05} is observing the LMC with a much larger field of view than previous campaigns. Towards M31, the Angrstom \\cite{ref:kerins06} and the PLAN \\cite{ref:loiano07} campaigns are currently underway. Both, in different ways, aims rather to the central M31 region, so to properly characterise the self-lensing signal. Furthermore, as opposed to previous analyses, there is an effort to get to a more suitable sampling of the observational data so to allow a better reconstruction of the microlensing event parameters, in particular of the Einstein time."
    },
    "0711/0711.4018_arXiv.txt": {
        "abstract": "We report the discovery of burst oscillations at 414.7 Hz during a thermonuclear X-ray burst from the low mass X-ray binary (LMXB) 4U 0614+091 with the Burst Alert Telescope (BAT) onboard {\\it Swift}. In a search of the BAT archive, we found two burst triggers consistent with the position of 4U 0614+091. We searched both bursts for high frequency timing signatures, and found a significant detection at 414.7 Hz during a 5 s interval in the cooling tail of the brighter burst. This result establishes the spin frequency of the neutron star in 4U 0614+091 as $\\approx 415$ Hz.  The oscillation had an average amplitude (rms) of $14 \\%$, These results are consistent with those known for burst oscillations seen in other LMXBs. The inferred ratio of the frequency difference between the twin kHz QPOs, and the spin frequency, $\\Delta\\nu / \\nu_s$, in this source is strongly inconsistent with either 0.5 or 1, and tends to support the recent suggestions by Yin et al., and Mendez \\& Belloni, that the kHz QPO frequency difference may not have a strong connection to the neutron star spin frequency. ",
        "introduction": "The detection of spin modulation pulsations during thermonuclear X-ray bursts, ``burst oscillations,'' has become an important method for measuring the spin rates of neutron stars in accreting binary systems (see Strohmayer \\& Bildsten 2006 for a review). To date, almost all of the {\\it a priori} detections of burst oscillations have been made with the Proportional Counter Array (PCA) onboard the {\\it Rossi X-ray Timing Explorer} (RXTE). Indeed, its combination of large collecting area, low background, high time resolution, and high telemetry capacity enabled the first detections of the phenomenon. Pulsations during a bright burst from SAX J1808.4-3568 were detected with the Wide Field Camera (WFC) on {\\it BeppoSAX} at the $\\approx 3\\sigma$ level, but in this case the spin frequency was already known (in 't Zand et al.  2003).  The Burst Alert Telescope (BAT) onboard {\\it Swift} has a combination of large collecting area (5,200 cm$^2$) and high time resolution (100 $\\mu$-sec) that should, in principle, make it sensitive to burst oscillations from particularly bright X-ray bursts, although the BAT does suffer a significant loss of sensitivity to such oscillations due to its much larger background compared to the PCA. The BAT has a sophisticated triggering system that identifies Gamma-ray bursts (GRBS), and other fast transients, including thermonuclear bursts from accreting neutron stars. When a trigger occurs, an estimate of the source's position is derived onboard the {\\it Swift} spacecraft, where it is checked against an onboard catalog of known objects. If the position is consistent with a source in the catalog, then the trigger is deemed to not be a GRB, and {\\it Swift} will not initiate a direct slew to the position. However, high time resolution data is still accumulated around the trigger time. These data can then be searched for burst oscillations.  Historically, some of the brightest X-ray bursts observed, with fluxes of order $2 \\times 10^{-7}$ ergs cm$^2$ s$^{-1}$, have been associated with the low mass X-ray binary (LMXB) 4U 0614+091. Based on the detection of a bright burst with {\\it Watch}, and Eddington limit arguments, Brandt et al. (1992) argued that the distance to 4U 0614+091 is likely $< 3$ kpc. Several lines of evidence point to the conclusion that 4U 0614+091 is an ultra-compact X-ray binary system. Juett, Psaltis \\& Chakrabarty (2001) inferred an enhanced neon to oxygen ratio for the system, and based on a comparison of the ratio with known ultra-compact systems, argued that 4U 0614+091 is a similar system. Sensitive optical spectroscopy by Nelemans et al. (2005) found carbon and oxygen emission lines, but no evidence for helium or hydrogen, further strengthening the ultra-compact classification.  The optical counterpart to 4U 0614+091, V1055 Ori, is also intrinsically faint, consistent with a compact orbit (Nelemans et al. 2005; van Paradijs \\& McClintock 1994).  The faintness of its persistent X-ray emission and nearby distance suggest a low accretion rate, again consistent with an orbital period $< 1$ hr (Deloye \\& Bildsten 2003). Finally, we note that in 't Zand et al. (2007) cite a report by O'Brien (2005) of a 50 min. optical modulation, that could represent the orbital period.  Although 4U 0614+091 has been extensively studied with RXTE, it has not yet observed any X-ray bursts from the source.  Ford et al. (1997) reported the discovery of kHz quasiperiodic oscillations (QPOs) from the source using RXTE data. They found a maximum QPO frequency of $\\approx 1,145$ Hz, and a mean frequency separation of $323 \\pm 4$ Hz between the upper and lower kHz oscillations (see van der Klis 2006, for a review of kHz QPOs). Based on this frequency separation they suggested a spin frequency of $\\approx 323$ Hz.  Recent studies, however, have suggested that the frequency separation between the kHz QPOs may be more or less independent of the neutron star spin frequency (Yin et al. 2007; Mendez \\& Belloni 2007).  In this Letter we report the discovery of 415 Hz burst oscillations in a burst from 4U 0614+091 observed with the BAT instrument onboard {\\it Swift}.  In \\S 2 we describe our timing study of bursts from 4U 0614+091 with BAT and summarize their basic properties. In \\S 3 we discuss our results in the context of recent efforts to understand the relationship (if any) between the kHz QPOs and the spin frequencies of neutron stars in LMXBs. ",
        "conclusions": "We have found strong evidence for a burst oscillation at 414.75 Hz in a thermonuclear burst from 4U 0614+091 with the BAT onboard {\\it Swift}. This is the first detection of burst oscillations with BAT and demonstrates its potential in this regard, particularly for relatively nearby (read bright) burst sources.  This detection implies a spin frequency of about 415 Hz for the neutron star in this system.  Since the discovery of kHz QPOs and burst oscillations, there has been a suggested link between the neutron star spin frequency, $\\nu_s$, (as indicated by burst oscillation frequencies, or persistent pulsation frequencies, or both) and the frequency difference of the twin kHz QPOs, $\\Delta\\nu / \\nu_s$. Initial discoveries suggested that the kHz QPO frequency separation was consistent with the spin frequency, but subsequent observations have revealed a more complex picture (Strohmayer et al. 1996; van der Klis 2006). The tentative consensus at present is that the frequency separation is related to the spin frequency (at least in some sources), in the sense that $\\Delta\\nu / \\nu_s$ can be close to either 1 or 0.5.  While this description fits reasonably well with several of the observed systems, it is not universally precise, in the sense that some objects show statistically significant deviations from 1 or 0.5, and that the frequency separation is known to vary in some objects. Recently, Mendez \\& Belloni (2007), and Yin et al. (2007) have re-examined the existing measurements of neutron star spin frequencies and kHz QPO, and suggest that there may be only a weak connection between the QPO frequency difference and the spin frequency.  Indeed, Mendez \\& Belloni (2007) hypothesize that the frequency difference may be more or less constant, essentially independent of the spin frequency.  With the detection of the spin frequency in 4U 0614+091 we now have an additional source in which to closely evaluate the relationship between kHz QPOs and spin. Interestingly, the burst oscillation frequency we have found is very significantly different from the frequency separation of $323 \\pm 4$ Hz of the twin kHz QPOs reported for the system (Ford et al. 1997).  In this case the QPO difference frequency was quite accurately determined, and since the spin frequency is not thought to differ from observed burst oscillation frequencies by more than a few Hz, this is a strong demonstration that the kHz QPO difference frequency and spin frequency in 4U 0614+091 do not follow the ``consensus'' model.  Assuming that the spin frequency is within 4 Hz of our measured burst oscillation frequency, then we arrive at $\\Delta\\nu / \\nu_s = 0.78 \\pm 0.02$, which is inconsistent with both 0.5 and 1 at very high confidence.  Interestingly, this value lies much closer to the two alternative relations suggested by Yin et al.  (2007, $\\Delta\\nu = 390 \\; {\\rm Hz} - 0.2\\nu_s$), and Mendez \\& Belloni (2007, $\\Delta\\nu = 308 \\; {\\rm Hz}$), although it stills fall about $2\\sigma$ away from these as well.  Since these relations were derived and or suggested prior to our discovery, the spin meausurement reported here can be thought of as an independent check on those models, and it appears they provide a much better description of the data than the ``consensus'' model, further supporting the notion that it may be past time to re-think the relationship between the kHz QPOs and neutron star spin, or to at least adopt the working hypothesis that the relationship does not hold for all sources, or at all times."
    },
    "0711/0711.1910_arXiv.txt": {
        "abstract": "{ The Milagro experiment observes the extensive air showers produced by very high energy $\\gamma$-rays impacting the Earth's atmosphere. Milagro uses 898 Hamamatsu R5912 Photomultiplier Tubes. To complete our Monte Carlo simulations, we tested the photocathode uniformity of our PMTs. The main finding was that the PMT gain and detection efficiency are a function of the distance from the center of the photocathode. Both quantities become considerably smaller as the illumination position nears the edge of the photocathode. }  \\email{vlasisva@umd.edu} \\begin{document} ",
        "introduction": "Milagro \\cite{milagro} is a water-Cherenkov detector at an elevation of 2650m at the Jemez Mountains in New Mexico. It comprises a central 60m x 80m x 8m pond surrounded by a 200m x 200m array of 175 ``outrigger'' tanks. The pond, covered with a light barrier, is instrumented with two layers of photomultiplier tubes (PMTs). The top ``air-shower'' layer consists of 450 PMTs while the bottom ``muon'' layer has 273 PMTs. Each outrigger tank contains $\\sim$4000l of water and one PMT. The PMTs collect the Cherenkov light produced by the air shower particles, as they transverse the detector's water volume. The AS and OR layers allow the accurate measurement of the air shower particle arrival times used for reconstructing the direction of the shower-initiating particle. The greater depth of the muon layer ($\\sim$17 radiation lengths) is used to distinguish deeply penetrating muons and hadrons, which are common in hadron induced air showers, from electrons and $\\gamma$-rays.  The motivation for the uniformity tests described here was a disagreement between the single-muon response of the muon-layer PMTs and a Monte Carlo simulation of this response. According to the PMT tests performed by the IceCube experiment\\footnote{http://www.ppl.phys.chiba-u.jp/research/IceCube/docs/presentations/index.html}, the photocathode response of their 10' PMTs is not uniform all over the photocathode's surface. Before the experimental work reported in this paper, the Milagro PMTs were simulated as having the same properties all over the face of the PMT. We thought it possible that the way the PMTs were simulated was the cause of the disagreement regarding the single-muon response of the muon-layer PMTs.  For that reason, an apparatus was constructed to examine the dependence of the gain and detection efficiency on the distance of the photon-detection position from the center of the photocathode. ",
        "conclusions": "The photocathode uniformity of the Hamamatsu R5912 PMTs used in the Milagro experiment was examined. The tests showed that the PMT gain and detection efficiency are a function of the photon-detection position. Both quantities are smaller when points near the edge of the photocathode are illuminated. This effect was present in all tested PMTs and was not negligible in magnitude.  Before this study, the Milagro PMTs were simulated as having the same properties all over the face of the PMT. After using the experimental results in the simulation of the PMTs, the simulation now predicts that an atmospheric muon produces $\\sim$110PEs (down from $\\sim$200PEs) in the PMTs of the muon layer compared with $\\sim$100PEs at the experiment."
    },
    "0711/0711.3124_arXiv.txt": {
        "abstract": "On the cosmological length scale, recent measurements by WMAP have validated $\\Lambda$CDM to a precision not see before in cosmology. Such is not the case on galactic length scales, however, where the `cuspy-core' problem has demonstrated our lack of understanding of structure formation. Here, we propose a solution to the 'cuspy-core' problem based on the observation that with the discovery of Dark Energy, $\\Lambda_{DE}$, there is now a universal length scale, $\\lambda_{DE}=c/(\\Lambda_{DE} G)^{1/2}$, associated with the universe. This length scale allows for an extension of the geodesic equations of motion that affects only the motion of massive test particles; the motion of massless test particles are not affected, and such phenomenon as gravitational lensing remain unchanged. An evolution equation for the density profile is derived, and an effective free energy density functional for it is constructed. We conjecture that the pseudoisothermal profile is preferred over the cusp-like profile because it has a lower effective free energy. A cosmological check of the theory is made using the observed rotational velocities and core sizes of 1393 spiral galaxies. We calculate $\\sigma_8$ to be $0.68_{\\pm0.11}$; this is within experimental error of the WMAP value $0.761_{-0.048}^{+0.049}$. We then calculate $R_{200}=270_{\\pm 130}$ kpc, which is in agreement with observations. We estimate the fractional density of matter that cannot be determined through gravity to be $0.197_{\\pm 0.017}$; this is nearly equal to the WMAP value for the fractional density of nonbaryonic matter $0.196^{+0.025}_{-0.026}$. The fractional density of matter that can be determined through gravity is then calculated to be $0.041_{-0.031}^{+0.030}$; this is nearly equal to $\\Omega_B=0.0416_{-0.0039}^{+0.0038}$. ",
        "introduction": "The recent discovery of Dark Energy (see \\cite{Ries1998, Perl1999} and references therein) has broadened our knowledge of the universe, and has demonstrated once again that it can hold surprises for us. The discovery has, most assuredly, also brought into sharp relief the degree of our understanding of it. Only a small fraction of the mass-energy density of the universe is made up of matter that we have characterized; the rest consists of Dark Matter and Dark Energy, and the precise properties of either one is not known. They are nevertheless needed to explain what is seen on an extremely wide range of length scales. On the galactic ($\\sim 100$ kpc parsec), galactic ($\\sim$ ~10 Mpc), and  supercluster ($\\sim$ 100 Mpc) scales, Dark Matter has been used to explain phenomena ranging from the formation of galaxies and their rotation curves, to the dynamics galaxies and the formation galactic clusters and superclusters. On the cosmological length scale, Dark Matter and Dark Energy is needed to explain the observed evolution of the universe from the Big Bang to the present, and will determine its fate in the future.  Observations thus tell us that over a vast range of length scales the dynamics and evolution of the observed universe is determined not by normal matter, but by Dark Matter and now Dark Energy (see \\cite{Cahi} for a quantum-cosmology approach that does not require Dark Matter or Dark Energy). Yet while the need and invocation of Dark Matter is ubiquitous on a wide range of length scales, our understanding of how matter determines dynamics at the galactic length scale is lacking. Recent measurements by WMAP \\cite{WMAP} have validated the $\\Lambda$CDM model of cosmology to an precision not seen before in cosmology. The situation on the galactic scale is not nearly as settled, however. Here, the cuspy-core problem for the density profile of matter in galaxies (\\cite{JNav, Krav, Moore}, and \\cite{PeeblesRev, Silk} for reviews) has demonstrated our lack of understanding of the formation of galactic rotation curves.  Current understanding of the structure formation is based on the work of Peebles \\cite{Peebles1984}, where the seeds of galaxies are due to local fluctuations in the density of matter that grow as the universe expands. Analytical solutions of this model have been done \\cite{Gunn, Fill, Hoff1985, Hoff1988} for a number of special cases, and have resulted in density profiles that are sensitive to initial conditions, and have a power-law dependence whose exponents vary over a range of values. More recently, numerical simulations \\cite{JNav, Krav, Moore} of galaxy formation have been done, and have consistently resulted in density profiles with a cusp-like structure \\begin{equation} \\rho_{\\hbox{\\scriptsize Simulation}} =\\frac{\\rho_i}{(r/R_S)^\\gamma(1+(r/R_S)^\\alpha)^{(\\beta-\\gamma)/\\alpha}}, \\label{cusp} \\end{equation} instead of the expected pseudoisothermal density profile. Here, $\\rho_i$ is a density parameter, $R_s$ is related to the radius of the galactic core, and the exponents $(\\alpha, \\beta, \\gamma)$ take a range of values from $(1.5, 2, 1.5)$ for Moore et.~al.~\\cite{Moore} to $(1.0, 3.0, 1.0)$ for Navarro, Frenk and Wilson (NFW) \\cite{JNav} (see \\cite{Silk} for review). This lead Moore to state in \\cite{Moore} that cold dark matter fails to reproduce the galactic rotation curves for dark-matter-dominated galaxies, one of the key reasons that dark matter was proposed in the first place. Soon afterwards, de Blok and coworkers \\cite{Blok-1, Blok-2, McGa} explicitly demonstrated that the NFW density profile does not fit the density profile observed for Low Surface Brightness (LSB) galaxies (see also \\cite{Gent-2007} for a recent analysis of cusp structure). Rather, the traditional pseudoisothermal profile, with $(\\alpha, \\beta, \\gamma)= (2,2,0)$, is the better fit. This demonstration is especially compelling as it is believed that dark matter dominates dynamics in LSB galaxies.  There have been a number of attempts to solving the cuspy-core problem within $\\Lambda$CDM \\cite{Bode, Dave, Somm, Sper2000}, and they have had varying degrees of success (see \\cite{PeeblesRev} for a review). This problem does not exist in Milgrom's Modified Newtonian Dynamics (MOND) \\cite{Mil-1, Mil-2, Mil-3} (see \\cite{Sand} for a review)\\textemdash a theory where Dark Matter is not needed\\textemdash but there are a number of theoretical and observational problems that MOND must overcome (see \\cite{Sell} for arguments in support of MOND, however).  Our approach to solving the cuspy-core problem, and to structure formation in general, is much more drastic; therefore, its reach is correspondingly broader. It is based on the observation that with the discovery of Dark Energy, $\\Lambda_{DE}$, we now have a \\text{universal} length scale, $\\lambda_{DE} = c/(\\Lambda_{DE}G)^{1/2}$, on hand \\footnote{It is possible to also construct the length scale $(\\hbar c/\\Lambda_{DE})^{1/4}\\approx 85$ $\\mu$ m. Experiments have been shown that this scale does not affect the Newtonian potential. \\cite{Adel2007}}. The geodesic equations of motion (GEOM)\\textemdash and thus the geodesic action\\textemdash is no longer unique, and extensions of it through the introduction of functions of $Rc^2/\\Lambda_{DE}G$ can be made. While there have been attempts at proving that the GEOM are the unique consequence of the Einstein field equations \\cite{Einstein, Gero1972, Gero2003}, such proofs assume that the background metric remains fixed under the passage of the test particle.  As our extension of the depends explicitly on the energy-momentum tensor of matter\\textemdash which includes the motion of the test particle\\textemdash these proofs do not preclude our extension of the GEOM.  In form, our extension of the GEOM preserves the equivalence principal, and through the choice of the function of $Rc^2/\\Lambda_{DE}G$, we can insure that their effects are not measurable on terrestrial scales. Physically, that this choice is possible is because $\\Lambda_{DE} = (7.21^{ +0.83}_{-0.84}) \\hbox{ x } 10^{-30}$ $g/cm^3$ (from \\cite{WMAP}) is so much smaller than any density of matter either presently achievable experimentally, or present in regions of space accessible to experiment. Correspondingly, $\\lambda_{DE} = 14020^{790}_{810}$ Mpc is much longer than the scale of any experiment that has been used to test general relativity. In fact, the issue is to make the theory relevant at \\textit{galactic} scales of a few kpc, and by doing so we arrive at an estimate for the exponent $\\alpha_\\Lambda$. This exponent is the only parameter in the theory, and it determines the power-law behavior of our extension of the GEOM. We also find that while affecting the motion of massive test particles, our extension does not affect the motion of massless test particles; photons still travel along null geodesics, and gravitational lensing and the deflection of light are left unchanged.  Applied to galaxy formation, the extended GEOM reduces to a nonlinear evolution equation for the density profile of a model galaxy in the nonrelativistic, linear gravity limit. This evolution equation minimizes a functional of the density, which is interpreted as an effective free energy functional for the system. Solutions to this equation is found using various velocity curves for galaxies as driving terms, and these solutions are then used to calculate the free energy associated with various profiles. We conjecture that like Landau-Ginzberg phenomenological theories in condensed matter physics, the system prefers to be in a state that minimizes this free energy. Showing that the pseudoisothermal profile is preferred over cusp-like profiles reduces to showing that it has the lower free energy.  In our model of a galaxy, the Hubble length scale $\\lambda_H = c/H$ (where $H = h H_0$, is the Hubble constant, $h = 0.732^{+0.031}_{- 0.032}$ and $H_0 = 100$ $km/s/$Mpc \\cite{WMAP}) naturally appears, \\textit{even though a cosmological model is not mentioned either in its construction or in its analysis}. What happens at galactic length scales are naturally tied to what happens  at cosmological length scales with our approach; the combination of the Dark Energy length scale and the nonlinear aspects of the extended GEOM link the two. This linkage allows us to extrapolate from the statistical properties of the observed universe the properties of a representative galaxy. These properties are then used to provide a cosmological check of the theory.  As with the Peebles model, the total density of matter for our model galaxy can be written as a sum of a background density $\\rho_{\\hbox{\\scriptsize asymp}}(\\mathbf{x})$ and a linear perturbation $\\rho_{II}^{1}(\\mathbf{x})$. As usual, $\\rho_{\\hbox{\\scriptsize asymp}}(\\mathbf{x})$ does not contribute to the motion of stars within the galaxy, while $\\rho_{II}^{1}(\\mathbf{x})$ does; in the absence of forces other than gravity, observations of the dynamics of the stars will determine only $\\rho_{II}^{1}(\\mathbf{x})$. But unlike the Peebles model, $\\rho_{\\hbox{\\scriptsize{asymp}}}$ is not a constant. With it, we are able to estimate $\\Omega_{\\hbox{\\scriptsize{asymp}}}$, the fractional density of matter that \\textit{cannot} be determined through gravity, to be $0.197_{\\pm0.017}$; this is nearly equal to the measured value of the fractional density of nonbaryonic (dark) matter in the universe $\\Omega_m-\\Omega_{B} = 0.196^{+0.025}_{-0.026}$ measured by WMAP \\cite{WMAP}. Correspondingly, we estimate $\\Omega_{\\hbox{\\scriptsize{Dyn}}}$, the fractional density of matter in the universe that \\textit{can} be determined through gravity, to be $0.041^{+0.030}_{- 0.031}$; this is nearly equal to the value of $0.0416^{+0.0038}_{- 0.0039}$ for $\\Omega_B$ measured by WMAP \\cite{WMAP}. We have also calculated $\\sigma_8$, the rms fluctuation in the fractional density of matter within a distance of $8 h^{-1}$ Mpc, as a direct check of our model. Using the average rotational velocity and core sizes of 1393 galaxies obtained through four different set of observations \\cite{Blok-1, Blok-2, McGa, Rubin1980, Rubin1982, Burs, Rubin1985, Cour, Math} that span 25 years, we obtain the value of $0.68_{\\pm 0.11}$ for $\\sigma_8$; this value is within experimental error of the measured value of $0.761^{+0.049}_{-0.048}$ by WMAP. Finally, we have calculated $R_{200}$, the radius of the galaxy at which the density equals 200 times that of the critical density, to be $270_{\\pm 130}$ kpc; this value also agrees well with observations.  Interestingly, $\\Omega_{\\hbox{\\scriptsize asymp}}/\\Omega_\\lambda$ depends only on the dimensionality and symmetry of spacetime, and the exponent $\\alpha_\\Lambda$. This suggests that there is an underlying coupling between Dark Energy and matter in the theory. Such a coupling has been dismissed before, primarily because it is believed that the coupling would result in a ``fifth force'' that would already have been observed \\cite{PeeblesRev}. The results of this paper suggest that with a suitable choice of this coupling, its effects will not be currently measurable.  A summary of the results here have appeared elsewhere \\cite{ADS}. Here, we provide the details of the theory and the calculations. ",
        "conclusions": "Given how sensitive of our expression for $\\sigma_8$ is dependent on $v_H^{*}$, $r_H^{*}$, and $\\alpha_\\Lambda$, that our predicted values of $\\sigma_8$ is within experimental error of its measured value is surprising. This is especially true as the data used in calculating $\\sigma_8$ was taken by four different groups over a period of 25 years, and for purposes that have no connection whatsoever with our analysis. Even in the absence of a direct experimental search for $\\alpha_\\Lambda$, this provides a compelling argument for the validity of our extension of the GEOM, and its impact on structure formation. This agreement also supports our free energy conjecture; the calculated $\\sigma_8^2$ would be very different if $\\beta=3$, say, were used in calculating $\\sigma_8^2$ instead of $\\beta=2$.  Direct detection and measurement of $\\alpha_\\Lambda$ through terrestrial experiments may be possible in the near future. As mentioned in Sections \\textbf{II.D} and \\textbf{III.B}, at a value of $1.51$ the exponent $\\alpha_\\Lambda$ is likely small enough that the effects of the additional terms in the extended GEOM may soon be detectable.  Interestingly, $\\Omega_m-\\Omega_{B} = 0.196_{-0.026}^{+0.025}$ is nearly equal to $\\Omega_{\\hbox{\\scriptsize asymp}}$ in value. Correspondingly, $\\Omega_B$ \\cite{WMAP} is nearly equal to $\\Omega_{\\hbox{\\scriptsize Dyn}}$. It would be tempting to identify $\\Omega_{\\hbox{\\scriptsize asymp}}$ with the fractional density of nonbaryonic (dark) matter in the universe, especially since matter in $\\rho_{\\hbox{\\scriptsize{asymp}}}(r)$ does not participate in the particle dynamics, and is not ``visible'' to measurements that inferred mass through particle motion under gravity. That $\\Omega_{\\hbox{\\scriptsize Dyn}}$ would then be identified with $\\Omega_B$ is consistent with the observation that most of the mass that has been inferred through gravitational dynamics indeed consists of baryons. We did not differentiate between normal and dark matter in our theory, however. Without a specific mechanism for funneling nonbaryonic matter into $\\rho_{\\hbox{\\scriptsize{asymp}}}$ and baryonic matter into $\\rho-\\rho_{\\hbox{\\scriptsize{asymp}}}$, we cannot at this point rule out the possibility that $\\Omega_m-\\Omega_B=\\Omega_{\\hbox{\\scriptsize{asymp}}}$ and $\\Omega_B\\approx\\Omega_{\\hbox{\\scriptsize Dyn}}$ is a numerical accident."
    },
    "0711/0711.2969_arXiv.txt": {
        "abstract": "A heuristic hypothesis about domination of Bose-Einstein statistics in the early Universe is suggested. The possibility of Bose-Einstein condensation (BEC) of primordial baryon-antibaryon pairs is considered. In accordance with this postulation enormous masses in the order of galactic mass may be accumulated within the cosmic scales. At the certain threshold value of the matter density the structural bosons decay into fermions and the sharp breakdown of quantum-mechanical symmetry of the particles wave functions occurs. Then, due to the Pauli principle of exclusion a large-scale phase transition occurs because of enormous pressure jump of the matter. This phenomenon might cause Cosmological Bang at the beginning stage of the Matter Era.  As a mechanism of accumulation of galactic mass much larger than the configuration with structural bosons, a hypothetical BEC of elementary bosons (gauge bosons $W^{\\pm }$\\ and $Z^{0})$ is discussed as well. ",
        "introduction": "As a possible version for observable evolution of the Universe an assumption may be put forward that the Universe has been started-up from the Radiation Era which probably may be discussed and commented in general arising from the energetic convenience of Bose-Einstein statistics. Based on this assumption, we predict that the evolution of the Universe has been started with the \\textquotedblleft Era of Bose-Einstein statistics\\textquotedblright . At the further period of Matter Era the Fermi-Dirac statistics began act, and since this period the \\textquotedblleft Era of Fermi-Dirac statistics\\textquotedblright\\ dominates. In the result, like to Big Bang phenomenon (because of abrupt breakdown of supersymmetry (SUSY) at the earliest stage of Universe), an alternative Cosmological Bang at the beginning stage of Matter Era might occur as a consequence of spasmodically breakdown of quantum statistics. The great likelihood of Bose-Einstein statistics is directly connected with the energetic convenience of the particles population in the quantum microstates with possible lowest energy without any restriction, in contrast to Fermi-Dirac statistics where the Pauli principle excludes the occupation of the same energetic state by other fermions. This is a well-known macro-scale quantum phenomenon of \\ Bose-Einstein condensation (BEC).  The stable baryonic configurations for considering problem were usually considered in assumption arising from the current state of the matter--already formed steady-state nucleosynthesis and existing nuclei within the model of Fermi-Dirac statistics. Such approach to the main cosmological problem, as the generation of cosmic objects is, seems to be a particular case of more general situation in the early Universe when before the thermodynamic equilibrium between the radiation and matter the Bose-Einstein and Fermi-Dirac statistics of primordial particles have been disjointed because of enormous values of corresponding thermodynamic potentials. This fact indicates in favor of initial large-scale phase transition within the evolution process of early Universe. On the base of above mentioned assumption concerning the Bose and Fermi statistics, and taking into account that the population of the same quantum microstates by the bosons stipulates incomparable less internal energy, a heuristic hypothesis about dominating of Bose-statistics in the early Universe is predicted. Such a physical assumption is in the accordance with the basic process of production of the structural Bose-pairs of primordial particles and antiparticles from the initial high energetic $\\gamma $-photons in the early Universe. The astrophysical conditions of the Universe at the beginning stage of Matter Era satisfy the physical requirements for BEC of such pairs.  For the beginning let \\textrm{us} leave these heuristic abstractions (detailed investigation will be done separately) and concretize on the possibility of generation of structural bosons by the high energetic primordial photons. The theoretical aspects and mechanisms developed for electron-positron pairs production from high energy photons in the superdense nuclear matter and multiphoton production by the nonlinear channels in the strong radiation field of incomparable low frequencies than the threshold one in diverse astrophysical cases (\\cite{1,2,3,4}; H.K.Avetissian, A.K.Avetissian, et. all.) principally can be generalized for the proton-antiproton pairs production. As to neutrons and antineutrons pairs, it is obvious that the mechanisms of their production in principle may be investigated through the other intermediate physical channels, because they do not involve in the electromagnetic interaction (\\cite{5}; W.Fauler, F.Hoyle; \\cite{6}; G.Feinberg, L.Lederman). Since the high energetic $\\gamma $-photons may generate baryon-antibaryon pairs, and the photons are electrically pure neutral, they could create just pure neutral elementary particles--bosons, or even structural bosons with the total zero electrical charge within appropriate elementary phase space. ",
        "conclusions": "It is physically clear that the possibility of generation of large-scale configurations in principle does not except just in one local cosmic space-time scales. Furthermore, this hypothesis and corresponding physical model, as well as its astrophysical applications might have cosmological consequences both within relatively small- and large-scales cosmic structures. Note that after predicted Cosmological Bang at the beginning stage of Matter Era the \\textquotedblleft islands\\textquotedblright\\ of similar configurations of degenerate matter might be accumulated, which might form then various stable cosmic objects. It is physically obvious that these objects principally might correlate in groups, clusters, or even associations.  May the predicted Cosmological Bang be candidate for a \\textquotedblleft generator of galactic-scale masses\\textquotedblright ? This physical question requires additional analysis in the scope of considered hypothesis and corresponding astrophysical model. Besides, the considered problems require further investigation of the discussed analytical model and development of the main ideas presented here in the scope of general relativity. In addition, it is essential to take into consideration the various empiric equations for BEC which might be more realistic from the cosmological point of view."
    },
    "0711/0711.4548_arXiv.txt": {
        "abstract": "Optical observations of the Type Ia supernova (SN~Ia) 2005bl in NGC~4070, obtained from $-6$ to $+66$\\,d with respect to the $B$-band maximum, are presented. The photometric evolution is characterised by rapidly-declining light curves ($\\Delta$m$_{15}(B)_\\mathrm{true}= 1.93$) and red colours at peak and soon thereafter. With M$_{B,\\mathrm{max}}=-17.24$ the SN is an underluminous SN~Ia, similar to the peculiar SNe~1991bg and 1999by. This similarity also holds for the spectroscopic appearance, the only remarkable difference being the likely presence of carbon in pre-maximum spectra of SN~2005bl. A comparison study among underluminous SNe~Ia is performed, based on a number of spectrophotometric parameters. Previously reported correlations of the light-curve decline rate with peak luminosity and $\\mathcal{R}$(Si) are confirmed, and a large range of post-maximum Si\\,{\\sc ii} $\\lambda6355$ velocity gradients is encountered. 1D synthetic spectra for SN~2005bl are presented, which confirm the presence of carbon and suggest an overall low burning efficiency with a significant amount of leftover unburned material. Also, the Fe content in pre-maximum spectra is very low, which may point to a low metallicity of the precursor. Implications for possible progenitor scenarios of underluminous SNe~Ia are briefly discussed. ",
        "introduction": "\\label{Introduction}  The history of underluminous SNe~Ia is a typical example of the ever-recurring pattern in which knowledge about nature is accumulated. Usually, in the beginning there is the observation of a phenomenon, followed by a successful theoretical explanation. However, as further experiments or observations are carried out in order to confirm the newly developed theoretical ideas, often an ever higher degree of diversity and ever more exceptions from the simple rules are found the closer the subject of interest is studied. The need for refined and more complex theories to obtain a realistic description of the involved processes becomes evident.  In the case of SNe~Ia, first a class of cosmic explosions apparently similar in absolute luminosity (``standard candles'') and spectroscopic appearance was identified. These events were explained as the disruptions of white dwarfs which had accreted matter until they reached their stability limit close to the Chandrasekhar mass ($M_\\mathrm{Ch}$). However, in 1991 the paradigm of SN~Ia homogeneity had to be relaxed a lot. This was triggered by the observation of two peculiar SNe~Ia, which thereafter served as prototypes of newly-defined SN~Ia subclasses with distinct spectrophotometric properties. One of these, SN~1991T (\\citealt{Filippenko92a}; \\citealt{Phillips92}; \\citealt{RuizLapuente92}; \\citealt*{Mazzali95}), was up to $0.6$ mag brighter than average SNe~Ia, and characterised by a hot early-time spectrum with strong Fe\\,{\\sc iii} features and weak or absent Si\\,{\\sc ii} and S\\,{\\sc ii} lines. The other one, SN~1991bg \\citep{Filippenko92b,Leibundgut93,RuizLapuente93,Turatto96,Mazzali97} was even more deviant, with low ejecta velocities and a cool spectrum dominated by intermediate-mass-element (IME) lines and particularly strong O\\,{\\sc i} and Ti\\,{\\sc ii}. Moreover, it had unusually red colours at early phases, and was underluminous by about $2$ mag at peak (hereafter we will refer to such appearance as 91bg-like). Hence, quasi instantaneously the luminosity range of SNe~Ia had increased to a factor of ten between the brightest and the faintest objects, proving that they were \\textit{no} standard candles. However, two years later \\citet{Phillips93} realised a tight correlation between peak luminosity and decline rate in the $B$ band. This relation and revised versions of it \\citep[e.g.][]{Phillips99} turned SNe~Ia into standardisable candles, and hence made them an extremely useful tool for precision cosmology\\footnote{Several years later even more peculiar objects such as SN~2002cx \\citep{Li03} were discovered which did not obey the Phillips relation, showing that not \\textit{all} SNe~Ia are standardisable.}.  In the following years, several more 91bg-like SNe~Ia were discovered, but the available data set grew much less rapidly than for ordinary SNe~Ia. From the results of the Lick Observatory Supernova Search (LOSS) and the Beijing Astronomical Observatory Supernova Survey (BAOSS), \\citet{Li01} estimated that about $16$\\,\\% of all SNe~Ia are of the 91bg-like variety. This may still be an underestimate, as their low intrinsic luminosity makes 91bg-like SNe prone to Malmquist bias; nevertheless \\citet{Li01} estimated this effect to be negligible in their sample. Statistical studies \\citep{Hamuy96a,Hamuy00,Howell01a} have shown that SNe~Ia occur in all host-galaxy types, but revealed a correlation between SN decline rate and host morphology, with a clear tendency for 91bg-like SNe to be associated with early-type hosts and hence old stellar populations.  While the single-degenerate (SD) Chandrasekhar-mass model has survived as the favoured scenario for the normal and indeed rather homogeneous SNe~Ia, a number of alternative models have been suggested for the 91bg-like subclass. Ideas include double-degenerate (DD) explosions of merging white dwarfs, sub-Chandrasekhar-mass explosions triggered by detonation of the accreted helium layer \\citep[cf.][for a review]{Hillebrandt00}, and deflagrations in strongly rotating white dwarfs, where the turbulent propagation of the flame front is suppressed by the differential rotation \\citep{Pfannes06}. Still, the notion that 91bg-like SNe are -- in terms of the underlying explosion model -- no different from ordinary SNe~Ia, and that the only discriminating parameter is the mass of synthesised $^{56}$Ni, has supporters in the SN~Ia community. No conclusive evidence for any of these ideas has been found so far.  In this paper we present the joint data set of SN~2005bl obtained by the European Supernova Collaboration (ESC)\\footnote{http:/$\\!$/www.mpa-garching.mpg.de/$\\sim$rtn/} and the Carnegie Supernova Project (CSP)\\footnote{http:/$\\!$/www.csp1.lco.cl/$\\sim$cspuser1/PUB/CSP.html}. Since these observations are among the earliest ever obtained for a 91bg-like SN, they may help to better constrain possible progenitor and explosion models. The observations and techniques applied for data reduction and calibration are discussed in Section~\\ref{Observations and data reduction}. In Section~\\ref{Distance and extinction} we estimate the distance of SN~2005bl and the extinction along the line of sight. Sections \\ref{Photometric evolution} and \\ref{Spectroscopic evolution} are devoted to the analysis of the light curves and spectra, respectively. Results of 1D spectrum synthesis calculations are presented in Section~\\ref{Spectral modelling}, and a comparison with other underluminous SNe~Ia is performed in Section~\\ref{Discussion}, where we also discuss the impact of SN~2005bl on our picture of SN~Ia explosions. A short summary of the main results is given in Section \\ref{Conclusions}. ",
        "conclusions": "\\label{Conclusions}  We have presented and analysed optical photometric and spectroscopic data of the underluminous Type Ia SN~2005bl from one week before to two months after maximum light in $B$. In this entire interval the evolution of SN~2005bl is substantially different from that of slow or intermediate decliners, but very similar to that of the prototypical fast decliners SNe~1991bg and 1999by. At peak, SN~2005bl appears slightly more luminous than the latter, but the difference is marginal and within the error bars, which for SN~2005bl are dominated by uncertainties in the amount of host-galaxy extinction. With respect to quantities like the $\\Delta$m$_{15}(B)_\\mathrm{true}$ of $1.93$, the ejecta velocities inferred from the minimum of the Si\\,{\\sc ii} $\\lambda6355$ line, and temperature indicators such as $\\mathcal{R}$(Si) or the time-evolution of various colour indices, SN~2005bl is in fact --  within the measurement uncertainties -- a clone of SNe~1991bg and 1999by.  We confirm that $\\Delta$m$_{15}(B)_\\mathrm{true}$ is a good luminosity indicator also for underluminous SNe~Ia, with a steeper dependence of the peak magnitudes on $\\Delta$m$_{15}(B)_\\mathrm{true}$ than for slow and intermediate decliners. From linear fits we obtain $\\frac{\\mathrm{d}(\\mathrm{M}_\\mathrm{peak})} {\\mathrm{d}(\\Delta\\mathrm{m}_{15}(B)_\\mathrm{true})} = 6.83 \\pm 0.32$ and $4.33 \\pm 0.31$ for the $B$ and $V$ band, respectively, considering objects with $\\Delta$m$_{15}(B)_\\mathrm{true} \\geq 1.69$.  We conducted an analysis analogous to that of \\citet{Benetti05}, focussing on the behaviour of the FAINT subclass and trying to identify trends within this group. In addition to the SNe in Benetti et~al. we not only added the measurements of SN~2005bl, but also included SN~1998de \\citep{Modjaz01, Matheson07} and updated some numbers on the basis of better data availability. We find a correlation between $\\Delta$m$_{15}(B)_\\mathrm{true}$ and $\\mathcal{R}$(Si)$_\\mathrm{max}$, nicely extending the observed linear trend among HVG and LVG SNe to larger decline rates. Because of the lack of early data, the time-evolution of $\\mathcal{R}$(Si) at phases where differences between the LVG and HVG groups become evident cannot be investigated. The Si\\,{\\sc ii} line velocities of FAINT SNe are mostly similar to those of LVG SNe near maximum light, but significantly lower a few weeks later. There is a large dispersion in post-maximum velocity gradients $\\dot v$ within the FAINT class, with numbers comparable to those of HVG SNe, but exceeding their range in both directions.  Synthetic spectra for SN~2005bl at four epochs before and soon after maximum light have been computed using a 1D Monte Carlo code. Despite uncertainties and shortcomings such as the poorly-constrained rise time, the possible inadequacy of the W7 density profile for 91bg-like SNe~Ia, and the fact that computations were performed without chemical stratification, a number of interesting results were obtained. The presence of carbon in the $-6$ and $-5$\\,d spectra, indicated by a visual comparison with SN~1999by at similar epoch, was confirmed, and an overall low burning efficiency was established. NSE elements, but also IMEs, are significantly less abundant than in ordinary SNe~Ia; instead, most of the ejecta above $6000$ km\\,s$^{-1}$ are made of oxygen, and are hence either unburned or the result of highly incomplete carbon burning. Furthermore, in the $-6$ and $-5$\\,d spectra only traces of Fe are found, one order of magnitude less than expected for unprocessed material with solar composition. Possible implications of this low metallicity on the nature of the progenitor were discussed in Section~\\ref{Element abundances from synthetic spectra}.  It is presently very difficult to decide whether 91bg-like SNe~Ia form a separate group, distinct from ordinary SNe~Ia and possibly descending from different progenitors or explosion mechanisms, or whether they are just the extreme end of a continuous distribution of objects, with their characteristic appearance owing to lower temperatures in conjunction with a different chemical composition, both caused by the overall low burning efficiency. To tackle this problem will require a new statistical analysis similar to that of \\citet{Benetti05} once a bigger sample of rapidly-declining SNe~Ia will be available, and more detailed spectral modelling of the data sets already published."
    },
    "0711/0711.4581_arXiv.txt": {
        "abstract": "The masses distribution inside the Earth governs the behaviour of the rotation axis in the Earth (polar motion), as well as the Earth rotation rate (or equivalently, length of day). This masses distribution can be measured from space owing to artificial satellites, the orbitography of which provides the Earth gravity field determination. Then, the temporal variations of the Earth gravity field can be related to the variations of the Earth Orientation Parameters (EOP) (with the Inertia Tensor). Nowadays, owing to the satellite laser ranging (SLR) technique and to the new gravimetric satellite missions (as CHAMP or GRACE), the temporal variations of the low degree coefficients of the Earth gravity field (i.e. Stokes coefficients) can be determined. This paper is one of the first study using gravity variations data in the equations already established (e.g. Lambeck 1988) and linking the variations of the length of day and of the $C_{20}$ Stokes coefficient (or, linking the polar motion and the $C_{21}$ and $S_{21}$ coefficients). This paper combines the Earth rotation data (mainly obtained from VLBI and GPS measurements) and the Earth gravity field variations ones (e.g. Lageos I and II data, or GRACE data), in order to complete and constrain the models of the Earth rotation. The final goal is a better Earth global dynamics understanding, which possible application can be the constraint on the couplings into the Earth system (e.g. core-mantle couplings). ",
        "introduction": "In this paper, we compare the mass terms of the Earth Rotation Parameters (ERP) series, length of day (LOD) and polar motion, obtained from various techniques. We use (i) on one hand the IERS (International Earth Rotation and refrence systems Service) C04 series (obtained combining mainly VLBI and GPS measurements), and (ii) on the other hand the gravity field variations data (CNES/GRGS gravity field degree~2-coefficients, obtained from Lageos~I and II positionning measurements: for further explanation see Bourda \\& Capitaine 2004). ",
        "conclusions": "Zonal tides, motion terms, oceanic tides and atmospheric angular momentum effects were removed from the IERS C04 series for valuable comparisons. We noticed a good agreement between LOD and polar motion exciation obtained from (i) gravity field variations data, and (ii) classical geodetic measurements. But it still remains some critical points. For example, the correlation between the IERS C04 LOD series and the one from $\\Delta C_{20}$ residual data (i.e. the measurements, without solid Earth tides, atmospheric pressure effects, and oceanic tides) is poor: 0.17. But this can come from the using of various models for the same effect (e.g. the atmospheric pressure part) in these two methods. Further investigations will focuss on the errors made removing atmospheric winds from LOD data (which is the main effect), and on the hypothesis $\\Delta Tr(I)=0$ (Bourda 2004a) assumed computing Eq.~(\\ref{eq:LOD}) (which is perhaps not valuable using the $\\Delta C_{20}$ residuals)."
    },
    "0711/0711.3474_arXiv.txt": {
        "abstract": "Recent observations indicate that fully convective stars can effectively build magnetic fields without the aid of a tachocline of shear, that those fields can possess large-scale components, and that they may sense the effects of rotation.  Motivated by these puzzles, we present global three-dimensional simulations of convection and dynamo action in the interiors of fully convective M-dwarfs of 0.3 solar masses.  We use the Anelastic Spherical Harmonic (ASH) code, adopting a spherical computational domain that extends from 0.08-0.96 times the overall stellar radius.  We find that such fully convective stars can generate magnetic fields of several kG strength, roughly in equipartition with the convective flows.  Differential rotation is established in hydrodynamic progenitor calculations, but essentially eliminated in MHD simulations because of strong Maxwell stresses exerted by the magnetic fields. Despite the absence of interior angular velocity contrasts, the magnetic fields possess strong mean (axisymmetric) components, which we attribute partly to the very strong influence of rotation upon the slowly overturning flows. ",
        "introduction": "Of all the many findings that have emerged from solar physics over the last two decades, the discovery of the tachocline of shear likely figures as the most significant for dynamo theory.  Revealed by helioseismology (e.g., Gough \\& Toomre 1991), the tachocline is a narrow boundary layer situated near the base of the convection zone, in which the solar angular velocity transitions from differential rotation in the unstable envelope to nearly solid-body rotation beneath it.  Its discovery motivated the now-prevalent ``interface dynamo'' paradigm, in which the tachocline plays a pivotal role -- both because of its shear, which helps stretch and organize toroidal fields, and because of the stable stratification, which may allow fields to be greatly amplified before becoming susceptible to magnetic buoyancy (e.g., Ossendrijver 2003).   Recently, simulations have helped affirm the likely importance of the tachocline in generating the Sun's organized magnetism: global simulations that include a tachocline (Browning et al. 2006; Browning et al. 2007) give rise to magnetic fields that appear more ``solar-like'' than do simulations of the convection zone alone (Brun, Miesch \\& Toomre 2004).  (See proceedings by Miesch et al. 2007, this volume.)  And if the tachocline is vital in building the Sun's orderly magnetism, then similar boundary layers are probably likewise central in the dynamo action of any star that possesses both a convective envelope and a radiative, stable core.  But not all stars are so configured.  Moving down the main sequence to stars of lower mass, the size of the inner stable layer gradually decreases.  By around spectral class M3, corresponding to a mass of about 0.35 solar masses, it has vanished entirely, and convection occurs throughout the interior.  In a fully convective star, there is no radiative-convective interface, and so presumably no interface dynamo -- implying that if fields are built in such stars, the dynamo responsible must operate rather differently than in the Sun.  A natural conclusion is that such fully convective stars should show surface magnetic field behavior quite distinct from that realized in their more massive cousins (e.g., Durney, De Young \\& Roxburgh 1993; Chabrier \\& Kuker 2006).  Surprisingly, observations of low-mass stars have partly confounded even this rather basic expectation.  In a recent breakthrough, Donati et al. (2006) used Zeeman Doppler imaging to constrain the surface magnetic field morphology of a very rapidly rotating fully convective M-dwarf.  The kG-strength fields they detected were predominantly axisymmetric and large-scale, in contrast to the purely small-scale fields predicted by some models (e.g., Durney et al. 1993).  Donati et al. (2006) also found that their target star did not exhibit any surface differential rotation. Separately, other observations have demonstrated that strong magnetic activity appears to be quite common in fully convective stars: indeed, the fraction of M-stars showing with detectable activity actually peaks at about spectral class M8 (West et al. 2004).  Further puzzles have come from study of the remarkable relation between stellar magnetic activity and rotation rate.  In stars like the Sun, observations indicate that chromospheric and coronal activity increase with rotation rate, then ``saturate'' above a threshold velocity (e.g., Noyes et al. 1984; Delfosse et al. 1998; Pizzolato et al. 2003).  Many authors have noted that this rotation-activity correlation is tightened considerably when rotation is expressed in terms of the Rossby number, estimated in these studies as just $P_{\\rm rot}/\\tau_{c}$, with $\\tau_c$ a typical convective overturning time; likewise, the threshold rotational period needed for ``saturated'' activity is apparently a function of stellar mass, varying from about a day in early-G stars to about ten days in early M-stars (Pizzolato et al. 2003).  This rotation-activity connection appears to persist in some fashion into the mid-M stars: Mohanty \\& Basri (2003) argued that a sample of stars ranging from M0 to M5 showed a common ``saturation-type''  rotation-activity relationship, with observed activity roughly independent of rotation rate above a threshold value. Because measuring the rotation rates of the slowest rotators is difficult, it remains unclear whether magnetic activity in these stars increases gradually with rotation rate as in solar-like stars, or instead transitions more abruptly.  In the late-M spectral classes and beyond, the rotation-activity relation does appear to break down (e.g., Reiners et al. 2007), but this may occur well beyond the point where stars become fully convective.  In all, it appears that fully convective stars can certainly act as magnetic dynamos, can (in at least some cases) build large-scale magnetic fields (Donati et al. 2006), and may sense the effects of rotation in roughly the same fashion as more massive stars.  Motivated by these puzzles, we have carried out global 3--D nonlinear simulations of convection and dynamo action in the interiors of fully convective M-dwarfs.  In the sections that follow, we briefly describe our modeling (\\S2), its principal results (\\S3), and the implications of our work (\\S4).  The simulations here are described in more detail in Browning (2007). ",
        "conclusions": "The simulation presented here is, we believe, the most faithful description yet achieved of dynamo action in a fully convective M-dwarf.  Although we have made many simplifications in our modeling, some of the main findings may turn out to be robust.  We summarize some of these findings and their implications here.  The strong stratification plays a major role in setting convective flow amplitudes and pattern scales, and thus indirectly magnetic field strengths and morphologies.  Although our model star is unstably stratified everywhere, there are still two conceptually distinct regions: one near the surface where convection is quite vigorous and small-scale, and a more quiescent deep interior.  Convection establishes differential rotation in hydrodynamic cases, but the angular velocity contrasts are essentially eliminated in MHD simulations. The strong Maxwell stresses ultimately yield an interior in nearly solid-body rotation.  The magnetic fields realized here possess structure on many spatial scales.  The axisymmetric mean component accounts for a surprisingly high percentage of the total magnetic energy ($\\sim$ 20\\%), despite the absence of persistent angular velocity contrasts.  The polarity of mean fields is remarkably stable, having flipped only once in the $\\sim$ 30 years we have so far evolved the simulation.  We attribute some of these effects to the strong influence of rotation upon the slowly overturning flows.  Although these simulations rotate at the solar angular velocity, the influence of rotation is much stronger than in the Sun because the luminosity, and hence the rms convective velocity, is much lower.  Comparison to prior simulations of the Sun and A-type stars (Browning et al. 2004; Brun et al. 2005; Brun et al. 2004; Browning et al. 2006) suggests that there are plausibly three different regimes of behavior, separated primarily by the Rossby number.  When rotation is weak, magnetic energies of less than $\\sim$ 30\\% KE are typically realized, and do not greatly modify the differential rotation established in hydrodynamic cases.  When rotation is somewhat stronger, yielding ME greater than 30\\% KE, cyclical variations of magnetism and differential rotation are possible; when it is stronger still (as in the simulation here, and the most rapidly rotating cases of Brun, Browning \\& Toomre 2005), ME can exceed KE without the aid of differential rotation, and any persistent angular velocity contrasts are strongly quenched.  These conclusions are also (tentatively) in keeping with calculations of magnetism in rapidly rotating Sun-like stars (Brown et al., this volume).  These ideas lead to some straightforward observational predictions -- namely that magnetic, rapidly rotating M-stars should seldom exhibit differential rotation, that less rapidly rotating M-stars should generally exhibit weaker fields and (for very slow rotation) may show differential rotation, and that the axisymmetric component of the field should generally increase with rotation rate.     \\begin{theacknowledgments} It is a pleasure to thank Juri Toomre, A. Sacha Brun, Andrew West, Mark Miesch, Benjamin Brown, and Nick Featherstone for helpful discussions and/or comments on this manuscript.  We also gratefully acknowledge Isabelle Baraffe and Gilles Chabrier for supplying the 1--D stellar model used here for initial conditions.  This work was supported by an NSF Astronomy \\& Astrophysics postdoctoral fellowship (AST 05-02413).   \\end{theacknowledgments}"
    },
    "0711/0711.1537_arXiv.txt": {
        "abstract": "Black holes are popping up all over the place: in compact binary X-ray sources and GRBs, in quasars, AGNs and the cores of all bulge galaxies, in binary black holes and binary black hole-neutron stars, and maybe even in the LHC! Black holes are strong-field objects governed by Einstein's equations of general relativity. Hence general relativistic, numerical simulations of dynamical phenomena involving black holes may help reveal ways in which black holes can form, grow and be detected in the universe. To convey the state-of-the art, we summarize several representative simulations here, including the collapse of a hypermassive neutron star to a black hole following the merger of a binary neutron star, the magnetorotational collapse of a massive star to a black hole, and the formation and growth of supermassive black hole seeds by relativistic MHD accretion in the early universe. ",
        "introduction": "Black holes are `sighted' everywhere in the universe these days. Originally located in compact binary X-ray sources in the 1970's, the cosmic presence of black holes has expanded considerably in recent decades. They now are believed to be the central engines that power quasars, active galactic nuclei (AGNs) and gamma-ray bursts (GRBs). They are identified in the cores of all bulge galaxies. They are presumed to form signficant populations of compact binaries, including black hole-black hole binaries (BHBHs) and black hole-neutron star binaries (BHNSs).  Black holes may even show up soon in the Large Hadron Collider!  Gravitationally, black holes are strong-field objects whose properties are governed by Einstein's theory of relativistic gravitation --- general relativity.  General relativistic simulations of gravitational collapse to black holes, BHBH mergers and recoil, black hole accretion, and other astrophysical phenomena involving black holes may help reveal how, when and where black holes form, grow and interact in the physical universe.  As a consequence, such simulations can help identify the ways in which black holes can best be detected.   To illustrate how our understanding of black hole phenomena is sharpend by large-scale simulations in general relativity, we summarize in this paper the results of several recent computational investigations. The first few involve the formation of black holes from stellar collapse, while the last one concerns supermassive black hole growth via disk accretion. Most of these simulations utilize the tools of numerical relativity to solve Einstein's field equations of general relativity. So we shall begin with a brief overview of this important, rapidly maturing field. ",
        "conclusions": "We have discussed a number of black hole scenarios that require numerical simulations in full general relativity for true understanding. Some of these processes involve the formation of black holes, others concern their subsequent growth and interactions. The important point is that numerical relativity is at last mature enough to probe these issues reliably. Numerical relativity can now serve as an important tool to simulate stellar collapse to black holes, the inspiral, merger and recoil of BHBHs, NSNSs and BHNSs, the generation of gravitational waves from stellar collapse and binary mergers, accretion onto black holes, and countless other phenomena involving black holes and their strong gravitational fields. With such a tool, these processes now can be investigated at a fundamental level without many of the ad hoc assumptions and approximations required in previous treatments. This should lead to an improved qualitative picture, as well as more reliable quantitative results.  We are grateful to T. Baumgarte, C. Gammie, Y-T Liu, and M. Shibata for useful discussions.  S.L.S is supported by NSF grants PHY-0205155, PHY-0345151, and PHY-0650377 and NASA Grants NNG04GK54G and NNX07AG96G to the University of Illinois."
    },
    "0711/0711.0330_arXiv.txt": {
        "abstract": "The least-squares (or Lomb-Scargle) periodogram is a powerful tool which is used routinely in many branches of astronomy to search for periodicities in observational data. The problem of assessing statistical significance of candidate periodicities for different periodograms is considered. Based on results in extreme value theory, improved analytic estimations of false alarm probabilities are given. They include an upper limit to the false alarm probability (or a lower limit to the significance). These estimations are tested numerically in order to establish regions of their practical applicability. ",
        "introduction": "Analysing astronomical time series, one often has to choose between at least two hypotheses, a base one $\\mathcal H$ and an alternative one $\\mathcal K$, based on the existing data array. In the signal detection problem, one should check whether the observations are consistent with some base model or they contain an extra deterministic signal. Under presence of random errors, such problem can be solved in a probabilistic sense only. We are never protected from mistakes of two kinds. They are the false retraction of $\\mathcal H$ (the `false alarm') and the false non-retraction of $\\mathcal H$ (the false non-detection). False alarms are generally believed to be more dangerous, hence the problem of estimation of the false alarm probability (hereafter $\\FAP$) associated with a candidate signal is very important. Given some small critical value $\\FAP_*$ (between $10^{-3}$ and $0.1$ usually), we could claim that the candidate signal is statistically significant (if its $\\FAP<\\FAP_*$) or is not ($\\FAP>\\FAP_*$).  For the \\citet{Lomb76}~-- \\citet{Scargle82} periodogram (hereafter also L--S), the base hypothesis is that the observations are pure zero-mean uncorrelated and Gaussian errors (also called the white Gaussian noise). The alternative one is that a sinuous harmonic is also present. Every single value of the L--S periodogram represents a test statistic for the corresponding problem of hypotheses testing. In routine practical cases, however, the period of a possible signal is not known a priory and we have to scan many periodogram values within a wide frequency range. In this case, the $\\FAP$ is provided by the probability distribution of the maximum periodogram value under the base hypothesis (i.e., without signal in the data). Existing methods of calculating this distribution for a continuous frequency range require time-consuming Monte-Carlo simulations. The aim of the present paper is to propose analytic approximations which could allow to avoid Monte-Carlo simulations (at least in many practical cases). Such approximations of the distribution of the maximum have already been constructed by mathematicians specializing in the field of extreme values of random processes. In the Section~\\ref{sec_Stat}, these results are adapted for and extended to the specific features of the periodogram analysis of astronomical time series. In the Section~\\ref{sec_Simul}, numerical simulations are used to explore the quality of the analytic results and to show regions of their practical applicability. ",
        "conclusions": "The problem of estimating statistical significance of periodogram peaks is discussed in the paper. The results published in the field of extreme values of random processes are adapted for and extended to the periodogram analysis of astronomical time series. For the Lomb--Scargle periodogram and its modifications the corresponding extreme value distributions are given by a closed formulae being ready for usage. If the spectral leakage cannot be neglected, the similar expressions provide upper limits to the false alarm probability (or lower ones to the significance).  It is established numerically that the region of validity of these approximations is large and has no sharp boundaries. Even if the aliasing is very strong, the error of the analytic estimation of false alarm probability does not favour false alarms and thus is not fatal. For strong aliases, the usage of this analytic approximation slightly decreases the sensitivity to low-amplitude signals. However, the respective increasing of amplitude thresholds should not exceed several per cent in the worst practical cases (like a strong aliasing enforced by lack of observations).  These results may be very useful in a wide variety of astronomical applications. They would be useful especially for systematic surveys that deal with large amounts of data consisting of many separate time series. Indeed, it would be very difficult and even impossible to perform Monte-Carlo simulation for every of such time series. Vice versa, it is easy to use simple analytic formulae~(\\ref{MaximaDistr},\\ref{MaximaDistrIneq}) or their analogs for the modified L--S periodograms. This will eliminate the need for Monte-Carlo simulations in the cases when the observed periodogram peak exceeds the adopted threshold and in the opposite cases when this peak is lower than this threshold by more than, say, $10\\%$. The rare intermediate cases are easy to be studied by means of Monte-Carlo simulations. It is also admissible not to use numerical simulations at all, especially for large datasets ($N\\gtrsim 100$). In this case, the number of undetected low-amplitude periodicities may be increased by a negligible quantity only."
    },
    "0711/0711.2429_arXiv.txt": {
        "abstract": "We analyse the dynamical properties of substructures in a high-resolution dark matter simulation of the formation of a Milky Way-like halo in a $\\Lambda$CDM cosmology. Our goal is to shed light on the dynamical peculiarities of the Milky Way satellites. Our simulations show that about 1/3 of the subhalos have been accreted in groups. We quantify this clustering by measuring the alignment of the angular momentum of subhalos in a group. We find that this signal is visible even for objects accreted up to $z \\sim 1$, i.e. 8 Gyr ago, and long after the spatial coherence of the groups has been lost due the host tidal field.  This group infall may well explain the ghostly streams proposed by Lynden-Bell \\& Lynden-Bell to orbit the Milky Way.  Our analyses also show that if most satellites originate in a few groups, the disk-like distribution of the Milky Way satellites would be almost inevitable.  This non-random assignment of satellites to subhalos implies an environmental dependence on whether these low-mass objects are able to form stars, possibly related to the nature of reionization in the early Universe.  With this picture, both the ``ghostly streams'' and the ``disk-like configuration'' are manifestations of the same phenomenon: the hierarchical growth of structure down to the smallest scales. ",
        "introduction": "In cold dark-matter (CDM) cosmologies, large galaxies are the result of the aggregation of smaller subunits. Some of these subunits may survive until the present day in the form of satellites, while some may be completely destroyed in the course of time, and contribute to the field.  In $N$-body CDM simulations of the formation of galaxy-size dark-matter halos, the satellites (often referred to as substructures or subhalos) show significantly different properties than those of the ``luminous'' satellites around galaxies like the Milky Way.  An example of this discrepancy is the ``missing satellite problem'': subhalos outnumber the bright satellites by factors of a hundred or more \\citep{kwg93,klypin99,moore99}. Furthermore, their spatial distribution is typically much shallower than observed for the luminous satellites \\citep{gao04a,tbs04}. Therefore, the relation between resolved substructures/subhalos in dark matter simulations and the luminous satellites in galaxy halos is still unclear. Attempts to reconcile these two populations, using semi-analytic models of galaxy formation \\citep{kwg93,benson02,kgk04} or full-fleshed SPH-simulations \\citep{maccio06,libeskind07} have produced interesting results, and helped us gain insight into the relevant processes on the smallest galactic scales.  In the past ten years new attention has been drawn to the properties of the satellite population in the Local Group. Starting with \\citet{lyndenbell95}, the existence of ghostly streams of satellites (dwarf galaxies and globular clusters) was proposed. These objects would share similar energies and angular momenta producing a strong alignment along great circles on the sky \\citep{palma02}. Recently, it has been argued that the MW satellites define a disk-like structure, that is so highly flattened (\\textit{rms} thickness of $10-30$ kpc) that may at first sight be inconsistent with CDM models \\citep{kroupa05,metz07}.  However, sophisticated modelling combining semi-analytic galaxy formation recipes with dark matter simulations has produced results that are consistent with observations \\citep{kang05,libeskind05,zentner05}.  In the past two years, about a dozen low surface brightness dwarf satellite galaxies around the MW have been found in the Sloan Digital Sky Survey \\citep{willman05a,willman05b,belokurov06,belokurov07,zucker06a,zucker06b,irwin07}. This is an increase by a factor of two in the number of known satellites. Because of the strong selection bias (SDSS focus is on the north galactic pole), it is unclear whether these satellites confirm the flattened disk-like structure. It is also unclear whether they help solve the missing satellite problem \\citep{sg07}. Satellites around M31 also show peculiarities in their distribution, such as an excess of objects on the side closest to the Milky Way \\citep{mcc06}, and possibly a similar degree of alignment \\citep{kg06}.  These facts have motivated us to revisit the distribution and properties of the subhalo population in CDM simulations. In particular, here we focus on the infall of substructures onto a Milky-Way like dark matter halo in a $\\Lambda$CDM cosmogony. We use a high-resolution dark-matter simulation that is a variant of the GA series \\citep{stoehr02,stoehr06} as described in Section~\\ref{simulation_sec}. We find evidence of group infall onto the MW-like halo, which may explain the ghostly streams proposed by \\citet{lyndenbell95}. This is discussed in Section \\ref{group_infall_sec}. In Section~\\ref{great_disk_sec} we show that the disky configuration of satellites is consistent with CDM if most satellites have their origin in a few groups. Thus both the ``ghostly streams'' and the ``planar configuration'' are manifestations of the same phenomenon: the hierarchical growth of structure down to the smallest galactic scales. We present our conclusions in Section~\\ref{discussion_sec}. ",
        "conclusions": "\\label{discussion_sec}  We have analysed a high-resolution cosmological simulation of the formation of a Milky Way-like halo, focusing on the properties of the satellite population at the present-day.  We have found that dark-matter subhalos are often accreted in groups in our simulations. Roughly $1/3$ of the surviving subhalos with mass $\\ge 2.9 \\times 10^6 \\msun$ at the present epoch share this property. This is clearly a lower limit since we are not be able to identify accompanying halos below our resolution limit (this is particularly severe at high-redshift).  This group infall is apparent as an enhancement in the number of subhalos whose angular momentum orientation is similar, particularly at the time of infall. This signal is measurable also from the present-day angular momentum of subhalos, even for those accreted 8 Gyr ago. These groups of subhalos share coherent orbits which can be traced back well before the accretion epoch. The differential group mass function follows a power-law distribution $dN/d \\log M|_{group} \\propto M^{n}$ with $n \\sim -0.5 \\pm 0.2$. This is reminiscent of the differential mass function of subhalos in both galaxy and cluster-size halos, albeit with a shallower slope (compared to $n \\sim -0.8$ as in e.g. \\citealt{delucia04,gao04b}).  We have also studied the degree of flattening of the spatial distribution of subhalos in our simulation. The mean minor-to-major axis ratio $c/a$ of the inertia tensor defined by the positions of $11$ randomly selected subhalos with $300$ kpc is $c/a \\sim 0.51 \\pm 0.12$. In comparison the $c/a$ of the $11$ ``classical'' MW satellites is $0.18 \\pm 0.01$.  Imposing the centrally concentrated MW satellite radial distribution leads to $c/a \\sim 0.4 \\pm 0.1$ and therefore somewhat alleviates the discrepancy with the observations \\citep[see also][]{kang05}.  We have explored also how this planar configuration may be obtained as a result of the infall of satellites in groups. The observed correlation in the angular momentum orientation of subhalos naturally gives rise to disk-like configurations. For example, we find that if all subhalos are accreted from just one group, it is almost impossible to avoid a disk-like distribution ($\\sim$80\\% probability), while for accretion from just two groups, the likelihood of obtaining a distribution as planar as observed is 40\\%.  These results may explain the origin of the ghostly streams proposed by \\citet{lyndenbell95}. Out of the streams originally proposed, only two appear to have survived the rigour of time, after modern and accurate measurements of proper motions have become available. \\citet{palma02} confirmed the LMC-SMC-UMi-Draco stream forms a kinematic group whose angular momentum separation is $< 18.5^{\\circ}$. More recently \\citet{piatek05} ruled out with $95\\%$ confident level UMi as a member using \\textit{HST} proper motions. The latest measurements of the Fornax proper motion by \\citet{piatek07} has apparently confirmed the Sculptor-Sextans-Fornax stream (although previous measurements led to conflicting results, see \\citealt{piatek02,dinescu04}).  If some of the luminous satellites are embedded in dark (sub)halos that fell in together, such coherent structures would be a naturally consequence of the hierarchical build-up of galaxies.  In our simulations, such groups remain coherent in angular momentum (i.e. they share similar orbital planes giving rise to great circle streams) for approximately 8 Gyr. This implies that these groups (or satellites) should have been accreted by the Milky Way at redshifts $z\\sim 1$ or below.  One of the possible implications of the reality of the ghostly streams is that its member galaxies formed and evolved in a similar environment before falling into the Milky Way potential. This would have implications on the (earliest) stellar populations of these objects, such as for example, sharing a common metallicity floor \\citep{helmi06}. On the other hand, this implies that there should be groups that have not been able to host any luminous satellites. This would hint at a strong dependence on environment on the ability of a subhalo to retain gas \\citep{scannapieco01}, or be shielded from re-ionization by nearby sources \\citep{mcb04,weinmann07}.  Recent proper motion measurements of the Large and Small Magellanic clouds by \\citet{smc-mu}, as well as the simulations by \\citet{bc05} suggest that these systems may have become bound to each other only recently. This would be fairly plausible in the context of our results.  The Clouds may well have been part of a recently accreted group \\citep[see also][for the link to Leo I]{sales07} and it may not even be necessary for them to ever have been a binary system. This may also have implications on the computations of the past trajectories of these systems, particularly if both are embedded in a larger common dark-matter envelope \\citep{besla07}.  Our analysis shows that the dynamical peculiarities of the Milky Way satellites can be understood in the context of the concordance cosmological model. Their properties must be a consequence of both the environment as well as of the hierarchical nature of the build up of galactic halos."
    },
    "0711/0711.2103_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\begin{table*} \\begin{center} \\caption{System parameters for \\tel} \\begin{tabular}{lccc} \\hline & \\vspace{-3mm} \\\\ \\multicolumn{1}{c}{Parameter} & Symbol & Strawman & Expansion \\\\ \\hline & \\vspace{-3mm} \\\\ Number of Dishes & $N$ & 30 & 45 \\\\ Dish Diameter (m) & & 12 & 12\\\\ Dish area (m$^2$) & $A$ & 113 & 113\\\\ Total collecting area (m$^2$) & & 3393 & 5089\\\\ Aperture Efficiency & $\\epsilon_a$ & 0.8 & 0.8 \\\\ System Temperature (K) & $T$ & 50 & 35 \\\\ Number of beams & & 30 & 30\\\\ Field-of-view (deg$^2$) & $F$ & 30 & 30 \\\\ Frequency range (MHz) & & 700 $-$ 1800 & 700 $-$ 1800\\\\ Instantaneous Bandwidth (MHz) & $B$ & 300 & 300\\\\ Maximum number of channels & & 16000 & 16000 \\\\ Maximum baseline (m) & & 2000 & 400, 8000\\\\ \\hline & \\vspace{-3mm} \\\\ \\end{tabular} \\label{intro:params} \\end{center} \\end{table*} The Australian SKA Pathfinder (\\tel) is a next generation radio telescope on the strategic pathway towards the staged development of the Square Kilometre Array (SKA). The \\tel\\ project is international in scope and includes partners in Australia, Canada, the Netherlands and South Africa. This paper, which concentrates on the science made possible with \\tel\\ was written jointly between Australian and Canadian research scientists following a collaborative agreement signed between the CSIRO, Australia and the National Research Council of Canada. \\\\  \\noindent \\tel\\ has three main goals: \\begin{itemize} \\item to carry out world-class, ground breaking observations directly relevant to the SKA Key Science Projects; \\item to demonstrate and prototype technologies for the mid-frequency SKA, including field-of-view enhancement by focal-plane phased-arrays on new-technology 12-m class parabolic reflectors; \\item to establish a site for radio astronomy in Western Australia where observations can be carried out free from the harmful effects of radio interference. \\end{itemize}  Following international science meetings held in April 2005 and March 2007, seven main science themes have been identified for \\tel. These are extragalactic \\HI\\ science, continuum science, polarization science, Galactic and Magellanic science, VLBI science, pulsar science and the radio transient sky. In this paper, we first give a general science overview, then outline the system parameters for \\tel, before allocating a section to a summary of each science theme. A larger, and more complete, version of the science case for \\tel\\ will be published elsewhere.  \\subsection{\\tel\\ and SKA science} The SKA will impact a wide range of science from fundamental physics to cosmology and astrobiology. The SKA Science Case ``Science with the Square Kilometre Array\" (Carilli \\& Rawlings 2004) identifies compelling questions that will be addressed as key science by the SKA : \\begin{itemize} \\item understanding the cradle of life by imaging the environments of the formation of earth-like planets, the precursors to biological molecules, and carrying out an ultra-sensitive search for  evidence of extra-terrestrial intelligence, \\item carrying out fundamental tests of the theory of gravity by using radio waves to measure the strong space-time  warp of pulsars and black holes and timing of arrays of pulsars over large areas of the sky to detect long-wavelength gravitational waves propagating through the Galaxy, \\item tracing the origin and evolution of cosmic magnetism by measuring the properties of polarized radio waves from galaxies over cosmic history \\item charting the cosmic evolution of galaxies and large-scale structure, the cosmological properties of the universe and dark energy, the imaging of atomic hydrogen emission from galaxies and the cosmic web from the present to time of the first galaxies, and \\item probing the dark ages and the epoch of reionization of the Universe when the first compact sources of energy emerged. \\end{itemize}  The technological innovation of \\tel\\ and the unique radio quiet location in Western Australia will enable a powerful synoptic survey instrument that will  make substantial advances in SKA technologies and on three of the SKA key science projects: the origin and evolution of cosmic magnetism, the evolution of galaxies and large scale structure, and strong field tests of gravity. The headline science goals for \\tel\\ are: \\begin{itemize} \\item The detection of a million galaxies in atomic hydrogen emission across 80\\% of the sky out to a redshift of 0.2 to understand galaxy formation and gas evolution in the nearby Universe. \\item The detection of synchrotron radiation from 60 million galaxies to determine the evolution, formation and population of galaxies across cosmic time and enabling key cosmological tests. \\item The detection of polarized radiation from over 500,000 galaxies, allowing a grid of rotation measures at $10'$  to explore the evolution of magnetic fields in galaxies over cosmic time. \\item The understanding of the evolution of the interstellar medium of our own Galaxy and the processes that drive its chemical and physical evolution. \\item The characterization of the radio transient sky through detection and monitoring of transient sources such as gamma ray bursts, radio supernovae and intra-day variables. \\item The discovery and timing of up to 1000 new radio pulsars to find exotic objects and to pursue the direct detection of gravitational waves. \\item The high-resolution imaging of intense, energetic phenomena through improvements in the Australian and global Very Long Baseline networks. \\end{itemize}  \\subsection{System Parameters} Table~\\ref{intro:params} gives the \\tel\\ system parameters. The first column gives the parameter with the second column listing the symbol used in the equations in this section. The strawman (or base model) parameters for \\tel\\ are given in the third column of Table~\\ref{intro:params} and these strawman assumptions have been used throughout this paper. Likely upgrade or expansion paths include the addition of further dishes and/or the cooling of the focal plane array elements to provide a lower system temperature. Parameters for this expansion path are listed in the column 4 of Table~\\ref{intro:params}.  \\tel\\ is designed to be a fast survey telescope. A key metric in this sense is the survey speed expressed in the number of square degrees per hour that the sky can be surveyed to a given sensitivity. Survey speeds and sensitivity for an interferometer like \\tel\\ have been derived elsewhere (e.g. Johnston \\& Gray 2006) and the full derivation will not be shown here. To summarise, the time, $t$, required to reach a given sensitivity limit for point sources, $\\sigma_s$, is \\begin{equation} t = \\left( \\frac{2\\,\\,\\, k\\,\\,\\, T}{A\\,\\,\\, N\\,\\,\\, \\epsilon_a \\,\\,\\, \\epsilon_c } \\right)^2 \\frac{1}{\\sigma_s^2\\,\\,\\, B\\,\\,\\, n_p} \\end{equation} where $B$ is the bandwidth (Hz), $n_p$ the number of polarizations, $A$ is the collecting area of a single element (m$^2$), $N$ is the number of elements and $\\epsilon_a$ and $\\epsilon_c$ represent dish and correlator efficiencies. The system temperature is $T$ with $k$ being the Boltzmann constant. The number of square degrees per second that can be surveyed to this sensitivity limit is \\begin{equation} SS_s = F\\,\\,\\,B\\,\\,\\,n_p \\left( \\frac{A\\,\\,\\, N\\,\\,\\, \\epsilon_a \\,\\,\\, \\epsilon_c \\,\\,\\, \\sigma_s}{2\\,\\,\\,k\\,\\,\\, T} \\right)^2 \\label{eq:ss} \\end{equation} where $F$ is the field of view in square degrees. The surface brightness temperature survey speed is given by \\begin{equation} SS_t = F \\,\\,\\,B\\,\\,\\,n_p  \\left( \\frac{\\epsilon_c \\,\\,\\, \\sigma_t}{T} \\right)^2 \\,\\,\\, f^2  \\,\\,\\, \\epsilon_s^{-2} \\label{eq:sb} \\end{equation} where now $\\sigma_t$ denotes the sensitivity limit in K and $f$ relates to the filling factor of the array via \\begin{equation} f = \\frac{A \\,\\,\\, \\epsilon_a \\,\\,\\, N \\,\\,\\, \\Omega\\,\\,\\, \\epsilon_s}{\\lambda^2} \\label{eq:ff} \\end{equation} Here, $\\epsilon_s$ is a `synthesised aperture efficiency' which is related to the weighting of the visibilities and is always $\\leq 1$.  There is interplay between these parameters when trying to maximise the survey speed for a given expenditure. For the majority of the science that will be considered here, the value of $SS_s$ and $SS_t$ are critical parameters, although the instantaneous sensitivity is also important, especially for pulsar science. Although these equations are useful, they are not the entire story. For example, the effects of good (u,v) coverage on the image quality and dynamic range do not appear in the equations. Furthermore, one should also not neglect the total bandwidth available for a spectral line survey. If the total bandwidth (or velocity coverage) is insufficient to cover the required bandwidth (velocity range) of a given survey, the survey speed suffers as a result of having to repeat the same sky with a different frequency setting.  In Table~\\ref{intro:speed} we list values of the sensitivity and survey speed for different `typical' surveys for both the strawman and the expansion parameters. The first entry gives a continuum survey where the entire 300~MHz of bandwidth is exploited and a desired 1-$\\sigma$ sensitivity of 100~$\\mu$Jy is required. The second entry gives a spectral line survey, with the third line listing a surface brightness survey needing to reach a 1-$\\sigma$ limit of 1~K over 5~kHz channel under the assumption of a $1'$ resolution. The final row lists the time necessary to reach 1-$\\sigma$ of 1~mJy across a 1~MHz bandwidth to a point source at the centre of the field. The expansion option for \\tel\\ offers a factor of almost 5 improvement over the strawman design.  \\subsection{Comparison with other instruments} The large field-of-view makes \\tel\\ an unprecedented synoptic radio telescope, achieving survey speeds not available with any other telescope. \\tel\\ will routinely image very large areas of the sky to sensitivities only achievable with current instruments over very small areas. \\begin{table*} \\begin{center} \\caption{Sensitivity and survey speeds for \\tel} \\begin{tabular}{lccc} \\hline & \\vspace{-3mm} \\\\ \\multicolumn{1}{c}{Parameter} & Strawman & Expansion\\\\ \\hline & \\vspace{-3mm} \\\\ Continuum survey speed (300 MHz, 100~$\\mu$Jy) & 250 & 1150 & deg$^2$/hr\\\\ Line survey speed (100~kHz, 5~mJy) & 209 & 960 & deg$^2$/hr\\\\ Surface brightness survey speed (5~kHz, 1~K, $1'$) & 18 & 83 & deg$^2$/hr\\\\ Point source sensitivity (1~MHz, 1~mJy) & 1290 & 280 & sec\\\\ \\hline & \\vspace{-3mm} \\\\ \\end{tabular} \\label{intro:speed} \\end{center} \\end{table*}  The survey speed of \\tel\\ exceeds that of the Parkes 20-cm multibeam receiver, the Very Large Array (VLA) and the Giant Metre Wave Telescope (GMRT) by more than an order of magnitude for spectral line surveys in the GHz band. In continuum, \\tel\\ can survey the sky some 50 times faster than the NVSS survey carried out by the VLA over a decade ago (Condon et al. 1998).  As an interferometer \\tel\\ provides low-frequency imaging not otherwise available at the other major southern hemisphere interferometric array, the Australia Telescope Compact Array (ATCA). For single pointings it exceeds the ATCA sensitivity at 1400~MHz and will also have better resolution and surface brightness sensitivity. In terms of survey speed however, it gains by large factors for both line, continuum and surface brightness sensitivity and will be comparable to other planned facilities such as the Karoo Array Telescope (KAT) in South Africa, the Allen Telescope Array (ATA; Deboer et al. 2004) in the USA and APERTIF in the Netherlands.  The survey speeds of \\tel\\ (with the strawman parameters), the ATA and APERTIF are almost identical although achieved in different ways. The ATA achieves its survey speed through a large number of small dishes each with a single pixel feed. APERTIF will have a focal plane array system on larger dishes. The presence of telescopes with similar survey speed in the northern hemisphere makes for excellent complementarity to \\tel.  \\subsection{Configuration of \\tel} A number of science projects (pulsar surveys, Galactic \\hi, low surface brightness mapping) require a highly compact array configuration in order to increase the surface brightness survey speed (see equations~\\ref{eq:sb} and \\ref{eq:ff}). On the other hand science such as continuum and transients require long baselines both to overcome the effects of confusion and to obtain accurate positions for identification at other wavelengths. In the middle is the extragalactic \\HI\\ survey which needs moderate resolution to avoid over-resolving the sources. With a total of only 30 dishes it is difficult to achieve all these requirements simultaneously.  It is envisaged that three array configurations will be available ranging from a very compact configuration (maximum baseline $\\lesssim$400~m) through a medium compact configuration (maximum baseline $\\sim$2~km) to an extended configuration (maximum baseline $\\sim$8~km). Configuration changes would occur by physically moving the antennas and would happen only infrequently.  A design study is currently underway to determine exact antenna positions for different configurations given the science case and the constraints of the local site topology and terrain.  \\subsection{Location of \\tel} The central core of \\tel\\ will be located at the Murchison Radio Observatory in inland Western Australia, one of the most radio-quiet locations on the Earth and one of the sites selected by the international community as a potential location for the SKA. The approximate geographical coordinates of the site are longitude 116.5 east and latitude 26.7 south. The southern latitude of \\tel\\ implies that the Galactic Centre will transit overhead and the Magellanic Clouds will be prominent objects of study. At least 30,000 square degrees of sky will be visible to \\tel. The choice of site ensures that \\tel\\ will be largely free of the harmful effects of radio interference currently plaguing the present generation of telescopes, especially at frequencies around 1~GHz and below. Being able to obtain a high continuous bandwidth at low frequencies is critical to much of the science described in this document.  \\subsection{\\tel\\ timeline} In early 2008, the \\tel\\ test bed antenna will be installed at the site of the Parkes radio telescope to allow testing of the focal plane array and beamforming systems. Following this, the goal is to have the first 6 antennas of \\tel\\ on-site in Western Australia in early 2010. Over the subsequent two years the remainder of the antennas would be deployed, with commissioning of the final system expected to take place in 2012. ",
        "conclusions": "\\tel\\ is a key step on the strategic pathway towards the SKA. The goals of \\tel\\ are to carry out world-class, ground breaking observations, to demonstrate and prototype technologies for the mid-frequency SKA and to establish a site for radio astronomy in Western Australia where observations can be carried out free from the harmful effects of radio interference.  In this paper we have outlined the main science themes to be tackled by \\tel, themes which derive from the major issues confronting astrophysics today such as understanding the evolution, formation and population of galaxies including our own, understanding the magnetic Universe, the nature of the transient radio sky and the direct detection of gravitational waves."
    },
    "0711/0711.2932_arXiv.txt": {
        "abstract": "{ The present status of Galactic Globular Clusters Database is briefly reviewed. The features implemented at the time writing are described, as well as plans for future improvements. ",
        "introduction": "The Galactic Globular Clusters Database (briefly, {\\it Gclusters}) \\footnote{http://snipurl.com/gclusters} is focused on presenting, in an organized way, a comprehensive list of bibliography, parameters and data for each of the known globular cluster of the Milky Way (GGCs). The need for a rational and organic assembly of these data is well described in a famous paper (\\citep{har96}):  {\\it \"The globular clusters in the Milky Way have proven thoughout this century to be irrepleaceble objects in an amazingly wide range of astrophysical studies... Year after year, it has proven important to have readily available up-to-date lists of parameters for these unique objects.\"}  After more than ten years, these words appear even more true: the amount of available data on GGCs has increased at a steady rate, following closely the enhanced capability of the technical instrumentations. Not only we have new and more reliable parameters for a great part of the known clusters, but - thanks to modern surveys conducted in bands different from the visible, such 2MASS (\\citep{skr06}) - several other objects keep going to increase our list of Milky Way clusters (e.g.,\\citep{fro07}, \\citep{bon07}). However, such data are inevitably scattered among the various papers, so what is needed is a simple way to have the relevant informations on a given cluster in a single source.  The Harris' catalogue of GGCs \\footnote{http://snipurl.com/gclusters2} is surely an unique resource for the researchers, in what it provides an extensive list of parameters for all the GGCs known at the time of its last revision (Feb. 2003): such a compilation is accessible online and is composed by tables of parameters available in form of flat text files.  \\begin{figure}[] \\caption{ \\footnotesize The main page of the Gclusters } \\end{figure}  Initially built around the Harris' compilation, {\\it Gclusters} is designed to allow a more flexible fruition of available data, so to make possible things such as ordering clusters according to the value of a given parameter, select objects whose parameters fall in a given interval, display related bibliography and colour magnitude diagrams, or even drop a note pertinent to that cluster, to be displayed in the website. Data are collected from a growing number of sources, such as NASA Astrophysics Data System (ADS), Clement's variable stars \\footnote{http://snipurl.com/gclusters3} pages, related websites, etc... ",
        "conclusions": ""
    },
    "0711/0711.2045_arXiv.txt": {
        "abstract": "We present new resolved near and mid-IR imaging and N-band spectroscopy of the two nuclei in the merger system NGC\\,3256, the most IR luminous galaxy in the nearby universe. The results from the SED fit to the data are consistent with previous estimates of the amount of obscuration towards the nuclei and the nuclear star formation rates.  However, we also find substantial differences in the infrared emission from the two nuclei which cannot be explained by obscuration alone. We conclude that the northern nucleus requires an additional component of warm dust in order to explain its properties. This suggests that local starforming conditions can vary significantly within the environment of a single system. ",
        "introduction": "The dominance of obscured star formation in the early Universe makes the dusty nearby systems interesting laboratories in which to probe their physics and evolution. Because of their high level of obscuration, such studies are well suited to observations at near and mid-IR wavelengths. Our understanding of these sources has improved considerably since the IRAS era of the 1980's. This is partly due to the results from ISO in the 1990's and most recently from Spitzer. However, high spatial resolution is still not atainable from space facilities. This situation is, however, rapidly changing with the advent of new ground based instrumentation which is a valuable complement to the space missions.  Because of its closeness ($D \\sim 39$ Mpc, for H$_{\\circ} = 70$ km s$^{-1}$ Mpc$^{-1}$) and luminosity ($L_{bol} \\sim L_{IR} \\sim 6 \\times 10^{11} L_{\\odot}$), NGC\\,3256 is one of the best studied Luminous Infrared Galaxies (LIRGs), note: it is almost a ULIRG!  Its powerful infrared emission is the result of a merger of two massive galaxies, whose nuclei have a projected separation of only $\\sim 1$ kpc, indicating an advanced stage of the collision.  NGC\\,3256 is also amongst the most X-ray luminous galaxies for which there is no clear evidence of the presence of a (luminous) active nucleus. In fact, the estimated luminosity of NGC\\,3256 ($L_{\\rm x} \\sim 10^{42}$ ergs s$^{-1}$ in the 0.5-10 keV band - Lira et al., 2002; Moran, Lehnert \\& Helfand, 1999) is consistent with the maximum limit normally adopted to separate AGN from starburst galaxies found in deep X-ray surveys (see Alexander et al. 2005, and references therein.)  In this paper we present a study of the near and mid-IR resolved emission from the two nuclei in the prototype LIRG NGC\\,3256. In section 2 we describe the high-spatial resolution, ground-based observations obtained in the N-band. In section 3 and 4 the details of our data analysis is presented. In Section 5 the SED of NGC\\,3256 is explored using our new N-band data, plus observations extracted from the literature and Spitzer and HST archival images. Finally, the discussion of our results is presented in Section 6. ",
        "conclusions": "The normalization of the starburst SED models allow us to infer a star formation rate (SFR) of $\\sim 15$ and $\\sim 6 M_{\\odot}/$yr for the northern and southern nuclei, respectively. This estimation does not take into account the Gray-body component added to the fit of the northern nucleus since its integrated luminosity corresponds to $\\la 25\\%$ of the total emission. These SFRs are in good agreement with the values obtained using the supernova rate of 0.3 yr$^{-1}$ per nucleus derived from radio observations (Norris \\& Forbes 1995) by assuming a conversion ${\\rm SNR} = 0.02 \\times {\\rm SFR}$ as inferred from the Starburst\\,99 models using a constant star formation episode (Leitherer et al., 1999).  One of the interesting results obtained from our new N-band data is establishing the photometric and spectroscopic differences between the northern and southern nuclei, which cannot be explained simply by differences in dust obscuration. Since the starburst activity in these nuclei must have been triggered by the merger of two progenitors, the age of the starbursts is likely to be comparable, but other physical parameters could be different in the two separate merging parent galaxies.  X-ray, optical and near-IR observations are indicative of a highly obscured southern nucleus which is completely hidden at visible wavelengths but which becomes increasingly prominent in the near-IR. The brightness of the two peaks are comparable at $\\sim 4 \\mu$m and at radio wavelengths. Hence, it might be anticipated that in the mid-IR, where extinction becomes negligible, the two nuclei would also have similar flux levels. But this is not observed. Instead, we find that the northern nucleus exhibits excess continuum emission at wavelengths $\\ga 10\\mu$m when compared with the southern nucleus and with our model SEDs.  Interestingly a similar behaviour was found by Alonso-Herrero et al.~(2006) when comparing the $8\\mu$m and Pa$\\alpha$ emission from a sample of LIRG nuclei, including NGC\\,3256. While the northern nucleus shows a high $8\\mu$m/Pa$\\alpha$ ratio, as was also seen in other LIRGs, the lower $8\\mu$m/Pa$\\alpha$ value of the southern nucleus is consistent with the ratios observed in HII regions. Alonso-Herrero et al.~(2006) explain this difference by the presence of large scale diffuse emission, which is characterised by strong PAH features in the mid-IR. As can be seen in Table 3, for NGC\\,3256 the EW of the 11.3 $\\mu$m PAH feature changes at most by a factor 3 between small (11.2--1.3\\arcsec) and large (10.5--24\\arcsec) apertures, while the line fluxes are greater by a factor $\\sim 6$. This suggests that the fractional increase in the continuum flux is also substantial (factor $\\ga 2$). In addition, the data presented in Figures 4 and 5 suggest that a significant fraction of the excess mid-IR flux is due to continuum emission produced by warm dust, heated in-situ by the starburst itself.  Using Spitzer IRS observations, Brandl et al.~(2006) has quantified the continuum fluxes at 6, 15 and 30 $\\mu$m for a sample of starburst galaxies, including NGC\\,3256. Brandl's observations show that the flux ratio $F_{6\\mu{\\rm m}}/F_{15\\mu{\\rm m}}=0.14$ determined for NGC\\,3256 is significantly lower than the average $<F_{6\\mu{\\rm m}}/F_{15\\mu{\\rm m}}>=0.26$ found for their starburst sample. As we have shown, the northern nucleus is more luminous that the southern nucleus by a factor of $\\sim 2$ at 5.8 $\\mu$m, while at 15$\\mu$m the flux should be dominated by the northern nucleus. Indeed, using the best fit SEDs we find $F_{6\\mu{\\rm m}}/F_{15\\mu{\\rm m}}=1.2$ for the southern nucleus, while for the northern nucleus we adopt the value tabulated from the IRS spectrum. The full range of observed flux ratios in Brandl's sample is $\\sim 0.1-0.6$, which probably reflects the range in starburst conditions, such as age and dust properties. We find that the nuclei in NGC\\,3256 are representative of the two extremes of this distribution, but are found in the same central region of a single galaxy.  Finally we compare the 10.5$\\mu$m and (unabsorbed, 2-10 keV) X-ray fluxes for both nuclei with the correlation derived by Krabbe, B\\\"oker \\& Maiolino (2001) using N-band MANIAC observations. The X-ray fluxes correspond to $9 \\times 10^{-17}$ and $4 \\times 10^{-17}$ W m$^{-2}$ (Lira et al., 2002), while the $10.5\\mu$m fluxes ($\\Delta \\lambda = 5 \\mu$m, assuming the best fit SEDs) are $\\sim 400$ and $\\sim 40$ mJy, respectively. It is found that both nuclei are X-ray deficient given their N-band photometry. In particular, because of its high mid-IR flux, the northern nucleus becomes a clear outlier in this correlation. But as shown by Krabbe's work, this is not the case for NGC\\,3256 when the fluxes are integrated over a large region, implying that the correlation holds for global starburst properties, but individual regions can show a much wider range of properties. In particular, our work shows that in-situ dust heating in LIRGs, whose emission can dominate the mid-IR emission, can be confined to very compact (few hundred-pc scale) regions."
    },
    "0711/0711.1877_arXiv.txt": {
        "abstract": "We study the effect of localized sound-speed perturbations on global mode frequencies by applying techniques of global helioseismology on numerical simulations of the solar acoustic wave field. Extending the method of realization noise subtraction (e.g. \\opencite{Hanasoge07a}) to global modes and exploiting the luxury of full spherical coverage, we are able to achieve very highly resolved frequency differences that are used to study sensitivities and the signatures of the thermal asphericities. We find that (1) global modes are almost twice as sensitive to sound-speed perturbations at the bottom of the convection zone as in comparison to anomalies well in the radiative interior ($r\\lesssim0.55 R_\\odot$), (2) the $m$-degeneracy is lifted ever so slightly, as seen in the $a$ coefficients, and (3) modes that propagate in the vicinity of the perturbations show small amplitude shifts ($\\sim 0.5\\%$). ",
        "introduction": "\\label{S-Introduction} Global helioseismology has proven very successful at inferring large scale properties of the Sun (for a review, see \\opencite{jcd.review}; \\opencite{jcd0}). Because they are very robust, the extension of methods of global helioseismology to study localized variations in the structure and dynamics of the solar interior has been of some interest (e.g. \\opencite{Swisdak}). However, the precise sensitivities of global modes to local perturbations are difficult to estimate through analytical means, especially in cases where the flows or thermal asphericities of interest possess complex spatial dependencies. To address questions relating to sensitivities and with the hope of perhaps discovering hitherto unknown phenomena associated with global modes, we introduce here for the first time a technique to study the effects of arbitrary perturbations on global mode parameters in the linear limit of small wave amplitudes.  Global modes attain resonant frequencies as a consequence of differentially sampling the entire region of propagation, making it somewhat more difficult (in comparison to local helioseismology) to pinpoint local thermal asphericities  at depth. Exactly how difficult is one of the questions we have attempted to answer in this article. Jets in the tachocline (e.g. \\opencite{jcd1}) are a subject of considerable interest since their existence (or lack thereof) could be very important in understanding the angular momentum balance of the Sun. Studying the sensitivities and signatures of waves to flows at depth may open up possibilities for their detection.   Forward modeling as a means of studying wave interactions in a complex medium like the Sun has become quite favoured (e.g. \\opencite{Hanasoge06}; \\opencite{Parchevsky}; \\opencite{Hanasoge07a}; \\opencite{Cameron}). The discovery of interesting phenomena, especially in the realm of local helioseismology (e.g. \\opencite{Hanasoge07b}; \\opencite{Birch07}), adds motivation to the pursuit of direct calculations. With the application of noise subtraction \\cite{Werne, Hanasoge07a}, we can now study the signatures of a wide range of perturbations in a realistic multiple source picture. Here, we attempt to place bounds on the detectability of thermal asphericities at various depths in the Sun. We introduce and discuss the method of simulation with a description of the types of perturbations introduced in the model in Section \\ref{simulations.sec}. The estimation of mode parameters can prove somewhat difficult due to restrictions on the temporal length of the simulation ($< 24$ hours; owing to the expensive nature of the computation). The data analysis techniques used to characterize the modes are presented in Section \\ref{peakbag.sec}. We then discuss the results from the analyses of the simulated data in $\\S$\\ref{results.sec} and summarize this work in $\\S$\\ref{conclusions.sec}. ",
        "conclusions": "\\label{conclusions.sec} We have introduced a method to systematically study the effects of various local perturbations on global mode frequencies. Techniques of mode finding and parameter fitting are applied to artificial data obtained from simulations of wave propagation in a solar-like stratified spherical shell. We are able to beat the issue of poor frequency resolution by extending the method of realization noise subtraction \\cite{Hanasoge07a} to global mode analysis. These methods can prove very useful in the study of shifts due to perturbations of magnitudes beyond the scope of first order perturbation theory; moreover, extending this approach to investigate systematic frequency shifts in other stars may prove exciting. We are currently studying the impact of complex flows like convection and localized jets on the global frequencies. Preliminary results seem to indicate that flows are stronger scatterers (larger $\\kappa$) than sound-speed perturbations although more work needs to be done to confirm and characterize these effects."
    },
    "0711/0711.1793_arXiv.txt": {
        "abstract": "} \\nc{\\eab}{ Quantum gravity effects are expected modify the primordial density fluctuations produced during inflation and leave their imprint on the cosmic microwave background observed today. We present a new analysis discussing whether these effects are detectable, considering both currently available data and simulated results from an optimal CMB experiment. We find that the WMAP data show no evidence for the particular signature considered in this work, but give an upper bound on the parameters of the model. However, a hypothetical experiment shows that with proper data, the transplanckian effects should be detectable through alternate sampling methods. This fuzzy conclusion is a result of the nature of the oscillations, since they give rise to a likelihood hyper-surface riddled with local maxima. A simple Bayesian analysis shows no significant evidence for the simulated data to prefer a transplanckian model. Conventional Markov Chain Monte Carlo (MCMC) methods are not suitable for exploring this complicated landscape, but alternative methods are required to solve the problem. This however requires extremely high-precision data. ",
        "introduction": "Quantum fluctuations in the scalar field responsible for driving inflation can give rise to a power spectrum of primordial density perturbations consistent with what is required to seed the large-scale structure (LSS) and the temperature anisotropies in the cosmic microwave background (CMB) radiation.  This fact is probably one of the main reasons for inflation having become a part of the concordance model of cosmology today.  The quantum fluctuations are imprinted during the early inflationary epoch, and become classical fluctuations in the gravitational potential as the fluctuations leave the causal horizon.  The standard way of calculating the power spectrum of the density fluctuations make use of the fact that space-time becomes Minkowskian in the distant past.  However, it has been argued \\cite{ martin-2001-63, easther-2001-64,ulf1, ulf2, jerome1, jerome2, jerome3, jerome4, bergstrom} that since a given length scale observable in the universe today shrinks and becomes smaller than the Planck length if followed sufficiently far back in time, effects of quantum gravity will at some point play a role in setting up the perturbations, and may potentially leave an observable signature. The primordial density fluctuations in the early universe may have small oscillations superimposed as a result of quantum gravitational effects. These effects are due to non-negligible curvature imposed by high energy densities in de Sitter space during the early stages of inflation. Several papers have suggested a generic shape of these effects \\cite{ulf1, ulf2, bergstrom} in the form of small oscillations superimposed on the standard, nearly scale-invariant primordial power spectrum of density fluctuations. A crucial questions is, of course, whether such a signature is observable in the universe today.  This question has been investigated in several papers \\cite{jerome1, jerome2, jerome3, jerome4, elgen, ulf1, ulf2}.  The conclusion in \\cite{elgen} was pessimistic: the so-called transplanckian effects were found to be unobservable in practice. However, this conclusion has been called into question in \\cite{jerome4}. One of the problems pointed out in that work is that the oscillations require more numerical accuracy than what is commonly employed in codes like CMBFAST \\cite{Seljak:1996is} and CAMB \\cite{Lewis:1999bs}.  In this paper, we revisit the question of the observability of transplanckian effects in the CMB, using a more accurate numerical treatment than in \\cite{elgen}.  Both present data and a hypothetical optimal CMB experiment are considered.  Furthermore, we mention alternative sampling algorithms to overcome the numerical problems encountered in exploring what turns out to be a particularly nasty likelihood hyper-surface. Section II gives a brief description of the particular model we consider and the data sets we use. In section III we discuss some technical details in the numerical computations undertaken.  Section IV then contains results obtained with currently available data sets, and in section V we investigate the possibilities with an ideal CMB data set. The exact likelihood function is described in section VI, while a Bayesian evidence analysis is performed in section VII. A discussion of the results and a suggestion for an improved approach to the search for transplanckian effects can be found in section VIII, and we conclude in section IX. ",
        "conclusions": " \\begin{itemize} \\item A low quantum gravity scale, for example the value  $\\xi = 0.0004$ chosen in our simulations, gives a very low amplitude of the oscillations in the primordial power spectrum. Hence, cosmic variance diffuses the (assumed noiseless) signal. \\item Evidence can maximally rise above $0.15$, showing that the transplanckian model is only very weakly preferred when assuming a Horava-Witten universe. \\item Even large (and unphysical) values of $\\xi$ will result in good fits for singular points due to the behaviour of the oscillations (they periodically reach zero when $\\sin x = 0$)\\\\ \\item The conventional MCMC-methods in CosmoMC fails, and we have no current computational methods implemented for solving the problem of transversing the extremely complicated likelihood surface. \\end{itemize}"
    },
    "0711/0711.4349_arXiv.txt": {
        "abstract": "Gravitational instability of a vertically thin, dusty sheet near the midplane of a protoplanetary disk has long been proposed as a way of forming planetesimals. Before Roche densities can be achieved, however, the dust-rich layer, sandwiched from above and below by more slowly rotating dust-poor gas, threatens to overturn and mix by the Kelvin-Helmholtz instability (KHI). Whether such a threat is real has never been demonstrated: the Richardson criterion for the KHI is derived for 2-D Cartesian shear flow and does not account for rotational forces. Here we present 3-D numerical simulations of gas-dust mixtures in a shearing box, accounting for the full suite of disk-related forces: the Coriolis and centrifugal forces, and radial tidal gravity. Dust particles are assumed small enough to be perfectly entrained in gas; the two fluids share the same velocity field but obey separate continuity equations. We find that the Richardson number $Ri$ does not alone determine stability. The critical value of $Ri$ below which the dust layer overturns and mixes depends on the height-integrated metallicity $\\Sigma_d/\\Sigma_g$ (surface density ratio of dust to gas). Nevertheless, for $\\Sigma_d/\\Sigma_g$ between one and five times solar, the critical $Ri$ is nearly constant at $\\sim$0.1. Keplerian radial shear stabilizes those modes that would otherwise disrupt the layer at large $Ri$. If the height-integrated metallicity is at least $\\sim$5 times greater than the solar value of 0.01, then midplane dust densities can approach Roche densities. Such an environment might be expected to produce gas giant planets having similarly super-solar metallicities. ",
        "introduction": "\\label{intro}  How do dust grains, known to permeate disks surrounding young stars, assemble into planets? One proposed stage of growth involves gravitational instability of a dust-rich layer at the disk midplane (Safronov 1969; Goldreich \\& Ward 1973). If dust grains are free to settle vertically out of gas, midplane dust densities eventually exceed Roche densities, and dust particles can begin to aggregate by self-gravity.  An objection to this means of forming planetesimals is that vertical velocities of gas may be too high for dust to settle (Weidenschilling 1980). Even apart from turbulence intrinsic to gas (e.g., turbulence driven by the magneto-rotational instability), dust layers that are too vertically thin can suffer from the Kelvin-Helmholtz instability (KHI). The KHI threatens to manifest because dust-rich gas at the midplane rotates at a rate different from that of dust-poor gas at higher altitude. Dust-poor gas experiences greater acceleration from a background radial pressure gradient, and so its rotation velocity in centrifugal balance must deviate more strongly from the Keplerian value.\\footnote{ This deviation from purely Keplerian flow is analogous to a ``thermal wind'' in the geophysical literature. See Knobloch \\& Spruit (1986) or Chapter 1 of Pedlosky (1979).} The deviation is smaller for gas laden with dust, since dust adds inertia but contributes no pressure. Usually it is assumed that the background radial pressure gradient $\\partial P/\\partial r < 0$ so that the vertical shear in the rotation velocity $\\partial v_{\\phi}/\\partial z <0$.  Whether the disk is KH-unstable is commonly assessed using the Richardson number (Chandrasekhar 1961, page 491):  $$ Ri \\equiv \\frac{ (g/\\rho) \\partial \\rho/\\partial z }{(\\partial v_{\\phi} / \\partial z)^2 } $$  \\noindent where $g$ is the vertical gravitational acceleration and $\\rho$ is the total mass density. In the limit that perturbations are incompressible, the numerator is the square of the Brunt-V\\\"{a}is\\\"{a}l\\\"{a} frequency of buoyant oscillations, while the denominator is the square of the vertical shearing rate. The Richardson number measures the amount of work required to overturn fluid elements that are originally in hydrostatic equilibrium, against the amount of free kinetic energy available in the background shear. For purely Cartesian flow (no rotational forces),  \\begin{equation} Ri < Ri_{\\rm crit} = 1/4 \\,\\,\\,\\,{\\rm is\\,\\,necessary\\,\\,for\\,\\, instability} \\label{onequarter} \\end{equation}  \\noindent (Miles 1961; Howard 1961; Chandrasekhar 1961; Drazin \\& Reid 1981; Li et al.~2003). Despite the fact that (\\ref{onequarter}) is formally not a sufficient criterion, laboratory experiments bear out its usefulness (Tritton 1988).  In the case of our disk, close to the midplane, $g = -\\Omega^2 z$ arises from the tidal gravity of the star, $z$ measures distance from the midplane, $\\Omega$ is the Keplerian angular frequency, and $\\rho = \\rho_g + \\rho_d$, where $g$ and $d$ denote gas and dust, respectively. Moreover, $|\\partial \\rho_g / \\partial z | \\ll |\\partial \\rho_d / \\partial z|$ within the thin dust layers of interest to us (see \\S\\ref{init}). Thus throughout this paper  $$ Ri = -\\frac{\\Omega^2 z}{\\rho_d + \\rho_g} \\frac{\\partial \\rho_d / \\partial z }{ \\left( \\partial v_{\\phi} / \\partial z \\right)^2} \\,. $$  \\noindent As dust settles vertically, $\\rho_d \\sim \\Sigma_d / z_d$ increases; the dust surface density $\\Sigma_d$ is assumed constant while the characteristic scale height $z_d$ of dust decreases. In the limit that $\\rho_d \\gg \\rho_g$, $Ri \\propto z_d^2$. Criterion (\\ref{onequarter}) for instability therefore places a lower bound on $z_d$ and a corresponding upper bound on $\\rho_d$. Unfortunately, for conventional solar nebular parameters---in particular for a height-integrated solar metallicity of $\\Sigma_d / \\Sigma_g = 0.01$---the maximum of $\\rho_d$ falls short of the Roche density by 1--2 orders of magnitude (e.g., Sekiya 1998; Youdin \\& Shu 2002).  But is it appropriate to apply criterion (\\ref{onequarter}) to a circumstellar disk where the flow is not Cartesian? One might guess that rotational forces merely modify $Ri_{\\rm crit}$ by a factor of order unity, since the Coriolis force and the radial shear ($d\\Omega/dr$) operate on a timescale $\\Omega^{-1}$---the same timescale characterizing vertical Brunt-V\\\"{a}is\\\"{a}l\\\"{a} oscillations and the vertical shear at $z\\sim z_d$. Yet when Coriolis forces are introduced, as in the analysis of G\\'{o}mez \\& Ostriker (2005, hereafter GO), the stability properties of the flow change dramatically. Numerical simulations by GO, performed under the assumption that dust is perfectly entrained in gas, reveal no well-defined threshold for instability. Unstable modes are detected for $Ri$ as high as 4, with growth rates that diminish with increasing $Ri$ but that show no sign of vanishing (see their Figures 9 and 10). Similar results obtain in numerical simulations by Johansen, Henning, \\& Klahr (2006, hereafter JHK), who relax the assumption of perfect coupling between dust and gas but who, like GO, retain only Coriolis forces and ignore radial shear. For their ``rocks'' and ``pebbles'' with momentum stopping times 10--50 $\\times$ shorter than $\\Omega^{-1}$, the dust distribution evolves to one where $Ri$ increases from $\\sim$1 at the midplane to $\\gtrsim 10$ at higher altitude (JHK, their Figures 4 and 5).  Here we conduct numerical simulations that account for the full complement of disk-related forces. We include not only the Coriolis force but also the centrifugal force and radial tidal gravity. The latter two forces combine to produce radial shear in rotational equilibrium. Thus we investigate the stability of doubly shearing flows: vertically shearing dust layers in near-Keplerian differential rotation. We expect our results to differ from those of GO and JHK. The modes emphasized by GO have growth rates $\\lesssim 0.1\\Omega$ at $Ri \\gtrsim 2$. Keplerian differential rotation, characterized by a strain rate of $3\\Omega/2$, should shear such modes apart before they have time to amplify.  Indeed this is the conclusion of Ishitsu \\& Sekiya (2003, hereafter IS), who numerically integrate the linearized equations of motion, including the full suite of disk-related forces, to compute the factors by which modes amplify. They find that amplification factors are sufficiently limited by radial shear that midplane dust-to-gas density ratios might reach values as high as $\\sim$2 (still too low for gravitational instability, unfortunately). They restrict, however, their linear analysis to odd-parity modes for which the vertical velocity $v_z(z) = -v_z (-z)$. G\\'{o}mez \\& Ostriker (2005) establish that even-parity modes grow faster and overturn the dust layer more effectively. Our simulations do not pre-select for either type of mode. Another difference between our work and IS is that we concentrate on $Ri= {\\rm constant}$ flows, whereas IS employ a simple, analytically tractable dust distribution. Of course, $Ri={\\rm constant}$ is merely a plausible condition to which the dust-gas mixture might relax. One of the goals of this study is to assess whether the Richardson number is a good predictor of stability even when all disk-related forces are accounted for. Finally, the analysis of IS is linear, while our numerical simulations enable access to nonlinear phenomena.  Our numerical simulations, like those of GO, assume that dust particles are perfectly entrained in gas, i.e., we assume that grain momentum stopping times $\\ll \\Omega^{-1}$. This approximation is valid for small particles, e.g., having sizes $\\ll 1$ m at disk radius $r = 1$ AU in a minimum-mass nebula (e.g., Weidenschilling 1977). The hydrodynamics code and initial conditions are described in \\S\\ref{method}. Results are presented in \\S\\ref{res}; there we determine $Ri_{\\rm crit}$ for our doubly shearing flows. A summary is provided in \\S\\ref{sum}.  While our entire study is rooted in the literature on the Kelvin-Helmholtz instability and its impact on planetesimal formation, and as such revolves around the Richardson criterion, it was brought to our attention by the referee that the KHI may not be the only instability afflicting the dusty midplane. Baroclinic instabilities, studied in disks by Cabot (1984), Knobloch \\& Spruit (1986), and Arlt \\& Urpin (2004) may also be relevant. These authors study baroclinity specifically in the context of a non-zero vertical shear. More generally, a baroclinic flow is one whose isobaric surfaces do not coincide with its isodensity surfaces. Our flows are baroclinic because of molecular weight gradients: dust adds to the density but not to the pressure. Baroclinic instability can strike even when the Richardson number is large. We will connect our findings to the baroclinic instability when appropriate; indeed, in \\S\\ref{cori} we propose that the instability discovered by GO in their simulations of non-radially-shearing disks is, in fact, the baroclinic instability discovered by Cabot (1984). And we will see in \\S\\ref{varsig} that the Richardson criterion alone does not determine stability under arbitrary conditions. ",
        "conclusions": "\\label{sum} We have performed numerical simulations of flows near the dust-rich midplanes of protoplanetary disks. Dust particles are assumed small enough that they are perfectly entrained in gas, and disk self-gravity is neglected. The midplane flow is doubly shearing; the rotational velocity varies with height according to the vertical gradient in dust density, and it also varies with radius according to the usual Keplerian shear. The simulations are necessarily three-dimensional and performed in a shearing box.  Despite the complications introduced by rotation and tidal gravity, the Richardson criterion for the onset of the Kelvin-Helmholtz instability still proves useful for solar to moderately super-solar height-integrated metallicities. Dust layers characterized by constant Richardson number and $\\Sigma_d/\\Sigma_g \\approx 0.01$--0.05 overturn and mix if $Ri \\lesssim 0.1$, but remain intact at larger $Ri$. This result contrasts with the situation when only the Coriolis force is accounted for and the radial shear is suppressed: in that case, unstable modes persist for $Ri$ as large as 16. But because these ``Coriolis-only'' (and likely baroclinic; Cabot 1984) modes at large $Ri$ grow at rates that are substantially slower than the Kepler strain rate of $3\\Omega/2$, they are stabilized and rendered impotent by the Kepler shear.  In hindsight, our guess (\\S\\ref{intro}) that the critical $Ri$ dividing stability from instability might only change by a factor of order unity proved correct, though only for solar to moderately super-solar metallicities: we find that $Ri_{\\rm crit}$ for our doubly shearing flows is about half that of the traditional value of 1/4, if $\\Sigma_d/\\Sigma_g \\approx 0.01$--0.05. This is sensible insofar as $Ri_{\\rm crit} \\approx 0.1$ gives a vertical shearing frequency $\\partial v_y / \\partial z$ that is a few times faster than any of the other frequencies of the problem: the Brunt-V\\\"{a}is\\\"{a}l\\\"{a} frequency for vertical oscillations (this reaches a maximum of $\\sim$$\\Omega$ at $z\\sim z_d$), the Kepler strain rate ($3\\Omega/2$), and the Coriolis turning frequency ($2\\Omega$). Deeper insights, including an understanding of why $Ri_{\\rm crit} \\ll 0.1$ for $\\Sigma_d/\\Sigma_g \\ll 0.01$, might be gained by study of the geophysical literature, where doubly shearing, rotating flows are commonplace (see, e.g., Chapter 7 of Pedlosky 1979, which discusses baroclinic instabilities).  We conclude that for a standard, height-integrated, solar metallicity of $\\Sigma_d/\\Sigma_g = 0.01$, the dust density at the midplane can be at most $\\sim$$1.5 \\times$ greater than the gas density. If we assume a gas density appropriate to a minimum-mass disk, then the maximum total density of gas and dust is still too low, by a factor of $\\sim$25, for the marginally stable layer to become gravitationally unstable. The literature discusses two remedies. The first is to find ways of enhancing $\\Sigma_d/\\Sigma_g$ by factors of 3--10, thereby increasing the dust-to-gas ratio at the midplane by factors of $\\sim$30 at fixed $Ri = 0.1$ (Sekiya 1998; Youdin \\& Shu 2002; see also our Figures \\ref{dens_sigma05} and \\ref{f9}, and run S8 in Table 2). Possible means of increasing the height-integrated metallicity include (i) decreasing $\\Sigma_g$ through photoevaporation of gas, (ii) increasing $\\Sigma_d$ by radial drift and pile-up of particles (Youdin \\& Shu 2002; Youdin \\& Chiang 2004) or (iii) increasing $\\Sigma_d$ by radiation blow-back of grains into the rim of a transitional disk (Chiang \\& Murray-Clay 2007). Gas giants that form in such a metal-rich environment might be expected to have metallicities enhanced above the solar value by similar factors of 3--10. Indeed some hot Jupiters have remarkably metal-rich interiors (Sato et al.~2005). Alternatively, one can relax the assumption of perfect coupling between dust and gas, and exploit drag instabilities that result from finite momentum stopping times and the backreaction of dust on gas (Youdin \\& Goodman 2005; Johansen et al.~2007)."
    },
    "0711/0711.2209_arXiv.txt": {
        "abstract": "{ High precision photometry and spectroscopy of low--mass stars reveal a variety of properties standard stellar evolution cannot predict. Rotation, an essential ingredient of stellar evolution, is a step towards resolving the discrepancy between model predictions and observations.  The first rotating stellar model, continuously tracing a low--mass star from the pre--main sequence onto the horizontal branch, is presented. The predicted luminosity functions of stars of globular clusters and surface rotation velocities on the horizontal branch are discussed. ",
        "introduction": "Rotation has to be included in realistic stellar models. Since recent high precision photometry and spectroscopy have become available, the effects of slower rotation on low--mass stars has become detectable in the colour--magnitude diagram (CMD) and Luminosity Functions (LFs). To predict them precisely, low--mass stellar models need to take rotation into account. \\subsection{Stellar Structure} The version of the stellar evolution code FRANEC, described in \\citet{Pietrinferni2004}, is used for our computations. Rotation was included in the stellar structure equations following the prescription of \\citet{Kippenhahn1970}. The required parameters $f_{\\rm P}$ and $f_{\\rm T}$ were derived analytically from the effective gravitational potential of the star approximating its equipotential surfaces as rotational ellipsoids.  We have eliminated the subatmospheric region used in the original code and increased the density of the outermost meshpoints in the stellar model. \\subsection{Transport Mechanisms} Chemical elements and angular momentum are transported within the star during its evolution. In a convective zone (CZ) the transport is assumed to be instantaneous. The abundance of chemical elements is averaged in a CZ and the angular momentum redistributed to achieve either specific angular momentum conservation or solid body rotation. Both distributions are limiting cases introduced by \\citet{Sweigart1970}. Solid body rotation is chosen for CZ until the the star reaches the turn off (TO), after which the specific angular momentum is conserved (see \\citealt{Palacios2006b} and \\citealt{Sills2000}).  In a radiative region the transport is a diffusive process, since \\citet{Palacios2006a} have shown that advective transport via meridional currents is less efficient than via diffusion resulting from the shear instability. This instability dominates rotational mixing \\citep{Zahn1992}. The diffusion coefficient $D_{v}$ is taken from \\citet{Denissenkov2004} introducing an efficiency parameter $F$. Chemical elements are transported by a combination of atomic diffusion and rotational mixing. Angular momentum is diffused by rotational mixing alone. \\subsection{Braking} The total amount of angular momentum of the star is changed by several processes. On the pre-main sequence (PMS) magnetic fields connect the star with its circumstellar disk. The disk allows angular momentum to be diffused beyond the influence of the magnetic field of the star, preventing the PMS spin-up. This process of disk braking \\citep{Shu1994} is active until a time $\\tau_{\\rm DB}$.  On the main sequence (MS) the stellar magnetic field increases the amount of angular momentum lost by its stellar wind. The angular momentum lost by this magnetic--wind braking is given by \\citet{Chaboyer1995}. We include a function determining the efficiency to generate a magnetic field $B_{\\rm eff}$ and  parametrise it by the mass of the convective envelope. The shape of $B_{\\rm eff}$ is determined by two parameters ($M_w$;$B_w$). The overall efficiency of the braking process is determined by $f_k$.  Angular momentum lost via mass loss is already included in the code, and is dominant on the red giant branch (RGB). ",
        "conclusions": "Given an isochrone without rotation that is consistent with observational constraints. The inclusion of rotation increases the helium core mass, this leads to a brighter ZAHB and distances from HB fitting that are no longer consistent with parallax based empirical MS fitting \\citep{Recio2004}.  A solution would be to reduce the rotation rate in the helium core on the RGB. This reduction may result from either more realistic angular momentum transport in CZs (see exploratory work by \\citealt{Palacios2006b}) or reducing $\\omega_0$.  We chose to reduce the initial rotation rate by adopting a different method to transpose solar metallicity rotation rates to lower metallicities. Instead of keeping the total amount of initial angular momentum constant, $\\omega_0$ is defined so that these stars reach a rotational period of $P_{\\rm 1\\,Myr}$=5\\,d. This reduces $\\omega_0$ by 60\\,\\% and leads to a SGB rotation velocity of 3\\,km\\,s$^{-1}$, consistent with observations. The ZAHB including rotation is only brighter by $\\Delta M_V$=0.122 (comparable to typical photometric errors for HB stars) and leads to ages younger by 1.5\\,Gyr. Therefore, an isochrone or LF with and without rotation can only be consistent with observational constraints, if the rotational isochrone or LF is 1.5\\,Gyr younger than without rotation.  A consistent comparison between a 10.5\\,Gyr old LF with the reduced $\\omega_0$ (solid line) and a 12.0\\,Gyr LF without (dashed line), is shown in Fig.~\\ref{lf}. The prior determined 11\\,\\% increase in number count is reduced to 5\\,\\%, but it is still detectable. \\begin{figure}[] \\resizebox{\\hsize}{!}{\\includegraphics[clip=true]{LFSGBcompareV3.eps}} \\caption{ \\footnotesize Three LFs are shown as a function of $M_V$ with an age of 12\\,Gyr without (solid line) and with rotation (dashed lines), as well as an age of 10.5\\,Gyr including the reduced initial rotation (dotted line). } \\label{lf} \\end{figure}"
    },
    "0711/0711.1163.txt": {
        "abstract": "\\vspace*{0.2in}  We analyzed the available LIGO data coincident with GRB~070201, a short duration hard spectrum $\\gamma$-ray burst whose electromagnetically determined sky position is coincident with the spiral arms of the Andromeda galaxy (M31). Possible progenitors of such short hard GRBs include mergers of neutron stars or a neutron star and black hole, or soft $\\gamma$-ray repeater (SGR) flares. These events can be accompanied by gravitational-wave emission. No plausible gravitational wave candidates were found within a 180~s long window around the time of GRB~070201. This result implies that a compact binary progenitor of GRB~070201, with masses in the range $1~M_\\odot < m_1 < 3~M_\\odot$ and $1~M_\\odot < m_2 < 40~M_\\odot$, located in M31 is excluded at $> 99\\%$ confidence. Indeed, if GRB~070201 were caused by a binary neutron star merger, we find that $D< \\checked{3.5} \\textrm{ Mpc}$ is excluded, assuming random inclination, at $90\\%$ confidence. The result also implies that an unmodeled gravitational wave burst from GRB~070201 most probably emitted less than $4.4\\times 10^{-4} M_{\\odot}c^2$ (7.9$\\times$10$^{50}$ ergs) in any 100~ms long period within the signal region if the source was in M31 and radiated isotropically at the same frequency as LIGO's peak sensitivity ($f \\approx$~150~Hz). This upper limit does not exclude current models of SGRs at the M31 distance.   \\vspace*{0.2in} ",
        "introduction": "\\label{sec:intro}   Gamma-ray bursts (GRBs) are intense flashes of $\\gamma$-rays which are observed to be isotropically distributed over the sky~\\cite[see, e.g.:][and references therein]{1973ApJ...182L..85K,ADM:Pir2005,ADM:ME2002}. The variability of the bursts on short time scales indicates that the sources are very compact. Combined observations, using $\\gamma$-ray and x-ray satellites such as \\href{http://heasarc.gsfc.nasa.gov/docs/heasarc/missions/vela5a.html}{Vela}, \\href{http://www.batse.msfc.nasa.gov/batse/}{CGRO}, \\href{http://bepposax.gsfc.nasa.gov/bepposax/index.html}{BeppoSax}, \\href{http://space.mit.edu/HETE/index.html}{HETE}, \\href{http://www.swift.psu.edu/}{Swift}, \\href{http://heasarc.gsfc.nasa.gov/docs/heasarc/missions/wind.html}{Konus-Wind}, and \\href{http://sci.esa.int/science-e/www/object/index.cfm?fobjectid=31083&fareaid_1=21&fareaid_2=21&farchive_objecttypeid=15&farchive_objectid=30995&fchoice=-1&startz=1&startpage=1}{INTEGRAL} \\cite[see][and references therein]{ 1973ApJ...182L..85K, 1992Natur.355..143M, 1999ApJS..122..465P, 2000ApJS..127...59F, 1981ApSS..75...47M, 2004ApJ...611.1005G}, missions, as well as by the Interplanetary Network (IPN), with follow-up by X-ray, optical and radio telescopes of the region around GRBs, have yielded direct observations of afterglows from $\\approx$350 GRBs. In turn, host galaxies were identified for many GRBs and redshifts were determined for $\\approx$125 bursts. The redshifts indicated that GRBs are of extra-galactic origin.  Two types of GRBs are distinguished by their characteristic duration~\\cite[see][]{grbs:general:discovery:multiple-classes, 2006Natur.444.1044G} and are understood to have different origins.  Long GRBs have duration $\\gtrsim 2 \\textrm{ s}$. Detailed observations of long GRBs demonstrate their association with star-forming galaxies ranging up to a redshift of $z\\approx 6.3$ \\cite[see][and references therein]{2006Natur.440..184K,2006ApJ...637L..69W,2006AA...447..897J}. Furthermore, several nearby long GRBs have been spatially and temporally coincident with supernovae \\cite[e.g.][]{Campana:2006qe,2004ApJ...609L...5M,Hjorth:2003jt,Galama:1998ea,2006ARA&A..44..507W}.  Short GRBs have duration $\\lesssim 2 \\textrm{ s}$.  The progenitors of short GRBs are not so well understood. While there are associations with distant galaxies of different types and different star formation histories, there are also powerful bursts of $\\gamma$-rays from Galactic sources, such as SGR~1806$-$20 \\citep{2006ApJ...640..849N,2005Natur.434.1098H,2005Natur.434.1107P}. However, statistical analyses indicate that at most 15\\% of known short GRBs can be accounted for as soft $\\gamma$-ray repeaters (SGRs) \\citep{2006ApJ...640..849N,2007arXiv0709.4640C}. Moreover, the spectral characteristics and energetics of some observed short GRBs and their afterglows seem to contradict this hypothesis in most cases \\citep{2006ApJ...640..849N}. The current leading hypothesis to explain most short GRBs is the merger of neutron star or neutron star $+$ black hole binaries \\cite[see for example][and references therein]{NakarReviewArticle2006,2007ApJ...654..878B}. However, to date no observations have definitively confirmed the association between short GRBs and binary mergers.  Therefore, given the candidate sources, it is plausible that GRB central engines are also strong gravitational wave (GW) emitters at frequencies accessible to ground-based detectors like LIGO, GEO-600, and Virgo \\citep{2005PhRvD..72d2002A,0264-9381-23-19-S01,Willke:2002,1993ApJ...417L..17K,LIGO-Bursts-s3,LIGO-Inspiral-s2-bns,2004ApJ...607..384F}. Bursts of gravitational waves are expected to be emitted during the GRB event, with a characteristic duration comparable to that of the associated GRB, though the amplitude and frequency spectrum of the gravitational-wave burst are unknown. In the case of short GRBs produced by compact binary mergers, gravitational waves with relatively well-modeled amplitude and frequency evolution will be emitted during the inspiral phase of the binary system, preceding the event that produces the GRB.   GRB~070201 was an intense, short duration, hard spectrum GRB, which was detected and localized by 4 IPN spacecraft (Konus-Wind, INTEGRAL and MESSENGER); it was also observed by Swift (BAT)  but with a high-intensity background as the satellite was entering the South Atlantic Anomaly~\\citep{2007GCN..6088....1G}. The burst light-curve exhibited a multi-peaked pulse with duration $\\sim0.15$\\,s, followed by a much weaker, softer pulse that lasted $\\sim0.08$\\,s. Using early reports, Perley and Bloom~\\citep{2007GCN..6091....1P} pointed out that the initial IPN location annulus of the event intersected the outer spiral arms of the Andromeda galaxy (M31). The refined error box, centered $\\approx$~1.1 degrees from the center of M31, was later reported~\\citep{2007GCN..6098....1P,2007GCN..6103....1H}, and it still overlaps with the spiral arms of M31 [see Figure~\\ref{fig:M31_ErrorBox} and~\\citep{HurleyPrivate,IPN}]. Based on the Konus-Wind observations~\\citep{HurleyPrivate, 2007GCN..6094....1G}, the burst had a fluence of $1.57 (-0.21, +0.06)\\times10^{-5}$~erg$\\cdot$cm$^{-2}$ in the 20~keV -- 1~MeV range.  \\begin{figure} \\includegraphics[width=\\columnwidth]{IPN_Error_Box_GRB070201.eps} \\caption{\\label{fig:M31_ErrorBox} The IPN3~\\citep{IPN} ($\\gamma$-ray) error box overlaps with the spiral arms of the Andromeda galaxy (M31). The inset image shows the full error box superimposed on an SDSS~\\citep{2006ApJS..162...38A,SDSS} image of M31. The main figure shows the overlap of the error box and the spiral arms of M31 in UV light~\\citep{Thilker_2005ApJ}. } \\end{figure}  It was also pointed out~\\citep{2007GCN..6094....1G} that if the burst source were actually located in M31 (at a distance of $\\simeq770$~kpc) the isotropic energy release would be $\\sim10^{45}$~erg, comparable to the energy release in giant flares of soft $\\gamma$-ray repeaters: e.g., the 5$^{\\rm  th}$ March 1979 event from SGR~0526$-$66 ($\\sim2\\times10^{44}$~erg in the initial pulse) and the 27$^{\\rm th}$ December 2004 event from SGR~1806$-$20 ($\\sim2\\times10^{46}$~erg). Conversely if the event had an isotropic energy release more typical of short hard GRBs, e.g., $\\sim10^{48}-10^{52}$~erg \\citep{berger-shortgrb-parameter-correlation-2007}, then it would have to be located at least $\\sim30$ times further than M31 (i.e., further than $\\sim23$~Mpc).  At the  time of GRB~070201, the Hanford detectors of the Laser Interferometer Gravitational-wave Observatory (LIGO)~\\citep{LIGOS1instpaper} were stable and recording science-quality data, while the LIGO Livingston, GEO-600, and Virgo detectors were not taking data. The LIGO data around GRB~070201 were searched for evidence of a gravitational wave signal from compact binary inspiral or the central engine of the GRB itself.  A standard measure of the sensitivity of a detector to gravitational waves is the distance to which an optimally oriented and located double neutron star binary would produce a response in the datastream that, when optimally filtered  for the inspiral waves, peaks at a signal to noise ratio of 8 \\citep[see, e.g.][and references therein]{LIGO-Inspiral-s2-bns}. At the time of GRB~070201, this distance was \\checked{35.7~Mpc} and \\checked{15.3~Mpc} for the Hanford 4~km and 2~km detectors, respectively. However, the sensitivity of a detector to a gravitational wave depends on the location of the source on the sky and on the polarization angle of the waves. In the case of compact binaries, it also depends on the inclination angle of the orbital plane relative to the line of sight. At the time of GRB~070201, the binary inspiral reach in the direction of M31 was only about $43\\%$ of this maximum. More details of the instrumental sensitivity can be found in Sec.~\\ref{sec:LIGO-Observations}.  The search for gravitational waves from a compact binary inspiral focused on objects with masses in the ranges $1~M_\\odot < m_1 < 3~M_\\odot$ and $1~M_\\odot < m_2 < 40~M_\\odot$. The core of the search is matched filtering, cross-correlating the data with the expected gravitational waveform for binary inspiral  and uses methods reported previously \\citep[see, e.g][and references therein]{LIGO-Inspiral-s2-bns}. Uncertainties in the expected waveforms can lead to decreased sensitivity of the search to the gravitational wave signal from the inspiral phase; this is particularly true of systems with higher masses and systems with substantial spin~\\citep{GrandclementKalogeraVecchio:2003}. This is accounted for by studying the dependence of sensitivity of the search to a variety of model waveforms based on different approximation methods.   The search for more generic transient gravitational waves coincident with the $\\gamma$-ray burst is based on cross-correlating data from two detectors and does not make use of a specific model for the gravitational wave signal.  This is an appropriate method when the gravitational wave signal is not well modeled theoretically, such as signals from the actual merger phase of a compact binary system or the core collapse phase of a supernova event.  The remainder of the paper is organized as follows. In Sec.~2, we discuss the LIGO detectors and the data taken around the time of GRB~070201.  In Sec.~3, we report on the inspiral gravitational wave search, briefly reviewing the methods and algorithms used, and concluding with the astrophysical implications of the search for the GRB~070201 event. In Sec.~4, we report on the search for burst-like gravitational wave signals and present the astrophysical implications of that search. The software used in this analysis is available in the LIGO Scientific Collaboration's data analysis code archives \\citep{lscsoft} \\footnote{The search for inspiraling binaries (Sec~3.) used LAL and LALAPPS with tag s5\\_grb070201\\_20070731 and the burst search (Sec.~4) used the MATAPPS package grbxcorr with tag grbxcorr\\_r1}. Since no plausible gravitational wave signal was detected above the background either in the inspiral or the burst search, we present the astrophysical implications of these results on the understanding of short GRBs in Sec.~5. ",
        "conclusions": "We analyzed the data from the LIGO H1 and H2 gravitational-wave detectors, looking for signals associated with the electromagnetic event GRB~070201. No plausible gravitational-wave signals were identified. Based on this search, a compact binary progenitor (neutron star + black hole or neutron star systems) of GRB~070201, with masses in the range $1~M_\\odot < m_1 < 3~M_\\odot$ and $1~M_\\odot < m_2 < 40~M_\\odot$, located in M31 is excluded at the $> 99\\%$ confidence.  Our model-independent search did not find correlated signatures inconsistent with the noise within the H1 and H2 data-streams that could be related to GRB~070201. Based on the sensitivity of our search and assuming isotropic gravitational-wave emission of the progenitor, an upper limit on the power emitted in gravitational waves by GRB~070201 was determined. A gravitational wave with characteristic frequency within the most sensitive range of the LIGO detectors ($f\\approx$~150~Hz) \\checked{most probably} emitted less than E$_\\mathrm{GW} <$7.9$\\times$ 10$^{50}$~ergs within any 100-ms-long time interval inside the on-source region if the source is in M31. This limit on radiated power is comparable to the emitted power of some GRBs. However, it is significantly higher than the associated electromagnetic emission of this particular GRB. Therefore the transient search only constrains the binary inspiral models for a short hard GRB in M31 and does not impose new limitations on magnetar-driven (SGR type) burst models.  As gravitational-wave observations continue and the sensitivity of the instruments improves, we look forward to the astrophysical insights that combined electromagnetic and gravitational observing campaigns can bring."
    },
    "0711/0711.2515_arXiv.txt": {
        "abstract": "Observational surveys of galaxies are not trivially related to single-epoch snapshots from computer simulations.  Observationally, an increase in the distance along the line-of-sight corresponds to an earlier cosmic time at which the properties of the surveyed galaxy population may change.  The effect of observing a survey volume along the light-cone must be considered in the regime where the mass function of galaxies varies exponentially with redshift.  This occurs when the halos under consideration are rare, that is either when they are very massive or observed at high-redshift.  While the effect of the light-cone is negligible for narrow-band surveys of Ly$\\alpha$ emitters, it can be significant for drop-out surveys of Lyman-break galaxies (LBGs) where the selection functions of the photometric bands are broad.  Since there are exponentially more halos at the low-redshift end of the survey, the low-redshift tail of the selection function contains a disproportionate fraction of the galaxies observed in the survey.  This leads to a redshift probability distribution (RPD) for the dropout LBGs with a mean less than that of the photometric selection function (PHSF) by an amount of order the standard deviation of the PHSF.  The inferred mass function of galaxies is then shallower than the true mass function at a single redshift with the abundance at the high-mass end being twice or more as large as expected.  Moreover, the statistical moments of the count of galaxies calculated ignoring the light-cone effect, deviate from the actual values. ",
        "introduction": "It has become common practice to analyse and interpret the observed abundance and distribution of high-redshift galaxies by approximating a limited survey volume to a single-epoch snapshot of the Universe \\citep{Ellis07}. The observed data is then compared to theoretical predictions which were calculated for an idealized snapshot of this nature.  However, in actual observations, an increase in the distance along the line-of-sight corresponds to an earlier cosmic time at which the properties of the surveyed galaxy population may change.  The ``snapshot approximation'' is adequate for galaxy surveys at low redshifts, when galaxy halos are common and their mass function is not evolving rapidly with cosmic time. At these low redshifts, a relatively small region of space spanning a narrow redshift range can still be sufficiently large to contain an adequate sample of these abundant objects. However, the validity of the approximation should be carefully examined at high redshifts when massive galaxies are rare and their abundance varies exponentially with redshift.  To illustrate the situation at high redshifts, let us consider two regions of the same shape centered at different redshifts and containing the same number of galaxy halos of a particular mass. The region centered at the higher redshift will span a larger range in redshift for two reasons. First, since halos of a given mass are rarer at a higher redshift, the higher redshift region must have a larger comoving size than the one at smaller redshift for each to contain the same number of halos.  Second, the same comoving distance corresponds to a larger redshift interval at higher redshift than at lower redshift (i.e. $dz = [H(z)/c] d\\chi$, where $\\chi$ is the comoving length and the Hubble parameter, $H(z)$, is an increasing function of $z$).  The difference between snapshot analysis (on a space-like hypersurface) and that along the light-cone is becoming increasingly relevant with purported discoveries of very massive galaxies near $z = 6$ \\citep{Mobasher05} and as new surveys probe redshifts up to $z=10$ \\citep{Bouwens06,Iye06,Stark07}.  Even at high-redshifts, narrow-band surveys of Lyman-alpha emitters (LAEs) span such a small range of redshifts that they are unaffected by the exponential change in the mass function of halos with redshift.  However, the photometric selection functions (PHSFs) of the bands used in drop-out surveys of Lyman-break galaxies (LBGs) can be fairly broad in redshift space \\citep{BI06}.  In this paper, we examine the significance of light-cone distortions on the inferred abundance and clustering properties of high-redshift galaxies in dropout surveys of LBGs.  First, we describe the PHSFs used in dropout surveys in \\S \\ref{selection}.  Subsequently, we derive analytic formulae for the first and second statistical moments of the count of halos in a given survey volume (\\S \\ref{moments}) and consider the two-point correlation function of halos on the light-cone (\\S \\ref{correlation}).  In \\S \\ref{model}, we review a simple model for high-redshift star forming galaxies by \\citet{SLE07}, that gives the luminosity of LBGs and LAEs contained in a halo of a given mass.  We then use this model to calculate the quantitative difference between our light-cone formalism and the standard snapshot approach for various survey volumes (\\S \\ref{results}), exploring the dependence on cosmological parameters. Finally, we discuss the significance of our results in \\S \\ref{discussion}.  Unless otherwise stated, we assume a flat, $\\Lambda$CDM cosmology with cosmological parameters $\\left( \\Omega_m, \\Omega_{\\Lambda}, \\Omega_b, h, \\sigma_8, \\alpha, r\\right) = \\left( 0.268, 0.732, 0.042, 0.704, 0.776, 0.947, 0.000 \\right)$ \\citep{Spergel07}.  All distance scales are comoving. ",
        "conclusions": "\\label{discussion}  The variation of the mass function of halos along the light-cone within the volume of a dropout survey results in a mass segregation effect that is also manifested in their hosted LBGs.  The observed sources are not at the mean of the photometric selection function (PHSF) but instead are distributed at lower redshifts (see Fig. \\ref{Nplot}).  This effect applies to true LBG dropouts and ignores possible interlopers from lower redshifts whose spectra mimic those of higher redshift LGBs. The mass segregation results in different measured statistics of LBGs from those expected from theory or simulations along a space-like slice (snapshot) through the universe at the ``mean survey redshift.\"  In particular, the mass/luminosity function extracted from such a survey is shallower than the underlying mass/luminosity function because of the different strengths of the light-cone effect on halos of different masses (LBGs of different luminosities).  This flattening of the mass function is particularly important in the J-band because of its larger width in redshift space (see Fig. \\ref{mfplot}).  Without spectroscopic measurements to confirm the redshifts of a sample of LBGs, the mass segregation effect presents an added complication in continuing efforts to determine the effect of high-redshift LBGs on reionization \\citep{Nagamine06,SLE07} and the microwave background \\citep{BL07}, or efforts to measure the high-redshift evolution of the star-formation rate \\citep{ST06,Ellis07}.  Reionization is a highly inhomogeneous process, and so the light-cone effect on the correlation function is also particularly relevant in that context \\citep{BL04,FL05}.  The lightcone effect on the measured correlation function of LBGs would also be important in attempts to use the clustering properties of LBGs to infer the masses of their host halos at higher redshift.  \\citet{Dow-Hygelund07} made an effort to follow up spectroscopically on LBGs from i-band dropout surveys near $z\\! \\simeq \\!6$ to measure their redshifts precisely.  However, they were only able to confirm redshifts on 6 LBGs in their sample, a number insufficient to trace the RPD in a statistically significant way.  \\citet{Ando04,Ando07} preform similar studies on LBGs near $z \\!\\simeq\\! 5$, but the sample size they obtained was also not sufficiently large for this purpose.  Ideally, one would like to perform this type of analysis on z- or J-band dropouts with the goal of tracing the RPD, but this would be very difficult both because of the present lack of candidates and because of the high integration times necessary with current technology."
    },
    "0711/0711.2386_arXiv.txt": {
        "abstract": "We investigate the deflections of UHE protons by Galactic magnetic field(GMF) using four conventional GMF models in order to discuss the positional correlation between the arrival distribution of UHECRs and their sources. UHE protons coming from the direction around the Galactic center are highly deflected above $8^{\\circ}$ by the dipole magnetic field during their propagation in Galactic space. However, in bisymmetric spiral field models, there are directions with the deflection angle below $1^{\\circ}$. One of these directions is toward Centaurus A, the nearest radio-loud active galactic nuclei that is one of possible candidates of UHECR sources. On the other hand, UHE protons arriving from the direction of the anti-Galactic center are less deflected, especially in bisymmetric spiral field models. Thus, the northern hemisphere, not including the Galactic center, is suitable for the studies of correlation with sources. The dependence on model parameters is also investigated. The deflection angles of UHE protons are dependent on the pitch angle of the spiral field. We also investigate distortion of the supergalactic plane by GMF. Since the distortion in the direction around Galactic center strongly depends on the GMF model, we can obtain information on GMF around Galactic center if Pierre Auger Observatory finds the significant positional correlation around the supergalactic plane. ",
        "introduction": "\\label{intro}  The origin of ultra-high-energy cosmic rays(UHECRs) is one of the most intriguing problems in astroparticle physics. In order to reveal their nature and sources, UHECR observatories with larger exposures have been constructed. Now, the total exposure of Pierre Auger Observatory(PAO) already exceeds that of Akeno Giant Air Shower Array(AGASA), which had been the largest detector before the PAO era. Expectations for elucidation of UHECR sources are raised since PAO can detect more than 100 events per year above $4 \\times 10^{19}$eV, which is an energy threshold that AGASA found small-scale anisotropy of observed UHECR arrival distribution from its 57 events.   The discussions toward UHECR astronomy have been recently began. There seems to be two standpoints for UHECR astronomy. One is study of direct correlation between UHECR sources and observed UHECR arrival directions. This can be expected if the Galactic magnetic field(GMF) and extragalactic magnetic field(EGMF) are weak to deflect UHECR trajectories weakly. However, the nature of the GMF and EGMF is poorly known observationally and theoretically. Recently, several simulations of the large-scale structure formation with magnetic field has described local magnetic structures \\citep*{sigl03,dolag05}. \\cite{sigl03} claimed that UHECR astronomy may not be possible if their source distribution model, magnetic structure and observer position are confirmed by future observation since UHE protons with $10^{20}$eV are even deflected $20^{\\circ}$ or more by their structured EGMF. On the other hand, \\cite{dolag05} showed that most of highest energy cosmic rays are very weakly deflected by EGMF since strong magnetic field ($\\sim \\mu$G) is highly localized at clusters of galaxies. Therefore, the deflection angles of UHECRs are controversial. The understanding of the GMF is known better than that of EGMF, but poor, as briefly reviewed in section \\ref{gmf}. Number density of UHECR sources has also a crucial role. If the source number density is large, the number of observed events to unveil UHECR sources or source distribution is large. Recently, PAO reports the positional correlation between the arrival directions of highest energy events and nearby AGN\\citep*{auger07}. AGASA also reported small-scale anisotropy of UHECR arrival distribution within its angular resolution\\citep*{takeda99}. These results are one of corroborating evidences for this approarch. The UHECR source number density is constrained at $\\sim 10^{-5}~{\\rm Mpc}^{-3}$ using small-scale anisotropy observed by AGASA, but this constraint has also large uncertainty due to small number of observed events \\citep*{blasi04,kachelriess05,takami07a}. However, this uncertainty is well reduced by nearly future observation \\citep*{takami07a}. In any case, these uncertainty prevents our understanding of the positional correlation study at present.   The other is purely statistical approach. On the arrival distribution of highest energy cosmic rays in data combined with results of several observatories before the PAO era, medium-scale ($\\sim 20^{\\circ}$) anisotropy was found \\citep{kachelriess06}. This angular scale corresponds a typical scale of event clusterings. The cumulative two-point auto-correlation function of the observed events can well reproduce that of mock data calculated by Monte-carlo simulation assuming that UHECR source distribution traces nearby large-scale structure within $z\\simeq 0.02$ \\citep*{cuoco07a}. They claimed that the medium-scale anisotropy reflects the local large-scale structure as UHECR sources. \\cite{cuoco07b} suggested that such a statistical method is a powerful tool to reveal UHECR source distribution even if the GMF is considered. This medium-scale anisotropy was also found in PAO data\\citep*{mollerach07}.   Our recent studies have been based on the former standpoint. In our previous study \\citep*{takami07b}, we discussed the possibility that future observations of UHECR arrival distribution can unveil large-scale structure of UHECR source distribution. We calculated the positional correlation between simulated arrival distribution of UHE protons and their source distribution which reproduces the local structures, taking into account their propagation in a structured EGMF with plausible strengths. Our EGMF model also constructed under simple assumptions so that it reproduces the observed local structures. Taking into account uncertainties on the EGMF strength and the source number density, we investigate the positional correlation in several numbers of their parameters. Our studies using source distribution and EGMF model based on astronomical observations just predict the recent PAO result. We concluded that five year observation by PAO can unveil nearby UHECR source distribution at a few degree scale if local number density of UHECR sources is $\\sim 10^{-5}~{\\rm Mpc}^{-3}$. More observation is needed if the local number density is larger. The best indicator of the correlation is UHE protons above $10^{19.8}$eV.   In the previous study, we considered EGMF as a first step towards the direct positional correlation study between UHECR arrival directions and source positions. The GMF was neglected. However, the GMF affects the arrival directions of UHECRs \\citep*{alvarez02,takami06,kachelriess07}. Thus, as a next step, we should discuss the deflections of UHECRs by the GMF.   In this study, we calculate trajectories of UHE protons in the Galactic space and investigate their deflections by the GMF. We adopt the backtracking method for calculating the trajectories. Trajectories of particles with the proton mass and the charge of -1 injected from the Earth can be regarded as those of extragalactic protons. We focus on UHE protons above $10^{19.8}$eV, which is best indicator of UHECR sources. We also investigate the distortion of extragalactic sky by the GMF and discuss its effect to the correlation study.  Compositions of UHECRs is essential for the positional correlation study. Heavier components of UHECRs are more strongly deflected by magnetic fields and disturb the positional correlation. One of observables for the composition measurements is the depth of shower maximum, $X_{\\rm max}$, which can be measured by fluorescence detectors. Its average value, $\\left<X_{\\rm max}\\right>$, is sensitive to UHECR composition. High Resolution Fly's Eye(HiRes) reported that compositions of cosmic rays above $10^{19}$eV are dominated by protons as a result of $\\left<X_{\\rm max}\\right>$ measurement \\citep*{abbasi05}. Recent result by PAO is compatible to the HiRes result within systematic uncertainty\\citep*{unger07}. Another observable is muon density from extensive air shower of UHECR, observed by the ground array. Recent studies of the muon content indicate some fraction of heavier components in highest energy cosmic rays \\cite{engel07,glushkov07}. However, the interpretation of these two observables is dependent on hadronic interaction models, which include uncertainties at UHE energy region. On the other hand, the correlation with active galactic nuclei found by PAO indicates light composition. For the positional correlation study, proton component is a powerful tool. Thus, despite such uncertainties on UHECR compositions, we consider only protons in this study.  This paper is organized as follows: in section \\ref{gmf} we explain GMF models used in this study. In section \\ref{distortion}, we investigate the deflection of UHE protons with $10^{19.8}$eV in four GMF models introduced in section \\ref{gmf}. Dependence on parameters in the GMF models is discussed in section \\ref{dependence}. In section \\ref{conclusion}, we summarize results of the deflection and discuss effects of the deflection  to the correlation study. ",
        "conclusions": "\\label{conclusion}  In this study, we investigate the deflections of UHE extragalactic protons with $10^{19.8}$eV in the Galactic space by calculating their propagation in the four different GMF models using the backtracking method in order to discuss the correlation between the arrival directions of UHECRs and their source positions.  Protons arriving from the direction of the Galactic center are highly deflected by the dipole field. Source positions of such protons are quite different from their arrival directions. Thus, the Galactic center region is not suitable for the correlation study. If the correlation is found in the direction of the Galactic center, we can obtain strong constraint on the GMF around the Galactic center. However, in the BSS models, there are directions around the Galactic center that the arrival directions are almost unchanged. One of the directions is toward near Cen A, which is the nearest radio-loud AGN which is one of strong candidates of UHECR sources. If Cen A is confirmed as one of UHECR sources, the BSS models with dipole field and one of UHECR sources are confirmed.   On the other hand, the deflection angles of protons coming from the direction of the anti-Galactic center are relatively small for the BSS models. Those become smaller by adopting smaller pitch angle, $p=-8^{\\circ}$ since the field reversal exterior to the solar system is nearer. Thus, the anti-Galactic center region with $|b|<30^{\\circ}$ is best region for the correlation study. This region is in the northern hemisphere, where PAO cannot detect UHECRs.  We also investigate the distortion of the supergalactic plane by the GMF. The supergalactic plane near the Galactic center is highly distorted and the pattern of the distortion is dependent on the spiral field model. \\cite{ide01} showed that we can check whether UHECR source distribution is correlated with the supergalactic plane by $O(10^3)$ event detection all the sky This claim is based on UHECR propagation in extragalactic space. If the positional correlation between the arrival distribution and the supergalactic plane are found in the future, we can obtain information on the global structure of the GMF. On the other hand, the supergalactic plane near the anti-Galactic center is less distorted than that near the Galactic center although the arriving cosmic rays are deflected along the longitude line in the ASS models. In the northern hemisphere which contains the anti-Galactic center, the correlation between the arrival distribution and the supergalactic plane can be tested, which is weakly dependent on the spiral field model.  As discussed above, the northern hemisphere is more suitable for the positional correlation study. The direction of the anti-Galactic center does not bother us with uncertainty of magnetic field around the Galactic center. In our previous study, we predicted the positional correlation between UHECR arrival directions and positions of local sources at a scale of $2^{\\circ}\\times2^{\\circ}$ taking into account only EGMF\\citep*{takami07b}. Adding the results in this study, we can predict the spatial correlation at a scale of $\\sim 4^{\\circ}$ for the BSS-S model with $p=-8^{\\circ}$ and $\\sim 6^{\\circ}$ for the same with $p=-10^{\\circ}$ in the direction around the anti-Galactic center. If the ASS model is a real situation, regions where we can find the positional correlation in small angular scale are highly restrained.  In the northern hemisphere, Telescope Array is under construction and will start observation in the near future\\citep*{fukushima07}. PAO also projects its northern site with larger exposure than its southern site\\citep*{nitz07}. Extreme Universe Space Observatory(JEM-EUSO), which can detect UHECRs from International Space Station with extremely large exposure, is also projected\\citep*{ebisuzaki07}. The southern site of PAO will strongly contribute to our understanding on UHECR nature and sources. As next step for UHECR astronomy, These detectors, which can observe in the northern hemisphere, will be requested."
    },
    "0711/0711.0984_arXiv.txt": {
        "abstract": "The orbital eccentricity of a single planet in a stellar binary system with a sufficiently large mutual inclination angle is known to oscillate on a secular timescale through the Kozai mechanism.  We have investigated the effects of the Kozai mechanism on double-planet systems in binaries.  The evolutionary sequence of a pair of planets under the influence of a binary companion is fairly complex.  Various dynamical outcomes are seen in numerical simulations.  One interesting outcome is the rigid rotation of the planetary orbits in which the planetary orbital planes secularly precess in concert, while the orbital eccentricities oscillate synchronously.  In such cases the outer planet acts as a propagator of the perturbation from the binary companion to the inner planet and drives the inner planetary orbit to precess at a rate faster than what is predicted by the Kozai mechanism. ",
        "introduction": "} The first decade of radial-velocity and photometric surveys has revealed that planetary systems commonly form among stellar binary systems.  As of July 2007, at least 37 planetary systems ($\\sim 20\\%$ of all the known systems) are members of multiple-stellar systems \\citep{raghavan06,desidera07}.  The true stellar multiplicity among planetary systems may be higher than $20 \\%$ since the photometric surveys in search for stellar companions around currently known planetary systems are not yet complete.  Because most of these stellar companions are in wide orbits (typically $a = 10^2$\\,AU or greater), their orbital inclinations are not correlated with the invariable plane of the planetary systems, implying that many binary companions orbit around planetary systems in largely inclined orbits \\footnote{If the distribution of the mutual inclination angle is completely isotropic, the probability of the mutual inclination angle being greater than the critical Kozai angle ($39.2\\deg$) is $77\\%$}.  \\begin{figure}[!ht] \\centerline{\\includegraphics[scale=.4]{sample_kozai.eps}} \\caption{An example of secular  evolution of a planet orbit through the Kozai mechanism.  The amplitude of the eccentricity oscillation ($\\emax$) is determined by the initial relative inclination angle.  The argument of periastron typically alternates between libration and circulation during the Kozai cycles. \\label{sample_kozai}} \\end{figure}  It has been known in celestial mechanics that in a hierarchical triple-body system with  sufficiently large mutual inclination angle ($I > I_{\\rm crit} = 39.2^\\circ$) between the inner and outer binaries, the inner binary's orbital eccentricity and inclination undergo large-amplitude oscillations on a secular timescale through the ``Kozai mechanism'' \\citep{kozai62}.  When adopted to planetary systems, the inner binary orbit can be replaced by the osculating orbit of a planet around the primary star.  The Kozai mechanism has following unique properties (see also Figure~\\ref{sample_kozai}): (1) the orbital eccentricity of the planet oscillates with an amplitude that is  dependent (in the lowest order) only on the initial mutual inclination angle $I_0$; $\\emax = \\sqrt{1 - 5/3 \\cos{I_0}^2}$.  Other orbital parameters only affect the eccentricity oscillation period, \\begin{equation} P_{\\rm e} \\simeq P_{\\rm b}^2 / P_{\\rm pl}\\, (1 - e_b^2)^{3/2} ,  \\label{Pkoz} \\end{equation} \\citep{ford00} where $P_{\\rm pl}$ and $P_{\\rm b}$ are the orbital periods of the planet and the binary, respectively, and $e_{\\rm b}$ is the binary eccentricity. (2) The orbital inclination $I$ of the planet with respect to the binary plane also oscillates on the same timescale as the eccentricity oscillation, conserving the Kozai integral $H_{\\rm K} = \\sqrt{1 - e^2} \\cos{I}$.  (3) The argument of the periastron $\\omega$ alternates between the libration mode around $\\pm 90^\\circ$ and the circulation mode \\citep{fabrycky07}.   An episode of the Kozai mechanism in an extrasolar system may leave a directly observable signature in the orbital eccentricity of the planet.  Theoretical models show that several highest planetary eccentricities observed are best explained by the Kozai mechanism \\citep{holman97,wu03}.  Also, the median eccentricity of the planets in binaries ($e_{\\rm med,binary} = 0.25$) is higher than that of the planets around single stars ($e_{\\rm med, single} = 0.20$), implying that the planetary eccentricities are commonly excited by stellar companions.  The Kozai eccentricity oscillation results from the torque exerted by the binary companion averaged over the planetary orbit.  If there is another secular perturbation at play that precesses the planetary orbit in a shorter timescale than the Kozai oscillation timescale (Eq.~\\ref{Pkoz}), the averaged Kozai torque is reduced, and the Kozai eccentricity oscillation will be suppressed.  Extra precessions are typically caused by general relativistic effects, tidal torques, or mutual interactions between planets \\citep{wu03,takeda05}. ",
        "conclusions": "The rigid rotation of the planetary orbits and synchronized eccentricity oscillations are frequently observed among our simulations.  The criterion for the up-libration in Equation~\\ref{up} can be easily satisfied if the Kozai oscillation amplitude of the outer planet is sufficiently large.  However, the system also needs to maintain its stability for sustainable Kozai oscillations.  One necessary stability requirement is the Hill stability \\citep{gladman93}.  The Hill stability needs to be maintained at all times including when $e_2$ is near its maximum.   Another case of instability arises when the initial conditions are chosen such that the system undergoes down-libration (Eq.~\\ref{down}) is satisfied.  During down-librations the apsidal vectors are anti-aligned ($\\Delta \\omega \\approx 180\\deg$), and $e_2/e_1 \\sim 1$ or smaller.  If the down-libration is maintained until the outer planetary orbit embarks on the first Kozai cycle, then $e_1$ is excited to a value larger than $e_2$.  Such an over-excitation leads to a fairly chaotic behavior of the inner orbital eccentricity and will most likely result in  collision of the inner planet with the central star or ejection from the system.   The synchronized eccentricity oscillation mechanism offers a few theoretical implications.  First, in this mechanism the outer planet helps the angular momentum transfer from the companion star to the inner planet.  Propagation of eccentricity disturbance by a stellar passerby through a system of planets has been previously studied by \\citet{zakamska04}.  In the case of two planets in a binary, the cyclic propagation of the stellar perturbation through the outer planet can effectively excite the inner planet's eccentricity to oscillate.  This eccentricity propagation mechanism may imply the presence of an undetected secondary planet around known single-planet systems in binaries.  Especially for single planets in wide binaries with the Kozai period too large compared to the estimated age of the system, yet with non-negligible observed eccentricities that are difficult to be explained by other dynamical processes, there is a possibility that there is an undetected secondary planet propagating the eccentricity excitation from the distant companion.  Another dynamically relevant outcome is the subsequent evolution of the planets after the Kozai mechanism induces instability to the system .  Many systems in our simulations have suffered close encounters between planets that resulted in losing a planet from the system by ejection or collision.  The orbital radius and the  eccentricity of the survived planet can take wide ranges of values, as investigated by planet scattering experiments \\citep{juric07, chatterjee07}.  The survived planet may resume the Kozai mechanism, but now the planet's orbital eccentricity is likely to have been already excited to a significantly larger value after ejecting the other planet.  Since the fraction of the time a planetary orbit remains highly eccentric during the Kozai oscillations is strongly dependent on its initial eccentricity \\citep{holman97}, single-planet systems in binaries with an episode of early scattering has a much higher probability to be observed at large eccentricities.  For example, known single planets in binaries with very large orbital eccentricities (e.g., HD~80606~b, $e=0.93$ and HD~20782~b, $e=0.92$) strongly suggest the past occurrence of planet ejections in these systems, enabling the eccentricities of the survived planets to oscillate at systematically larger values."
    },
    "0711/0711.3455_arXiv.txt": {
        "abstract": "The space astrometry mission GAIA will construct a dense optical QSO-based celestial reference frame. For consistency between the optical and radio positions, it will be important to align the GAIA frame and the International Celestial Reference Frame (ICRF) with the highest accuracy. Currently, it is found that only 10\\% of the ICRF sources are suitable to establish this link, either because they are not bright enough at optical wavelengths or because they have significant extended radio emission which precludes reaching the highest astrometric accuracy. In order to improve the situation, we have initiated a VLBI survey dedicated to finding additional high-quality radio sources for aligning the two frames. The sample consists of about 450 sources, typically 20 times weaker than the current ICRF sources, which have been selected by cross-correlating optical and radio catalogues. This paper presents the observing strategy and includes preliminary results of observation of 224 of these sources with the European VLBI Network in June 2007. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.1666_arXiv.txt": {
        "abstract": "\\normalsize{ Numerical simulations of strongly compressible MHD corresponding to a stellar atmosphere with substantial gravity and near force-free magnetic fields show that the current sheet collapses (its width decreasing substantially). As a result, the reconnection rate increases dramatically.} ",
        "introduction": "\\label{sec0} It is well known that magnetic configurations are dissipating rather slowly in astrophysical conditions. Obviously, the Ohmic reconnection time, $$ t_O=\\frac{L^2}{\\eta}, $$ where $L$ is the characteristic magnetic scale, and $\\eta$ is resistivity, is huge, comparable with cosmological times.  The dissipation becomes much more efficient in the presence of current sheets. In particular, the Sweet-Parker (SP) flow in the vicinity of the current sheet results in faster reconnection rate, so that it can be estimated as \\begin{equation} t_{SP}=\\frac{t_O}{S^{1/2}}=t_AS^{1/2}, \\label{tSP} \\end{equation} where $S$ is Lundquist number, $$ S=\\frac{C_AL}{\\eta}, $$ $C_A$ is Alfv\\'en velocity, $$ C_A=\\frac{B}{\\sqrt{4\\pi\\rho}}, $$ and $t_A$ is Alfv\\'en time, \\begin{equation} t_A=\\frac{L}{C_A}, \\label{tA} \\end{equation} see \\cite{SP1}, \\cite{SP2}. As $S\\gg 1$ in astrophysical conditions, the SP mechanism (\\ref{tSP}) provides magnetic dissipation which is much faster than Ohmic. Still, this mechanism is rather slow compared with the time of stormy events like solar flares, or flares in the stars.  More efficient mechanism was suggested by Petschek \\cite{petschek}, with reconnection time, $$ t_P=t_A\\frac{8\\ln{S}}{\\pi}. $$ The  Alfv\\'en time (\\ref{tA}) is usually small enough, so that the characteristic time of stormy and violent events is comparable or even slightly larger than it. For this reason, the time $t_P$ is sufficient to explain these stormy events. It was pointed out however, that there should be some additional nontrivial conditions satisfied in order this mechanism to work. In particular, there are difficulties for this mechanism  in simple two-dimensional geometry, without additional effects like anomalous resistivity, Hall currents, localized resistivity etc., see e.g. \\cite{biskamp1}, \\cite{biskamp_book}, \\cite{drake}, \\cite{biskamp2}. ",
        "conclusions": ""
    },
    "0711/0711.4111_arXiv.txt": {
        "abstract": "Measuring the albedo of an extrasolar planet provides insights into its  atmospheric composition and its global thermal properties, including heat dissipation and weather patterns.  Such a measurement requires very precise photometry of a transiting system sampling fully many phases of the secondary eclipse.  Spacebased optical photometry of the transiting system HD 209458 from the MOST (Microvariablity and Oscillations of STars) satellite, spanning 14 and 44 days in 2004 and 2005 respectively, allows us to set a sensitive limit on the optical eclipse of the hot exosolar giant planet in this system.  Our best fit to the observations yields a flux ratio of the planet and star of 7 $\\pm$ 9 ppm (parts per million), which corresponds to a geometric albedo through the MOST bandpass (400-700 nm) of $A_g$ = 0.038 $\\pm$ 0.045.  This gives a 1$\\sigma$ upper limit of 0.08 for the geometric albedo and a 3$\\sigma$ upper limit of 0.17.  HD 209458b is significantly less reflective than Jupiter (for which $A_g$ would be about 0.5). This low geometric albedo rules out the presence of bright reflective clouds in this exoplanet's atmosphere.  We determine refined parameters for the star and exoplanet in the HD 209458 system based on a model fit to the MOST light curve. ",
        "introduction": "Among the roughly 250 planets discovered to date around nearby stars\\footnote{The Extrasolar Planets Encyclopedia: http://www.exoplanet.eu}, about two dozen are observed to transit.  Multicolour optical photometry of the transits provides accurate measurements of the orbital inclination, and hence the mass and radius of the planet with respect to the parent star \\citep[e.g.,][]{knu07}. Infrared measurements of the eclipses in thermal emission \\citep{dem05a,cha05}  give brightness temperatures at specific wavelengths to constrain the atmospheric properties.  Measurement of the optical eclipse of such a system provides a geometric albedo, which when combined with the thermal emission data, sensitively probes the exoplanetary atmosphere to test for high-altitude clouds (such as silicate condensates) and possibly even weather patterns \\citep{row06a}.  Models of the interior structure of giant exoplanets require appropriate boundary conditions -- planetary mass, radius and surface temperature -- which can be set by optical photometry of transiting hot extrasolar giant planets (EGPs).  HD 209458b, which orbits its host star at a distance of only 0.047 AU every 3.5 days, the thermal equilibrium temperature for the dayside photosphere can be estimated as \\begin{equation}\\label{eq:teq} T_{eq}=T_* (R_*/2a)^{1/2} [f(1-A_B)]^{1/4}, \\end{equation} where the star and planet are considered blackbodies, $T_*$ and $R_*$ are the effective temperature and radius of the host star, the planet at distance $a$ with a Bond albedo of $A_B$ and $f$ is a proxy for atmospheric thermal circulation.  The Bond albedo, $A_B$, is the fraction of total power incident on a body scattered back into space.  The geometric albedo $A_g$ is the ratio of brightness at zero phase angle compared to an idealized flat, fully reflecting and diffusively scattering disk.  From $A_g$, one can derive the ratio of the stellar and planetary fluxes \\begin{equation}\\label{eq:ag} \\frac{F_p}{F_*} = A_g \\left(\\frac{R_p}{a}\\right)^2, \\end{equation} where $R_p$ is the planetary radius and $a$, the separation between star and planet.  If the composition of HD 209458b is similar to Jupiter, then the surface temperature of the exoplanet is the dominant variable driving the atmospheric chemistry, and hence, the albedo. This means that the albedo should be a function of both the star-planet separation, and the luminosity of the host star \\citep{mar99}.  The relationship between albedo and surface temperature is a good tracer of chemical processes (such as the formation of molecules) in planetary atmospheres.  HD 209458 is a nearby (47 pc), bright ($V = 7.65$) Sun-like star (G0V) hosting a transiting hot EGP \\citep[e.g.,][]{cod02}.  Spitzer observations at 24 $\\mu m$ detected an eclipse at a depth of $0.26\\% \\pm 0.046\\%$ which gives a brightness temperature (1130 $\\pm$ 150 K) for the observed band \\citep{dem05a}.  Estimating the equilibrium temperature is highly model-dependent \\citep{sea05,for05,bar05,bur07b}.  Atmospheric models by \\citet{for05} predict an effective temperature of 1442 K, assuming 4$\\pi$ reradiation and solar metallicity.  The brightness of the star and the short orbital period ($\\sim$3.5 d) of the exoplanet made this planetary system an ideal candidate for photometry by the MOST (Microvariablity and Oscillations of STars) satellite.  The objectives of the observations were: (1) to measure accurately times of several consecutive transits to search for undetected Earth-mass planets in close orbits \\citep{mil07}; (2) to search for transits of nearly-Earth-radius planets at other periods in the nearly continuous data sets \\citep{cro07}; (3) and to measure or set a limit on the optical eclipse of the exoplanet HD 209458b.  The first results of our eclipse analysis \\citep{row06a}, based on a MOST trial run of 14 days in August 2004 sampling 4 consecutive transits and 4 secondary eclipses, set upper limits on the eclipse depth of 53 ppm (parts per million) and on the geometric albedo of 0.25.  We present here improved measurements of the eclipse based on the original photometry and a second 44-day MOST run in 2005, sampling an additional 12 consecutive transits and eclipses.  The combined datasets give  1$\\sigma$ upper limits on the eclipse depth of 16 ppm and on the geometric albedo of 0.08. ",
        "conclusions": "HD 209458b is less reflective than Solar System giant planets such as Jupiter.  Our measurements place a 1$\\sigma$ upper limit on $F_p/F_*$ of $1.6 \\times 10^{-5}$. The inferred low Bond albedo $<0.12$ rules out presence of highly reflective clouds in the atmosphere of HD 209458b and is also consistent with non-cloudy atmospheric models.  It also constrains the equilibrium temperature to be between $1400$ and $1650K$   The largest unknown in determining $T_{eq}$ is the efficiency with which the planet distributes heat from the dayside to nightside of the planet.  Further infrared measurements at different phases of the planet orbit will place stronger constrains on this parameter.  The MOST satellite re-observed the HD 209458 system in August - September 2007 for approximately 4 weeks.  A new observing mode has been added that allows images to be stacked onboard the satellite before downloading.  This dramatically increases the number of photon counts that can be recorded.  Instead of obtaining a 1.5-second exposure every 10 seconds one can obtain 1 exposure every 15 seconds composed of 15 stacked 1.5-second exposures.  This has decreased scatter in the data by a factor of $\\sim$3.  MOST will either detect the secondary eclipse or place a significance limit of a few percent on the albedo.  Analysis of these data is underway."
    },
    "0711/0711.1502_arXiv.txt": {
        "abstract": "We discuss implications of recent precision measurements for the ${}^{93}$Rh proton separation energy for the production of the lightest molybdenum isotopes in proton-rich type II supernova ejecta.  It has recently been shown that a novel neutrino-induced process makes these ejecta a promising site for the production of the light molybdenum isotopes and other ``p-nuclei'' with atomic mass near 100. The origin of these isotopes has long been uncertain.  A distinguishing feature of nucleosynthesis in neutrino-irradiated outflows is that the relative production of ${}^{92}$Mo and ${}^{94}$Mo is set by a competition governed by the proton separation energy of ${}^{93}$Rh. We use detailed nuclear network calculations and the recent experimental results for this proton separation energy to place constraints on the outflow characteristics that produce the lightest molybdenum isotopes in their solar proportions. It is found that for the conditions calculated in recent two-dimensional supernova simulations, and also for a large range of outflow characteristics around these conditions, the solar ratio of ${}^{92}$Mo to ${}^{94}$Mo cannot be achieved.  This suggests that either proton-rich winds from type II supernova do not exclusively produce both isotopes, or that these winds are qualitatively different than calculated in today's supernova models. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.1358_arXiv.txt": {
        "abstract": "We analyse brightness profiles for 143 early-type galaxies in the Virgo and Fornax Clusters, observed with the {\\it Advanced Camera for Surveys} on the {\\it Hubble Space Telescope}. S\\'{e}rsic models are found to provide accurate representations of the global profiles with a notable exception: the observed profiles deviate systematically inside a characteristic ``break\" radius of $R_b \\approx 0.02^{+0.025}_{-0.01}R_e$, where $R_e$ is the effective radius of the galaxy. The sense of the deviation is such that bright galaxies ($M_B \\lesssim -20$) typically show central light {\\it deficits} with respect to the inward extrapolation of the S\\'{e}rsic model, while the great majority of low- and intermediate-luminosity galaxies ($-19.5 \\lesssim M_B \\lesssim -15$) show central light {\\it excesses}; galaxies of intermediate luminosities ($-20 \\lesssim M_B \\lesssim -19.5$) are generally well fitted by S\\'{e}rsic models over all radii. We show that the slope, $\\gamma^{\\prime}$, of the central surface brightness profiles, when measured at fixed fractions of $R_e$, varies smoothly as a function of galaxy luminosity in a manner that depends sensitively on the choice of measurement radius. We find no evidence for a core/power-law dichotomy, and show that a recent claim of strong bimodality in $\\gamma^{\\prime}$ is likely  an artifact of the biased galaxy selection function used in that study.  To provide a more robust characterization of the inner regions of galaxies, we introduce a parameter, ${\\Delta}_{0.02} = \\log{({\\cal L}_g/{\\cal L}_s)}$ --- where ${\\cal L}_g$ and ${\\cal L}_s$ are the integrated luminosities inside $0.02R_e$ of the observed profile and of the inward extrapolation of the outer S\\'ersic model --- to describe the central luminosity deficit (${\\Delta}_{0.02} < 0 $) or excess (${\\Delta}_{0.02} > 0 $). We find that ${\\Delta}_{0.02}$ varies smoothly over the range of $\\approx$ 720~in luminosity spanned by the sample galaxies, with again no evidence for a dichotomy. We argue that the central light excesses in $M_B \\gtrsim -19$ galaxies may be the analogs of the dense central cores predicted by some numerical simulations to form via gas inflows. ",
        "introduction": "\\label{sec:intro}  Pioneering HST imaging studies of the centers of early-type galaxies suggested an apparently abrupt transition in central stellar density at $M_B \\sim -20.3$ mag --- the so-called ``core/power-law dichotomy\" (e.g., Ferrarese \\etal 1994; Lauer \\etal 1995). These findings prompted the widely held view that the bright (``core\") and faint (``power-law\") galaxies follow distinct evolutionary routes (e.g., Faber \\etal 1997). However, the evidence for such a dichotomy has lessened (although not entirely disappeared) following more recent studies that identified a population of galaxies with intermediate properties (Rest \\etal 2001; Ravindranath \\etal 2001). The slope, $\\gamma^{\\prime}$, of the central surface brightness profile --- usually parameterized as a ``Nuker\" law (essentially two power-laws that merge at a characteristic ``break\" radius; Lauer \\etal 1995) and measured at the angular distance corresponding to the instrumental resolution --- has traditionally been taken as a diagnostic of this behavior.  However, using ACS profiles for 100 early-type galaxies belonging to the Virgo Cluster, Ferrarese \\etal (2006a) showed that S\\'{e}rsic models provide more accurate parameterizations of the global brightness profiles than do Nuker models (see also Graham 2004; Ferrarese \\etal 2006c) and argued that the core/power-law dichotomy is an artifact introduced in part by the use of an inappropriate (i.e., power-law) parameterization of the outer profiles, combined with a tendency in previous work (which usually relied on HST brightness profiles of limited radial extent) to not properly account for the compact stellar nuclei found in low- and intermediate luminosity galaxies (e.g., Graham \\& Guzm\\'an 2003; Grant \\etal 2005; C\\^ot\\'e \\etal 2006).  In this paper, we use the best available imaging dataset --- in terms of depth, radial coverage, angular resolution, completeness and homogeneity --- to re-examine the central structure of early-type galaxies. Our analysis relies on HST/ACS imaging for 100 early-type members of the Virgo (previously discussed in Ferrarese \\etal 2006ab and C\\^ot\\'e \\etal 2006) and new HST/ACS imaging for 43 early-type members of the Fornax cluster (Jord\\'an \\etal 2007). Our principle finding is clear evidence for a continuous, systematic progression from central luminosity deficit ($M_B \\lesssim -20$) to excess ($M_B \\gtrsim -19$) within a characteristic radius, approximately equal to 2\\% of the galaxy effective radius. We find no evidence for a ``core/power-law\" dichotomy. ",
        "conclusions": "\\label{sec:discussion}  The central deficits in bright ellipticals have been traditionally interpreted as evidence for core evacuation by coalescing supermassive black hole (SBH) binaries (Ebisuzaki \\etal 1991; Faber \\etal 1997; Milosavljevi\\'c \\& Merritt 2001), although at least one recent study suggests that the dissipationless collapse of initially cold stellar distributions in existing dark matter haloes can also give rise to such features (Nipoti et al. 2006). The discovery of a systematic transition from central deficit to excess should provide an important new clue to the formation of both high- and low-luminosity early-type galaxies and the connection to SBHs. It is worth noting that the mass fraction contributed by the central nuclei in low- and intermediate-luminosity galaxies is indistinguishable ($\\sim 0.2$\\%) from that of SBHs in bright galaxies, suggestive of a shared formation path (C\\^ot\\'e \\etal 2006, Ferrarese \\etal 2006ab; Wehner \\& Harris 2006; Rossa \\etal 2006). On the other hand, McLaughlin \\& King (2006) point out that feedback by stellar winds and supernovae may also be capable of reproducing the observed scaling relations without a direct connection between the nuclei and SBHs. And, as Ferrarese \\etal (2006b) note,  there are at least some galaxies (such as M32) that contain both a central excess (Worthey 2004) and a SBH (Verolme \\etal 2002).  A lengthy discussion of the selection effects that prevented the ubiquitous nature of these nuclei from being fully appreciated with ground-based imaging was presented in C\\^ot\\'e \\etal (2006). Not surprisingly, the foremost advangtage of HST/ACS in this regard is its exceptional angular resolution. One might think, then, that the detection of nuclei in low- and intermediate-luminosity would have been straightforward in early high-resolution HST imaging, but this often proved not to be the case. For instance, the left panel of Figure~\\ref{fig:v1440} compares our ACS profile for VCC1440 --- a low-luminosity elliptical previously classified as a ``power-law\" galaxy --- to the WFPC1 profile of Lauer \\etal (1995, 2007).\\footnote{Lauer \\etal (2007) report $\\gamma^{\\prime} = 0.89$ for this galaxy. We measure $\\langle\\gamma^{\\prime}\\rangle = 1.08, 1.40, 0.79$ and $1.31$ at $R/R_e = 0.005, 0.01, 0.05$ and $0.3$, respectively. These radii are indicated by the tickmarks in Figure~\\ref{fig:v1440}.} In this case, the limited radial extent of the WFPC1 profile, and the adopted parameterization of the galaxy profile as a power-law, allowed the full extent of the central excess to go unrecognized, although its signature is evident in the {\\tt S}-shaped residuals of the fitted Nuker model (see also Ferrarese \\etal 2006c).  In hindsight, the realization that compact stellar mass concentrations reside at the centers of most low- and intermediate galaxies should perhaps have come as no surprise. Many theoretical studies have shown that such mass concentrations may be a generic outcome of the galaxy formation process. Proposed formation routes include repeated mergers of star clusters drawn to the galaxy center via dynamical friction, the growth of $\\rho \\propto r^{-7/4}$ density cusps due to two-body relaxation of stars orbiting a central SBH (Bahcall \\& Wolf 1976) and/or central star formation driven by gas inflows --- likely modulated by stellar evolution, mergers, IGM confinement, and feedback from stellar winds, supernovae or pre-existing AGN (e.g., Barnes \\& Hernquist 1991; Mihos \\& Hernquist 1994; McLaughlin \\etal 2006; Li \\etal 2007; Hopkins \\etal 2007).  A comparison of the properties (e.g., sizes, luminosity functions, mass fractions, colors) of nuclei and star clusters in ACSVCS program galaxies argues against the first scenario as the {\\it dominant} formation mechanism (e.g., see \\S5.2 of C\\^ot\\'e \\etal 2006). Likewise, Bahcall-Wolf ``cusps\" generated by SBHs do not seem viable since most of the nuclei are {\\it resolved} and thus much more extended than the predicted cusps (e.g., Merritt \\& Szell 2006). On the other hand, gas inflow models --- either internally or externally modulated --- seem more attractive as they have long predicted the formation of ``dense stellar cores\" in early-type galaxies. Consider the center and right panels of Figure~\\ref{fig:v1440} which compare the brightness profile of VCC1440 (the same galaxy shown in the left panel) to the simulations of Mihos \\& Hernquist (1994). These authors showed that compact central starbursts are found in the remnants of disk/halo mergers and went on to argue that this central component would lead to ``a break in the mass profile at small radii ($\\sim 2\\%R_e$)\". As Figure~\\ref{fig:v1440} and the discussion in \\S\\ref{sec:radii} show, this may turn out to be a remarkably prescient prediction. The fact that the deviations from the fitted S\\'ersic models occur on roughly the same ($\\sim 2\\%R_e$) scale for both the central excesses {\\it and} central deficits (which are usually attributed to binary SBH evolution; see above) presents something of a puzzle since it is unclear why these two unrelated physical mechanisms would give rise to deviations on the same (fractional) scale.  In any case, it will be important to test the gas inflow scenario envisaged by Mihos \\& Hernquist (1994) by measuring star formation and chemical enrichment histories for the central regions of bright (deficit) and faint (excess) galaxies alike. Both the measured color profiles (Ferrarese \\etal 2006a) and the model-dependent integrated colors (C\\^ot\\'e \\etal 2006) of the nuclei and galaxies in the ACSVCS suggests that stellar population differences between the two components may indeed exist, with the nuclei in the faintest systems generally being bluer than the underlying galaxy. On the other hand, at least some intermediate-luminosity galaxies show evidence for surprisingly red central components (e.g., $g-z \\gtrsim 1.5$).  Finally, we point out that the open symbols in {\\it panel (e)} of Figure~\\ref{fig:delta} denote galaxies with ``dE/dIrr transition\" morphologies: i.e., dust, young stellar clusters and/or evidence of young stellar populations from blue integrated colors (Ferrarese \\etal 2006a; C\\^ot\\'e \\etal 2006). Such galaxies typically lack conspicuous central nuclei at their photocenters. The absence of a central excess may suggest that the process of nucleus-building --- perhaps through delayed gas inflows and subsequent star formation --- has yet to occur in these apparently young galaxies."
    },
    "0711/0711.4782_arXiv.txt": {
        "abstract": "We present a theoretical model for Type Ib supernova (SN)~2006jc. We calculate the evolution of the progenitor star, hydrodynamics and nucleosynthesis of the SN explosion, and the SN bolometric light curve (LC). The synthetic bolometric LC is compared with the observed bolometric LC constructed by integrating the UV, optical, near-infrared (NIR), and mid-infrared (MIR) fluxes. The progenitor is assumed to be as massive as $40\\Msun$ on the zero-age main-sequence. The star undergoes extensive mass loss to reduce its mass down to as small as $6.9\\Msun$, thus becoming a WCO Wolf-Rayet star. The WCO star model has a thick carbon-rich layer, in which amorphous carbon grains can be formed. This could explain the NIR brightening and the dust feature seen in the MIR spectrum. We suggest that the progenitor of SN~2006jc is a WCO Wolf-Rayet star having undergone strong mass loss and such massive stars are the important sites of dust formation. We derive the parameters of the explosion model in order to reproduce the bolometric LC of SN~2006jc by the radioactive decays: the ejecta mass $4.9\\Msun$, hypernova-like explosion energy $10^{52}$ ergs, and ejected \\Nifs\\ mass $0.22\\Msun$. We also calculate the circumstellar interaction and find that a CSM with a flat density structure is required to reproduce the X-ray LC of SN~2006jc. This suggests a drastic change of the mass-loss rate and/or the wind velocity that is consistent with the past luminous blue variable (LBV)-like event. ",
        "introduction": "\\label{sec:introduction}  On 9th October 2006, \\cite{nak06} reported K. Itagaki's discovery of a possible supernova (SN) in UGC~4904. Although the SN was discovered after the peak, an upper limit of the $R$ magnitude ($M_R>-12.2$) was obtained at $\\sim$20 days before the discovery \\citep{pas07}. Interestingly, \\cite{nak06} also reported that an optical transient had appeared in 2004 close to the position of SN~2006jc. The transient was as faint as $M_{\\rm R}\\sim-14$ and its duration was as short as $\\sim10$ days. Since the event was faint and short-lived, they speculated that the transient was a luminous blue variable (LBV)-like event. The spatial coincidence between the LBV-like event and SN~2006jc is confirmed by \\cite{pas07}. Because of such an intriguing association with the LBV-like event, many groups performed follow-up observations of SN~2006jc in various wavebands: X-ray, ultra violet (UV), optical, infrared (IR), and radio.  Spectroscopic observations showed many broad features and strong narrow \\ion{He}{1} emission lines. According to the He detection, SN~2006jc was classified as Type Ib \\citep{cro06,fes06a,ben06,mod06a,mod06b}. However, strange spectral features and their evolutions were reported. A bright blue continuum was prominent in the optical spectrum at early epochs \\citep{fol07,pas07,smi07}. Such a bright blue continuum had also been observed in Type II SN~1988Z \\citep{tur93}, but the origin of this feature is still unclear. As the blue continuum declined, the red wing brightened and the optical spectra showed ``U''-like shapes \\citep{smi07,kaw07b}. This is a distinguishing feature of SN~2006jc in contrast to the spectra of usual SNe that have a peak in optical bands.  Photometric observations in optical and IR bands were performed continuously. The optical light curve (LC) showed a rapid decline from 50 days after the discovery, as in the case of SN~1999cq \\citep{mat00}. At the same epoch, near infrared (NIR) emissions brightened \\citep{ark06,dic07}. The NIR brightness increased from $\\sim40$ days to $\\sim70$ days after the discovery and then declined \\citep{dic07}. The epoch of the NIR brightening corresponds to that of the development of the red wing in the optical spectra \\citep{smi07}.  The NIR brightening, as well as the fact that the redder side of the He emission profile declined faster than the bluer side, has been interpreted as an evidence of an ongoing dust formation \\citep{smi07}. Additionally, on 29th April 2007 (200 days after the discovery), the {\\sl AKARI} satellite performed NIR and mid-infrared (MIR) photometric and spectroscopic observations \\citep{sak07} and the {\\sl MAGNUM} telescope obtained the NIR photometries \\citep{min07}. They report the formation of amorphous carbon dust: another piece of evidences of the dust formation.  \\begin{figure} \\epsscale{1.} \\plotone{f1.eps} \\caption{Multicolor available observations of SN~2006jc. (a) X-ray and UV luminosities obtained with the {\\sl Swift} and {\\sl Chandra} satellites ({\\it black}: X-ray, {\\it red}: UVW2-band, {\\it green}: UVM2-band, {\\it blue}: UVW1-band, \\citealt{imm07}). X-ray LC is shown in a unit of \\ergs\\ (right axis). (b) Optical luminosities obtained with the {\\sl HCT} telescope ({\\it filled}, \\citealt{anu07}) and the {\\sl SUBARU} telescope ({\\it open}, \\citealt{kaw07a,kaw07b}). The upper limit is taken from \\cite{pas07}. The color represents the wavelengths (U-band: {\\it red}, B-band: {\\it green}, V-band: {\\it blue}, R-band: {\\it magenta}, I-band: {\\it cyan}). (c) IR luminosities obtained with the {\\sl AZT-24} telescope ({\\it filled circles}, \\citealt{ark06,dic07}), the {\\sl MAGNUM} telescope ({\\it open circles}, \\citealt{min07}), and the {\\sl AKARI} satellite ({\\it squares}, \\citealt{sak07}). The color of circles represents the wavelengths (J-band: {\\it red}, H-band: {\\it green}, and K-band: {\\it blue}). The contributions to the IR luminosities from the hot dust ({\\it filled square}) and with the hot and warm dust ({\\it open square}) are shown in a unit of \\ergs\\ (right axis). \\label{fig:LCobsall}} \\end{figure}  X-ray and UV emissions have also been observed by the {\\sl Swift} and {\\sl Chandra} satellites \\citep{bro06,imm06,imm07,hol07}. X-ray observations were performed at seven epochs and showed a brightening from $\\sim20$ days to $\\sim100$ days after the discovery \\citep{bro06,imm06,imm07}. The X-ray detection suggests an interaction between the SN ejecta and the circumstellar matter (CSM). On the contrary, the radio emission was not detected by Very Large Array (VLA) \\citep{sod06}.  We present a SN explosion model of a Wolf-Rayet star that explains the bolometric and X-ray LCs. Hydrodynamics, nucleosynthesis, and LC synthesis calculations are performed assuming the spherical symmetry. In this study, we assume the explosion date of SN~2006jc to be 15 days before the discovery ($t=0$) and the energy source of the light to be the \\Nifs-\\Cofs\\ decay.  The paper is organized as follows: in \\S~\\ref{sec:bol}, we describe how we derive the bolometric LC from observations in the various wavebands, in \\S~\\ref{sec:preSN}, we briefly discuss the presupernova evolutionary properties of the progenitor star; in \\S~\\ref{sec:hyd}, hydrodynamical and nucleosynthesis calculations are described; in \\S~\\ref{sec:LC}, LC synthesis calculations are presented; in \\S~\\ref{sec:CSM}, we calculate the X-ray emission due to the ejecta-CSM interaction; in \\S~\\ref{sec:conclude} and \\S~\\ref{sec:discuss}, conclusions and discussion are presented. ",
        "conclusions": "\\label{sec:conclude}  We present a theoretical model for SN~2006jc whose properties are summarized as follows.  (1) {\\bf WCO progenitor and Dust Formation}: The progenitor is a WCO Wolf-Rayet star whose total mass has been reduced from $\\Mms=40\\Msun$ to as small as $\\Mpre=6.9\\Msun$. The WCO star model has a thick C-rich envelope and CSM. This is consistent with the formation of amorphous carbon grains in the SN ejecta and the CSM suggested by {\\sl AKARI} observations \\citep{sak07}. \\cite{noz07} have calculated dust formation in the WCO star explosion model and shown that carbon dust is formed in the C-rich layer at $t\\sim50$ days. This is much earlier than the dust formation after $t\\sim1$ yr in Type II SNe (SNe~II), because of the much smaller $\\Mej$ in the WCO star than SNe~II. According to the models in \\cite{lc06}, the stars with $\\Mms>40\\Msun$ typically become WCO stars to form a thick C-rich layer and CSM. This limiting mass, however, is still uncertain and strongly depends on many details of the stellar evolution. The early dust formation in the SN ejecta and the CSM suggests that the progenitor of SN~2006jc is a massive star becoming a WCO Wolf-Rayet star.  (2) {\\bf Explosion and Bolometric Light Curve}: The multicolor LCs of SN~2006jc show peculiar evolutions, \\eg a rapid decline of the optical LC and brightening of the IR LC. These can be interpreted as an ongoing dust formation. Assuming the absorbed optical light is re-emitted in the IR band, the bolometric LC is constructed as a summation of $L_{\\rm UV}$ ($=0.8L_{\\rm opt}$), $L_{\\rm opt}$, and $L_{\\rm IR}$. By calculating the hydrodynamics, nucleosynthesis, and the bolometric LC for the SN explosion with the various explosion energies, $E_{51}=1$, 5, 10, and 20, we find that the hypernova-like SN explosion model with $\\Mej=4.9\\Msun$, $E_{51}=10$, and $\\Mni=0.22\\Msun$ best reproduces the bolometric LC of SN~2006jc with the radioactive decays. Also, the temperature evolution of the carbon dust heated by the \\Nifs-\\Cofs\\ decay reasonably well explains the IR observations for $t\\sim50-220$ days \\citep{noz07}.  (3) {\\bf CSM Interaction and X-ray Light Curve}: Applying the model with $\\Mej=4.9\\Msun$ and $E_{51}=10$, we calculate the ejecta-CSM interaction and the resultant X-ray LC. We derive the CSM density structure to reproduce the X-ray LC of SN~2006jc as $\\rho =2.75\\times10^{-19}$ g cm$^{-3}$ for $r < 2.2\\times10^{16}$ cm and $2.75\\times10^{-19} (r / 2.2\\times10^{16} {\\rm cm})^{-6}$ g cm$^{-3}$ for $r > 2.2\\times10^{16}$ cm. The flat density distribution in the inner CSM indicates a drastic change of the mass-loss rate and/or the wind velocity that is consistent with the LBV-like event two years before the explosion.  \\vspace{1.5cm}"
    },
    "0711/0711.3993_arXiv.txt": {
        "abstract": "A test case comparison is presented for different dust cloud model approaches applied in brown dwarfs and giant gas planets.  We aim to achieve more transparency in evaluating the uncertainty inherent to theoretical modelling. We show in how far model results for characteristic dust quantities vary due to different assumptions. We also demonstrate differences in the spectral energy distributions resulting from our individual cloud modelling in 1D substellar atmosphere simulations. ",
        "introduction": "Spectra are windows through which we access the physics and the chemistry of substellar atmospheres down into the planetary regime. Consequently, model atmospheres are needed to produce synthetic spectra. The complexity of such simulations necessitates the use of approximations and/or analytic simplifications to make them numerically feasible. The input data needed can be frequently incomplete or of limited accuracy, translating into uncertainties in the model results.  The formation of clouds in their atmospheres is one of the most fundamental challenges in studies of planets and brown dwarfs.  Element depletion and gravitational settling are only two feedback mechanisms arising.  We therefore have set out to compare our approaches of cloud formation in substellar atmosphere simulations to achieve more transparency in evaluating the uncertainty inherent to theoretical modelling. We show in how far model results for characteristic dust quantities vary due to different assumptions. We also demonstrate differences in the spectral energy distributions resulting from our individual cloud modelling in 1D substellar atmosphere simulations. ",
        "conclusions": "Our test case studies show that the results of our individual cloud models are comparable regarding general feature like the location of cloud base and the maximum dust-to-gas ratio. However, the cloud-model results differ if studied in more detail. It is therefore no surprise that the spectral energy distributions produced from our different 1D atmosphere simulations for a given parameter combination (T$_{\\rm eff}$, $\\log$g, [M/H]) differ in almost the entire wavelength range. It remains to quantify these differences in order to provide a general range of applicability for e.g. stellar parameter determinations inside and beyond the substellar regime.   \\medskip \\noindent {\\small\\bf Acknowledgement:}\\\\ We thank all the participants of the workshop {\\sf From Brown Dwarfs to Planets: Chemistry and Cloud formation} which was supported by the Lorentz Center of the University Leiden, Nederlandse Organisatie voor Wetenschappelijk Onderzoek, The Netherlands research School for Astronomy, the Scottish University Physics Alliance, and European Space Agency. ChH acknowledges an IAU travel grant."
    },
    "0711/0711.4057_arXiv.txt": {
        "abstract": "% In young star clusters, the density can be high enough and the velocity dispersion low enough for stars to collide and merge with a significant probability. This has been suggested as a possible way to build up the high-mass portion of the stellar mass function and as a mechanism leading to the formation of one or two very massive stars ($M_\\ast > 150\\,{\\rm M}_\\odot$) through a collisional runaway.  I quickly review the standard theory of stellar collisions, covering both the stellar dynamics of dense clusters and the hydrodynamics of encounters between stars. The conditions for collisions to take place at a significant rate are relatively well understood for idealised spherical cluster models without initial mass segregation, devoid of gas and composed of main-sequence (MS) stars. In this simplified situation, 2-body relaxation drives core collapse through mass segregation and a collisional phase ensues if the core collapse time is shorter than the MS lifetime of the most massive stars initially present. The outcome of this phase is still highly uncertain. A more realistic situation is that of a cluster still containing large amounts of interstellar gas from which stars are accreting. As stellar masses increase, the central regions of the cluster contracts. This little-explored mechanism can potentially lead to very high stellar densities but it is likely that, except for very rich systems, the contraction is halted by few-body interactions before collisions set in. A complete picture, combining both scenarios, will need to address many uncertainties, including the role of cluster sub-structure, the dynamical effect of interstellar gas, non-MS stars and the structure and evolution of merged stars. ",
        "introduction": "We first consider the timescale for direct collisions between stars. We assume two populations of stars, with stellar masses $m_1$ and $m_2$, stellar radii $r_1$ and $r_2$, and Maxwellian velocity distributions of (1D) dispersions $\\sigma_1$ and $\\sigma_2$. The mean time for a star of type 1 to collide with a star of type 2 is \\begin{equation} t_{\\rm coll} = \\left\\{\\sqrt{8\\pi}n_2\\sigma_{\\rm rel}(r_1+r_2)^2 \\left[1+\\frac{G(m_1+m_2)}{\\sigma_{\\rm rel}^2(r_1+r_2)}\\right]\\right\\}^{-1}, \\label{eq.tcoll} \\end{equation} with $\\sigma_{\\rm rel}^2=\\sigma_1^2+\\sigma_2^2$ \\citep{BT87_MFreitag}. The second term in the square brackets originates in the mutual gravitational attraction of the stars. This gravitational focusing dominates over the purely geometrical cross section (represented by the term $1$ in brackets) unless the velocity dispersion is similar or higher than the escape velocity from the surface of a star, i.e.\\ $500-1000\\,{\\rm km}\\,{\\rm s}^{-1}$. Such relative velocities are presumably only reached in the vicinity of massive black holes \\citep[e.g.,][]{FB02b_MFreitag,FB05_MFreitag}. For other situations, the collision time can be written \\begin{equation} t_{\\rm coll} \\simeq 5\\,{\\rm Gyr}\\,\\frac{10^6 {\\rm pc}^{-3}}{n_\\ast} \\frac{\\sigma_{\\rm rel}}{10\\,{\\rm km}\\,{\\rm s}^{-1}} \\frac{2\\,{\\rm R}_\\odot}{r_1+r_2} \\frac{2\\,{\\rm M}_\\odot}{m_1+m_2}. \\label{eq.tcoll2} \\end{equation} In Fig.~\\ref{fig.DensColl}, I show (with dashed lines) the stellar densities required for a massive star to experience on average one collision during the MS. Typical values are in excess of $10^7\\,{\\rm pc}^{-3}$ or $10^6\\,{\\rm pc}^{-3}$ for a star of $10\\,{\\rm M}_\\odot$ or $120\\,{\\rm M}_\\odot$, respectively.  Stars passing each other within a few stellar radii at low velocities can dissipate enough orbital energy in tides to form a bound binary \\citep*[e.g.,][]{FPR75_MFreitag,KL99_MFreitag}. Such a ``tidal binary'', being born very tight, is likely to merge quickly, or, at least, to enter a common-envelope phase \\citep{DD06_MFreitag}. The merger can also be induced by the perturbation of a passing star \\citep{BB05b_MFreitag}. In any case, because tidal dissipation decreases steeply with increasing closest-approach distance, formation and merging of tidal binary can reduce the effective value of the collision time by a factor $\\sim 3-5$ at most.   A more promising way to get interesting collision rates is to consider binary interactions. The timescale over which a binary with separation $a$ and mass $m_{\\rm bin}$ suffers from a close interaction with a field star is simply \\begin{equation} t_{\\rm bin,inter} \\simeq 30\\,{\\rm Myr}\\,\\frac{10^6 {\\rm pc}^{-3}}{n_\\ast} \\frac{\\sigma_{\\rm rel}}{10\\,{\\rm km}\\,{\\rm s}^{-1}} \\frac{1\\,{\\rm AU}}{a} \\frac{3\\,{\\rm M}_\\odot}{m_{\\rm bin}+m}, \\label{eq.tbininter} \\end{equation} where $m$ is the individual mass of the field stars and $n_\\ast$ their number density. For a sufficiently hard binary, i.e.\\ when the total orbital energy of the binary and impactor system is negative, most close interactions are ``resonant''. This means that the three stars orbit each other in a chaotic fashion for several dynamical times within a volume of order $a^3$ until this meta-stable system decays into a binary and single star if stars were point masses\\footnote{Binary-binary interactions occur in a similar fashion.} \\citep[e.g.,][]{HH03_MFreitag}. There is in fact a high probability that two of the stars collide and merge during a resonant interaction. For instance, \\citet{FregeauEtAl04_MFreitag} showed that, in the hard regime, the cross section for a collision to occur during an encounter of a circular, $a=1\\,$AU binary consisting of two $1\\,{\\rm M}_\\odot$ MS stars with a similar star is about $\\frac{3}{2}\\pi a G{\\rm M}_\\odot v_{\\rm rel}^{-2}$, corresponding to a hard sphere of radius $a/4$. If massive binaries are born with eccentricities $\\approx 0.9$, as a result of accretion of interstellar gas, more distant, non-resonant, interactions with passing stars can suffice to cause a merger, as suggested by \\citet{BB05b_MFreitag}.  At relative velocities below about $5\\,{\\rm km}\\,{\\rm s}^{-1}$, encounters involving two massive stars, one or both of them surrounded by a massive disc, can result into a bound binary if the closest approach distance $d$ is smaller than a few tenth of the disc radius, which corresponds typically to $d\\approx 10-100\\,$AU \\citep{MB07a_MFreitag}. With $d$ substituted for $a$, Eq.~\\ref{eq.tbininter} suggests that this mechanism might be responsible for the high proportion of massive stars with a companion of similar mass in clusters such as the ONC, despite the short life time of massive discs \\citep{MB07b_MFreitag,PO07b_MFreitag}. ",
        "conclusions": ""
    },
    "0711/0711.2392_arXiv.txt": {
        "abstract": "Two methods are used to observationally determine the solar radius: One is the observation of the intensity profile at the limb, the other one uses f-mode frequencies to derive a 'seismic' solar radius which is then corrected to optical depth unity. The two methods are inconsistent and lead to a difference in the solar radius of $\\sim$0.3\\,Mm. Because of the geometrical extention of the solar photosphere and the increased path lengths of tangential rays the Sun appears to be larger to an observer who measures the extent of the solar disk. Based on radiative transfer calculations we show that this discrepancy can be explained by the difference between the height at disk center where $\\tau_{\\mathrm{5000}}=1$ ($\\tau_{\\mathrm{Ross}}=2/3$) and the inflection point of the intensity profile on the limb. We calculate the intensity profile of the limb for the MDI continuum and the continuum at 5000\\,{\\AA} for two atmosphere structures and compare the position of the inflection points with the radius at $\\tau_{\\mathrm{5000}}=1$ ($\\tau_{\\mathrm{Ross}}=2/3$). The calculated difference between the 'seismic' radius and the inflection point is $0.347\\pm 0.06$\\,Mm with respect to $\\tau_{\\mathrm{5000}}=1$ and $0.333\\pm 0.08$\\,Mm with respect to $\\tau_{\\mathrm{Ross}}=2/3$. We conclude that the standard solar radius in evolutionary models has to be lowered by $0.333\\pm 0.08$\\,Mm and is 695.66~Mm. Furthermore, this correction reconciles inflection point measurements and the 'seismic' radii within the uncertainty. ",
        "introduction": "The solar radius is a key parameter for understanding the Sun's evolution and many astrophysical applications. Being our closest star, the Sun can be considered as our laboratory, allowing us to determine its radius with high precision. Over the past 25 years various groups were dedicated to determine the solar radius by applying basically two approaches. First, there is the direct measurement of the photospheric radius. In earlier times its visual determination was carried out by an observer. Later, employing CCD-imaging devices, the visual radius determination was replaced by deriving the inflection point to the intensity profile of the limb. The error of both methods has been investigated in detail by \\cite{Laclare1996} who derive a solar radius of 959''.60, which agrees well with 959''.63 $\\pm$ 0''.10 or 695.99 $\\pm$0.07\\,Mm \\citep{Allen1976}, commonly used for calibrating solar evolutionary models \\citep{Brown1998,Turck2000}. Second, \\cite{Schou1997,Antia1998} determine the helioseismic radius by means of f-mode frequencies obtained from the Michelson Doppler Imager \\citep[MDI]{Scherrer1995} on board the ESA/NASA spacecraft SOHO. Due to the high quality of the MDI observations the helioseismic determination is very precise and leads to a radius that is approximately 0.3~Mm smaller than the standard value by \\cite{Allen1976}. \\cite{Schou1997} conclude that the standard solar radius has to be decreased by this value. Finally, \\cite{Brown1998} determine the solar radius from observations obtained with the High Altitude Observatory's Solar Diameter Monitor \\citep[HAO SDM]{Brown1982} by applying a model of the data reduction procedure. The authors point out that there are significant uncertainties associated with the currently used radius value which are related to the definition of the solar limb adopted in the radius determinations and the reduction of the measured value to the photosphere. The authors indicate a height of about 0.3\\,Mm which 'perhaps' explains the radius correction inferred from the f-mode frequencies. In the present paper we determine the difference between the height at $\\tau_{\\mathrm{5000}}=1$ at disk center and the distance of the inflection point from radiative transfer calculations for two wavelengths and atmosphere structures. ",
        "conclusions": "From radiative transfer calculations with the spherical code COSI we determine a distance of 0.333$\\pm$0.008~Mm ($0.347\\pm 0.06$\\,Mm) between the height where $\\tau_{\\mathrm{Ross}}=2/3$ ($\\tau_{\\mathrm{5000}}=1$) at disk center and the position of the inflection point of the 5000\\,{\\AA} intensity profile at the limb. This correction explains the differences between the f-mode frequencies derived from models calculations and observations. Therefore, we conclude that standard solar radius currently used in evolutionary models has to be corrected by 0.333~Mm and is 695.66~Mm.  Furthermore, the correction of 0.333~Mm reconciles the radius values derived from inflection point measurements and the 'seismic' radius by 0.09$\\pm$0.17~Mm. Due to the increasing precision of the instruments it becomes necessary that the community agrees on a generally binding definition of the solar radius and that values derived via different techniques are to be corrected accordingly. We propose that the solar radius refers to $R_\\odot=R(\\tau_{\\mathrm{Ross}}=2/3)$. Thus, we suggest that the value of the solar radius determined from inflexion point measurements has to be lowered by 0.333~Mm to be comparable with 'seismic' radii. A detailed study of the position of the wavelength and solar cycle dependent inflection point will be crucial for upcoming observations with the SODISM instrument on board PICARD \\citep{ThuiDewSchm2006ASpR}."
    },
    "0711/0711.0504_arXiv.txt": {
        "abstract": "In this paper it is given a brief review of the current limits on the magnitude of CPT and Lorentz Invariance violations, currently predicted in connection with quantum gravity and string/M-theory, that can be derived from astrophysical and cosmological data. Even if not completely unambiguous, these observational tests of fundamental physics are complementary to the ones obtained by accelerator experiments and by ground or space based direct experiments, because potentially can access very high energies and large distances. ",
        "introduction": "\\label{sect:intro} \\cite{Einstein1905AnP322891E} introduced the postulate of the constancy of the velocity of light in empty space, justifying it on the bases of the negative result of the Michelson-Morley experiment \\citep{michelson1887}. Since then the covariance of physical laws under the group of Lorentz transformations, usually nicknamed the ``Lorentz invariance''(LI), has been proven to hold locally within a very high accuracy in tests done on the Earth or nearby space. A modern version of the Michelson-Morley experiment has been performed comparing the frequency of solid state resonant cavities, in two orthogonal mode while in rotation respect to the Cosmic Microwave Background Radiation (CMBR) reference frame. Nevertheless the accuracy of the measurements is of the order of $10^{-7}$ for resonators in motion with the Earth \\citep{SaathoffPhysRevLett.91.190403} and of $10^{-10}$ for rapidly spinning ones \\citep{Stanwix:2006jb}. As we will discuss briefly in the remaining part of this introduction, this level of accuracy cannot exclude that Special Relativity Theory could be only an approximation of the physical reality, that could be violated by tiny perturbations introduced by conceivable mechanism.\\par Around the middle of the last century, the development of a fully Relativistic Quantum Field Theory (RFQT) has pointed out  that there is a strong logical connection between the Lorentz invariance and the matter and antimatter duality. This is not unexpected since the very existence of antimatter is a consequence of the relativistically covariant Dirac equation \\citep{Dirac1933nl}. Following an early indication by \\cite{SchwingerQFT-I}, \\cite{Pauli1955book} introduced the idea that for every process occurring in nature there is an allowed one, occurring with the same probability, in which each particle is replaced with its corresponding antiparticle, having reversed spin and following a trajectory that is the reflection under space and time inversion of the original ones. This is now called the ``Schwinger-Pauli-Luders CPT theorem'' because \\cite{Luders1957} showed that assuming \\begin{enumerate} \\item Lorentz Invariance; \\item Locality; \\item Unitarity; \\item Correct connection between spin and statistics \\end{enumerate} it is possible to prove that the successive application of charge conjugation C, parity reflection P and time reversal T in any order leave the Hamiltonian invariant. As a complement \\citet{Greenberg2002prl} has also proved the \\quotes{inverse-CPT theorem}, showing that if under any circumstances CPT invariance is violated, then also Lorentz invariance will be.\\par In certain sense the CPT invariance is more fundamental than Lorentz invariance, because, as we will show later, it is possible to conceive modifications of the Hamiltonian of the fundamental interactions which violates Special Relativity, but are invariant under CPT transformations, acting equally on particles and antiparticles.\\par The CPT theorem points out what could be the physical origin of Lorentz violations. In particular the second hypothesis that leads to it is \\quotes{locality}, which demands that space-like separated events should not influence each other \\citep{Wess:1988dm}. There are at least three good theoretical reasons to suspect that locality could not hold for arbitrary small distances: \\begin{enumerate} \\item  The unavoidable singularities of the General Theory of Relativity (GRT) \\citep{Hawking1982} that makes the structure of space-time very complex at distances of the order of the Planck length $\\ell_P=\\sqrt{\\hbar G/ c^3}\\simeq1.66\\times 10^{-35}\\hbox{ m}$. The fabric of space-time has been vividly described as a \\quotes{foamy} structure (see \\emph{e.g.} \\citealt{Wheeler1998book}), turbulently perturbed by black holes continuously popping out of the vacuum and evaporating in times of the order of $\\Delta t= c\\ell_P$.\\bigskip \\item String/M-theories  are intrinsically non local theories \\citep{Amati:1988tn} where the ordinary concept of point-like components of matter is substituted by two dimensional object with a finite, even if very small, dimensions . The space-time structure of the string theory is discontinuous on a scale $\\Delta t\\Delta x\\ge c \\ell_s^2$ , being $\\ell_s$ the size of the string \\citep{Yoneya:1989ai}. If string theory has incorporates gravity, one of the characteristic length of the theory should be $\\ell_s^{grav}=\\ell_P$, but might exist other characteristic lengths, corresponding to different type of interactions, with $\\ell_s>\\ell_P$ \\citep{Lykken:1996fj}. \\bigskip \\item In string theory, gravity is just one of the many possible excitations of a string (or other extended object) living over some background metric space. The existence of such background metric space, over which the theory is defined, is needed for the formulation and for the interpretation of the theory. The \\quotes{Loop Quantum Gravity} theory \\citep{lrr-1998-1} is an attempt to eliminate this background space-time. In this theory the space-time has a kind of \\quotes{polymeric} structure with minimal space cell with volumes $\\Delta V\\sim \\ell_p^3$ \\citep{Rovelli:1994ge}. \\end{enumerate} In any case the indeterminacy principle states that to resolve a length scale $\\ell$ we need energies of the order of $\\Lambda=\\hbar c/\\ell=(\\ell_P/\\ell)M_P$ where $M_P=\\sqrt{\\hbar c/G}=1.22\\times 10^{19}\\;\\mathrm{GeV}$ is the Planck mass. Na\\^{\\i}ve expectations of the orders of magnitude for $CPT$ violations, motivated by dimensional considerations, will be \\begin{equation}\\label{eq:intro1} \\Delta\\mathcal{H}\\approx \\frac{E^2}{\\Lambda} \\end{equation} \\par At low energies the Hamiltonian is of the order of magnitude of the mass of the particle, therefore CPT symmetry can be tested in the laboratory measuring the difference of mass between particles and antiparticles. The best constraint has been obtained from the limits on the mass difference between the neutral strange mesons $K^0-\\overline{K}^0$ recently obtained from the KLOE experiment at the Daphne $\\varphi$-factory (LNF-INFN)\\citep{Ambrosino:2006ek} \\begin{equation}\\label{eq:intro2} \\frac{|m_{K^0}-m_{\\overline{K}^0}|}{m_{K^0}}<1.26 \\times 10^{-18}\\quad \\mathrm{(95\\%\\; C.L.)} \\end{equation} that from Eq.\\ \\eqref{eq:intro1} is predicted to be $\\approx 0.5\\times 10^{-19}$.\\par Astrophysical and cosmological test of the CPT/LI violations, that I will discuss in the rest of this paper are complementary to local tests, because test modifications of physical law over large spatial scale $D\\sim10^{26}\\hbox{ m}$ and long times $t_0\\sim13.7\\hbox{ Gy}$. ",
        "conclusions": "Astrophysical tests show that Einstein's Special Relativity Theory is in quite good health, far beyond the Planck scale. The final assessment of the evidence for GZK cut-off in the primary UHE cosmic ray spectrum would secure an upper limit to CPT-even violations $\\lesssim 10^{-25}$.\\par The Crab Nebula promises to be an excellent particle physics laboratory for the search of CPT-odd effects, that could allow the exploitation of electron beam energies up to $\\approx 2500\\hbox{ TeV}$, but unfortunately the interpretation of data is far from being lacking in ambiguity. For example if it could be demonstrated that synchrotron emission is produced by negative electrons only, a limit to CPT-odd violations of the order of $\\lesssim 10^{-8}\\ell_P$ could be assessed.\\par The existence of a preferred direction in space-time, possibly connected with a quintessential dark-energy, is constrained by the optical polarimetry of far distant galaxies to be very small $<5\\times 10^{-5}\\ell_P$, but the scale of anisotropy estimated from the dark energy density is $<2.5\\times 10^{-5}\\ell_P$, close but not conclusive. It is intriguing that in the future X and $\\Gamma$-ray polarimetry of bright objects at cosmological distances, like AGN and GRB, could improve the present limits, by order of magnitudes, if emission models are also improved.\\par The polarization of the CMBR is another interesting source of data on possible CPT-odd violations in the photon sector, but at present the results of WMAP allow to constraint the scale in the range $\\sim 0.1\\ell_P$. The predicted sensitivity of the Planck satellite, to be launched about one year from now, can improve significantly the above limit in the next decade."
    },
    "0711/0711.1264_arXiv.txt": {
        "abstract": " ",
        "introduction": "Most of the matter in the Universe is \"dark\", i.e. it can not be observed directly  by astronomical observations, as the particles that form it are not baryons \\cite{peebles}, \\cite{kamionkowski}. A lot of types of new particles were suggested by the theorists to explain the Dark Matter (DM) problem. One of the best candidates with a strong theoretical background is axion -- a light neutral pseudoscalar particle, which appears in spontaneous breaking of Peccei-Quinn symmetry \\cite{pq} as a solution of the strong CP problem.  In fact, axions are not really dark. In external electromagnetic field they can couple to  virtual photons and produce  real photons (so-called Primakoff effect \\cite{prim})  \\begin{equation} \\label{primakoff} a+\\gamma_{\\rm virt} \\rightarrow\\gamma \\end{equation}  Several experiments that utilize the Primakoff effect are underway now. One direction of such studies is to look for an effect related to conversion of high energy axions from interiors of  the Sun in strong magnetic fields in a laboratory \\cite{CAST}, \\cite{Tokyo}. Also solar axions can be searched in X-ray observations because of  conversion of these particles in the Earth magnetic field \\cite{ziotas},\\cite{davoud}. All these experiments are searching for   \"hot\" and \"young\" axions, not \"cold\" cosmological particles originated in the Early Universe. Here we propose an astronomical method to detect emission due to conversion of cosmological axions. Direct experimental searches for such axions are in progress now in laboratories, too \\cite{sik2}.  The  strongest magnetic fields known can be found in the surroundings of neutron stars (NS), so under certain conditions axions that constitute cosmological DM can experience the Primakoff effect in magnetospheres of NS. Possible use of such super-strong magnetic fields for searching light pseudo-scalar bosons in processes of photon conversion was previously discussed in several papers \\cite{gnedin1,gnedin2}. In what follows, a possibility of detection of photons from such conversion  is studied. ",
        "conclusions": "We suggest to use radio observations of close-by X-ray dim cooling isolated NSs to search for observational appearances of DM axions. Whether the axion-photon conversion in NS's magnetic field takes place, the radio spectrum of the object might have a very distinct feature -- a narrow spike at a frequency corresponding to the rest mass of axion. If the coupling constant of the photon-axion interaction is $M=10^{10}$ GeV, the density of DM axions is $\\rho=10^{-24} \\, {\\rm g \\, cm^{-3}}$ and a NS with $B\\sim 10^{14}$~G is located at the distance of 300 pc from the Solar system,  then the flux density of  signal for axions with rest mass of $5 \\, {\\rm \\mu eV}$ is as large as several tenths of mJy at the frequencies $\\approx 1200$ MHz in the bandwidth $\\approx 3$ MHz.  Acknowledgements. M.P. wishes to thank  Dr. Mikhail Sazhin for useful comments. S.P. thanks Alexander Gvozdev and Igor Ognev for discussions. The work of S.P. was supported by INTAS and by the RFBR grant 07-02-00961."
    },
    "0711/0711.1863_arXiv.txt": {
        "abstract": "We present the galaxy cluster autocorrelation function of 277 galaxy cluster candidates with $0.25 \\le z \\le 1.5$ in a 7 deg$^2$ area of the IRAC Shallow Cluster Survey.  We find strong clustering throughout our galaxy cluster sample, as expected for these massive structures.  Specifically, at $\\meanz = 0.5$ we find a correlation length of $r_0 = 17.40^{+3.98}_{-3.10}$ \\hmpc, in excellent agreement with the Las Campanas Distant Cluster Survey, the only other non--local measurement.  At higher redshift, $\\meanz = 1$, we find that strong clustering persists, with a correlation length of $r_0=19.14^{+5.65}_{-4.56}$ \\hmpc.  A comparison with high resolution cosmological simulations indicates these are clusters with halo masses of $\\sim 10^{14} M_\\odot$, a result supported by estimates of dynamical mass for a subset of the sample.  In a stable clustering picture, these clusters will evolve into massive ($10^{15} M_\\odot$) clusters by the present day. ",
        "introduction": "The clustering amplitude of massive galaxy clusters, and in particular its dependence on richness or cluster mass, is a strong function of the underlying cosmology.  Recent theoretical work \\citep[\\eg][]{majumdar&mohr03,younger06} has demonstrated how cluster surveys can be ``self--calibrated,'' providing precise simultaneous constraints on both the cosmology and cluster evolution models.  Given redshift and approximate mass information for each cluster, reliable cosmological parameter estimation is feasible even in the presence of significant, and potentially unknown, evolution in cluster physical parameters \\citep[\\eg][]{gladders07}.  The ultimate goal of these cosmological studies, a measurement of the equation of state of dark energy with an accuracy competitive with SNe Ia methods, also requires knowledge of the cluster power spectrum or autocorrelation function \\citep{majumdar&mohr04,wang04}. Such extensive sample characterization naturally emerges from mid--infrared selected photometric redshift cluster surveys like the IRAC Shallow Cluster Survey \\citep[ISCS;][hereafter E07]{eisenhardt07}.  Previous studies \\citep[\\eg][and references therein]{bahcall03} have produced largely local ($z\\la 0.3$) galaxy cluster clustering measurements with limited baselines for evolutionary studies. The ISCS provides the opportunity to address these issues in one of the largest statistical samples of high redshift clusters to date. In this Letter we present the first measurement of the galaxy cluster autocorrelation function extending over more than half the cosmic age of the Universe. Such measurements also allow evolutionary connections between galaxy clusters and high redshift, highly clustered galaxy populations to be explored.  We use a concordance cosmology throughout, with $\\Omega_M = 0.3$, $\\Omega_\\Lambda = 0.7$, and $H_0 = 70$ km s$^{-1}\\!$ Mpc$^{-1}$. For consistency with previous studies we report distances, including correlation lengths, in units of comoving \\hmpc, with $H_0 = 100 \\,h$ km s$^{-1}$ Mpc$^{-1}$.  % ",
        "conclusions": "\\label{Sec: conclusion}  By deprojecting the angular correlation function measured in redshift bins spanning $z=0.25$ to $z=1.5$, we have determined the real--space clustering amplitudes for ISCS clusters at $\\left<z\\right>=0.53$ and $\\left<z\\right>=0.97$ to be $r_0 = 17.40^{+3.98}_{-3.10}$ and $r_0=19.14^{+5.65}_{-4.56}$ \\hmpc, respectively.  These measurements are consistent with the relation between correlation amplitude and mean intercluster distance predicted by LCDM.  The ISCS clusters have total masses exceeding $10^{14} M_\\odot$ and will evolve into very massive clusters by the present day."
    },
    "0711/0711.2267_arXiv.txt": {
        "abstract": "Rotation has become an important element in evolutionary models of massive stars, specifically via the prediction of rotational mixing. Here, we study a sample of stars, including rapid rotators, to constrain such models and use nitrogen enrichments as a probe of the mixing process. Chemical compositions (C, N, O, Mg and Si) have been estimated for 135 early B-type stars in the Large Magellanic Cloud with projected rotational velocities up to $\\sim$300\\,km\\,s$^{-1}$ using a non-LTE {\\sc tlusty} model atmosphere grid. Evolutionary models, including rotational mixing, have been generated attempting to reproduce these observations by adjusting the overshooting and rotational mixing parameters and produce reasonable agreement with 60\\% of our core hydrogen burning sample. We find (excluding known binaries) a significant population of highly nitrogen enriched intrinsic slow rotators ($v \\sin i \\lesssim$50\\,km\\,s$^{-1}$) incompatible with our models ($\\sim$20\\% of the sample). Furthermore, while we find fast rotators with enrichments in agreement with the models, the observation of evolved ($\\log g < 3.7$dex) fast rotators that are relatively unenriched (a further $\\sim$20\\% of the sample) challenges the concept of rotational mixing.  We also find that 70\\% of our blue supergiant sample cannot have evolved directly from the hydrogen burning main-sequence. We are left with a picture where invoking binarity and perhaps fossil magnetic fields are required to understand the surface properties of a population of massive main sequence stars. ",
        "introduction": "\\label{s_intro}  Stellar rotation in massive stars has been adopted in recent theoretical models to, for example, predict the correct blue to red supergiant ratio \\citep{mae01}, the progenitors of gamma-ray bursts through homogeneous evolution \\citep{yoo05} and Wolf-Rayet populations as a function of metallicity (\\citealt{mey05}; \\citealt{vin05}). However, the consequent surface enrichment of helium and nitrogen through rotational mixing is poorly constrained by observations \\citep{daf01}, with the mixing typically calibrated from either studies of main-sequence stars with low projected rotational velocities \\citep{kor02} or evolved supergiant stars \\citep{ven99}. The former are biased towards slow rotators while the latter have evolved beyond the core hydrogen burning phase. The analysis of an unbiased sample of fast rotating core-hydrogen burning stars in order to properly calibrate the predicted rotational mixing efficiency was a strong driver for a large survey of O- and early B-type stars in our Galaxy and the Magellanic Clouds \\citep{eva05}. \\citet{hun07a} and \\citet{tru07}, hereafter Paper I and II respectively, have presented chemical abundances for the narrow lined objects (predominately slow rotators) in the LMC sample. Here we extend the study with the chemical analysis of the faster rotating stars. ",
        "conclusions": "In Fig.~\\ref{f_nvsini} the nitrogen abundances are plotted against the projected rotational velocities for our sample stars, and compared with the evolutionary models. Given the large number of objects in our sample, it is reasonable to assume a random distribution of inclination angles, and hence the rotational velocity of the tracks has been scaled by $\\pi$/4. Although we have attempted to contrain the evolutionary tracks to the observations, we are able to reproduce agreement for only 60\\% of the data (excluding known radial velocity variables).   However, two groups of core hydrogen burning stars in Fig.~\\ref{f_nvsini}(a) stand out as being in conflict with the evolutionary models. Group~1 contains fast rotators which have little chemical mixing. In particular, this group includes fast rotating single stars with surface gravities indicating that they are near the end of core hydrogen burning (excluding radial velocity variables this is $\\sim$20\\% of the core hydrogen burning sample.) The gravities used to discriminate these stars have not been corrected for the fact that we may observed them almost equator-on, i.e. their true polar gravities could be larger, meaning that they are less evolved than our derived gravities imply. However, applying the surface gravity corrections described in \\citet{hua06b} would increase the gravities by up to $\\sim$0.3\\,dex for our faster rotators. As the zero age main-sequence gravity of our 13\\,M$_{\\sun}$ models is about 4.3\\,dex, and since most of the stellar expansion on the main sequence occurs towards the end of core hydrogen burning, this gravity uncertainty can not reconcile the situation.   The second discrepant group of stars in Fig.~1a are the 15 slow rotators ($v \\sin i < $50\\,km\\,s$^{-1}$) that show significant nitrogen enrichment (Group~2). This group also forms $\\sim$20\\% of the non-binary core hydrogen burning sample. Although it could be argued that they may be fast rotators observed pole-on, statistically this is unlikely. In order to reproduce their mean nitrogen abundance requires a rotational velocity of $\\sim$200\\,km\\,s$^{-1}$. Using simple geometry to calculate the solid angle restriction that we must impose, i.e. $4\\pi(1-cos\\theta)$, and assuming that we need 15 stars to have this velocity, we expect less than one star to appear with such a low rotational velocity due to random inclinations. The simple corollary of this is that if we have 15 stars with velocities $\\sim$200\\,km\\,s$^{-1}$ populating Group~2 through random angles of inclination, over 300 similary enriched stars with $v\\sin i$ values between 150 and 200\\,km\\,s$^{-1}$ would be expected. Such a population clearly does not exist.  Evolved rapid rotators with low nitrogen enrichment (Group~1) may be produced by close binaries with an initial period small enough to ensure tidal locking and slow rotation but large enough for highly non-conservative mass transfer, i.e. with an amount of mass accreted which is sufficient to spin-up the star but insufficient to enrich it significantly \\citep{pet05}. However, for many of the evolved stars in Group~1, there is no indication of binarity. The single star nature for these objects would pose a serious challenge to the theory of rotational mixing.  While binaries might be able to populate Group~2 \\citep{lan08}, another explanation may be more likely. \\citet{mor06} have analyzed ten slowly rotating Galactic early $\\beta$-Cephei B-type stars and found a highly enriched group (four out of ten) of which three have detected magnetic fields. Although pulsations in these stars may cause an apparent surface enrichment, \\citet{bou07} report that only slight nitrogen enrichments would be expected. As such, assuming that the effect of pulsations is negligible, there appears to be a correlation between magnetic fields and nitrogen enrichment. Indeed, \\citet{wol07} have attributed the large number of slow rotators often seen in massive star populations to magnetic locking of the star to the accretion disk during the star formation process and \\citet{ale07} also show a correlation between slow rotation and magnetic fields for magnetic A- and B-type stars. Hence we postulate that the highly enriched slow rotators in the LMC are analogs of the enriched Galactic magnetic stars. This would imply that, independent of metallicity, a significant fraction of early B-type stars are intrinsically magnetic, analogous to the well known situation for lower mass stars. In this context, the identification of three He-rich slowly rotating OVz stars in NGC\\,346 \\citep{mok06} suggests that the phenomenon of intrinsic magnetic fields in massive stars may not be confined to the B-star regime. If these magnetic fields are of fossil origin, i.e. possessing long-term stability \\citep{bra04}, in contrast to the fields produced by the \\citep{spr02} mechanism, one may speculate that stars in Group~2, and their Galactic counterparts, might be the progenitors of magnetars, analogous to the suggestion of \\citet{fer06} that magnetic white dwarfs evolve from  magnetic A, F and late B stars. While the fossil field hypothesis might well explain the slow rotation, the physical process which leads to the enrichment of nitrogen in the stars of Group~2 remains to be identified.   The nitrogen abundances of the supergiants (Fig.~\\ref{f_nvsini}(b)) appear to fall into two distinct groups, one group having a level of nitrogen enhancement consistent with that seen for the majority of the core hydrogen burning objects ($<$7.2\\,dex; Group 3), and a second having a much greater level of enrichment ($>$7.6\\,dex; Group 4). It should be noted that although the Group~4 objects appear to be well fitted by evolutionary tracks with initial rotational velocities of 200-300\\,km\\,s$^{-1}$, such a rotational velocity distribution for supergiants is inconsistent with that for core-hydrogen burning objects and rotation does not increase the blue-supergiant lifetime. Hence, the enrichments of the supergiants in Group~4  (70\\% of the non-binary supergiant sample) are incompatible with the theory of rotational mixing. Additionally it should be noted that the lowest gravity stars ($<$2.8\\,dex) are all enriched in nitrogen, implying that some evolutionary process not accounted for in the models is responsible for the enrichment.  The simplest way to interprete this would be to characterize Group~3 as pre-red supergiant objects, and Group~4 as post-red supergiant objects. The nitrogen abundances and rotational velocities in Group~4 are indeed consistent with the predictions for a blue loop stage \\citep{heg00}. However, we note that our models do not return to the high effective temperatures of the stars in Group~4 and hence their position on the H-R diagram cannot be reproduced. As several of the enriched objects show evidence of binarity, mass transfer may also be important \\citep{wel01}.  To conclude, our study can not provide unambigous evidence for rotational mixing acting in massive stars. The fast rotators in our sample can be interpreted in two ways. Firstly as rotationally mixed stars, if the relatively non-enriched stars (Group~1) can be explained by binary effects. However, if it were confirmed that these Group~1 stars are single stars, then the enriched rapid rotators may need to be understood as binary products --- and rotational mixing would be inefficient, or much more complex than described by the present-day shellular models. Additionally, the population of intrinsically slow rotators with nitrogen enhancements of up to a factor of $\\sim$6 implies that studies of nitrogen enhancements with stellar samples which are restricted to low projected rotational velocities --- which applies to most previous works --- are not suited for constraining, or motivating, the adoption of rotational mixing in massive stars. Finally our supergiant sample implies that these stars cannot be considered as representative of the amount of mixing that occurs during the hydrogen burning main-sequence and hence should not be used to constrain this process. In summary, our study provides a challenge to rotational mixing, and provides evidence for two other processes affecting the rotation and the surface abundances of populations of massive stars, likely binarity and magnetic fields."
    },
    "0711/0711.2321_arXiv.txt": {
        "abstract": "Anomalies in the large-scale CMB temperature sky measured by WMAP have been suggested as possible evidence for a violation of statistical isotropy on large scales.  In any physical model for broken isotropy, there are testable consequences for the CMB polarization field.   We develop simulation tools for predicting the polarization field in models that break statistical isotropy locally through a modulation field. We study two different models: dipolar modulation,  invoked to explain the asymmetry in power  between northern and southern ecliptic hemispheres, and quadrupolar modulation, posited to explain the alignments between the quadrupole and octopole. For the dipolar case, we show that predictions for the correlation between the first 10 multipoles of the temperature and polarization fields can typically be tested at better than the 98\\% CL.  For the quadrupolar case, we show that the polarization quadrupole and octopole should be moderately aligned. Such an alignment is a generic prediction of explanations which involve the temperature field at recombination and thus discriminate against explanations involving foregrounds or local secondary anisotropy.  Predicted correlations between temperature and polarization multipoles out to $\\ell = 5$ provide tests at the $\\sim$ 99\\% CL or stronger for quadrupolar models that make the temperature alignment more than a few percent likely. As predictions of anomaly models, polarization statistics move beyond the {\\it a posteriori} inferences that currently dominate the field. ",
        "introduction": "Several anomalies observed in the CMB data suggest a possible breaking of statistical isotropy on the largest angular scales. Broken isotropy would represent a radical revision of the standard cosmological model as statistical isotropy underlies  all cosmological inferences.  As such, it is important to test such claims by all means possible.  Specifically, the measurements from the {\\it Wilkinson Microwave Anisotropy Probe} (WMAP) \\cite{bennett, spergel, wmap3yr, page} indicate an asymmetric distribution in power between the northern and southern ecliptic hemispheres \\cite{eriksen_hansen}. Also, the data exhibit alignments in the temperature field; in particular, the observed quadrupole and octopole of the CMB are planar and mutually aligned, and this plane is oriented roughly orthogonal to the CMB dipole direction \\cite{oliveira,land_and_magueijo}.  The statistical significance of these preferred directions has been investigated by several authors \\cite{oliveira, tegmark, copi,schwarz,axis, eriksen_banday,Copi:2005ff,Copi:2006tu,Wiaux:2006zh,Land:2004bs}, but the {\\it a posteriori} choice of statistics makes their interpretation difficult \\cite{Magueijo:2006we,Land:2006bn,Gordon:2007xm}.   Indeed several generic tests of isotropy yield no evidence for statistical anisotropy \\cite{Hajian:2003qq,Hajian:2006ud,Donoghue:2004gu,ArmendarizPicon:2005jh}. Since measurements of the temperature field have already reached the cosmic variance limit, further information from it will not be forthcoming, barring improvements in foreground and systematics modeling. On the other hand, large-scale CMB polarization measurements are still in their infancy and provide fertile grounds for future tests of isotropy.  While generic tests of isotropy have also been proposed for polarization \\cite{Basak:2006ew,Pullen:2007tu} and unrelated statistical anomalies may be found, any physical model that purports to explain the temperature anomalies provides testable predictions for the statistics of the polarization field.   Matching anomalies in polarization for such models can provide a means of going beyond both {\\it a posteriori} inferences and blind searches for isotropy anomalies.   In this {\\it Paper}, we develop the radiative transfer tools required to investigate polarization statistics.  Large-angle polarization arises from rescattering of CMB radiation from recombination during reionization.  Given a model that predicts the statistical properties of the temperature field at recombination, Monte-Carlo realizations of the polarization field can be generated and analyzed for signatures of broken isotropy.  We apply these techniques to models that break statistical isotropy with a multiplicative modulation field \\cite{wayne}.  These models posit that the observed temperature field is the product of a modulating field with superhorizon wavelength fluctuations and a second field that carries fluctuations at the observed scales.   Gradients in the modulation field naturally pick out a preferred direction within our horizon. Such models have been proposed to explain both the hemisphere asymmetry \\cite{spergel,gordon,eriksen_banday} and the quadrupole-octopole alignment \\cite{wayne}.   The structure of the paper is as follows.  In \\S~\\ref{sec: isotropy_breaking} we discuss the general form of the modulation model and develop the simulation tools to predict the polarization field. The analytic framework for polarization predictions and explicit relations for the simplest modulation statistics are given in the Appendix. In \\S~\\ref{sec: dipolar}, we study the polarization predictions for the dipolar modulation models that are designed to explain the hemisphere asymmetry. We continue in  \\S~\\ref{sec: quadrupolar} by studying quadrupolar modulation models that explain the quadrupole-octopole alignment.  We discuss our results in \\S~\\ref{sec: discussion}. ",
        "conclusions": "\\label{sec: discussion}  Statistical anomalies in the large-scale CMB temperature sky measured by WMAP have been suggested as possible evidence for the violation of statistical isotropy on large scales.  In any physical model for such a violation, there are testable consequences for the CMB polarization field.  We have developed a radiative-transfer simulation-based approach to predicting these consequences.  We have applied this approach to modulation models of broken isotropy.  In this class of models, a long-wavelength modulation of shorter-wavelength fluctuations breaks isotropy locally, while the universe remains globally isotropic and homogeneous.  A linear or dipolar modulation can be used to explain the observed power asymmetry between the northern and southern ecliptic hemispheres at multipole moments $\\ell \\sim 20-40$ \\cite{eriksen_hansen,spergel,gordon, eriksen_banday}. Linear modulation produces a dipole coupling that correlates multipole moments $\\ell$ to $\\ell \\pm 1$ in all of the two point functions of the temperature and $E$ polarization field.   Using an analytic framework developed in the Appendix, we derive an estimator of the modulation amplitude from these fields from these correlations.  We find that the addition of polarization information substantially improves the ability to test for dipole modulation at the lowest multipoles $\\ell \\lesssim 10$. Under a wide range of assumptions, the dipole modulation explanation of the hemisphere asymmetry can be tested  at the 98\\% CL or greater. These tests will require near cosmic variance limited polarization measurements at the low multipoles.   Such measurements are within the statistical reach of the {\\it Planck} satellite but systematic contamination from foregrounds will have to be sufficiently controlled.  A quadratic or quadrupolar modulation in the temperature field can  be used to explain the alignment of the quadrupole and octopole \\cite{oliveira,wayne}. We generalize the considerations of Ref.~\\cite{wayne} and develop full 3D models that enhance the probability of alignments by a factor of $\\sim 20$.  These models predict that the quadrupole and octopole of the $E$-polarization will tend to align as well.  However alignment of the polarization mainly tests that the temperature alignment is cosmological and originates from recombination, not the specifics of the modulation model.  As such its main use will be to reject explanations of the alignments due to foregrounds or purely local secondary anisotropy. In fact a very high alignment, comparable to the temperature field, is not expected in the modulation model due to the smaller strength of the modulation at the distance to reionization.   If such an alignment is seen then it is more likely that it arises from a chance configuration of the temperature field at recombination.  Quadrupolar modulation models can be distinguished in a manner similar to the dipole modulation but through a correlation between multipoles $\\ell$ and $\\ell \\pm 2$ in the temperature and polarization fields.  We find that the combination of temperature and polarization information out to $\\ell \\le 5$ can strongly constrain the level of modulation in the low multipoles.  As a consequence, polarization measurements can in principle rule out models  that make the temperature alignments more than a few percent likely $\\sim 99\\%$ CL or stronger.  In general, matching signatures of broken isotropy in the polarization provide a powerful test of competing explanations of the temperature anomalies. Explanations that involve foregrounds (e.g. \\cite{Slosar:2004xj}) or local gravitational effects \\cite{Inoue:2006rd,Vale:2005mt} would tend to predict different signatures.   This is especially true since matching anomalies in the $E$ polarization are non-local in the Stokes parameters. Possible experimental systematics would be specific to a given experiment, and would thus have to be analyzed in a case by case basis. Therefore we do not consider them in this paper.  The examples we have studied illustrate that polarization in the future will provide a fruitful testing ground for explanations of the observed temperature anomalies.  As predictions of anomaly models, polarization statistics move beyond the {\\it a posteriori} inferences that currently dominate the field."
    },
    "0711/0711.0002_arXiv.txt": {
        "abstract": "We report on the detection of \\fka emission in F04103$-$2838, an ultraluminous infrared galaxy (ULIRG; log[L$_{\\rm IR}$/L$_\\odot$] $\\ge$ 12) that is optically classified as a LINER.  Previous {\\it Chandra} observations suggested the presence of both a starburst and an AGN in this source.  A deeper ($\\sim$20~ksec) {\\it XMM-Newton} spectrum reveals an \\fka line at rest frame energy $\\sim$6.4~keV, consistent with cold neutral iron.  The best-fit spectral model indicates the \\fka line has an equivalent width of $\\sim$1.6~keV.  % The hard X-ray emission is dominated by a Compton-thick AGN with intrinsic 0.2--10~keV luminosity $\\sim$10$^{44}$ ergs~s$^{-1}$, while the soft X-ray emission is from $\\sim$0.1~keV gas attributed to the starburst.  The X-ray spectrum of this source bears a striking resemblance to that of the archetypal luminous infrared galaxy NGC~6240 despite differences in merger state and infrared properties. ",
        "introduction": "\\label{sec:intro}  The primary energy source (AGN versus starburst activity) of ultraluminous infrared galaxies (ULIRGs; log[L$_{\\rm IR}$/L$_\\odot \\ge$ 12]) is still a matter of debate.  Optical and infrared emission-line spectra suggest that the energy output of most local ULIRGs is dominated by starbursts, but the ``warm'' infrared colors and quasar-like spectra of the more luminous ULIRGs indicates that black-hole driven activity plays an increasingly important role in these objects (e.g., Veilleux et al. 1995, 1997, 1999a, 1999b; Genzel et al. 1998; Surace \\& Sanders 1999; Tran et al. 2001).  The dusty, gas-rich nature of ULIRGs implies, however, that observations in energy bands other than radio and X-ray may not always probe the true nuclear energy source of these objects.  Since the luminosity of ULIRGs in the radio is insignificant, X-ray observations remain arguably the best option to solve this energy source mystery.  Unresolved hard X-ray emission is in principle a telltale sign of a dominant AGN. However, if a large column density of gas ($\\ga 10^{24}$ cm$^{-2}$) is located in front of the nucleus, then directly viewed X-rays from the AGN will be strongly attenuated.  For such cases, \\fka lines with large equivalent widths ($\\ga$ 1 keV) are expected due to scattering off circumnuclear material (e.g., Ghisellini et al. 1994; Krolik 1994). Thus, the discovery of such \\fka lines may be the best evidence for energetically dominant AGNs in highly obscured ULIRGs.   In recent years, three X-ray surveys have added considerably to our knowledge of ULIRGs. \\citet{ptak} performed a volume-limited ($z < 0.045$) survey of ULIRGs with {\\it Chandra}.  On the basis of their dust temperatures (25-to-60~$\\mu$m flux ratio) and X-ray luminosities, three of the eight ULIRGs sampled by \\citet{ptak} were classified as AGN-dominated (Mrk~231, Mrk~273, and F05189-2524).  In the same year, \\citet{frances} published the results of a similar survey with {\\it XMM-Newton} that focused on the brightest local ULIRGs (only one ULIRG in their sample had $z > 0.082$).  Of the ten ULIRGs sampled by \\citet{frances}, three were AGN-dominated (Mrk~231, F19254-7245, and F20551-4250) and two had X-ray signatures of both a starburst and an AGN (F20100-4156 and F23128-5919).  All of the AGN-dominated ULIRGs showed strong Fe~K emission lines (Maloney \\& Reynolds 2000; Braito et al. 2003; Ptak et al. 2003; Franceschini et al. 2003; Braito et al. 2004).  These two pioneering surveys proved the viability of using the X-ray emission as a diagnostic for AGN activity in ULIRGs.  However, these two surveys only studied a small set of the nearest and brightest ULIRGs, and therefore were not able to draw general conclusions on the issue of the energy source among ULIRGs as a class.  In an attempt to expand this type of study to a more characteristic sample of ULIRGs, our group (Teng et al. 2005, hereafter \\teng) conducted a snapshot (10 ksec/target) survey of 14 ULIRGs from the 1-Jy sample\\footnote{The 1~Jy sample of ULIRGs is comprised of {\\it IRAS} galaxies with fluxes at 60~$\\mu$m exceeding 1~Jy, L$\\rm{_{IR} \\ge 10^{12}}$~L$\\rm{_\\odot}$, galactic latitude $\\mid$b$\\mid$ $> 30^\\circ$, f(60~$\\mu$m) $>$ f(12~$\\mu$m) (to avoid stars), {\\it IRAS} color log[$f_{60}/f_{100}$] $>$ --0.3 (to favor luminous infrared systems), and redshift 0.018 $< \\rm{z} <$ 0.268 (Kim \\& Sanders, 1998).}. These sources were carefully selected to sample the full range of infrared luminosities and infrared colors that characterize the entire class of local ULIRGs. All 14 galaxies were detected by {\\it Chandra}, though most (11/14) had less than 40 counts.  The analysis showed that the two brightest galaxies in the sample have optical and X-ray spectral characteristics of Seyfert~1 nuclei.  Most others have X-ray photon indices (estimated using hardness ratios) and hard X-ray to far-infrared flux ratios which are similar to those of starbursts.  One exception, F04103$-$2838, had a hardness ratio (deduced from only 30 counts) that suggested the presence of a starburst coexisting with an AGN.  The low signal-to-noise data could not distinguish between a Compton-thick AGN or an intrinsically faint nuclear source.  This object is optically classified as a LINER (Veilleux et al. 1999a). F04103$-$2838 has one of the largest 25-to-60~$\\mu$m flux ratios of all 1~Jy ULIRGs ($f_{25}/f_{60}$ = 0.30).  In fact, this is the warmest of all {\\it IRAS} 1~Jy ULIRGs with optical LINER or H~II classification.  This source is even warmer than some of the Seyfert galaxies in the 1~Jy sample (see Fig. 1 of \\teng).  Recent {\\it Spitzer} detection of [Ne~V] and [O~IV] lines from this source has confirmed the existence of an AGN in this system (Veilleux et al. 2007, in prep.), making it a rare example of a ULIRG optically classified as a LINER that is not classified as a starburst on the basis of mid-infrared spectroscopy (Lutz, Veilleux, \\& Genzel 1999; Taniguchi et al. 1999).  In this paper, we present an {\\it XMM-Newton} observation of F04103$-$2838 which delves deeper into the nature of this AGN.   F04103$-$2838 is an interacting galaxy system in the late stages of a merger as indicated by the presence of a single nucleus with distinct tidal tails \\citep{vei02, dasyra}.  This object has an infrared (8--1000 $\\mu$m) luminosity of 10$^{12.15}$L$_\\odot$ and a cosmological redshift of 0.118.  Assuming H$_0$ = 75~km~s$^{-1}$~Mpc$^{-1}$ and q$_0$ = 0 (used throughout this paper), the luminosity distance of this object is 497~Mpc.  At this distance, 1$\\arcsec$ corresponds to $\\sim$2.4~kpc. In \\S 2 of this paper, we describe our new {\\em XMM-Newton} observation of F04103$-$2838 and the methods we used to reduce these data. In \\S 3, we present the analysis of these data, emphasizing the results on the X-ray morphology, the lack of flux variability, and the spectral decomposition of the X-ray emission. The implications of these results are discussed in \\S 4.  The main conclusions are summarized in \\S 5. ",
        "conclusions": "\\label{sec:dis}  \\subsection{The Soft Component} \\label{sec:soft}   The results from the spectral fitting suggest that the soft X-ray (0.2--2~keV) flux is best described as thermal emission from hot gas with kT $\\sim$ 0.1~keV (T$\\sim$1.2$\\times 10^6$~K).  This is somewhat lower than the range of gas temperatures (0.6--0.8~keV) found in LINERs \\citep{gonzalez}.  The results for F04103--2838 is also somewhat lower than the results from \\citet{grimes} who performed a {\\it Chandra} archival study of the soft X-ray emission from starburst galaxies ranging in luminosity from dwarf galaxies to ULIRGs.  The authors found that the soft X-ray thermal emission of these starburst galaxies tends to fall in the temperature range kT $\\sim$0.25--0.8~keV with ULIRGs occupying the upper end of this temperature range. These large temperatures can all be attributed to powerful starbursts.  The soft X-ray emission in F04103$-$2838 is likely the result of thermal bremsstrahlung from a hot gas produced by the merger-induced starburst or by intrinsically extended soft X-ray emission heated by the AGN.  % If the ion density equals that of the electrons, the relationship between the electron density ($n_e$) and luminosity of an emitting region of a given volume ($V$) is \\begin{equation} L_{ff} \\approx 1.7 \\times 10^{-25} n_e^2 f V~{\\rm ergs~s^{-1}}, \\label{eq:brems} \\end{equation} where $f$ is the filling factor for the hot gas\\footnote{Equation~\\ref{eq:brems} is based on equation (5.14b) and Figure~5.2 of \\citet{rybicki} for T = $10^6$~K in the energy range of 0.2--2~keV.}.  The non-AGN contribution of the nominal 0.2--2~keV luminosity from the best-fit model (model C) for F04103$-$2838 is $1.6 \\times 10^{41}$~ergs~s$^{-1}$. Assuming the emitting region is spherical with a diameter of $\\leq$5$\\arcsec$\\footnote{While the selection of a $\\leq$5$\\arcsec$ emitting region is based on the spatial resolution of the telescope, it should be noted that the linear diameter of 5$\\arcsec$ at the distance of F04103--2838 is less than a factor of two larger than the soft X-ray (0.5--2.5~keV) emitting region of NGC~6240 \\citep{komossa}.  Therefore, the assumption of a $\\leq$5$\\arcsec$ diameter is reasonable, even though it was chosen based on the instrument PSF.}, the average electron density has a lower limit of $\\sim$0.19$f^{-1/2}$~cm$^{-3}$.  This value is consistent with simulation results for the warm (10$^{5.5} \\lesssim$ T $\\lesssim 10^{6.5}$~K) component in the wind models of \\citet{winds}.  Observationally, this hot gas component is difficult to probe because of its low density and emissivity.  \\citet{winds} performed hydrodynamic simulations of starburst-driven galactic winds with various ISM models.  The authors found that, in general, the soft X-ray emission comes from gas with low filling factors ($10^{-3} < f < 10^{-1}$; see also Cecil, Bland-Hawthorn, \\& Veilleux 2002; and Strickland et al. 2004a, 2004b for observational constraints).  Using these values for $f$, the electron density of the hot gas in F04103$-$2838 is $\\sim$0.6--5.9~cm$^{-3}$, consistent with values derived by \\citet{netzer} in NGC~6240.  The soft X-ray emission detected in F04103--2838 may be thus the result of superwinds from the starburst.  X-ray superbubbles have been observed in Arp~220 \\citep{arp220} and NGC~6240 \\citep{netzer}. Furthermore, powerful outflow events are now thought to take place in most ULIRGs (e.g., Rupke et al. 2002, 2005a, b, c, though their sample did not include F04103--2838).  \\subsection{The Iron Feature} \\label{sec:linedis}  F04103$-$2838 joins the growing list of ULIRGs with Fe~K detections [e.g. Arp~220 \\citep{arp220}, Z11598-0112 (\\teng), F19254-7245 \\citep{frances, braito03}, Mrk~231 \\citep{malreynolds, ptak, braito04}, F05189-2524 \\citep{ptak}, Mrk~273 \\citep{ptak}, and UGC~05101 \\citep{iman, ptak}], supporting the view that an obscured AGN exists in many of these objects. The presence of an AGN in F04103$-$2838 was first suggested by \\teng\\ based on the large hard X-ray to far-infrared flux ratio; the {\\em XMM} detection of Fe~K now indicates that the luminosity of this AGN has probably been underestimated.  Few LINERs have detected \\fka lines.  \\citet{tera} studied a sample of 53 LINERs and low-luminosity Seyfert galaxies using {\\it ASCA}.  Of the 21 LINERs in their sample, Fe emission lines were detected in only five galaxies (NGC~1052, NGC~3998, NGC~4261, NGC~4579, and NGC~4736). Of these five objects, only four (i.e. those excluding NGC~4261) have centroid line energies consistent within the uncertainties of the measurements with \\fka emission due to neutral iron (E $\\sim$ 6.4~keV).  Three other LINERs have known Fe~K detections; all three are powerful luminous or ultraluminous infrared galaxies.  These galaxies are Arp~220 \\citep{arp220}, NGC~6240 \\citep{ptak, komossa}, and UGC~5101 \\citep{iman, ptak}.  {\\it Chandra} observations of Arp~220, the archetypal ULIRG, show an iron line at 6.7$\\pm$0.1~keV. This is consistent with emission due to Fe~XX up to Fe~XXVI, but not neutral iron at 6.4~keV \\citep{arp220}.  \\citet{komossa} detected Fe~K emission from each of the two nuclei in NGC~6240.  Their analysis showed that the iron lines in each nucleus are consistent with Fe K$\\alpha$ and Fe K$\\beta$ emissions.  In Figure~\\ref{fig:eqwid} we show the distribution of published Fe~K equivalent widths of all LINERs and ULIRGs known to have line emission.  Arp~220, NGC~6240, and F04103--2838 appear to have iron emission with the greatest EW measurements of all the LINERs and ULIRGs.  These large Fe~K features could be results of the blending of multiple narrower lines.  \\citet{komossa} did not publish the EWs of the lines from each of the nuclei in NGC~6240.  The result quoted here is from \\citet{ptak}.   The authors did not distinguish \\fka emission from \\fkb emission and the EW measurement is likely dominated by the brighter southern nucleus alone.  The large equivalent widths of the ULIRGs are telltale signs of obscured AGNs where line-of-sight columns of material exceeding $10^{24}$ cm$^{-2}$ prevent a direct view of the AGN; the 2 -- 10 keV flux is dominated by light scattered off dust or electrons (e.g., Ghisellini et al. 1994; Krolik et al. 1994).  The large amount of molecular gas ($\\sim$ 10$^4$ M$_\\odot$ pc$^{-2}$) within 400 pc from the nuclei of NGC~6240 (e.g., Bryant \\& Scoville 1999) is sufficient to cause this obscuration. A similar explanation likely applies to F04103$-$2838, although we are not aware of any CO measurements in this system.   Interestingly, the Fe~K complex in NGC~6240 breaks up into a number of narrow lines.  Both \\citet{netzer} and \\citet{boller03} detected Fe~K lines due to neutral iron (6.41$\\pm$0.2~keV), Fe~XXV (6.68$\\pm$0.02~keV), and Fe~XXVI (7.01$\\pm$0.04~keV) in NGC~6240.  \\citet{komossa} also detected lines at 6.4 and 6.95~keV.  The centroid energies of the lines due to neutral iron and Fe~XXV in NGC~6240 are consistent with the respective centroid energies suggested by the doublet in the F04103$-$2838 3 counts bin$^{-1}$ data.  Although simulations suggest the two line model is only significant at the $\\sim$60\\% level, a FWHM of $\\sim$30000~km~s$^{-1}$ ($\\sigma \\sim$0.3~keV) seems too broad and the two component interpretation may be more likely.  The Fe~XXVI line in NGC~6240 is much fainter than the other lines so it is not surprising that we were unable to detect this feature in the modest signal-to-noise ratio data of F04103$-$2838.  Despite their X-ray similarities, F04103$-$2838 is $\\sim$2.5 times more infrared luminous than NGC~6240.  These objects also differ in terms of {\\it IRAS} $f_{25}/f_{60}$ ratios (0.15 for NGC~6240 and 0.30 for F04103$-$2838) and merger state (NGC~6240 is in a pre-merger phase with a nuclear separation of $\\sim$ 1.3 kpc while F04103$-$2838 is in the post-merger stage with a single coalesced nucleus). There is growing observational evidence (e.g., Veilleux et al. 2002, 2006; Ishida 2004; Dasyra et al. 2006a, b, 2007) and theoretical motivation (e.g., Hopkins et al. 2005) that mergers of gas-rich galaxies often produce ``cool'' ($f_{25}/f_{60} < 0.2$) luminous infrared galaxies that evolve into ``warm'' ($f_{25}/f_{60} \\ge 0.2$) ULIRGs before becoming optical quasars.  If this evolutionary sequence applies to NGC~6240 and F04103$-$2838, the first object may actually be the precursor to the latter.  \\subsection{Energy Source of the ULIRG} \\label{sec:energy}  The lack of of short timescale variability (see \\S \\ref{sec:var}) is to be expected if most of the primary X-ray flux is being absorbed or reprocessed.  As discussed in \\S 4.2, the large equivalent width of the iron line in F04103$-$2838 implies the presence of a highly obscured AGN.  It is very difficult in such cases to estimate the intrinsic luminosity of the AGN without measurements of the $>$ 10 keV flux from the buried AGN (e.g., Mrk~231; Braito et al. 2004).  Here we follow the method of \\citet{malreynolds} to estimate the intrinsic luminosity of F04103--2838.  In their analysis of an {\\it ASCA} observation of Mrk~231, they discussed two ways of estimating the intrinsic AGN flux.  The observed X-ray flux is due to a combination of two effects: reflection and scattering.  \\citet{malreynolds} estimated the intrinsic AGN flux from the reflection and the scattering components separately.  In their geometry, the observer has an obstructed view of the nucleus so the observed flux must be either scattered or reflected into the line of sight along which there is some amount of absorbing material.  The reflected component is light from the central engine reflected off of the circumnuclear torus; the amount of reflection depends on the size of the reflecting surface and the composition of the torus.  On the other hand, the scattered component is light from the central engine (unobstructed by the torus) scattered into the line of sight. Based on their spectral fitting of the {\\it ASCA} data, \\citet{malreynolds} found that the X-ray flux of Mrk~231 is scattering-dominated with 75\\% scattered and 25\\% reflected light.  Due to the low signal-to-noise ratio of our data on F04103--2838, the same spectral fitting as performed by \\citet{malreynolds} could not be done.  The large equivalent width of the Fe~K$\\alpha$ line ($\\sim$1.6~keV) above 1~keV suggests a reflection-dominated spectrum. However, the width of the line implies it could be a blend of narrower Fe~K$\\alpha$ and ionized iron emission lines (as suggested by the 3~counts bin$^{-1}$ data).  If this were the case, the \\fka EW may be more consistent with a scattering dominated spectrum. Therefore, we will consider two cases: (1) the majority of the observed flux is due to reflection and (2) the majority of the observed flux is due to scattering to estimate the intrinsic X-ray luminosity of the AGN.  After correction for absorption, the nominal 0.2--10~keV flux of the buried AGN in F04103--2838 derived from our best-fit model (model C) is $1.83 \\times 10^{42}$ ergs~s$^{-1}$.  In the first scenario, we will assume the reflection component is 75\\% and the scattering component is 25\\% of the total observed flux.  This implies that  $L_{\\rm scattered} = 0.45 \\times 10^{42}$ ergs~s$^{-1}$ and $L_{\\rm reflected} = 1.38 \\times 10^{42}$ ergs~s$^{-1}$ for the AGN in F04103$-$2838.  In \\citet{malreynolds}, the luminosity from the reflected portion is scaled up by a factor of 25 in their modeling of the reflection process.  The reflection process differs for different galaxies; it depends on the ionization state of the mirror and the steepness of the photon index of the central black hole. \\citet{malreynolds} assumed reflection from neutral material, a reflecting fraction of 10\\%, and the canonical value of the photon index due to an AGN ($\\Gamma$ = 1.8). The scaling factor used by \\citet{malreynolds} corrects for the flattening of a spectrum with $\\Gamma$ = 1.8 to $\\Gamma \\sim$1.1 (based on a single absorbed power-law model) for Mrk~231 due to reflection. The correction factor of 25, therefore, is a maximum correction factor. The minimum scaling factor is 10 to simply correct for a reflecting surface fraction of 10\\%.  % We will conservatively assume this minimum scaling factor of 10 for the reflection component. For the scattering component, we will assume the same scattering fraction as \\citet{malreynolds} (i.e. 1\\% for electron scattering).  After the corrections, the intrinsic 0.2--10~keV luminosity of the AGN in this scenario, where reflection dominates the observed flux, is $1.4 \\times 10^{43}$ ergs~s$^{-1}$ (from reflection) and $4.5 \\times 10^{43}$ ergs~s$^{-1}$ (from scattering). Therefore, the total reflection- and scattering-corrected luminosity in the 0.2--10~keV band is $5.9 \\times 10^{43}$~ergs~s$^{-1}$ if we assume reflected light dominates the observed spectrum.  Similarly, for the second scenario where the majority of the observed flux was scattered into the line of sight, we will assume the reflection component is 25\\% and the scattering component is 75\\% of the total observed flux.  This implies that the intrinsic 0.2--10~keV luminosity of the AGN is $4.5 \\times 10^{42}$ ergs~s$^{-1}$ (from reflection) and $1.4 \\times 10^{44}$ ergs~s$^{-1}$ (from scattering). Hence, the corrected 0.2--10~keV luminosity is $1.4 \\times 10^{44}$~ergs~s$^{-1}$.  It should be noted that in both cases we considered, the luminosity from the scattering portion dominated the total after the corrections.   Thus, with the assumptions made above, the intrinsic 0.2--10~keV luminosity of the AGN ranges from 0.6 to $1.4 \\times 10^{44}$ ergs~s$^{-1}$.  This range in luminosity overlaps with that of quasars ($\\sim$10$^{44}$~ergs~s$^{-1}$; e.g., Elvis et al. 1994; Piconcelli et al. 2005) and is similar to that of NGC~6240 ($\\sim$ 0.7 $-$ 2 $\\times$ 10$^{44}$ ergs s$^{-1}$, after correction for an HI column density of 1 $-$ 2 $\\times$ 10$^{24}$ cm$^{-2}$; Vignati et al. 1999).  % The ratio log[L$_{\\rm 2-10~keV}$/L$_{\\rm IR}$] for F04103--2838 corrected for scattering and reflection is --2.2 to --1.7.  These values fall precisely within the range found in radio-quiet PG quasars (--3 to --1; Sanders et al. 1989). Assuming F04103--2838 has the same X-ray to bolometric luminosity ratio as radio-quiet QSOs \\citep[L$_{\\rm x}$/L$_{\\rm bol} \\sim$3\\%]{elvis}, the AGN contribution to the bolometric luminosity of F04103--2838 is $\\sim$15--38\\%.  Therefore, within the large uncertainties, the AGN in F04103--2838 does not dominate the total energy output of the galaxy.    The results from our analysis of the {\\it XMM-Newton}\\ spectrum of the 1-Jy ULIRG/LINER F04103$-$2838 can be summarized as follows:  \\begin{enumerate}  \\item The soft (0.2--2~keV) X-ray flux of F04103$-$2838 is attributed to hot gas with kT $\\sim$0.1~keV.  This temperature is similar to that derived in other starburst galaxies and LINERs.  % The electron density in F04103--2838 % is $\\sim$0.6--5.9~cm$^{-3}$, consistent with theoretical predictions and observational estimates in wind systems.  \\item An \\fka line located at $\\sim$ 6.4 keV with an equivalent width of $\\sim$1.6~keV is detected in F04103$-$2838.  The line could be intrinsically broad or could be made up of two narrow lines located at rest frame energies of $\\sim$ 6.3 and 6.7~keV % but this decomposition is only significant at the $\\sim$60\\% level, so it needs to be verified with higher resolution spectra.  \\item The large equivalent width of the \\fka line suggests that the AGN is Compton-thick. Using simple assumptions, we estimate that the intrinsic 0.2--10~keV luminosity of this AGN is 0.6 -- 1.4 $\\times 10^{44}$ ergs~s$^{-1}$.  If these assumptions are correct and the galaxy has a QSO-like X-ray to bolometric luminosity ratio, the AGN detected by our observations does not dominate the bolometric luminosity of F04103$-$2838.  \\item The X-ray spectral characteristics of F04103$-$2838 are strikingly similar to those of the local luminous infrared galaxy NGC~6240. Given the similarities in X-ray properties but differences in merger state and in infrared color and luminosity, objects like NGC~6240 could conceivably be the precursors of ULIRGs like F04103$-$2838.  \\end{enumerate}"
    },
    "0711/0711.0658_arXiv.txt": {
        "abstract": "We study the mechanism of particle production in the world-volume of a probe $\\overline{D6}$-brane (or $D6$ with SUSY breaking) moving in the background created by a fixed stack of $D6$-branes. We show that this may occur in a regime of parametric resonance when the probe's motion is non-relativistic and it moves at large distances from the source branes in low eccentricity orbits. This leads to an exponential growth of the particle number in the probe's world-volume and constitutes an effective mechanism for producing very massive particles. We also analyze the evolution of this system in an expanding universe and how this affects the development of the parametric resonance. We discuss the effects of transverse space compactification on the probe's motion, showing that it leads to the creation of angular momentum in a similar way to the Affleck-Dine mechanism for baryogenesis. Finally, we describe possible final states of the system and their potential relevance to cosmology. ",
        "introduction": "One of the most important aspects of string theory with potential applications to beyond-the-standard-model physics and cosmology is the existence of additional extended objects.  A particularly interesting class of these, $Dp$-branes, occur as $(p+1)$-dimensional surfaces arising as supersymmetric solutions of the low-energy supergravity field equations but have also a high-energy interpretation as surfaces confining the endpoints of open strings \\cite{Polchinski}.   By now many different configurations of $Dp$-branes have been studied, including parallel and intersecting sets of branes \\cite{Berkooz,SMIntersectingBranes}, the combination of branes with different dimensionalities \\cite{Douglas} and also several possible embeddings in the compactified extra-dimensions \\cite{embedding}. These models have an enormous potential in the construction of gauge theories with matter and several new extensions of the Standard Model have been obtained in this way \\cite{SMIntersectingBranes}.  Branes have also been intensively studied in their application to the problems of early universe cosmology, such as inflation and reheating. In the context of effective field theory treatments of brane-world model building, an immense amount of work has been performed on the astrophysical and cosmological implications of these theories for both flat and warped extra dimensions \\cite{earlybranecosmo}, while many studies in string theory have also been undertaken \\cite{Quevedo,Martineau,braneinflation}. One important property of $Dp$-branes is the fact that they interact at large distances via the exchange of closed string modes, which allows one to study them as massive point-like charged objects moving along the dimensions transverse to their world-volumes.  An interesting application of this was proposed in \\cite{Burgess}, where it was shown that branes with opposite charges may form bound states similar to electron-positron states, despite the different origin of the central potentials involved. This analogy has motivated the name of \\textit{branonium} to designate this type of systems. Several properties of these bound states have been analyzed in \\cite{Burgess}, including the classical orbits, the associated quantum dynamics and possible compactification schemes. This analysis assumed a probe (anti)brane moving in the background spacetime created by a fixed central stack of branes, thus neglecting the gravitational backreaction of the probe. In the case of $D6$-branes, the probe's trajectory follows closed elliptical orbits similar to those found in planetary systems. It was also suggested that the probe's motion would lead to radiation of particles into both the bulk space and the probe's world-volume, making otherwise stable orbits decay. The stability of this system was further studied in \\cite{Lukas}, where it was shown that orbital decay necessarily occurs when geometric moduli of the background spacetime are allowed to vary.  In this work, we explore some properties of branonium systems, namely the mechanism leading to particle production in the world-volume of the probe brane. For the particular case of a scalar field living in the world-volume of a probe $D6$-brane, we show that particle production may occur in a regime of parametric resonance similar to the one behind the phenomenon of preheating after inflation \\cite{preheating,Kofman}, leading to an exponential growth of the particle number in the probe's world-volume.\\footnote{Discussions of preheating in brane systems other than branonium can be found in \\cite{otherbranepreheating}.} We determine how the orbital parameters of the system constrain this mechanism and analyze the possible properties of the resonantly produced particles. Also, we study how the probe's trajectory is modified by compactifying three of the dimensions parallel to the branes, so that a 4-dimensional spacetime with a flat Friedmann-Robertson-Walker geometry is obtained. In this case, we analyze the effects of Hubble expansion on the orbital parameters and how it modifies the parametric resonance regime. Other instability sources such as radiation into bulk modes are discussed as well. We also consider the effects of transverse space compactification on the probe's motion in the case of a compact 3-torus and analyze the associated production of angular momentum. Finally, we explore some of the cosmological applications of the resonant particle production mechanism and relate them to possible final configurations of the system, including stable states or brane-antibrane annihilation.  In the next section, we begin by introducing some of the general properties of $Dp-\\overline{Dp}$ branonia, with particular emphasis on the $p=6$ case, and discuss the effects of Hubble expansion on the orbital parameters. In Section III, we study the evolution of scalar particle-modes confined to the probe brane and analyze the properties of the associated parametric resonance mechanism in both the non-expanding and expanding universe cases. We distinguish between the production of massless and massive particles and determine the significance of energy damping into these modes on the probe's motion. We also consider, in this section, the effects of radiation into bulk modes on the resonant production of brane-bound particles. In Section IV, we analyze how compactifying the directions transverse to the branes modifies the probe's trajectory and dynamically generates the required angular momentum. Possible final states of the system are described in section V, where we also analyze the relevance of the particle production mechanism to cosmology in some particular scenarios. We summarize the main results and conclusions of this work in section VI. ",
        "conclusions": "In this work we have analyzed the mechanism of particle production in $D6-\\overline{D6}$ branonium systems, as well as in $D6-D6$ systems with SUSY breaking. We have shown that, in the limit of large distances and small velocities and for small eccentricity orbits, a parametric resonance will develop, producing scalar particles confined to the probe brane's world-volume. Massless particles are produced in a narrow resonance band centered at the comoving momentum corresponding to one half of the probe's orbital frequency $\\Omega$, as if resulting from the direct decay of the particles associated with the interbrane distance field. The strength of the resonance is given by the exponent $\\mu_k$ defined in Eq. (\\ref{eq72}). Massive particles are produced in the same way, with energies of order $\\Omega/2$. These energies are small compared to the string energy scale but, as the latter should be close to the Planck scale, very massive particles can be produced by the resonance mechanism, according to Eq. (\\ref{eq84}). For both massless and massive particles, the associated particle number grows exponentially in time and a large number of particles can be produced by this mechanism, just as in the case of preheating after inflation \\cite{preheating,Kofman}. However, the resonance surely requires the probe to complete a large number of orbits before the particle-modes can be significantly amplified, due to the small value of the exponent $\\mu_k\\simeq{q\\over2}\\ll1$.  Hubble expansion of the 4-dimensional effective flat FRW spacetime makes the probe's orbit decay, redshifting its energy and angular momentum. This alters the development of the parametric resonance significantly if the probe's motion does not occur at sufficiently late times. If, however, its motion begins late in the history of the universe, at least compared to typical string times, the resonance will still develop with the associated resonance band being shifted towards higher momentum modes. Each mode will be excited during the finite period it spends inside the resonance band, after which the associated particle number becomes an adiabatic invariant. A significant number of particles, given appproximately by Eq. (\\ref{eq101}), can be produced if this period is sufficiently long for the resonance to develop, according to the condition $q_0^2\\gsim\\gamma$, where $\\gamma$ quantifies the effects of the expansion and was defined in Eq. (\\ref{eq93}).  We have also concluded that a large number of brane-bound particles can be produced without affecting the probe's orbital motion, as can be seen in Eq. (\\ref{eq112}). Radiation into bulk closed string modes will also have a negligible effect if the probe moves at sufficiently large distances, as assumed for the validity of our study.  A realistic implementation of this mechanism requires particular initial configurations with large interbrane distances, non-relativistic velocities and almost circular orbits. We have shown that angular momentum creation may result from the effects of compactification of the directions transverse to the branes, illustrating this for the case of a compact 3-torus, where low eccentricity orbits can be produced for some range of initial conditions.  In this work, we have also discussed how, for a probe $D6$-brane, an interbrane potential may arise from supersymmetry breaking, if e.g. the relevant RR-form and the dilaton, $\\Phi$, gain a mass in the process. Stabilization of the probe is then possible at late times and at distances of order $m_{RR}^{-1}$, if $m_{\\Phi}\\gg m_{RR}$. Such a stabilization avoids annihilation with the central stack, as necessarily happens for a probe antibrane, making the parametric resonance regime potentially relevant for baryon number or dark matter generation, for example.  In our analysis, the parametric resonance arises only for $D6-\\overline{D6}$ or $D6-D6$ systems, where the $1/r$ potential allows the probe to move in closed elliptical orbits in a non-expanding universe. It is crucial for the development of the resonance that the interbrane distance field $\\rho$ exhibits an oscillating behavior, although the orbits do not need to close, as we have seen for the case of an expanding universe. Thus, other $Dp-\\overline{Dp}$ or $Dp-Dp$ branonium systems with $p<6$, where the probe trajectories are necessarily unbounded from below, cannot exhibit such resonant particle production. Additional stabilization potentials may, however, provide the required oscillatory behavior, so that it may be possible to find a resonant particle production mechanism in other cases.  It is nevertheless clear that resonant particle production in branonium systems may play an important role in our universe's evolution and we hope with this work to motivate future exploitation of their properties and applications.    \\vskip 0.2cm      \\centerline{\\bf {Acknowledgements}} \\vskip 0.2cm The authors would like to thank Graham Ross, Martin Schvellinger, Andre Lukas and Yang-Hui He for useful discussions. JGR is supported by FCT (Portugal) under the grant SFRH/BD/23036/2005. JMR is partially supported by the EC Network $6^{th}$ Framework Programme Research and Training Network ``Quest for Unification\" (MRTN-CT-2004-503369). This work was partially supported by the EU FP6 Marie Curie Research and Training Network ``UniverseNet\" (MRTN-CT-2006-035863).  \\vskip 0.5cm  \\vfill  \\appendix"
    },
    "0711/0711.1142_arXiv.txt": {
        "abstract": "We report high-resolution radio imaging of the recurrent nova RS Ophiuchi (RS Oph) during the first month of the 2006 outburst, using the Very Long Baseline Array (VLBA).  Observations made on days 20.8 and 26.8 of the outburst show a synchrotron-emitting partial shell that is much brighter to the east than to the west.  Assuming the broad component of the infrared lines corresponds to the outermost part of the shell seen by the VLBA, the distance to the source is $2.45\\pm0.4$\\,kpc. The circular shape and spectral indices of the shell emission challenge simple models for the radio structure immediately after the outburst.  The second epoch also shows an additional, resolved, synchrotron-emitting component well to the east of the shell.  Its inferred velocity is comparable to the escape speed from the surface of a high-mass white dwarf.  This component was not seen in the first epoch.  Its appearance may be related to the outflow reaching the edge of the nebula created by the red giant wind, which had been re-filling the void left by the last outburst in 1985.  This eastern component is likely related to the jets previously seen in this and other symbiotic stars, and represents the earliest clear detection of such a jet, as well as the best case yet for synchrotron emission from a white dwarf jet. ",
        "introduction": "The most recent major outburst of the symbiotic star and recurrent nova RS Ophiuchi (\\object{RS Oph}) was discovered near the optical peak on 2006 February 12.83 \\citep{Nar06}, hereafter taken as time zero. RS Oph is a binary system in which a red giant deposits material through its wind onto a white dwarf (WD) companion \\citep[see, e.g.,][]{ken86}.  Approximately every 20 years, the system erupts as a nova, as 'runaway' thermonuclear burning in the hydrogen-rich white-dwarf envelope leads to a visual brightness increase of more than 6 mag in less than 2 days, a sustained period with WD luminosity near the Eddington limit, and the ejection of material from the WD surface \\citep[e.g.,][]{Bod87}. The last such eruption was in 1985. The frequency of these outbursts suggests that the mass of the WD in RS Oph is very close to the Chandrasekhar limit \\citep[e.g.,][]{Hac01,Sok06e}, making this a prime candidate for a Type Ia supernova progenitor system.  Whereas most novae detected in the radio emit via thermal free-free bremsstrahlung from the shell of ejecta \\citep[e.g.,][]{Sea89}, post-outburst radio emission from RS Oph instead emanates primarily from the dense circumstellar material (CSM) that is swept up and shocked by the fairly low-mass ejected shell.  Observations after the previous outburst in 1985 with the European Very Long Baseline Interferometry Network (EVN), the Multi-Element Radio Linked Interferometer Network (MERLIN), and the Very Large Array (VLA) suggested that the radio emission was initially non-thermal at low frequencies, with a possible thermal component becoming dominant later in the outburst. The radio source consisted of a resolved central component plus possible bi-polar jets expanding at 1.5 mas/day \\citep{pad85,Hje86,Tay89,dav87,por87}.  Although jets have been observed from an increasing number of accreting white dwarfs in symbiotic stars and supersoft X-ray sources \\citep{Sok04}, and even several novae \\citep{Kaw06,Iij03}, these jets are generally not spatially well resolved, and basic questions such as the relationship between jet production and outbursts remain.  Most jets from white dwarfs have velocities of hundreds to thousands of km s$^{-1}$, and when jet emission can be spatially isolated, it usually appears to be thermal \\citep[e.g.,][]{Hol97,Kel01,Bro04,Gal04}.  In a few cases, however, observations have hinted that particles can be accelerated sufficiently that non-thermal radio and/or X-ray emission can also be produced \\citep{Cro01,Nic07}.  Here we report the detection of an expanding, synchrotron-emitting ring just weeks after the nova explosion, and, between three and four weeks after the outburst, the appearance of a synchrotron-emitting feature to the east of this ring aligned with the putative jet from 1985. We describe our observations of the expanding ring and the eastern component in \\S2.  We discuss these observations in more detail in \\S3, and summarize our conclusions in \\S4. ",
        "conclusions": "Dual-frequency VLBA observations made 20.8 and 26.8 days after the February 2006 outburst of the recurrent nova RS Oph have shown that the bulk of the radio emission at this stage comes from an expanding synchrotron shell.  The shell's increasing angular size matches the shock velocities inferred from X-ray observations and measured through infrared spectroscopy, confirming the basic picture of a decelerating shock between the material flung off in the explosion from the surface of the white dwarf and the circumstellar material produced by the wind of its red giant companion.  This gives a direct distance estimate of $\\sim2.45\\pm0.05\\pm0.37\\rm\\,kpc$, the most accurate yet determined for this source.  This evolving shock is far from simple.  The shell evolution is {\\it not} self-similar, with the shell thickness increasing from $\\sim15\\%$ to $\\sim26\\%$ between the two epochs.  In addition, the brightness distribution around the shell is highly asymmetric, with the eastern side many times brighter than the western side.  The eastern side of the ring also has a much flatter radio spectrum.  To our knowledge, such a spectrum has been seen only once before for a high-brightness resolved source, in the radio images of the expanding shell associated with the peculiar X-ray outburst of the emission-line star CI\\,Cam \\citep{Mio04}. Finally, the flux density of the western part of the shell is {\\it increasing} with time, in contrast to the optically-thin behavior of the ejecta-driven shocks in physically similar systems such as supernovae \\citep[e.g.,][]{Wei02}.   Our VLBA images also reveal an additional eastern component that is well separated from the dominant ring.  `Jets' (generally just elongations or separate components in barely-resolved sources) have been seen in previous outbursts of this and other symbiotic stars.  The VLBI observations of RS Oph constitute the first unambiguous evidence for a synchrotron jet emanating from a white dwarf system.  The inferred speed is close to the escape speed from a high-mass white dwarf, in accord with the generally good agreement between directed outflow and escape speeds for physical systems ranging from low-mass protostars to black hole binaries and active galactic nuclei\\footnote{In supernovae and gamma-ray bursts, the relevant escape speeds are those of the compact objects which are formed, not those of the progenitor stars.}. Similarly, synchrotron emission is characteristic of fast jets, where emission far from the powering source is generally assumed to derive from an interaction between the flow and the surrounding medium.  The association of a highly directed and a quasi-spherical outflow is more unusual, but certainly not unprecedented; for instance, the jet model for gamma-ray bursts, and their association with core-collapse supernovae \\citep[e.g.,][]{Mac99}, has brought the physical connection between quasi-spherical and highly directed outbursts into much greater prominence.  The eastern component that was so prominent on day 27 appeared to ``turn on'' around day 21. Although this turn-on could be due to decreasing absorption by the circumstellar medium as the eastern component moved outwards, there are several difficulties with this picture.  Intrinsic brightening, possibly associated with the component reaching the edge of the region re-filled by the red giant wind since the last nova outburst in 1985, may also play a role in the turn-on.   The preferred east-west axis for the observed asymmetries is consistent with VLBI observations on day 77 of the last (1985) outburst, which showed a strong east/west extension \\citep{Tay89}. The persistent east-west asymmetry suggests that this preferred axis reflects an intrinsic asymmetry in the binary system, rather than (for instance) some statistical irregularity in a particular explosive event.  The binary orbit provides one possible source of the asymmetry. Although a highly anisotropic density structure around the binary could collimate the outflow, with the `jet' naturally focused along the axis of least resistance \\citep[cf.][]{Llo93}, one would then expect the observed ring to correspond to the high-density `waist' of this structure.  The ring is in fact circular in projection and brightest along the trajectory of the eastern component. Whatever the collimation mechanism in RS Oph, we may for the first time be seeing the formation of a white-dwarf jet, as well as the most compelling demonstration to date of both its high initial speed and the physical emission mechanism (synchrotron emission)."
    },
    "0711/0711.3976_arXiv.txt": {
        "abstract": "Utilizing spatially resolved VLT/FORS spectroscopy and HST/\\-ACS imaging, we constructed a sample of over 200 field spiral galaxies at redshifts $0.1 < z < 1.0$. We find that the ratio between stellar and total mass remains roughly constant over the observed epochs, in compliance with the framework of hierarchical structure growth. However, the stellar mass--to--light ratios evolve more strongly in low--mass spirals than in high--mass spirals, indicating an anti-hierarchical evolution of their stellar populations (aka ``down-sizing''). ",
        "introduction": "Observations of the properties of distant galaxies at various cosmic epochs are a powerful tool to test the predictions of cosmological simulations in the framework of the hierarchical Cold Dark Matter paradigm. Combining high-resolution HST/ACS imaging and deep VLT/FORS spectroscopy and imaging, we have observed a sample of 202 disk galaxies at redshifts $0.1<z<1.0$ that represent a mean look-back time of $\\sim$\\,5 Gyr. Such a data set allows - via a comparison to local reference samples - to study the evolution of fundamental parameters of galaxies, like luminosity, size, mass, $M/L$ ratio etc., as a function of cosmic time. By applying models that fully account for observational effects like seeing and the influence of the slit width, we were able to extract spatially resolved rotation curves from the spectra and derive the galaxies' maximum rotation velocities as well as the total masses for 124 galaxies in our data set. ",
        "conclusions": ""
    },
    "0711/0711.4598_arXiv.txt": {
        "abstract": "We probe the evolution of globular clusters that could form in giant molecular clouds within high-redshift galaxies.  Numerical simulations demonstrate that the large and dense enough gas clouds assemble naturally in current hierarchical models of galaxy formation.  These clouds are enriched with heavy elements from earlier stars and could produce star clusters in a similar way to nearby molecular clouds. The masses and sizes of the model clusters are in excellent agreement with the observations of young massive clusters.  Do these model clusters evolve into globular clusters that we see in our and external galaxies?  In order to study their dynamical evolution, we calculate the orbits of model clusters using the outputs of the cosmological simulation of a Milky Way-sized galaxy.  We find that at present the orbits are isotropic in the inner 50 kpc of the Galaxy and preferentially radial at larger distances.  All clusters located outside 10 kpc from the center formed in the now-disrupted satellite galaxies.  The spatial distribution of model clusters is spheroidal, with a power-law density profile consistent with observations.  The combination of two-body scattering, tidal shocks, and stellar evolution results in the evolution of the cluster mass function from an initial power law to the observed log-normal distribution. ",
        "introduction": "The outcomes of many proposed models of globular cluster formation depend largely on the assumed initial conditions.  The collapse of the first cosmological $10^6\\ \\Msun$ gas clouds, or the fragmentation of cold clouds in hot galactic corona gas, or the agglomeration of pressurized clouds in mergers of spiral galaxies could all, in principle, produce globular clusters, but only if those conditions realized in nature.  Similarly, while observational evidence strongly suggests that all stars and star clusters form in molecular clouds, the initial conditions for cloud fragmentation are a major uncertainty of star formation models.   \\begin{figure} \\centering \\includegraphics[width=3.1in]{gnedin_fig1.ps} \\caption{A massive gaseous disk with prominent spiral arms, seen face-on at redshift $z=4$ in a process of active merging.  The gas density is projected over a $3.5$~kpc slice.  In our model star clusters form in giant gas clouds, shown by circles with the sizes corresponding to the cluster masses.  From Kravtsov \\& Gnedin (2005).} \\label{fig:gc} \\end{figure}   The only information that we actually have about the initial conditions comes from the early universe, when primordial density fluctuations set the seeds for structure formation.  These fluctuations are probed directly by the anisotropies of the cosmic microwave background radiation.  Cosmological numerical simulations study the growth of these fluctuations via gravitational instability, in order to understand the formation of galaxies and all other structures in the Universe.  The simulations begin with tiny deviations from the Hubble flow, whose amplitudes are set by the measured power spectrum of the primordial fluctuations while the phases are assigned randomly.  Therefore, each particular simulation provides only a statistical description of a representative part of the Universe, although current models successfully reproduce major features of observed galaxies.  Kravtsov \\& Gnedin (2005, \\apj, 623, 650) attempted to construct a first self-consistent model of star cluster formation, using an ultrahigh-resolution gasdynamics cosmological simulation with the Adaptive Refinement Tree code.  They identified supergiant molecular clouds in high-redshift galaxies as the likely formation sites of globular clusters.  These clouds assemble during gas-rich mergers of progenitor galaxies, when the available gas forms a thin, cold, self-gravitating disk.  The disk develops strong spiral arms, which further fragment into separate molecular clouds located along the arms as beads on a string (see Fig. \\ref{fig:gc}).  In this model, clusters form in relatively massive galaxies, with the total mass $M_{\\rm host} > 10^{9}\\ \\Msun$, beginning at redshift $z \\approx 10$.  The mass and density of the molecular clouds increase with cosmic time, but the rate of galaxy mergers declines steadily. Therefore, the cluster formation efficiency peaks at a certain extended epoch, around $z \\approx 4$, when the Universe is only 1.5 Gyr old.  The host galaxies are massive enough for their molecular clouds to be shielded from the extragalactic UV radiation, so that globular cluster formation is unaffected by the reionization of cosmic hydrogen.  As a result of the mass-metallicity correlation of progenitor galaxies, clusters forming at the same epoch but in different-mass progenitors have different metallicities, ranging between $10^{-3}$ and $10^{-1}$ solar.  The mass function of model clusters is consistent with a power law $dN/dM \\propto M^{-\\alpha}$, where $\\alpha = 2.0 \\pm 0.1$, similar to the observations of nearby young star clusters.   \\begin{figure} \\centering \\includegraphics[height=2.25in]{gnedin_fig2a.ps} \\includegraphics[height=2.25in]{gnedin_fig2b.ps} \\includegraphics[height=2.25in]{gnedin_fig2c.ps} \\includegraphics[height=2.25in]{gnedin_fig2d.ps} \\caption{Spatial distribution of surviving model clusters in the Galactic frame.  Dashed circles are to illustrate the projected radii of 20, 50, and 150 kpc.  The number density profile ({\\it bottom right}) can be fit by a power law, $n(r)\\propto r^{-2.7}$. The distribution of model clusters is similar to that of surviving satellite halos ({\\it dashed line}) and smooth dark matter ({\\it dotted line}).  It is also consistent with the observed distribution of metal-poor globular clusters in the Galaxy ({\\it solid line}), plotted using data from the catalog of Harris (1996, \\aj, 112, 1487).} \\label{fig:density} \\end{figure}   \\begin{figure} \\centering \\includegraphics[height=2.25in]{gnedin_fig3a.ps} \\includegraphics[height=2.25in]{gnedin_fig3b.ps} \\caption{{\\it Left:} Average eccentricity distribution of the surviving model clusters.  {\\it Right:} Anisotropy parameter $\\beta$ as a function of radius.  Vertical errorbars represent the error of the mean for each radial bin, while horizontal errorbars show the range of the bin.  Horizontal dashed lines illustrate an isotropic ($\\beta=0$) and a purely radial ($\\beta=1$) orbital distributions.} \\label{fig:ecc} \\end{figure} ",
        "conclusions": ""
    },
    "0711/0711.4621_arXiv.txt": {
        "abstract": "The GLAST satellite mission will study the gamma ray sky with considerably greater exposure than its predecessor EGRET. In addition, it will be capable of measuring the arrival directions of gamma rays with much greater precision. These features each significantly enhance GLAST's potential for identifying gamma rays produced in the annihilations of dark matter particles. The combined use of spectral and angular information, however, is essential if the full sensitivity of GLAST to dark matter is to be exploited. In this paper, we discuss the separation of dark matter annihilation products from astrophysical backgrounds, focusing on the Galactic Center region, and perform a forecast for such an analysis. We consider both point-like and diffuse astrophysical backgrounds and model them using a point-spread-function for GLAST.  While the results of our study depend on the specific characteristics of the dark matter signal and astrophysical backgrounds, we find that in many scenarios it is possible to successfully identify dark matter annihilation radiation, even in the presence of significant astrophysical backgrounds. ",
        "introduction": "Dark matter constitutes an essential element of modern cosmology. Evidence of its existence come from a wide variety of observations including the rotational speeds of galaxies~\\cite{rotationcurves}, the orbital velocities of galaxies within clusters~\\cite{clusters}, gravitational lensing~\\cite{lensing}, the cosmic microwave background~\\cite{wmap}, the light element abundances~\\cite{bbn} and large scale structure~\\cite{lss}. But despite these many observational indications of dark matter, we remain ignorant of the particle nature of this substance. To reveal the identity of dark matter, it will be crucial to measure its non-gravitational couplings. Efforts in this direction include direct detection experiments, which are designed to observe the elastic scattering of dark matter particles with nuclei, and indirect detection experiments, which aim to detect the annihilation products of dark matter. Such annihilations could potentially produce observable fluxes of gamma rays, neutrinos~\\cite{neutrinos}, positrons~\\cite{positron}, antiprotons~\\cite{antiproton}, antideuterons~\\cite{antideu}, synchrotron radiation~\\cite{syn} and X-rays~\\cite{Bergstrom:2006ny}. Of these various channels, gamma rays have the important advantage of retaining directional information. This is a feature that can be used, together with the peculiar energy spectra expected from dark matter annihilations, to disentangle dark matter annihilation products from astrophysical backgrounds~\\cite{angular}.  In this paper, we study the ability of gamma ray experiments to identify dark matter annihilation radiation from the Galactic Center region (for earlier work on this subject, see Refs.~\\cite{gchist}) by using both spectral and angular information. This is a goal of existing Atmospheric Cerenkov Telescopes (ACTs) including HESS~\\cite{HESSurl} , MAGIC~\\cite{MAGICurl}, VERITAS~\\cite{VERITASurl} and CANGAROO-III~\\cite{CANGAROOIIIurl}, as well as of the forthcoming satellite mission, GLAST~\\cite{GLASTurl}. Motivated by its sensitivity to gamma rays over an energy range well suited to dark matter searches (0.1-300 GeV), we focus here on the GLAST experiment.   The remainder of this paper is structured as follows: In Secs.~\\ref{backgrounds} and~\\ref{darkmatter} we describe the model used for the point-like and diffuse astrophysical backgrounds and the expected characteristics of a dark matter annihilation signal, respectively. In Sec.~\\ref{results} we describe our analysis method and assess the ability of GLAST to either place limits on or identify dark matter annihilations in the Galactic Center region. In Sec.~\\ref{measure} we extend our method to assess GLAST's ability to measure the annihilation cross section, mass and distribution of dark matter. In Sec.~\\ref{conclusions} we summarize and present our conclusions. ",
        "conclusions": "The most challenging task for indirect dark matter searches is not detecting the products of dark matter annihilations, but in confidently identifying those particles as such. In particular, any signal must be separated from astrophysical backgrounds if it is to be reliably claimed to be a detection of dark matter annihilation products. This is certainly true in the case of gamma ray telescopes hoping to observe dark matter annihilations in the region of the Galactic Center, where astrophysical backgrounds are especially foreboding.  In this paper, we have studied quantitatively the ability of GLAST to identify gamma rays from dark matter annihilations in the Galactic Center region, using all information encoded in the spectrum and angular distribution of the observed events. Due to the complex nature of the Galactic Center's backgrounds and dark matter distribution, it is important for GLAST to make full use of this information to best separate dark matter annihilation products from any astrophysical backgrounds which are present.  When GLAST begins its mission in 2008, it will detect a number of astrophysical sources in the region of the sky around the Galactic Center. These include the point sources identified by HESS and EGRET, and perhaps other, thus far unknown, sources. A diffuse gamma ray background will also likely be present. In this article, we have modeled these backgrounds based on known spectral information, and using a point-spread-function for GLAST, in the hope of identifying statistically significant departures from this model resulting from dark matter annihilations.  The spectrum of gamma rays from dark matter annihilations is quite distinctive, and does not resemble the power-law form observed from typically astrophysical sources. Furthermore, the angular distribution of dark matter annihilation products is not a simple point source, nor is it isotropic or trace the Galactic Ridge, as the backgrounds are expected to. By exploiting these differences, it may be possible to  identify the products of dark matter annihilation observed from the inner galaxy, even they are a subdominant component of the total emission.   If dark matter annihilation products are identified by GLAST, then it may also be possible to measure or constrain the properties of dark matter, including its mass, annihilation cross section and spatial distribution.  We find that GLAST is unlikely to determine the WIMP's mass with high precision, however. For example, for the case of a 100 GeV WIMP with an annihilation cross section of $3 \\times 10^{-26}$ cm$^3$/s and distributed with an NFW halo profile, the mass could be determined to lie within approximately 50-300 GeV. In the same benchmark model, the inner slope of the dark matter halo profile could be determined to $\\sim 10\\%$ precision by GLAST.  {\\it Note added:} After we had completed this study, we became aware of updated estimates for the performance of GLAST~\\cite{lat}. In particular, the point spread function that we have used is a factor $3 \\div 4$ more optimistic than was found in the current estimate. As a consequence, a reduction of the ability of GLAST to identify dark matter and to measure its properties is expected, but not to an overwhelming extent: we checked for example that for a NFW profile and a DM mass of 100 GeV the bound worsens by less than 30\\%. Also, additional effects not taken into account here (like effective area degradation due to orbit inclination and spacecraft rocking; non-negligible dead-time and a high-particle-background associated to orbital passage over the South-Atlantic Anomaly) may reduce the exposure by up to a factor $\\sim 2$ compared to the present estimate. We would like to thank the referee for bringing this to our attention.    \\bigskip {\\bf Acknowledgments} This work has been supported by the US Department of Energy and by NASA grant NAG5-10842. Fermilab is operated by Fermi Research Alliance, LLC under Contract No.~DE-AC02-07CH11359 with the United States Department of Energy."
    },
    "0711/0711.3237_arXiv.txt": {
        "abstract": "We investigate the signal from supernova relic neutrinos in future large scale observatories, such as MEMPHYS (UNO, Hyper-K), LENA and GLACIER, at present under study. We discuss that complementary information might be gained from the observation of supernova relic electron anti-neutrinos and neutrinos using the scattering on protons on one hand, and on nuclei such as oxygen, carbon or argon on the other hand. When determining the relic neutrino fluxes we also include, for the first time, the coupling of the neutrino magnetic moment to magnetic fields within the core-collapse supernova. We present numerical results on both the relic $\\nu_e$ and $\\bar{\\nu}_e$  fluxes and on the number of events for $\\nu_e + ^{12}$C, $\\nu_e + ^{16}$O, $\\nu_e + ^{40}$Ar and $\\bar{\\nu}_e +$p for various oscillation scenarios. The observation of supernova relic neutrinos might provide us with unique information on core-collapse supernova explosions, on the star formation history and on neutrino properties, that still remain unknown. ",
        "introduction": "\\noindent The observation of neutrinos from past core-collapse supernovae would represent an exciting and important discovery. Supernova relic neutrinos (SRN) bring crucial information on the explosion phenomenon, on the star formation rate and on neutrino properties. The understanding of how massive stars explode constitutes one of the main future challenges in astrophysics. While the various stages of the stellar evolution are quite well established, it still remains a mistery how, in these stars, the shock wave formed in the iron core, after collapse, manages to eject the mantle. Essentially all the gravitational energy is emitted in a short burst of neutrinos of all flavours. This luminosity curve follows closely the collapse, accretion phase and the cooling of the proto-neutron star left \\cite{Raffelt:wa}.  The first time such neutrinos have been observed is from the SN1987A located in the Large Magellanic Cloud, thanks to the IMB \\cite{Bionta87}, Baksan \\cite{Alexeyev87} and Kamiokande detectors \\cite{Hirata87}. The few electron anti-neutrino events measured have roughly confirmed the theoretical expectations, even though their energy and angular distributions still have unsolved problems. For the SRN fluxes, the best upper limits at 90 $\\%$ C.L. of 6.8 $\\times$ 10$^3$ $\\nu_e$ cm$^{-2}$s$^{-1}$  (25 MeV $< E_{\\nu_e} <50$ MeV) and 1.2 $\\bar{\\nu}_{e}$ cm$^{-2}$s$^{-1}$ ($E_{\\bar{\\nu}_e} > 19.3$ MeV) come from the LSD \\cite{Aglietta:1992yk} and the Super-Kamiokande detectors \\cite{Malek:2002ns}. The latter appears extremely close to most of theoretical predictions. Indeed it seems quite likely that the supernova relic neutrino background could be observed either by using an improved technology, e.g. the use of Gadolinium doped water \\v{C}erenkov detectors \\cite{Beacom:2003nk}, or by scaling the present technology at megaton sizes, like in MEMPHYS \\cite{deBellefon:2006vq}, UNO \\cite{uno}, and Hyper-K \\cite{itow}. These large-size observatories could provide us with the complementary information from a possible future (extra-)galactic explosion and from the SRN background.  Many authors have estimated the expected SRN flux using different approaches to determine the supernova rate \\cite{Malaney:1996ar,Hartmann:1997qe,Totani:1995dw,Kaplinghat:1999xi,Strigari:2003ig,Fukugita:2002qw,Wurm:2007cy} (see \\cite{Ando:2004hc} for a review). In fact, the latter can be deduced either from direct observations \\cite{Cappellaro:1996cc} or from star formation rates (see e.g. \\cite{Strigari:2005hu}). The effects of neutrino masses and mixings on the SRN flux have been included recently (see e.g. \\cite{Ando:2002ky}), \\cite{Ando:2002zj}). In \\cite{Lunardini:2005jf} a pragmatic approach has been used: to predict the SRN flux using available data, namely direct supernova rate measurements and supernova neutrino fluxes that give the best fit of the SN1987A events. Most of the available literature presents calculations for the $\\bar{\\nu}_e$ scattering on protons, which is the dominant detection channel in water \\v{C}erenkov (see e.g. \\cite{Ando:2004sb}) or scintillator detectors \\cite{Wurm:2007cy}. In particular, a study of what could be learned on the neutrino spectra and/or the core-collapse supernova rate at next generation water \\v{C}erenkov detectors is realized in \\cite{Lunardini:2006pd}. In Ref. \\cite{Cocco:2004ac} the potential of liquid argon detectors to observe supernova relic electron neutrinos is analyzed, while the specific case of $\\nu_e$ detection in SNO is considered in \\cite{Beacom:2005it,Lunardini:2006sn}. Ref. \\cite{Lunardini:2006sn} also obtains an upper limit on the relic $\\nu_e$ flux from the $\\bar{\\nu}_e$ upper limit in Super-Kamiokande. Finally, in \\cite{Ando:2003ie,Fogli:2004gy} the authors have considered the possibility of using the SRN to gather information on more exotic unknown neutrino properties such as neutrino decay.  The study presented in this paper has several new features compared to the available literature. First we study the possible signal from both the supernova relic $\\nu_e$ and $\\bar{\\nu}_e$ in future large-scale observatories. We consider the three technologies under investigation, water \\v{C}erenkov, scintillators and liquid argon, having in mind the proposed MEMPHYS  (UNO, Hyper-K), LENA and GLACIER detectors \\cite{Autiero:2007zj}. We also include for the first time in the neutrino propagation within the supernova the interaction with matter (the Mikheev-Smirnov-Wolfenstein effect \\cite{Wolfenstein:1977ue,Mikheev:1986wj}) as well as the coupling of the neutrino magnetic moment to the magnetic fields (the spin-flavour conversion phenomenon). Clearly our present understanding of supernovae explosion is still limited from several respects and, in particular, for the knowledge of the stellar magnetic fields and their evolution during the explosion. We think however that, even at the present stage, it is interesting to perform a first calculation including this process, to start investigating under which conditions some effects might be seen. We make here the assumption of a static supernova density profile, neglecting in this first work the changes induced by the shock wave. We also do not consider earth matter effects that have been discussed in \\cite{Ando:2002zj}. To determine the SRN flux, we integrate the neutrino spectra from one supernova explosion over the redshift by convoluting it with the supernova explosion rate. We predict the expected event rates and analyse the effects of the presence of a neutrino magnetic moment both on the relic number fluxes and on the number of events. Note that  we follow previous investigations for the choice of the neutrino fluxes at the neutrinosphere, of the supernova rate as well as for the discussion of the possible backgrounds in the detectors. The paper is organized as follows. The formalism is described in Section II. The results and conclusions are presented in Sections III and IV. ",
        "conclusions": "\\noindent In this work we have analysed both the relic supernova electron anti-neutrino and neutrino fluxes and the associated signals due to scattering on protons and nuclei respectively, in large-scale observatories such as MEMPHYS (UNO, Hyper-K), LENA or GLACIER. Our calculations of the neutrino propagation in the star include both the matter interaction and the resonant spin-flavour conversion. One of our goals has been to try to investigate under which conditions (profile, $\\mu_{ij}$ and/or $B_0$ values) RSF effects could be seen. In all cases studied, the results show that RSF phenomenon might increase/decrease significantly the number of events, for strong magnetic fields and/or large values of the neutrino magnetic moment. It is interesting to note that these effects are maximal for electron anti-neutrinos, when no resonant conversion is present due to matter efffects (normal hierarchy and a large or small value of $\\theta_{13}$) due to the RSF-$h$ resonance. In the case of electron neutrinos, the RSF tend to decrease the events, in most cases. Our results are also sensitive to the magnetic field profiles. To obtain thes results with the spin-flavour conversion, we have made the assumption of a static density supernova profile and also neglected possible effects of the shock wave on the stellar magnetic fields. Our present knowledge in this respect is in fact still too limited to include such aspects. From this point of view, our analysis should be taken as a first step in this direction and should be updated once progress is performed on the knowledge of the supernova magnetic fields and on the shock wave impact.  Our main goal has been to perform a numerical study of the relic fluxes and their signal in the next generation observatories. We have obtained that $\\nu_e$ fluxes are larger than the $\\bar{\\nu}_e$ up to a factor 2.5 depending on the oscillation scenarios (normal or inverted hierarchy and $\\theta_{13}$ value). We have shown that, after 10 years operation, between 100 and 250 $\\bar{\\nu}_e + p$ events are expected in the 19.3-27 MeV neutrino energy range for MEMPHYS, while 45-65 are expected in LENA in the 9.3-25 MeV energy range. We find that neutrino-carbon scattering will produce 2-3 events, while neutrino-oxygen will give 10-20 events and the neutrino-argon interaction will induce 100 to 200 events. (These number are obtained by taking the normalization factor for the star formation rate $f_{\\star}=1$, which implies a local formation rate of 0.7$\\times 10^{-2} h_{70} M_{\\odot}$ year$^{-1}$Mpc$^{-3}$, consistent with mildly dust-corrected UV data at low redshift). The possible measurement of both $\\nu_e$ and $\\bar{\\nu}_e$ could help us determining the neutrino fluxes at the neutrinosphere and/or still unknown neutrino properties, such as the hierarchy or the knowledge of the third neutrino mixing angle. This will depend on the progress made in supernova modelling or neutrino physics by the time such observatories might come into operation. In conclusion, future large-scale observatories based on water, scintillator or liquid argon have a great discovery potential for relic supernova neutrinos. Their observation would clearly provide us with essential information on core-collapse supernova explosions, the supernova explosion rate and important neutrino properties that still remain unknown."
    },
    "0711/0711.1232_arXiv.txt": {
        "abstract": "We present a kinematic analysis of the globular cluster (GC) system in M31. Adding new velocity data for 150 GCs in our wide-field survey to those in the literature, we increase the number of M31 GCs with measured velocities by 42 per cent, to 504. Using the photometric and spectroscopic database of these 504 GCs, we have investigated the kinematics of the M31 GC system. We find that the all GC system shows strong rotation, with rotation amplitude of $v_{\\rm rot}\\sim190$ km s$^{-1}$, and that a weak rotation persists even for the outermost samples at $|Y|\\ge 5$ kpc where $Y$ represents the projected distance from the major axis. The rotation-corrected velocity dispersion for the GC system is estimated to be $\\sigma_{p,r}\\sim130$ km s$^{-1}$, and it increases from $\\sigma_{p,r}\\sim120$ km s$^{-1}$ at  $|Y|< 1$ kpc to $\\sigma_{p,r}\\sim150$ km s$^{-1}$ at  $|Y|\\ge 5$ kpc. These results are very similar to those for the metal-poor GCs. This shows that there is a dynamically hot halo in M31 that is rotating but primarily pressure-supported. We have identified 50 ``friendless\" GCs, and they appear to rotate around the major axis of M31, unlike M31's disk rotation around the minor axis. For the subsamples of metal-poor and metal-rich GCs, we have found that the metal-rich GCs are more centrally concentrated than the metal-poor GCs, and both subsamples show strong rotation. For the subsamples of bright and faint GCs, it is found that the rotation for the faint GCs % is stronger than that for the bright GCs. % We have identified 56 GCs and GC candidates with X-ray detection including 39 genuine GCs with measured velocities. It is found that the majority of X-ray emitting GCs follow the disk rotation, and that the redder, more metal-rich, and brighter GCs are more likely to be detected as X-ray emitting GCs, as seen for GCs   in early-type galaxies. We have derived a rotation curve of M31 using the GCs at $|Y|\\le 0.6$ kpc, and it agrees well for the range of $R=20-45\\arcmin$ to that  based on other tracers except for the planetary nebulae. We have estimated the dynamical mass of M31 using `Projected Mass Estimator (PME)' and `Tracer Mass Estimator (TME)' as $M_{\\rm PME}= 5.5_{-0.3}^{+0.4} \\times 10^{11} M_\\odot$ out to a radius of $\\sim55$ kpc and $M_{\\rm TME}= 19.2_{-1.3}^{+1.4} \\times 10^{11} M_\\odot$ for a radius of $\\sim100$ kpc, respectively. We finally discuss the implication of these results and compare the kinematics of GCs with that of planetary nebulae in M31. ",
        "introduction": "Globular clusters (GCs) are the oldest stellar systems in the universe and thus maintain a fossil record of the early history of galaxies. They are much brighter than stars and the spatial distribution of the GCs is much more extended than that of the halo stars in galaxies. Therefore they are an excellent tool to investigate the kinematics of the main components (stellar and dark matter halo, bulge, and disk) of galaxies and to estimate the mass of their host galaxies.  M31 is a unique object as               the nearest spiral galaxy in the Local Group, providing an excellent place to study GCs. It has the largest GC system of any galaxy within several Mpc. Therefore, there are numerous studies of GCs in M31 since the first study by \\citet{hub32} (e.g., \\citep{van69,sar77,bat87,huc91,bar00,per02,mor04,bur04,hux05,bea05,puz05,gal06,kim07} and see extensive lists of earlier references in \\citet{hod92,van00}). In addition, there were  several surveys and studies of the M31 GCs with X-ray detection, finding that the X-ray emitting GCs are commonly associated with low-mass X-ray binaries \\citep{sup01,kong02,di02,kaa02,wil04,tru04,pie05,fan05,tru05,tru06}.  Kinematics is one of the notable features for the M31 GCs, and has long been a focus of study. It is known that the kinematics of the M31 GCs is much different from that of the Galactic GCs in the sense that the former is dominated by rotation, while the latter shows weak rotation \\citep{per02} %  The first detailed kinematic analysis of the M31 GC system was given by \\citet{har74} using radial velocities for 44 GCs in \\citet{van69}. They estimated the total mass of M31, M$_{\\rm tot}=3.4\\pm1.4\\times10^{11} M_\\odot$, and found that the metal-poor GCs have a larger velocity dispersion than the metal-rich GCs. \\citet{huc82} obtained digital, low-dispersion spectra of 61 GCs, and found that the GC system rotates with an amplitude of $160\\pm40$ km s$^{-1}$. They also found  a measurable but small metallicity gradient and no significant spectroscopic difference between the GCs with X-ray detection and those without X-ray detection. \\citet{fed90} using the velocity data of 81 GCs obtained a total mass of $3.2\\times 10^{11} M_\\odot$ for a radius of 16 kpc using a projected mass estimator. A comparison of the data with the model of \\citet{ric84} showed that the GCs have circular or isotropic orbits, and lead to a total mass of $\\sim 7\\times 10^{11} M_\\odot$ based on circular orbits. Later, \\citet{fed93} updated the total mass of M31 to be $5-8\\times 10^{11} M_\\odot$ for a radius of 30 kpc using velocity data for 176 GCs.  In the 1990's, metallicity information for a large sample of GCs ($N>100$) became available, and it was possible to investigate the kinematics of GC subsamples based on the metallicity. \\citet{huc91} estimated the radial velocities and metallicities for 150 GCs, and presented the detailed kinematics of the GC system using   expanded subsamples of GCs based on metallicity. They identified a weak metallicity gradient as a function of projected radius. They found that, in the inner region ($<2$ kpc), the metal-rich ([Fe/H]$\\ge-0.8$) GCs exhibit a rapid rotation of $100-200$ km s$^{-1}$, while the metal-poor GCs show no  significant rotation. However, the rotation of the two subsamples could  not be easily measured at larger radii. Later, \\citet{per02} showed, using higher precision data of over 200 GCs ( typical velocity measurement error of 12 km s$^{-1}$), that the entire M31 GC system is highly rotation-supported. This is in strong contrast to the case of the Milky Way where the GC system composed mostly of halo GCs is pressure-supported, although a minor system of the bulge GCs shows some rotation. They also found that the centrally concentrated metal-rich GCs appear not to be flattened, but are   analogous to the bulge population with a strong rotation, and that the less concentrated metal-poor GCs also show a rapid rotation. \\citet{gal06} obtained velocity data for 42 genuine GCs, and produced a large data set of 349 M31 GCs with up-to-date velocities by combining all information in the literature.  Recently, large datasets of velocities for M31 GCs and several deep photometric studies of the halo fields in M31 have provided important clues for understanding the evolution of M31. \\citet{bro03} presented very deep photometry of stars in a halo field along the minor axis of M31 at a projected distance of 11 kpc (corresponding to a deprojected distance of 53 kpc) and found that the M31 halo includes a major intermediate-age metal-rich population and a significant old metal-poor population. They suggested from this result that M31 may have undergone a major merger $6-8$ Gyr ago. Later \\citet{bro07} suggested, from photometry of another minor-axis halo field at a projected distance of 21 kpc, that the outer spheroid is dominated by the debris of earlier merging events that might have occurred more than 8 Gyr ago.  \\citet{mor04} suggested the existence of an old thin disk composed of GCs in M31, using the data for 321 GCs presented in \\citet{per02}. Since the GCs in the thin disk have a metallicity distribution from [Fe/H] $\\sim-2.2$ to above solar, and the metal-poor GCs are regarded as being as old as their Galactic counterparts, they concluded that the thin disk has remained undisturbed for $\\sim10$ Gyr or more. This is contradictory with the scenario of \\citet{bro03} that M31 may have undergone a major merger $6-8$ Gyr ago. Later \\citet{bea04,bea05} suggested that the majority of M31 thin disk GCs may be significantly younger than 10 Gyr, up to $\\sim 1$ Gyr and metal-rich ($-0.20\\le$[Fe/H]$\\le0.35$). The existence of young and metal-rich GCs in the disk is also confirmed in other studies as well \\citep{bur04,puz05}.  On the other hand \\citet{ash93} first tried to identify subgroups in the M31 GC system using the positional and velocity information for 144 GCs in \\citet{huc91}. They suggested that the clustering of GCs in position and velocity space may be a surviving signature of gaseous clumps from which the galaxy halo of M31 formed. Later, \\citet{per03}, using 301 GCs with measured velocity and metallicity in \\citet{per02}, searched for the subclustering of M31 GCs in the parameter space of projected radius, radial velocity, and metallicity. Interestingly \\citet{mor03} found a possible tidally distorted satellite of M31, Andromeda VIII, from the identification of a group of $5-12$ planetary nebulae (PNe), $1-3$ GCs, and two HI clouds. Recent kinematic studies of M31 streams \\citep{iba01,mcc04} started examining the GC groups including Andromeda VIII as a possible progenitor of the stream (e.g., \\citealt{fon06a,far06a}).  Although there were numerous studies to obtain the radial velocities and the metallicities for M31 GCs, there are still many GC candidates for which spectroscopic information is needed. For example, there are 569 GC candidates for which radial velocities are not available in a combined sample of M31 GCs in \\citet{gal06}. Recently we have performed a homogeneous photometric and spectroscopic, wide-field survey of GCs in M31 covering a field of $\\sim$ 3\\arcdeg $\\times$ 3\\arcdeg~ centered on M31. We presented the first result of our survey: a catalog of new GCs in M31 (\\citealt{kim07}, Paper I). This is the second in the series of papers on this survey.  In this paper, we present the results of a kinematic study of the GC system in M31 using the velocity data of 150 GCs newly obtained in our survey and 354 GCs previously known in the literature. Section 2 gives a brief description of the data used in this analysis, and section 3 investigates the kinematic properties for various subsamples of M31 GCs. We discuss primary results % in \\S 4. A summary is given in the final section.  We adopted a distance to M31  of $780$ kpc derived by \\citet{mcc05} based on the method of the tip of the red giant branch, which is similar to the value given by \\citet{fre90}, 770 kpc. At this distance one arcmin in the sky corresponds to 0.227 kpc. We adopted the center position of M31 as $\\alpha_0=00^{\\rm h}42^{\\rm m}44.^{\\rm s}3$, $\\delta_0=+41^\\circ16\\arcmin09\\arcsec$ (J2000; \\citealt{cra92}). Following the geometric transformation of \\citet{huc91}, we define the $X$ and $Y$ coordinates, which are the projected distances along the major and minor axis, respectively. $X$ and $Y$ increase towards the northeast and the northwest, respectively, and are given by, \\begin{eqnarray} x^\\prime=\\sin{(\\alpha - \\alpha_0)} \\cos{\\delta}, \\\\ y^\\prime=\\sin{\\delta}\\cos{\\delta_0} - \\cos{(\\alpha - \\alpha_0)} \\cos{\\delta}\\sin{\\delta_0},  \\\\ X = x^\\prime\\,\\sin{({\\rm PA})} + y^\\prime\\,\\cos{({\\rm PA})},  \\\\ Y = -x^\\prime\\,\\cos{({\\rm PA})} + y^\\prime\\,\\sin{({\\rm PA})}, \\end{eqnarray} where $x'$ and $y'$ are Cartesian coordinates for $\\alpha$ and $\\delta$ with a center at ($\\alpha_0$, $\\delta_0$), and PA$=37.7^\\circ$ is the position angle for the major axis \\citep{dev58}. ",
        "conclusions": "\\label{discussion}   \\subsection{A dominant bulge versus a wimpy halo in M31?}  The existence of a halo in M31 has been controversial. While \\citet{huc91} suggested that the metal-poor GCs in M31 show little evidence of rotation, indicating that they belong to the halo of M31, other studies based on surface photometry, counts of red giant branch stars, and counts of PNe along the minor axis of M31 found an extended bulge, but no evidence of any halo component in M31 \\citep{pri94,mou86,hur04,irw05,mer06}. However, recent studies covering a very wide field of M31 are finding evidence for a huge halo in M31 \\citep{irw05, guh05, cha06, kal06, gil06, iba07} as well as a vast extended disk reaching about 40 kpc from the center of M31 \\citep{iba05}.  \\citet{irw05} presented a minor-axis surface brightness profile reaching out to a projected radius of $4^\\circ$ ($\\approx 55$ kpc), combining surface photometry and faint red giant star counts. They found that the minor-axis profile can be fit by a de Vaucouleurs law with effective bulge radius  $R_{\\rm eff} = 0^\\circ.10$ (1.4 kpc) out to a projected radius of $1^\\circ.4$ ($\\approx 20$ kpc), and that there is a separate component beyond $1^\\circ.4$ that can be fit by either a power law of index about --2.3 or an exponential law of scale length 14 kpc. One thing interesting about this component is that it has the same color as the inner region, indicating that it has a similar metallicity to that of the bulge  (assuming the two components have the same age). In addition, the photometry of remote fields along the minor axis of M31 showed that most stars in these fields have unexpectedly high metallicity, [m/H]$\\sim -0.5$ ([Fe/H]$\\sim -0.8$) while there are some metal-poor stars as well \\citep{mou86,dur01}. Thus there are two components in the outer region of M31 that are distinguished in structure, but not in metallicity. We may say from the results based on stars that M31 has a huge dominant bulge reaching out to $4^\\circ$, or that M31 has two bulges: an inner bulge (out to $1^\\circ.4$) and an outer bulge (out to $4^\\circ$). Even the inner bulge is much larger than the central bulge in our Galaxy, and the outer bulge may be called a halo.  The progenitors of PNe are low to intermediate mass stars so that they are often considered as a good tracer for old stellar populations. For example, the radial number density profiles of PNe in M31 show excellent agreement with the $R$-band surface brightness profiles of the stellar bulge and disk \\citep{wal87,mer06}. In particular, the number density profile of PNe along the minor axis of M31 is consistent with the $R^{1/4}$ profile out to one degree \\citep{mer06, hur04}. \\citet{mer06} pointed out that there is  little excess of PNe over the  $R^{1/4}$ profile even at $\\sim$10 $R_{eff}$ (one degree), indicating no separate halo component. This is also consistent with the result based on the surface brightness profile along the minor axis \\citep{pri94, irw05}. Considering the rotating kinematics of the outer PNe at $|Y|>6$ kpc in M31 as well as the above points, \\citet{hur04} suggested that the bulge is the most dominating component in M31, while a nonrotating, pressure-supported halo is, if any, only a very minor component of M31.  In general, or at least in our Galaxy, metal-poor GCs are considered as a tracer of the halo, while metal-rich GCs belong to the bulge or thick disk like the PNe, although the bulge does have some metal-poor GCs \\citep{har96}. Note that the kinematics of 17 old star clusters in M33, irrespective of their metallicity, show evidence for the existence of a hot halo in M33 \\citep{cha02}. Therefore we may use metal-poor GCs to investigate the properties of any halo in M31. The kinematics of the metal-poor GCs along the major axis at $|Y|<1$ kpc show a rotation as strong as the metal-rich GCs. The rotation amplitude of the metal-poor GCs decreases with increasing distance from the major axis, and the metal-poor GCs show some rotation ($60\\sim100$ \\kms) even far from the major axis at $|Y|>5$ kpc. These results are consistent with those of \\citet{per02}, but not with \\citet{huc91}.  However, the rotation-corrected velocity dispersion for the metal-poor GCs is about twice as large as    the rotation amplitude in the outer region at $|Y|>5$ kpc. The rotation-corrected velocity dispersion for the metal-poor GCs at $|Y|>5$ kpc in M31, $\\sigma_{p,r}=145^{+11}_{-11}$ km s$^{-1}$, is comparable to  that of the metal-poor GCs in the Galactic halo at $R_{GC} >4$ kpc, $\\sigma =123$ km s$^{-1}$ given by \\citet{cot99}. Considering that the spatial distribution of the metal-poor GCs in M31 is much more extended than that of the metal-rich GCs and that the velocity dispersion of the metal-poor GCs in the outer region of M31 is about two times larger than the rotation amplitude, we conclude that there is a dynamically hot halo in M31 that is rotating but primarily pressure-supported. This is consistent with the discovery of a pressure-supported metal-poor stellar halo based on the observations of red giant branch stars at radii from 10 kpc to 165 kpc \\citep{cha06,gil06, kal06}. \\citet{cha06} found from the analysis of a large sample of red giant branch stars: (a) there is seen a stellar halo component with low metallicity ( [Fe/H]$\\sim -1.4$ ) at radii from 10 to 70 kpc, (b) there is little radial gradient of metallicity, (c) there is no significant evidence for rotation, and (d) the velocity dispersion of this metal-poor halo can be represented by $\\sigma_v (R_p ) = 152 - 0.90 R_p $ km s$^{-1}$ where $R_p$ is a projected galactocentric distance in units of kpc.  In Figure 25 we display the radial variation of the mean metallicity for the metal-poor GCs in comparison with that for the red giant branch stars given by \\citet{kal06}. A few features are noted in Figure 25: (a) the mean metallicity for the metal-poor GCs is much lower than that for the red giant branch stars at the same distance from the M31 center; (b) the mean metallicity for the metal-poor GCs in the inner region is similar to that for the red giant branch stars much further out at $R_p \\approx 80$ kpc that was considered to be the part of the stellar halo by \\citet{kal06}; and (c) the mean metallicity of the metal-poor GCs shows little radial dependence: [Fe/H]$=-1.48\\pm0.27$ at $R=7.1$ kpc to [Fe/H]$=-1.62\\pm0.25$ at $R=25.5$ kpc (if we use all GCs, we obtain [Fe/H]$=-1.24\\pm0.27$ at $R=6.7$ kpc to [Fe/H]$=-1.62\\pm0.25$ at $R=25.5$ kpc). These show that the metal-poor GCs are tracing the metal-poor halo, even in the inner region of M31.  Recently \\citet{iba07} found a faint, smooth and extremely extended halo component, reaching out to 150 kpc, about three times larger than the limit studied by \\citet{irw05}. In addition, they found, from the photometry of red giant branch stars, that the majority of the stars in this extended halo are metal-poor, in contrast to the \\citet{irw05} finding that its integrated color is similar to the inner bulge. In spite of its huge size, the total luminosity of this halo derived assuming that it is symmetric is found to be $\\sim 10^9$ L$_\\odot$, similar to the stellar halo of our Galaxy. Therefore it will be interesting to find GCs and investigate their kinematics in this outer halo of M31.   \\subsection{Comparison of GC and PN kinematics}  PNe are an excellent tool for studying the kinematics of the bulge and halo of a galaxy. Recently huge data sets for the PNe in M31 were published by \\citet{mer06, hal06}. How do  the kinematics of the GCs compare with that of the PNe in M31? We have derived the kinematics of about 2500 PNe using the same method to the data given by \\citet{mer06}, as used for the analysis of the GC data, % which are summarized in Table 2.  Figure 26 displays radial velocities (dots) versus the projected distances along the major axis ($X$) and the velocity histogram for the PNe in M31 based on the data given by \\citet{mer06}, like Figure 8 for GCs. We plot also the mean radial velocity and its dispersion of the PNe in a distance bin of 10 $\\arcmin$ along the major axis. Velocity histograms of PNe are displayed with Gaussian fits as in Figure 8. Figure 26 shows that the rotational properties of the PNe are very similar to those of the GCs. The rotational amplitudes of all PNe are $v_{\\rm rot}\\simeq187^{+12}_{-13}$ km s$^{-1}$ from the rotation curve, and $v_{\\rm rot}\\simeq198^{+58}_{-33}$ km s$^{-1}$ from the velocity histogram, which are similar to those of GCs derived in this study, $v_{\\rm rot}\\simeq188_{-28}^{+34}$ km s$^{-1}$, and $v_{\\rm rot}\\sim188_{-33}^{+39}$ km s$^{-1}$, respectively. The rotation amplitudes of the PNe decrease as the distance along the minor axis of M31 increases, but the rotation amplitude is not zero even for the outermost samples at $|Y|\\ge 5$ kpc, which are consistent with the results based on GCs derived in this study.  The rotation-corrected velocity dispersion for  all PNe is estimated to be $\\sigma_{p,r}=94^{+1}_{-1}$ km s$^{-1}$, which is much smaller than that for all GCs, $\\sigma_{p,r}\\sim 134 \\pm 5$ km s$^{-1}$. The rotation-corrected velocity dispersion for the PNe at $|Y|<1$ kpc is   $\\sigma_{p,r}=97^{+2}_{-2}$ km s$^{-1}$, and it becomes smaller and shows little variation with $|Y|$: $\\sigma_{p,r}=75^{+2}_{-2}$ km s$^{-1}$ for  $1\\le|Y|<3$ kpc, $\\sigma_{p,r}=88^{+4}_{-4}$ km s$^{-1}$ for  $3\\le|Y|<5$ kpc, and $\\sigma_{p,r}=79^{+3}_{-3}$ km s$^{-1}$ for  $|Y|\\ge5$ kpc. Thus the rotation-corrected velocity dispersion for the PNe at $|Y|<1$ kpc is consistent  with that for the metal-rich GCs at $|Y|<1$ kpc, $\\sigma_{p,r}=107^{+24}_{-22}$ km s$^{-1}$. We cannot compare the kinematics of the PNe at $|Y|>1$ kpc with that of the metal-rich GCs at $|Y|>1$, because there are too few such GCs in this region.  The rotation-corrected velocity dispersions for the PNe at $|Y|>1$ kpc are significantly different from those for the metal-poor GCs. That is, the rotation-corrected velocity dispersion for the PNe at $|Y|>1$ kpc are smaller than that for the metal-poor GCs at $|Y|>1$ kpc, and the former shows little variation depending on $|Y|$ for  $|Y|>1$ kpc, while the latter increases with $|Y|$. These differences show that the PNe are tracing well the bulge, while the metal-poor GCs are tracing better the metal-poor halo, although their spatial coverages in the sky are similar. Therefore we conclude again that there do exist both an extended rotating bulge and a dynamically hot metal-poor halo in M31.  \\subsection {Formation and Evolution of M31}  Main features related with the kinematics of the M31 GC system revealed in this study can be summarized as follows. Most of the GCs show a strong rotation around the minor axis. The rotation velocity decreases with increasing distance from the major axis. It is still as large as $60\\sim100$ \\kms ~even at $|Y|>5$ kpc, indicating the bulge of M31 is large and rotation-dominated, which is consistent with the finding based on the PNe by \\citet{hur04} and \\citet{mer06}. The rotation velocity of the metal-rich GCs is  similar to % that of the metal-poor GCs in the region close to the major disk. The spatial distribution of metal-rich  GCs shows a stronger central concentration than that of the metal-poor GCs. The rotation-corrected velocity dispersion of the metal-poor GCs in the outer region of M31 is about twice larger than the rotation amplitude. There are several substructures in the spatial distribution of both metal-poor GCs and metal-rich GCs. There are about 50 friendless GCs all over the disk of M31, and most of them show retrograde motion. Interestingly it appears that these GCs show a rotation around the major axis, while most of M31 GCs show a rotation around the  minor axis.  What do these results in addition to the known properties of M31 tell us about the formation and evolution of M31 in comparison with our Galaxy? The most distinguishable kinematic difference between the M31 GCs and the Galactic GCs is a difference in the bulge and halo kinematics. In our Galaxy metal-poor halo GCs are the most dominant population among the GCs, and metal-rich GCs in the bulge (or thick disk) are a minor population . The bulge GCs show significant rotation ($v_{rot}=193\\pm29$ \\kms), while the halo GCs are pressure-supported ($v_{rot}=43\\pm29$ \\kms and $\\sigma_v = 123$ km s$^{-1}$) \\citep{arm89, cot99}. In contrast, the disk and bulge GCs are the most dominant population in M31, and the halo GCs may be a minor population compared with the Galactic halo GCs.  The current picture of the formation and evolution of the halo in our Galaxy involves basically two classical scenarios: (a) The monolithic collapse model according to which the stellar halo is formed by the rapid collapse of a monolithic protogalactic cloud in a dynamical time scale of about 0.1 Gyr% \\citep{egg62}, and (b) the accretion/merging model according to which the stellar halo is formed via successive accretion and merging of protogalactic fragments \\citep{sea78, whi78}. The latter is related with the current paradigm for explaining the formation of the cosmological large scale structure: a hierarchical merging scenario in the $\\Lambda$CDM universe. It appears that both scenarios are involved in forming the halo of our Galaxy \\citep{chi00}. While the inner part of our Galaxy can be explained mainly by the first, the outer part of our Galaxy is explained better by the latter.  Recently, \\citet{ren05a} suggested from a semi-analytic multizone chemical evolution model that the observed higher metallicity in the M31 halo compared with the Galactic halo can be explained by two scenarios: (a) a higher star formation efficiency for the halo that is expected from the accretion of smaller satellites, or (b) a larger reservoir of infalling halo gas with a longer halo phase, producing the intermediate-age population in the halo. These were also confirmed in the semi-cosmological spiral galaxy simulation by \\citet{ren05b}, indicating that the M31 halo may have a more protracted assembly history compared with the halo of our Galaxy. On the other hand, \\citet{fon06b} found from the simulation of the stellar halo in our Galaxy based on the hierarchical scenario: (a) the stellar halo of our Galaxy formed inside out via accretion of satellites; (b) the inner halo at $R<20$ kpc assembled rapidly, and most of its mass were in place more than 8-9 Gyr ago; (c) satellites surviving today were accreted within the past few Gyr; and (d) a major fraction of the present-day stellar halo was originally from a few massive satellites with $M\\sim 10^8 - 10^{10} $ $M_\\odot$. They also pointed out that the much lower metallicity of the Galactic halo compared with the M31 halo indicates that our Galaxy could have formed mainly through smaller satellites. Noting that our Galaxy has a stellar mass and angular momentum that are 2 to 3 times lower than M31, which is a typical spiral galaxy, \\citet{ham07} suggested that our Galaxy has an exceptionally quiet formation history involving no major merger during the last 10-11 Gyr. However, it is noted that the term 'M31 halo' used in these studies corresponds to the extended bulge, not to the metal-poor halo in M31. Therefore these studies explain the formation of the extended bulge in M31 and the halo in our Galaxy.  M31 might have formed and evolved in a way  similar to our Galaxy, with some differences. The existence of the rotating metal-rich central bulge, the radial gradient of the metallicity and the extended smooth metal-poor halo in M31 can be explained by a rapid dissipative collapse of a proto-Galactic cloud, while the extended bulge and its several substructures including giant streams and friendless GCs (\\citet{iba05, gil07, iba07} and this study), are evidence for accretion and merging of small scale fragments. Differences between M31 and our Galaxy are that the degree of merging was larger, the scale of merging fragments was larger, and the duration of merging (or the frequency of merging) was larger in M31 compared with our Galaxy. Remaining questions are: (a) Why does M31 have an extended metal-rich bulge in M31, which is not seen in our Galaxy?; (b) Why is the M31 extended bulge a strong rotator?; (c) Are there metal-poor GCs that do not show systematic rotation in the halo of M31?; (d) Why do the metal-poor GCs in M31 show some rotation while those in our Galaxy do not?; and (e) Is there indeed a population of young disk GCs in M31 as posited by \\citet{mor04}, which has no counterpart in the Galaxy?"
    },
    "0711/0711.3001_arXiv.txt": {
        "abstract": "In this contribution we discuss how neutron stars are produced and retained in globular clusters, outlining the most important dynamical channels and evolutionary events that affect the population of mass-transferring binaries with neutron stars and result in the formation of recycled pulsars. We confirm the importance of electron-capture supernovae in globular clusters as the major supplier of retained neutron stars. By  comparing the observed millisecond pulsar population and the results obtained from simulations, we  discuss several constraints on the evolution of mass-transferring systems. In particular, we find that in our cluster model the following mass-gaining events create populations of MSPs that do not match the observations  (with respect to binary periods and companion masses or the number of produced systems) and therefore likely do not lead to NSs spun up to millisecond periods: (i) accretion during a common envelope event with a NS formed through accretion-induced collapse, and (ii) mass transfer from a WD donor. By restricting ourselves to the evolutionary and dynamical paths that most likely lead to neutron star recycling, we obtain good agreement between our models and the numbers and characteristics of observed millisecond pulsars in the clusters Terzan 5 and 47 Tuc. ",
        "introduction": "Millisecond pulsars (MSPs) are present in globular clusters (GCs) in great numbers: about 140 GC MSPs have been detected\\footnote{See http://www.naic.edu/~pfreire/GCpsr.html for an updated list.}, with more than a dozen in several GCs -- in 47~Tuc \\citep{Camilo00}, M28 \\citep{Stairs06},  and Terzan~5 \\citep{Ransom05}. It is estimated that of order 1000 potentially detectable MSPs are present in the Galactic GC system \\citep{Heinke05a}, as at present the pulsar searches have reached the bottom of the pulsar luminosity function only in a few clusters (e.g., 47~Tuc and M15).  Per unit mass, the number of MSPs in GCs greatly exceeds their numbers in the Galaxy. This was, however, an expected discovery: MSPs are thought to be descendants of LMXBs (for a review see, e.g., Bhattacharya and van den Heuvel 1991), and the abundance of  X-ray binaries per unit mass was known to be $\\sim$100 times greater in GCs than in the Galaxy as a whole \\citep{Clark75}. The high formation rate of X-ray binaries, in itself, is widely accepted to be a consequence of the high stellar density of GCs, which may lead to the creation of compact NS binaries in close stellar encounters.  Although the numbers of MSPs and X-ray binaries per unit mass in globulars are much higher than those in the rest of the Galaxy, only a few X-ray binaries or a few dozen MSPs are present per fairly massive (more than a million stars) and dense cluster. This makes the problem computationally very challenging. E.g., the target time for direct $N$-body methods to address the million-body problem is as far in the future as 2020 (Hut 2006). In our studies, we use a modified encounter rate technique method, described in detail in \\citet{Ivanova05_bf}, and with the updates described in \\citet{Ivanova07}. ",
        "conclusions": "We considered in detail the problem of the formation of MSPs in GCs, which includes NS formation, retention and recycling. Our simulations showed that most of the retained neutron stars in GCs must come from an electron capture supernova formation channel (c.f. the field, where most NSs are from core-collapse supernovae). A typical GC could contain at present (where the adopted cluster age is 11 Gyr) as many as $\\sim 200$ NSs where about half of them are located outside the core; a massive GC like 47 Tuc could have more than a thousand NSs.  Our simulations produce LMXBs in numbers comparable to observations if indeed as many as a few hundred NSs are retained per GC. If, contrary to our assumptions, AIC does not lead to the formation of NSs, then the number of formed NSs is reduced only by $\\sim 20\\%$, but the number of appearing LMXBs is decreased by 2-3 times (per Gyr, at the cluster age of 11 Gyr). Given the large scatter in the simulations, it may still be consistent with the observations.  We find that, in case of a very dense cluster, up to half of NSs could gain mass after their formation through mass transfer, hyper-accretion during a common envelope, or physical collision. It is likely that most of these mass-gaining events do not lead to  NS spin-up, and that only a few percent of all NSs appear eventually as MSPs, implying that there is a large underlying population of unseen NSs in GCs. In particular, we find that if we assume that a NS would accrete and spin up during a common envelope, we overproduce bMSPs in the cluster halos from primordial binaries of intermediate masses. Such bMSPs would be present in low-density clusters and have not yet been seen.  Excluding the systems discussed above, as well as those which are still actively accreting their donor's material and are seen instead as LMXBs, we obtained the number of ``detectable'' bMSPs. The predicted numbers -- about 2 dozen -- are in good agreement with the observations. The fraction of isolated pulsars is comparable with that observed and is bigger in Terzan 5 than in 47 Tuc. Comparing the population census of our models with the observations  of all detected pulsars to date in GCs, we find good agreement for all types of pulsars.  In conclusion we outline several important questions that must be addressed for further progress in studies of NSs in globular clusters:  \\begin{itemize}  \\item What is the final fate of a mass-transferring NS-WD binary -- the post-MT period, the NS luminosity in its quiescent state, and how does it lose its companion?  \\item What is the result of mass accretion on a NS during a merger or a physical collision?  \\item For a common envelope event in primordial binaries of intermediate masses that produce a NS via ECS -- does the NS accrete the material and spin up?  What is the common envelope efficiency?  \\end{itemize}"
    },
    "0711/0711.2627_arXiv.txt": {
        "abstract": "In the standard galaxy formation scenario plasma clouds with a high thermal energy content must exist at high redshifts since the proto-galactic gas is shock heated to the virial temperature, and extensive cooling, leading to efficient star formation, must await the collapse of massive halos (as indicated by the massive body of evidence, referred to as {\\it downsizing}). Massive plasma clouds are potentially observable through the thermal and kinetic Sunyaev-Zel'dovich effects and their free-free emission. We find that the detection of substantial numbers of galaxy-scale thermal SZ signals is achievable by blind surveys with next generation radio telescope arrays such as EVLA, ALMA and SKA. This population is even detectable with the 10\\% SKA, and wide field of view options at high frequency on any of these arrays would greatly increase survey speed. An analysis of confusion effects and of the contamination by radio and dust emissions shows that the optimal frequencies are those in the range 10--35 GHz. Predictions for the redshift distributions of detected sources are also worked out. ",
        "introduction": "\\label{sec:intro}  A satisfactory theory of galaxy formation requires a good understanding of the complex physical processes governing the collapse of primordial density perturbations and the early galaxy evolution. Measurements of the galaxy luminosity and stellar-mass functions up to substantial redshifts have highlighted that these functions show conspicuous differences with respect to the halo mass functions predicted by the cold dark matter (CDM) theory with the \"concordance\" cosmological parameters. At the low-mass end, the halo mass functions is much steeper than the galaxy luminosity function. As discussed by many authors (Larson 1974: Dekel \\& Silk 1986; Cole 1991; White \\& Frenk 1991; Lacey \\& Silk 1991; Kauffmann et al. 1993; Cole et al. 1994; Somerville \\& Primack 1999; Granato et al. 2001; Benson et al. 2003), the relative paucity of low-luminosity galaxies may be attributed to the quenching of star formation in low-mass halos by energy injections from supernovae and stellar winds, and by photoionization of the pre-galactic gas. This leads to the conclusion that efficient star formation must await the collapse of massive halos. On the other hand, the above processes have little effect on very massive halos, which, in the absence of additional relieving mechanisms, would convert too large fractions of gas into stars, yielding too many bright galaxies, with wrong metallicities  (Thomas et al. 2002; see Benson et al. 2003 and Cirasuolo et al. 2005 for discussions of the effect of quenching mechanisms). An effective cure for that is the feedback from active nuclei (AGNs), growing at the galaxy centers (Granato et al. 2001, 2004; Bower et al. 2006; Croton et al. 2006)\\footnote{Note that the AGN feedback invoked by Granato et al. is radically different from that advocated by Bower et al. and Croton et al.. The former is a property of all AGNs and is attributed to a combination of radiation pressure (especially line acceleration) and gas pressure. The latter is associated to the radio active phase of quasars (`radio mode' feedback).}.  During their very early evolutionary phases, massive proto-galaxies are expected to contain large amounts of hot gas, but the gas thermal history is obscure. According to the standard scenario (Rees \\& Ostriker 1977; White \\& Rees 1978), the proto-galactic gas is shock heated to the virial temperature, but this view has been questioned  (Katz et al. 2002; Binney 2004; Birnboim \\& Dekel 2003; Kere{\\v s} et al. 2005), on the basis of independent approaches: analytic methods, a high-resolution one-dimensional code, smoothed particle hydrodynamics simulations. The general conclusion is that only a fraction, increasing with halo mass, of the gas heats to the virial temperature. The hot gas is further heated by supernova explosions and by the AGN feedback, and may eventually be pushed out of the halo. Kere{\\v s} et al. (2005) and Dekel \\& Birnboim (2006) find that there is a critical shock heating halo mass of $\\sim 10^{11.4}\\hbox{--}10^{12}\\,M_{\\sun}$, above which most of the gas is heated to the virial temperature, while most of the gas accreted by less massive halos is cooler.  The large thermal energy content of the hot proto-galactic gas in massive halos makes this crucial evolutionary phase potentially observable by the next generation of astronomical instruments through its free-free emission and the thermal and kinetic Sunyaev-Zel'dovich effects (Oh 1999; Majumdar, Nath, \\& Chiba 2001; Platania et al. 2002; Oh, Cooray, \\& Kamionkowski 2003; Rosa-Gonz{\\'a}lez et al. 2004; De Zotti et al. 2004). In this paper we investigate the detectability of this proto-galactic gas exploiting an up to date model. For the purposes of the present analysis, the adopted model can be taken as representative of the most popular semi-analytic models (White \\& Frenk 1991; Kauffmann et al. 1993; Cole et al. 1994; Somerville \\& Primack 1999; Benson et al. 2003), that all adopt a similar mass function of dark matter halos, a cosmological gas to dark matter ratio at virialization, and assume that all the gas is heated to the virial temperature.  Even if some single dish telescopes have the required theoretical sensitivity, especially at mm and sub-mm wavelengths (e.g. LMT/GTM, GBT at 3 mm, Rosa-Gonz{\\'a}lez et al. 2004) these continuum observations will be hampered by fluctuations in tropospheric emission. Interferometric array observations offer a better trade off between angular resolution, sensitivity and control of systematics (Birkinshaw \\& Lancaster 2005) together with larger fields of view, and allow us to work at lower frequencies where the sources of contaminations from backgrounds and foregrounds are lower and may be better estimated. For this reason we focus mainly on the capabilities of next generation interferometers: the Square Kilometer Array (SKA)\\footnote{http://www.skatelescope.org/}, the Atacama Large Millimiter Array (ALMA)\\footnote{http://www.alma.info/}, the Expanded Very Large Array (EVLA)\\footnote{http://www.aoc.nrao.edu/evla/}, the new 7 mm capability of the Australia Telescope Compact Array (ATCA). We note that the situation for galaxy scale SZ detection is quite different from that for cluster SZ detection.  Cluster SZ signals are stronger and have much larger angular scales than optimum for the interferometer arrays and are best observed with single dishes at high quality sites (Carlstrom, Holder \\& Reese 2002).  The outline of the paper is the following: in $\\S\\,$\\ref{sec:model} we briefly describe our reference model; in $\\S\\,$\\ref{sec:counting}, we present our predictions for the counts of proto-galaxies seen through their free-free emission and their thermal and kinetic Sunyaev-Zel'dovich effects; in $\\S\\,$\\ref{sec:survey} we analyze the potential of the next generation interferometers for detecting such signals, describe possible survey strategies, and discuss possible contaminating emissions and confusion effects; in $\\S\\,$\\ref{sec:conclusions}, we summarize our main conclusions.  We adopt a $\\Lambda$-CDM cosmology with $h=0.71$, $\\Omega_m=0.27$, $\\Omega_\\lambda=0.73$, $\\Omega_b=0.04$, $\\sigma_8= 0.8$, consistent with the results from the Wilkinson Microwave Anisotropy Probe (WMAP) (Spergel et al. 2006). ",
        "conclusions": "\\label{sec:conclusions}  In the standard scenario for galaxy formation, the proto-galactic gas is shock heated to the virial temperature. The observational evidences that massive star formation activity must await the collapse of large halos, a phenomenon referred to as {\\it downsizing}, suggest that proto-galaxies with a high thermal energy content existed at high redshifts. Such objects are potentially observable through the thermal and kinetic Sunyaev-Zel'dovich effects and their free-free emission. The detection of this phase of galaxy evolution would shed light on the physical processes that govern the collapse of primordial density perturbations on galactic scales and on the history of the baryon content of galaxies.  As for the latter issue, the standard scenario, adopted here, envisages that the baryon to dark matter mass ratio at virialization has the cosmic value, i.e. is about an order of magnitude higher than in present day galaxies. Measurements of the SZ effect will provide a direct test of this as yet unproven assumption, and will constrain the epoch when most of the initial baryons are swept out of the galaxies.  As mentioned in \\S\\,\\ref{sec:intro},  almost all semi-analytic models for galaxy formation adopt halo mass functions directly derived or broadly consistent with the results of N-body simulations, and it is commonly assumed that the gas is shock heated to the virial temperature of the halo. They therefore entail predictions on counts of SZ effects similar to those presented here. On the other hand, the thermal history of the gas is governed by a complex interplay of many astrophysical processes, including gas cooling, star formation, feedback from supernovae and active nuclei, shocks. As mentioned in \\S\\,\\ref{sec:intro}, recent investigations have highlighted that a substantial fraction of the gas in galaxies may not be heated to the virial temperature. Also, it is plausible that the AGN feedback transiently heats the gas to temperatures substantially above the virial value, thus yielding SZ signals exceeding those considered here. The gas thermal history may therefore be substantially different from that envisaged by semi-analytic models, and the SZ observations may provide unique information on it.  We have presented a quantitative investigation of the counts of SZ and free-free signals in the framework of the Granato et al. (2004) model, that successfully accounts for the wealth of data on the cosmological evolution of spheroidal galaxies and of AGNs (Granato et al. 2004; Cirasuolo et al. 2005; Silva et al. 2005; Lapi et al. 2006).  We find that the detection of substantial numbers of galaxy-scale thermal SZ signals is achievable by blind surveys with next generation radio interferometers. Since the protogalaxy thermal energy content increases, for given halo mass, with the virialization redshift, the SZ ``fluxes'' increase rather strongly with $z$, especially for the thermal SZ effect, partially compensating for the rapid decrease of the density of massive halos with increasing redshift. The redshift distributions of thermal SZ sources are thus expected to have substantial tails up to high $z$.  There are however important observational constraints that need to be taken into account. The contamination by radio and dust emissions associated to the star formation activity depends on mass and redshift of the objects, but is expected to be stronger than the SZ signal at very low and very high frequencies. We conclude that the optimal frequency range for detecting the SZ signal is from 10 to 35 GHz, where such signal dominates over the contamination at least for the most massive objects. It must be noted however that contaminating emissions have typical scales of the order of those of the stellar distributions, i.e. $< 1''$ at the redshifts of interest here (see Fomalont et al. 2006), while the SZ effects show up on the scale of the dark matter halo, which is typically ten times larger. Therefore arcsec resolution images, such as those that will be provided by the SKA, will allow to reconstruct the uncontaminated SZ signal.  The coexistence of the hot plasma halo, responsible for the SZ signal, with dust emission implies that the scenario presented in this paper may be tested by means of pointed observations of high-$z$ luminous star-forming galaxies detected by (sub)-mm surveys.  Confusion noise is a very serious limiting factor at mm wavelengths. Contributions to confusion come on one side from radio sources and on the other side from dusty galaxies. At 10 GHz only radio sources matter; a modest extrapolation of the 8.4 GHz $\\mu$Jy counts by Fomalont et al. (2002) gives a $5\\sigma_{\\rm conf}$ detection limit $S_{\\rm d} \\simeq 4\\,\\mu$Jy, for a SZ signal subtending a typical solid angle of $100\\,\\hbox{arcsec}^2$. Fluctuations due to dust emission from high-$z$ luminous star-forming galaxies may start becoming important already at 20 GHz; at this frequency, quadratically summing them with those due to radio sources we find again $S_{\\rm d} \\simeq 4\\,\\mu$Jy, for the same solid angle. On the other hand, the high resolution of the SKA and of the EVLA will allow us to effectively detect and subtract out confusing sources, thus substantially decreasing the confusion effects. Beating confusion will be particularly important for searches of the weaker kinetic SZ signal."
    },
    "0711/0711.2761_arXiv.txt": {
        "abstract": "{ We review the available empirical evidence for the presence of ``gaps'' and multimodal distributions among horizontal branch (HB) stars, along with some of the theoretical scenarios that have been proposed to explain these features. While gaps along the HB have become increasingly less prominent and frequent as more and better color-magnitude diagram data have been obtained for Galactic globular clusters, the evidence for multimodal HBs has instead become stronger. In addition, different HB modes have recently started to be traced down to multiple components that have been detected among subgiant branch and main sequence stars, thus suggesting that their origin lies in the complex physical processes that took place at the earliest stages in the history of massive stellar clusters. ",
        "introduction": "``Gaps'' and ``multimodality'' are two terms that one will invariably come across, when studying the literature that addresses the horizontal branch (HB) morphology of resolved globular star clusters (GCs). In this sense, it is instructive to review the following authoritative definitions of the key terms ``gap'' and ``mode,'' as provided by the Cambridge dictionary:  \\begin{center} {\\bf Gap:} {\\em an empty space or opening in the middle of something or between two things} \\smallskip   {\\bf Mode:} {\\em the number or value which appears most frequently in a particular set} \\end{center}  According to these definitions, a multimodal distribution may clearly exist {\\em without} any gaps being present~-- all that is required, in this case, is for the probability distribution to present two or more statistically significant peaks. Indeed, as we shall soon see, the empirical evidence has increasingly been suggesting that ``empty spaces'' or ``forbidden regions'' do not actually exist along the HB. The evidence for HB multimodality, on the other hand, {\\em is} becoming increasingly stronger, especially among some of the most massive globulars, driving renewed interest in theoretical interpretations of the observed features.  It is interesting to note that the first reported gaps on observed color-magnitude diagrams (CMDs) were not located on the HB, but rather either along the main sequence (MS) or the red giant branch (RGB) of both open \\citep{mj57,es64,es69} and globular \\citep*{asea68,hr74} clusters.  Not all such gaps have withstood the test of time. In particular, RGB gaps, once viewed as a ``major significant feature'' \\citep{asea68}, were soon attacked on a statistical basis \\citep{by72}. Later, \\citet{rfp88} pointed out that ``the tendency has been for such gaps to get filled with increasing sample size''; it is indeed unusual for one to find recent CMD studies in which significant RGB gaps are claimed to be present. On the other hand, some of the gaps along the MSs of open clusters \\citep[the ``B\\\"ohm-Vitense gaps'';][]{ebv70,bvc74} are still widely thought to be real \\citep[e.g.,][]{rc00}, and caused by the change in behavior of convection as a function of MS mass \\citep{fdea02b}. Will any of the widely reported HB gaps similarly withstand the test of time, or will they share the same fate as the RGB gaps? ",
        "conclusions": "The empirical evidence for HB gaps has become weaker over the past several years, with several previously reported such features having become at least partly filled as more and better data have become available. It does not appear very likely at present that {\\em all} GCs have HB gaps. Still, plausible mechanisms for the production of some HB gaps have been advanced, generally involving a sharp discontinuity in surface chemical abundances with temperature. Such a discontinuity may be brought about by radiative levitation, which leads to the ``Grundahl jump'' at $T_{\\rm eff} \\simeq 11,\\!500$~K, or by a late hot flasher, which leads to a predicted gap in ZAHB temperatures beyond the end of the canonical EHB. Such phenomena may lead to significantly different CMD features depending on the bandpasses used: in some planes, ``jumps'' and ``gaps'' may become much more apparent than in others.  In contrast, the empirical evidence for multimodal HBs has become significantly stronger lately. As a matter of fact, in several cases HB multimodality has been successfully traced to multiple sequences on the SGB or/and the MS, and in some such cases He enrichment, not necessarily accompanied by an increase in metallicity, is strongly suspected to be the culprit. HB multimodality does not appear to have a single origin though, and more data and theoretical studies are required before we are in a position to claim that we fully understand the origin of multimodal HB distributions."
    },
    "0711/0711.2282_arXiv.txt": {
        "abstract": "{ We discuss the basic difficulties in understanding the origin of the highest energy particles in the Universe - the ultrahigh energy cosmic rays (UHECR). It is difficult to imagine the sources they are accelerated in. Because of the strong attenuation of UHECR on their propagation from the sources to us these sources should be at cosmologically short distance from us but are currently not identified. We also give information of the most recent experimental results including the ones reported at this conference and compare them to models of the UHECR origin. }  \\begin{document} ",
        "introduction": "More than forty years ago, in 1963, John Linsley~\\cite{Linsley63} published an article about the detection of a cosmic ray of energy 10$^{20}$ eV. The article did not go unnoticed, neither it provoked many comments. The few physicists that were interested in high energy cosmic rays were at that time convinced that the cosmic ray energy spectrum will continue forever. The fact that cosmic rays may have energies exceeding 10$^6$ GeV (10$^{15}$ eV) was established in the late thirties by Pierre Auger and his collaborators. Showers of higher and higher energy were detected in the mean time - seeing a 10$^{20}$ eV shower seem to be only a matter of time and exposure. Already in the fifties there was a discussion about the origin of such ultra high energy cosmic rays and Cocconi~\\cite{Cocconi56} reached the conclusion that they must be of extragalactic origin since the galactic magnetic fields are not strong enough to contain such particles.  How exclusive this event is became obvious three years later, after the discovery of the microwave background radiation (MBR). Almost simultaneously Greisen in US~\\cite{G} and Zatsepin\\&Kuzmin~\\cite{ZK} in the USSR published papers discussing the propagation of ultra high energy particles in extragalactic space. They calculated the energy loss distance of nucleons interacting in the microwave background and reached the conclusion that it is shorter than the distances between powerful galaxies. The cosmic ray spectrum should thus have an end around energy of 5$\\times$10$^{19}$ eV. This effect is now known as the GZK cutoff.  The experimental statistics of such events grew with the years, although not very fast. The flux of UHECR of energy above 10$^{20}$ eV was estimated to be 0.5 to 1 event per square kilometer per century per steradian. Even big detectors of area tens of km$^2$ would only detect  few events for ten years of work. The topic became one of common interest during the last decade of the last century when ideas appeared for construction of detectors with effective areas in thousands of km$^2$~\\cite{Auger,TA}. Such detectors would detect tens of events per year and finally solve all mysteries surrounding UHECR. The Auger observatory (3,000 km$^2$) is now almost fully deployed and the Telescope Array (TA, 1,000 km$^2$)  is being constructed in Utah, U.S.A. The expectations for the flux of UHECR are now smaller. The High Resolution Fly's Eye (HiRes) and Auger have shown that the rate of events above 10$^{20}$ eV is about 10 times smaller than previously thought.  Cosmic rays are defined as charged nuclei that originate outside the solar system. They come on a featureless, power law like, $F(E) = K \\times E^{-\\alpha}$, spectrum that extends beyond 10$^{11}$ GeV per particle. There are only two distinct features in the whole spectrum. At energy above 10$^6$ GeV the power law index $\\alpha$ steepens from 2.7 to about 3.1. This is called {\\em the knee} of the cosmic ray spectrum. At energy above 3$\\times$10$^9$ GeV the spectrum flattens again at {\\em the ankle}.  The common wisdom is that cosmic rays below the knee are accelerated at galactic supernovae remnants. Cosmic rays above the knee are also thought to be of galactic origin, although there is no clue of their acceleration sites. Cosmic rays above the ankle are thought to be extragalactic.  More recently, with the improved accuracy and exposure of the modern experiments, several theoretical models that fit the measured ultra high energy cosmic rays (UHECR) spectrum have been developed. We will discuss the experimental data and compare them to some of these models.  \\subsection{Air Shower Detection Methods.}  Cosmic rays of energy above 10$^{14}$ eV are detected by the showers they generate in the atmosphere. The atmosphere contains more than ten interaction lengths even in vertical direction and is much deeper for particles that enter it under higher zenith angles. It is thus a deep calorimeter in which the showers develop, reach their maximum, and then start being absorbed. There are generally two types of air shower detectors: air shower arrays and optical detectors. Air shower arrays consist of numerous particle detectors that cover large area. The shower triggers the array by coincidental hits in many detectors. The most numerous particles in an air shower are electrons, positrons and photons. The shower also contains muons, that are about 10\\% of all shower particles, and hadrons.  The direction of the primary particle can be reconstructed quite well from the timing of the different hits, but the shower energy requires extensive Monte Carlo work with hadronic interaction models that are extended orders of magnitude above the accelerator energy range. The main composition sensitive variable is the ratio of the number of muons in the shower N$_\\mu$ to the number of electrons N$_e$, or the ratio of muon to electron densities at a certain distance from the shower axis. The type of the primary particles can only be studied in statistically big samples because of the fluctuations in the individual shower development. Even then it is strongly affected by the differences in the hadronic interaction models.  The optical method uses the fact that part of the particle ionization loss is in the form of visible light. All charged particles emit in air Cherenkov light in a narrow cone around their direction. In addition to that charged particles excite Nitrogen atoms in the atmosphere, which emit fluorescence light. The output is not large, about 4 photons per meter, but the number of shower particles in UHECR showers is very large, and the shower can be seen from distances exceeding 30 km. The fluorescence detection is very suitable for UHECR showers because the light is emitted isotropically and can be detected independently of the shower direction. Since optical detectors follow the shower track, the direction of the primary cosmic ray is also relatively easy. The energy of the primary particles is deduced from the total number of particles in the shower development or from the number of particles at shower maximum. The rough number is that every particle at maximum carries about 1.5 GeV of primary energy. The mass of the primary cosmic ray nucleus is studied by the depth of shower maximum $X_{max}$, which is proportional to the logarithm of the primary energy per nucleon.  \\subsection{The Highest Energy Cosmic Ray Event} \\begin{figure}[thb] \\includegraphics[width=0.48\\textwidth]{highest_b.ps} \\caption{The shower profile of the highest energy cosmic ray shower detected by the Fly's Eye. } \\label{ts:fig1} \\end{figure}  The highest energy cosmic ray particle was detected by the Fly's Eye experiment~\\cite{Bird_high}. We will briefly describe this event to give the reader an idea about these giant air showers. The energy of this shower is estimated to be 3$\\times$10$^{20}$ eV. This is an enormous macroscopic energy. 3$\\times$10$^{20}$ eV is equivalent to 4.8$\\times$10$^{8}$ erg, 7.2$\\times$10$^{34}$ Hz or the energy of 290 km/h tennis ball in the favorite units of Alan Watson. Fly's Eye was the first air fluorescence experiment, located in the state of Utah, U.S.A. Fig.~\\ref{ts:fig1} shows the shower profile of this event as measured by the Fly's Eye. Note that the maximum of this shower contains more than 2$\\times$10$^{11}$ electrons and positrons. Both the integral of this shower profile and the number of particles at maximum give about the same  energy.   The errors of the estimates come from the errors of the individual data points, but mostly from the uncertainty in the distance between the detector and the shower axis. The minimum energy of about 2$\\times$10$^{20}$ eV was calculated in the assumption that the shower axis was much closer to the detector than the data analysis derived. ",
        "conclusions": "The cosmic ray spectrum steepens at the approach of 10$^{20}$ eV according to the HiRes and the new Auger data. Is this steepening the long expected GZK feature or is it of a different origin? An example for a different origin could be the inability of the sources to accelerate cosmic rays of energy higher than 5$\\times$10$^{19}$ eV. Solving this would require higher statistics than is currently available.  The current statistics is, however, enough to establish the fact that 1-2$\\times$10$^{19}$ eV cosmic rays are not $\\gamma$-rays. This fact supports the acceleration {\\em bottom-up} scenarios for cosmic ray production. At the higher energies more statistics is needed to set limits of a higher quality. Thus it is still possible (although not likely) that the highest energy events could result from {\\em top-down} models.  It is not currently obvious what the cosmic ray nuclear composition is in the energy range above 10$^{18}$ eV. The Auger data points at a composition that is somewhat heavier than the one derived by the HiRes experiment. Differences are not statically significant yet. Composition measurements are hurt by the insufficient understanding of the shower development. Current hadronic interaction models should improve after comparisons with the Large Hadronic Collider (LHC) data that will become available soon.  We still can not see the UHECR sources. Studies of anisotropy, in addition of revealing the UHECR sources, are complimentary to the composition studies. If these particles were heavy nuclei they would scatter stronger in the magnetic fields and show smaller correlation with their sources. Protons would more or less point at their sources. The identification of the UHECR sources would also contribute to our knowledge of magnetic fields in the Galaxy and the Universe.  We are now in a period when the UHECR statistics is increasing very fast and will help the solution of the problems listed above. In addition to Auger and TA the work on a satellite based UHECR Observatory continues. While the previous two projects, EUSO~\\cite{EUSO} and OWL~\\cite{OWL} barely exist anymore, the Japanese JEM/EUSO project, that is to be installed at the Japanese module of the International Space Station is doing well and may be launched in 2013. Such an experiment would increase the statistics by another factor of 10 over the surface arrays.  Finally, the neutrino astronomy projects are also in the process of fast development. Auger is looking for horizontal air showers initiated by neutrinos, the construction of IceCube~\\cite{IceCube} is going extremely well, and the radio detection of UHE neutrinos is also moving forward. The possible detection of cosmogenic neutrinos will help the solution of the UHECR origin when compared to the direct observations."
    },
    "0711/0711.1762_arXiv.txt": {
        "abstract": "We summarize a novel approach which has been recently proposed for direct detection of low energy neutrino backgrounds such as the cosmological relic neutrinos, exploiting neutrino/antineutrino capture on nuclei that spontaneously undergo $\\beta$ decay. ",
        "introduction": "The electron {\\bf N}eutrino {\\bf C}apture from a nucleus $A$ which spontaneously decays via {\\bf B}eta decay to a daughter nucleus $B$ (in the following NCB) \\begin{equation} \\nnbe + A \\rightarrow B + e^\\pm \\label{main} \\, , \\end{equation} shows the remarkable property that it has no energy threshold on the value of the incoming neutrino energy. In the limit of vanishing value for neutrino mass $m_\\nu$ and energy the neutrino contributes to (\\ref{main}) uniquely via its lepton flavor quantum number and in this case the electron in the final state has exactly the $\\beta$ decay endpoint energy $Q_\\beta$. However, for finite $m_\\nu$ the electron kinetic energy is $Q_\\beta+E_\\nu \\geq Q_\\beta + m_\\nu$, while electrons emerging from the corresponding $\\beta$ decay has at most an energy $Q_\\beta - m_\\nu$, neglecting nucleus recoil energy.  The idea of using NCB to measure the cosmological relic neutrino background predicted in the framework of the Hot Big Bang model was already suggested many years ago in \\cite{weinberg}. The original idea was that if relic neutrinos have a large chemical potential $\\mu$, the electron (positron) energy spectrum for $\\beta$ decays and NCB would get quite a typical signature in a interval of order $\\mu$ around the zero neutrino mass endpoint $Q_\\beta$. However, Big Bang Nucleosynthesis constrains relic neutrino--antineutrino asymmetry to the small value $\\mu/T_\\nu \\leq 0.1$, see e.g. \\cite{cuoco,hansen}. This implies that, unless more exotic scenarios are considered, as a larger amount of relativistic degrees of freedom in the Early Universe, the effect of neutrino degeneracy in $\\beta$ decays and NCB is too small to be detected experimentally. However, for massive neutrinos a gap around $Q_\\beta$ is expected of the order of twice the neutrino mass, which for $m_\\nu \\sim 1$ eV is several orders of magnitude larger than the corresponding effect due to neutrino-antineutrino asymmetry. At least in principle, this allows to distinguish between $\\beta$ decay and NCB interaction. In this paper we briefly summarize the results of \\cite{cocco} where it is argued that if $m_\\nu$ is in the eV range, future NCB experiments could represent an almost unique way to detect cosmological neutrinos. ",
        "conclusions": "In this paper we have summarized the analysis of NCB performed in \\cite{cocco}. These processes have the remarkable property of having no energy threshold on the incoming neutrino energy and thus they might represent a good and numerous class of interactions suitable for low energy neutrino flux measurements. The possibility to pursue the ultimate goal of cosmological relic neutrino background detection via a future experimental implementation of this approach depends upon two crucial issues, a high value for the expected order of magnitude of NCB event rate, as well a very good energy resolution of the outgoing electron or positron, of the order of neutrino mass. Both these aspects should be optimized by a careful choice of the $\\beta$ decaying nuclei."
    },
    "0711/0711.1281_arXiv.txt": {
        "abstract": "The presence of dilute hot cavities in the intracluster medium (ICM) at the cores of clusters of galaxies changes the relation between gas temperature and its X-ray emission properties. Using the hydrostatic equations of a porous medium we solve for the ICM density for a given temperature as a function of the filling factor of dilute bubbles. We find that  at a given temperature, the  core X-ray luminosity increases with the filling factor. If the frequency of AGNs in clusters were higher in the past, then the filling factor could  correspondingly be significant, with implications for the cluster scaling relations at high redshifts. This is especially important for the core properties, including the temperature-luminosity ($L_X-T$) relation and estimates of the core gas mass. The results imply an epoch-dependent sensitivity of the $L_X-T$ relation in the core to the porosity of the ICM. Detection of such an effect would give new insights into AGN feedback. ",
        "introduction": "X-ray observations of the intracluster medium (ICM) of massive clusters indicate the presence of cavities seen as depressions in the X-ray surface brightness (e.g. Boehringer \\etal 1993; Fabian \\etal 2000; McNamara \\etal 2000; B\\^irzan \\etal 2004).  These cavities or bubbles are believed to have been generated by AGN activity of central galaxies and possibly other galaxies in the cluster (Nusser, Silk \\& Babul 2006; Best et al. 2007; Martini, Mulchaey, \\& Kelson 2007).  The observed properties of the x-ray emitting gas in clusters of galaxies provide strong constraints on the thermodynamic properties of the ICM. The most  satisfactory model which accounts for the observed self-similar scaling properties of outer regions  invokes  galaxy formation by cooling gas (Suginohara \\& Ostriker 1998; Lewis et al. 2000; Pearce et al. 2000; Dav\\'e, Katz \\& Weinberg 2002; Ettori et al. 2004; Nagai, Kravtsov \\& Vikhlinin 2007), but this in turn leads to a major puzzle. Approximately twice as much cold gas is predicted as compared with the mass in stars. AGN feedback in the core can  resolve this problem (e.g. Quilis \\etal 2001, Babul \\etal 2002, Dalla Vecchia \\etal 2004, Ruszkowski \\etal 2004,  Voit \\& Donahue 2005, Br\\\"uggen \\etal 2005, Roychowdhury \\etal 2005, Sijacki \\& Springel 2006, Nusser, Silk \\& Babul 2006, Peterson \\& Fabian 2006, McNamara \\& Nulsen 2007, Cattaneo \\& Teyssier 2007). Related discussions of the role of AGN heating of the ICM are given in semi-analytic modelling of galaxy formation (Croton et al. 2006; Bower et al. 2006). We point out here that AGN feedback drives porosity in the ICM and modifies the hydrostatic properties of cluster cores. A higher AGN activity in the past leads to a   redshift evolution of the scaling relations. ",
        "conclusions": "\\label{sec:sum}  \\begin{table*} \\centering \\begin{minipage}{140mm} \\caption{Filling factors of radio lobes in a few rich clusters (data taken from B\\^irzan et al. 2004).} \\begin{tabular}{@{}lllllll@{}} & Perseus     & Hydra A   & Cygnus A  & MKW 3S    &  RBS 797     & A2199 \\\\ F($R<30\\rm kpc$)  & 0.006 -0.036&  0.07-0.14& 0.11-0.25 & 0.12-0.23 &  0.023 -0.14& 0.006-0.03\\\\ F($R<40\\rm kpc$)  & 0.002-0.015 &  0.06-0.11& 0.12-0.26 & 0.09-0.19 &  0.01 -0.066&  0.002-0.01 \\end{tabular} \\end{minipage} \\end{table*}   Our results may be interpreted as follows: \\begin{enumerate} \\item For a given observed X-ray emission and spectroscopic temperature of the residual ICM gas (outside the bubbles), hydrostatic equilibrium can yield erroneous mass estimates if the porosity of the medium is not properly included. \\item  The $L_X-T$ relation  will depend on the porosity.   The effect is enhanced if analysis is restricted to the core where the porosity is highest. This is in addition to  changes resulting  from gas  cooling in the core (Allen \\& Fabian 1998,  Markevitch 1998). \\item The ICM porosity is expected to be larger in the past  due to an enhanced frequency of AGN activity. This leads to redshift dependence of a core  $L_X-T$ relation which is different from that seen for the outer parts of the cluster (Kotov \\& Vikhlinin 2005). \\item Porosity enhances the X-ray emission and hence shortens the cooling time. This leads to more AGN activity and the effect is only quenched when the porosity becomes sufficiently large, when only little amount of gas remains available for feeding the AGN. At this point the cold gas is all driven out of the core until there is time for a new episode of cooling and associated AGN activity. \\end{enumerate}   Predictions of the modified hydrostatic equation should be compared  against observations by the analysis of hydrostatic equilibrium of clusters. Porosity of individual clusters is needed for such a study. But reliable estimates of the porosity is not always possible using current data. One could alternatively proceed by solving the modified equation for assumed porosity as a function of cluster temperature and mass. The scatter in the relevant relations could then be minimised with respect to the assumed porosity values.  Rough estimates of porosity could be derived from  the results of B\\^irzan et al. (2004) (see also Dunn \\& Fabian 2004 and Dunn, Fabian \\& Taylor 2005, Dunn \\& Fabian 2006). Using their table 2 for the sizes of prominent active bubbles (radio lobes), we compute the filling factors corresponding to some of their clusters. The filling factors within a distance  of $30$ and $40\\rm kpc$ from the center are listed in table 1 here. The range of the filling factors corresponds to the uncertainty in projection effects (cf. B\\^rzan et al.). Each cluster in  table 1 contains 2 radio lobes and the upper (lower) limiting value of the  filling factor is obtained by summing the maximum (minimum) values of the sizes of both lobes. Table 1 demonstrates that the porosity resulting from active prominent bubbles could be significant. Note that, low contrast bubbles remain undetected (Dunn \\& Fabian 2006) so that the porosity could be much higher than the estimate based on observed bubbles. Therefore, more rigorous estimates of the porosity, especially at larger radii and at different redshifts,  by detailed observational studies of higher S/N are needed in order to test the evolutionary implications of the AGN feedback hypothesis.  The modified hydrostatic equation is valid for any two-component fluid   where the density of one of the components  is negligible. The generalization to multi-component fluids of any densities could easily be done using the derivation  described in Nusser, Silk \\& Babul (2006)."
    },
    "0711/0711.4976_arXiv.txt": {
        "abstract": "We investigate the conformations  of a semiflexible polymer confined to a square box. Results of Monte Carlo simulations show the existence of a shape transition when the persistence length of the polymer becomes comparable to the dimensions of box. An order parameter is introduced to quantify this  behavior. A simple mean-field model is constructed to study the effect of the shape transition on the effective persistence length of the polymer. ",
        "introduction": "Biological macromolecules live in the crowded, confined environment of a cell.   They are often packed into spaces that are much smaller than their natural size and adopt conformations which are unlikely to occur in free space.  For example, viral DNA is packaged in a capsid whose dimensions are comparable to the persistence length of DNA \\cite{Gelbart,Odijk01,Yeomans01,Jiang,Katzav}, which obliges the DNA to adopt a tightly bent shape. Other examples of strong confinement include actin filaments in eukaryotic cell \\cite{Sarah}, and protein encapsulated in E.Coli \\cite{Marrison}. In such environments, the macromolecule is forced to optimize the competing demands of configurational entropy, excluded volume and bending energy.  The packaging of DNA in viral capsids has been investigated extensively and illustrates the effects of these competing interactions and the additional electrostatic forces arising from its charged nature \\cite{Gelbart,Odijk01}. Experimental and theoretical studies \\cite{Yeomans01,Muthukumar01,Frenkel} show that the static, and dynamic properties such as threshold force for packaging and packaging time,  change drastically with changes in confining geometry.  Another example is provided by driving DNA through a nanochannel by electrical force \\cite{Clementi,Kolomeisky}, where the dimensions of the channel determine the translocation time and threshold voltage. Understanding the statistics of polymer in confinement is therefore essential for exploring  both the static and dynamic properties of DNA \\cite{Sakaue}.  In this paper, we investigate the shape and rigidity of a single, semiflexible polymer under confinement. Using the Bond-Fluctuation-Model (BFM) \\cite{Kremer01}, we simulate a single, semiflexible  polymer chain confined to a box in two dimensions. One advantage of  the two-dimensional (2D) simulations is that they can be compared directly to experimental measurements, since conformations of DNA on a substrate can be visualized through atomic force microscopy (AFM) which characterize the statistics of fluctuations of DNA\\cite{Arneodo}. Our simulations show that there is a shape transition as a function of increasing rigidity.  We analyze the nature of this transition by defining an order parameter associated with the transition.  As expected, the flexibility of the chain, measured by the tangent-tangent correlation function, evolves with the shape change.  We construct a theory for this correlation function by considering Gaussian fluctuations around the ``average shape''. ",
        "conclusions": "We have investigated the conformational and elastic properties of a single semiflexible polymer confined in 2D square box using numerical simulations based on a widely used lattice model of polymers, the BFM. Since effects of the competition among configurational entropy, bending energy and excluded volume are included, this simplified model leads to non-trivial results. By mapping the polymer configurations onto the tangent space, we visualized  changes in conformation of the polymer as the bending rigidity was increased from a disordered shape to an intermediate elliptical shape and eventually to a spiral shape (Fig. \\ref{lpsmall}, \\ref{lplarge}). We introduced an order parameter (Eq. \\ref{orderparameter}) to characterize the shapes and the distribution of the order parameter provided further insight into the nature of the shape changes. We used a meanfield approach to calculate the tangent-tangent correlation function and the effective persistence length.  The combination of numerical and analytical results provides a detailed picture of the conformational and elastic properties of a polymer under extreme confinement.   The variation of the persistence length with confinement has been studied previously under tube confinement where the polymer is unconstrained in one dimension.   Our results make specific predictions about the more extreme form of confinement where the polymer is constrained in all dimensions.   The definition of the order parameter, introduced to characterize shapes in two dimensions, can be easily generalized to higher dimensions and should prove useful for characterizing shapes of polymers under confinement. \\\\"
    },
    "0711/0711.1554_arXiv.txt": {
        "abstract": "Our new study of the two central compact object pulsars, \\epsr\\ ($P = 424$~ms) and PSR J1852$+$0040 ($P = 105$~ms), leads us to conclude that a weak natal magnetic field shaped their unique observational properties.  In the dipole spin-down formalism, the 2-sigma upper limits on their period derivatives, $< 2 \\times 10^{-16}$ for both pulsars, implies surface magnetic field strengths of $B_s < 3 \\times 10^{11}$~G and spin periods at birth equal to their present periods to three significant digits.  Their X-ray luminosities exceed their respective spin-down luminosities, implying that their thermal spectra are derived from residual cooling and perhaps partly from accretion of supernova debris. For sufficiently weak magnetic fields an accretion disk can penetrate the light cylinder and interact with the magnetosphere while resulting torques on the neutron star remain within the observed limits. We propose the following as the origin of radio-quiet CCOs: the magnetic field, derived from a turbulent dynamo, is weaker if the NS is formed spinning slowly, which enables it to accrete SN debris. Accretion excludes neutron stars born with both $B_s < 10^{11}$~G and $P > 0.1$~s from radio pulsar surveys, where such weak fields are not encountered except among very old ($> 40$~Myr) or recycled pulsars. We predict that these birth properties are common, and may be attributes of the youngest detected neutron star, the CCO in Cassiopeia~A, as well as an undetected infant neutron star in the SN~1987A remnant. In view of the far-infrared light echo discovered around Cas~A and attributed to an SGR-like outburst, it is especially important to determine via timing whether Cas~A hosts a magnetar or not.  If not a magnetar, the Cas~A NS may instead have undergone a one-time phase transition (corequake) that powered the light echo. ",
        "introduction": "The discovery in recent years of many isolated neutron stars (NSs) at the centers of supernova remnants (SNRs) confirms the long-held notion that these ultra-dense stellar remnants are born in supernova explosions \\cite{baa34}.  Most NSs are identified as pulsars, whose emission derives either from rotational energy loss, as for the rapidly spinning pulsars in the Crab ($P = 33$~ms) and Vela ($P = 89$~ms) remnants, or from magnetic field decay, as posited for the highly magnetized $(10^{14} \\simlt B_s \\simlt 10^{15}$~G) anomalous X-ray pulsars (AXPs) and soft gamma-ray repeaters (SGRs).  However, the nature of a significant fraction of the young ($\\simlt 10^4$~yrs) NSs in SNRs is uncertain. These so-called central compact objects (CCOs) are seemingly isolated NSs, distinguished by their steady flux, predominantly thermal X-ray emission, lack of optical or radio counterparts, and absence of a pulsar wind nebula (see review by Pavlov \\etal\\ \\cite{pav04}).  The properties of eight confirmed and proposed CCOs are summarized in Table~1. (We omit the unique source 1E 161348$-$5055 at the center of RCW 103 because its large-amplitude variability \\cite{got99,del06} violates the adopted definition.) CCO luminosities are typically $10^{33}$ erg s$^{-1}$, similar to the younger pulsars; however, their spectra are best characterized as hot blackbody emission of $kT_{BB} \\sim 0.4$~keV, or two such components.  This is significantly hotter than the surfaces of radio pulsars or other radio-quiet NSs and corresponds to only a small fraction of the NS surface area, with $R_{BB} \\sim 0.7$~km.  Here we summarize new results on two unique pulsars that were initially identified as CCOs.  Their properties, listed in Table 2, are strikingly different from both the traditional radio pulsars and the magnetars.  The lack of a measurable period derivative requires that they have: \\begin{itemize} \\item small natal magnetic fields ($\\simlt 10^{11}$~G), \\item slow initial spin periods very near to their present values ($>0.1$~s), and \\item X-ray luminosities powered by internal cooling, and possibly by accretion of supernova (SN) debris. \\end{itemize} \\noindent We predict that these properties will come to characterize the CCO class of young neutron stars in general. This can be tested with more sensitive searches for pulsations from those objects listed in Table~1.  \\begin{table} \\begin{tabular}{llcccrrl} \\hline\\hline {\\hfil CCO \\hfil} & {\\hfil SNR \\hfil} & {Age} & {$d$}   & {$P$}  & {\\hfil $f_p$\\tablenote{Upper limits on pulsed fraction are for a search down to $P=12$~ms or smaller.}\\hfil } & {\\hfil $L_x$ \\hfil} & {\\hfil References \\hfil} \\cr {   } & {   } & {(kyr)} & {(kpc)} & {(ms)} & (\\%) & {(erg s$^{-1}$)} & \\\\ \\hline RX~J0822.0$-$4300 & Puppis~A           & 3.7  &  2   & \\dots\\   & $<5$ & $5 \\times 10^{33}$ & \\cite{hui06}\\cr CXOU~J085201.4$-$461753 & G266.1$-$1.2    & 1    &  1  & \\dots\\   & $<13$    & $2.5\\times 10^{32}$ & \\cite{sla01,kar02,bam05,iyu05}\\cr 1E 1207.4$-$5209   & PKS~1209$-$51/52& 7   &  2 & 424      & 9         & $2.1 \\times 10^{33}$   & \\cite{zav00,mer02a,big03,del04}\\cr CXOU~J160103.1$-$513353 & G330.2$+$1.0     & $\\simgt 3$ & 5 & \\dots\\ & \\dots\\ & $1.0 \\times 10^{33}$ & \\cite{par06}\\cr 1WGA~J1713.4$-$3949 & G347.3$-$0.5    & 1.6   &  1.3    & \\dots\\   & $< 25$     & $1\\times 10^{33}$ & \\cite{laz03,cas04} \\cr CXOU~J181852.0$-$150213 & G15.9$+$0.2       & $1-3$ & (8.5) & \\dots\\   &  \\dots\\   & $1 \\times 10^{33}$ & \\cite{rey06}\\cr CXOU~J185238.6$+$004020 & Kes~79          & $7$  & 7   & 105      & 80        & $3.3 \\times 10^{33}$ & \\cite{sew03,got05,hal07}\\cr CXOU~J232327.9$+$584842 & Cas~A           & 0.33  & 3.4  & \\dots\\   & $<27$    & $2.6\\times 10^{33}$ & \\cite{pav00,cha01,mur02,mer02b}\\cr \\hline \\end{tabular}\\\\ \\caption{Central Compact Objects in Supernova Remnants} \\end{table}  \\subsection{\\epsr\\ in SNR \\esnr}  The central source \\esrc\\ in the SNR \\esnr\\ (Fig.~1) is the most intensively studied of the CCOs.  It acquired special importance as the first CCO detected as a pulsar \\cite{zav00,pav02}. It was distinguished again as the first isolated NS to display strong absorption lines in its X-ray spectrum \\cite{san02,mer02a,big03}. We have made a comprehensive study of all X-ray timing data on \\esrc\\ (\\epsr) \\cite{got07}, showing that its rotation is highly stable (Fig. 2), with $P = 424.130749(4)$~ms and $\\dot P = (6.6 \\pm 9.0) \\times 10^{-17}$ ($1\\sigma$ errors), superseding previous claims of large period changes in the same data \\cite{zav04,woo07}. In the dipole spin-down formalism, the $2\\sigma$ upper limit on $\\dot P$ implies a spin-down luminosity $\\dot E \\equiv -I\\Omega\\dot \\Omega < 1.3 \\times 10^{32}$~erg~s$^{-1}$, surface magnetic field strength $B_s < 3.3 \\times 10^{11}$~G, and characteristic age $\\tau_c \\equiv P/2\\dot P> 27$~Myr.  This lower limit on $\\tau_c$ exceeds the SNR age by at least $3$ orders of magnitude, requiring that the pulsar was born spinning at its present period to three significant digits.  The X-ray luminosity of \\esrc, $L_{\\rm bol} \\approx 2.1 \\times 10^{33}\\,(d/2\\ {\\rm kpc})^2$ erg~s$^{-1}$, exceeds its $\\dot E$, implying that $L_{\\rm bol}$ derives from residual cooling, and perhaps partly from accreting SN debris.  \\subsection{\\kpsr\\ in SNR Kes 79}  We also discovered pulsations from a second CCO, in the SNR Kes~79 \\cite{got05} (Fig.~3). Our follow-up program to time \\kpsr\\ produced a remarkable result: no change in its 105~ms period over 2.4 yr \\cite{hal07}.  From the data shown in Figure~4, we derived a $2\\sigma$ upper limit of $\\dot P < 2.0 \\times 10^{-16}$, which leads to a spin-down luminosity $\\dot E < 7 \\times 10^{33}$~erg~s$^{-1}$, surface magnetic field strength $B_s < 1.5 \\times 10^{11}$~G, and characteristic age $\\tau_c > 8$~Myr. Again, this implies that the pulsar was born spinning at its current period, with a weaker $B$-field than that of any other young pulsar.  These are the only two NSs whose initial rotation periods are so precisely inferred.  They are longer than what was once thought typical, but in fact are consistent with recent statistical analyses of the radio pulsar population, e.g., Faucher-Gigu\\`ere \\& Kaspi \\cite{fau06} who favor a wide distribution of birth periods (Gaussian mean $P \\sim 300$~ms and $\\sigma_P \\sim 150$~ms). The absence of spin-down for both CCO pulsars then highlights the difficulty of existing theories to explain the high luminosities and temperatures of CCO thermal X-ray spectra. Their X-ray luminosities are a large fraction of their $\\dot E$, challenging the rotation-powered assumption, and greater than their reservoirs of $B$-field energy, refuting the magnetar hypothesis.  \\begin{table} \\begin{tabular}{lcc} \\hline\\hline {\\hfil Parameter \\hfil} & {PSR J1209$-$5226} & {PSR J1852$+$0040} \\cr \\hline $P$ (ms)                 & 424.1307 & 104.9126 \\cr $\\dot P$ ($10^{-17}$ s s$^{-1}$)   & $6.6 \\pm 9.0$ & $-34 \\pm 27$  \\cr $\\dot E$ ($10^{32}$ erg s$^{-1}$)\\tablenote{$2\\sigma$ limit assuming magnetic dipole spin-down.} & $<1.3$ & $<70$ \\cr $L({\\rm bol})/\\dot E$ & $> 15$ & $>0.5$ \\cr $B_s$ ($10^{11}$ G)$^*$              & $<3.3$   & $<1.5$   \\cr $\\tau_c$ (Myr)$^*$ & $>27$   & $>8$     \\cr SNR age (kyr)                       & $\\sim 7$ & $\\sim 7$ \\cr Reference & \\cite{got07} & \\cite{hal07} \\cr \\hline \\end{tabular}\\\\ \\caption{Spin Parameters of CCO Pulsars} \\end{table} ",
        "conclusions": ""
    },
    "0711/0711.1624_arXiv.txt": {
        "abstract": "{ We have analysed near-infrared NTT/SofI observations of the starburst cluster Westerlund 1, which is among the most massive young clusters in the Milky Way. A comparison of colour-magnitude diagrams with theoretical main-sequence and pre-main sequence evolutionary tracks yields improved extinction and distance estimates of A$_{\\rm Ks}$ =  1.13$\\pm$0.03\\,mag and d = 3.55$\\pm$0.17\\,kpc (DM  = 12.75$\\pm$0.10\\,mag). The pre-main sequence population is best fit by a Palla \\& Stahler isochrone for an age of 3.2\\,Myr, while the main sequence population is in agreement with a cluster age of 3 to 5\\,Myr. An analysis of the structural parameters of the cluster yields that the half-mass radius of the cluster population increases towards lower mass, indicative of the presence of mass segregation. The cluster is clearly elongated with an eccentricity of 0.20 for stars with masses between 10 and 32\\,M$_\\odot$, and 0.15 for stars with masses in the range 3 to 10\\,M$_\\odot$. We derive the slope of the stellar mass function for stars with masses between 3.4 and 27\\,M$_\\odot$. In an annulus with radii between 0.75 and 1.5\\,pc from the cluster centre, we obtain a slope of  $\\Gamma = -1.3$. Closer in, the mass function of Westerlund 1 is shallower with $\\Gamma = -0.6$. The extrapolation of the mass function for stars with masses from 0.08 to 120\\,M$_\\odot$ yields an initial total stellar mass of $\\approx$52,000\\,M$_\\odot$, and a present-day mass of 20,000 to 45,000\\,M$_\\odot$ (about 10 times the stellar mass of the Orion Nebula Cluster, and 2 to 4 times the mass of the NGC 3603 young cluster), indicating that Westerlund\\,1 is the most massive starburst cluster identified to date in the Milky Way. ",
        "introduction": "Clustered star formation is the dominant mode of star formation in the Universe with about 80\\% of all stars originating in giant molecular clouds (Lada \\& Lada \\cite{lada}). The globular cluster systems in the Milky Way or M\\,31 might be remnants of an epoch of Super Star Cluster (SSC) formation early in the history of these galaxies. Low-mass stars lost in the course of the dynamical evolution of the SSCs could explain the majority of Pop II field stars found in galaxy halos (Kroupa \\& Boily \\cite{kroupa02}). While the Galactic globular clusters formed more than 12\\,Gyr ago, 100s of young SSCs with masses comparable to or greater than Galactic globular clusters; i.e., present-day masses $\\ge 100,000\\,{\\rm M}_\\odot$, have been identified in interacting galaxies like, e.g., the \\object{Antennae galaxies} (Whitmore \\& Schweizer \\cite{whitmore95}).  At extragalactic distances, however, individual cluster members are often unresolved, and physical properties of SSCs  -- such as their initial mass function (IMF) and total mass, degree of binarity, or density profile -- must be  estimated from their integrated spectra and photometry alone, which in turn are  dominated by the most massive, luminous stars present. The contribution of lower mass stars to the total cluster mass, which is crucial for the dynamical evolution and the long-term survival of a cluster (e.g., de Grijs \\& Parmentier \\cite{degrijs07}, and references therein), remains uncharted. Indirect estimates of the total cluster mass can be derived by extrapolating the IMF towards lower masses. Another possibility is to determine the dynamical mass of a cluster from its half mass radius and velocity dispersion under the assumption of virial equilibrium. Studies by different groups either agreed with a Kroupa (\\cite{kroupa02a}) IMF over the whole range of stellar masses (e.g. Larsen \\& Richtler \\cite{larsen}) or found a truncated IMF towards lower stellar masses (e.g., Rieke et al.\\ \\cite{rieke93}; McCrady et al.\\ \\cite{mccrady05}). Young, massive clusters, however, are not necessarily in virial equilibrium, as photoionization and winds from O-stars in their centres rapidly expel any remnant gas left over from the formation process (Kroupa \\& Boily \\cite{kroupa02};  Bastian \\& Goodwin \\cite{bastian06}).  The determination of the IMF slope as well as the potential presence or absence of a low mass stellar component, is critical for constraining the long term evolution of such clusters. This addresses the question of whether (proto-) Globular Clusters  are still forming in the Universe today. Moreover, as the majority of low mass field stars might have formed in clusters, the uncertainty in the initial ratio of high to low mass stars in clusters  results in a corresponding uncertainty in the star formation and chemical enrichment history of galaxies like our Milky Way.  Therefore we examine starburst clusters in the Milky Way and its companion galaxies, where SSCs can be resolved into individual stars, and cluster properties like the IMF can be determined directly. As a {\\it starburst} cluster we consider any cluster which is massive enough to house several early O-type stars with initial masses in excess of 50\\,M$_\\odot$ in its centre. Among the most massive Galactic starburst clusters studied to date are \\object{NGC\\,3603YC} (e.g., Brandl et al.\\ \\cite{brandl}, Stolte et al.\\ \\cite{stolte06}) and \\object{Westerlund 2} (Wd 2, Ascenso et al.\\ \\cite{ascenso07}), both in the Carina spiral arm,  and the two massive clusters \\object{Arches} (Stolte et al.\\ \\cite{stolte02,stolte05}; Figer et al. \\cite{figer02}; Kim et al.\\ \\cite{kim06}) and \\object{Quintuplet} (e.g., Cotera et al.\\ \\cite{cotera96}, Figer et al.\\ \\cite{figer04})   in the Galactic centre (GC) region.  One questions of particular interest is whether or not the presence of high mass stars hinders or even quenches the formation of low mass stars. Intense UV radiation fields and strong stellar winds of nearby OB stars might lead to the destruction  of protostellar cores before low-mass stars can form. Furthermore, in extreme environments like the GC region, variations in the physical conditions leading to higher thermal and/or turbulent  pressure, as well as the presence of strong magnetic fields might preferentially yield higher mass prestellar cores than in more benign environs (Morris \\cite{morris93}, Stolte et al. \\cite{stolte05}, Larson \\cite{larson06}, Klessen et al.\\ \\cite{klessen07}).  Along with the Arches and Quintuplet clusters in the GC region, and the NGC 3603YC in the Sagittarius-Carina spiral arm, \\object{Westerlund 1} (Wd 1), located in the Scutum-Crux spiral arm, is one of the very few young, massive starburst clusters in the Milky Way.  Discovered by Westerlund (\\cite{west61}), subsequent optical observations of the galactic open cluster  Westerlund 1 (Wd~1) suggested an unusually rich population of cool and hot  supergiants (Westerlund \\cite{west87}), from which a large cluster mass could be inferred. Recent optical spectroscopic and photometric observations have confirmed these assertions (Clark \\& Negueruela \\cite{cn}; Clark et al. \\cite{clark05} (henceforth C05); Negueruela \\& Clark \\cite{nc}). Located close to the Galactic plane (b = $-0.35^\\circ$) at a distance of $\\sim$3--5\\,kpc, Wd\\,1 is subject to significant foreground extinction. Thus optical studies of the cluster were limited to the most massive stars  of Wd1. C05 spectroscopically identified about 50 cluster members with masses in excess of 30\\,M$_\\odot$. Adding 150 probable cluster members identified by means of photometry, they estimate that the high-mass stellar content of Wd\\,1 alone amounts to 6000\\,M$_\\odot$. Assuming a Kroupa IMF to extrapolate from this population, C05 then infer a likely total mass of $\\sim 10^5$\\,M$_{\\odot}$  for Wd1; directly comparable  to the masses of SSCs observed in other galaxies.  Up to now, however, only the evolved, high-mass stellar population of the cluster has been characterised. The presence of Wolf-Rayet stars (e.g., Crowther et al.\\ \\cite{crowther06}; Skinner et al.\\ \\cite{skinner06a}), a pulsar indicative of a recent supernova (Skinner et al.\\ \\cite{skinner06b}, Muno et al.\\ \\cite{muno06a}), as well as O supergiants, suggest an age between 3 and 5\\,Myr for the cluster. Because of a visual foreground extinction of A$_{\\rm V} \\approx 10$\\,mag, and the uncertainties in the intrinsic luminosity of the evolved, massive stars, distance determinations to Wd\\,1 range from 2 to 5.5\\,kpc (C05).  Because of its proximity, Wd1 provides the unique opportunity to directly determine the spatially resolved IMF of a SSC and, given that it is located well away  from the GC in a supposedly `quiescent' region of the Galactic disc, potentially study physical properties imprinted by the formation process from the parental molecular cloud on the cluster. With the aim to better determine basic cluster properties such as age, foreground extinction and distance, to directly measure the IMF, and to study the structural properties of Wd\\,1, we have analysed deep near IR observations. This paper describes the analysis of  NTT/SofI observations of Wd1 covering a $\\sim 5' \\times 5'$ field of view centred on the cluster core, which  are sensitive to $\\sim$ solar mass stars at a distance of 4\\,kpc; a second  paper details subsequent VLT/NAOS-CONICA adaptive optics observations of selected core regions sensitive to $\\sim$0.2\\,M$_{\\odot}$.  The structure of the paper is as follows: in section 2 we present the observations and the data analysis. The basic astrophysical quantities extinction, distance, and age of Wd~1 are derived in section 3. Section 4 discusses the mass function and the total stellar mass of Wd~1. Dynamical and structural parameters of Wd~1 are addressed in section 5, followed by a comparison of Wd~1 to other well studied starburst clusters in section 6. Section 7 summarises the results.  \\begin{figure*}[htb] \\centerline{ \\hbox{ \\includegraphics[width=8.5cm,angle=0]{figures/wd1_sofi_jhk_rgb.jpg} \\includegraphics[width=8.6cm,angle=0]{figures/field_sofi_jhk_rgb.jpg} } } \\caption{Composite colour JHKs images of Westerlund 1 (left) and the comparison field (right) obtained with NTT/SofI. The field of view is $4' \\times 4'$. North is up, and east is to the left.} \\label{image_rgb} \\end{figure*} ",
        "conclusions": "We have analysed near-infrared NTT/SofI observations of the starburst cluster Westerlund 1, which is among the most massive young clusters in the Milky Way. A comparison of colour-magnitude diagrams with theoretical main-sequence and pre-main sequence evolutionary tracks yields improved extinction, distance and age estimates of A$_{\\rm Ks}$ =  1.13$\\pm$0.03\\,mag, d = 3.55$\\pm$0.17\\,kpc (DM  = 12.75$\\pm$0.10\\,mag) and t = 3.6$\\pm$0.7\\,Myr, respectively. We derive the slope of the stellar mass function for stars with masses between 3.4 and 27\\,M$_\\odot$. In an annulus with radii between 0.75 and 1.1\\,pc from the cluster centre, we get a slope of  $\\Gamma = -1.3$, i.e.\\ the Salpeter slope. Closer in, the mass function of Westerlund 1 is shallower with $\\Gamma = -0.6$, while at larger separations from the cluster it is getting steeper, reaching $\\Gamma = -1.6$ for separations larger than 1.5\\,pc. This is in good agreement with the change in mass function slope found in the starburst cluster NGC\\,3603YC (Stolte et al.\\ \\cite{stolte06}).   Only considering stars with masses between 3.0 and 32\\,M$_\\odot$, we derive a half-mass radius of 1.0\\,pc for the cluster. Wd\\,1 exhibits clear deviations from spherical symmetry. The distribution of stars with masses between 10 and 32\\,M$_\\odot$ has an eccentricity of 0.20 with the major axis aligned roughly in the North-South direction. The distribution of stars with masses between 3.0 and 10\\,M$_\\odot$ is elongated in the same direction with an eccentricity of 0.15. The flattening of the cluster might be explained by rotation of the cluster along its minor axis.  By extrapolation of the observed mass function slopes for different annuli, we derive an upper limit of the total initial stellar cluster mass of $\\approx 52,000$M$_\\odot$. Stellar evolution at the high mass end should have reduced the total initial cluster mass by $\\approx$6,000\\,M$_\\odot$ over the past 4\\,Myr. By adding up the individual initial masses of stars directly detected in the cluster, we derived a lower limit for the total present cluster mass of m $\\ge 2\\times 10^4$\\,M$_\\odot$.  With a present-day mass of 20,000 to 45,000\\,M$_\\odot$ and an initial stellar mass of $\\approx$52,000\\,M$_\\odot$, Wd\\,1 is the most massive starburst cluster identified in the Milky Way to date, and about 10 times as massive as the ONC, and 2 to 4 times as massive a NGC\\,3603\\,YC. Additional cluster members located outside the field of view of the SofI data might push the mass of Wd\\,1 even higher.  In a following paper, we will report of the analysis of high-angular resolution adaptive optics data, which trace the pre-main sequence population of the cluster down to $\\approx$0.2\\,M$_\\odot$.  Further studies should aim at determining proper motions for individual cluster members as well as more radial velocity measurements. From this precise models of the cluster kinematics and dynamics could be derived, which also might enable one to trace back the dynamical evolution of the cluster since its formation $\\approx$4\\,Myr ago. Wd\\,1, being the most massive, and one of the most nearby starburst clusters, could turn into the prime example for studies on starburst cluster formation and evolution."
    },
    "0711/0711.3417_arXiv.txt": {
        "abstract": "{We study the propagation of bubbles of new vacuum in a radially inhomogeneous Lema\\^{i}tre-Tolman-Bondi background that includes a cosmological constant.  This exemplifies the classical evolution of a tunneling bubble through a metastable state with curvature inhomogeneities, and will be relevant in the context of the Landscape. We demand that the matter profile in the LTB background satisfy the weak energy condition. For sample profiles that satisfy this restriction, we find that the evolution of the bubble (in terms of the physically relevant coordinates intrinsic to the shell) is largely unaffected by the prsence of local inhomogeneities. Our setup should also be a useful toy model for capturing the effects of ambient inhomogeneities on an inflating region. } ",
        "introduction": "Although bubble propagation in homogeneous cosmological backgrounds has been amply studied \\cite{Israel:1966rt}-\\cite{Lake:1984pn}, bubble propagation in inhomogeneous backgrounds is virtually unexplored \\cite{Khakshournia:2002jr}. Since tunneling is of enormous importance \\cite{Vilenkin:1983xq, Hawking:2006ur} in populating the string Landscape \\cite{Kachru:2003aw, Susskind:2003kw, Aguirre:2007an, Chang}, we expect that the evolution of a new Universe through an old one, is worthy of study.  One reason for the lack of emphasis on inhomogeneous backgrounds is the cosmic no-hair theorem \\cite{Gibbons:1977mu, Hawking:1981fz, Starobinsky:1982mr}. This theorem states that the evolution of an expanding universe that contains a cosmological constant will eventually be dominated by vacuum energy, and will asymptote to a de Sitter space. It will imply in particular, that inhomogeneous backgrounds containing a cosmological constant (like the ones constructed on Landscape vacua) will evolve toward `ironing out' such inhomogeneities, and therefore it will make sense to restrict the study of bubble propagation to homogeneous backgrounds. But fluctuations can still clump locally even in an asymptotically de Sitter space, and therefore the question of evolution of vacuum bubbles in an inhomogeneous spacetime is still relevant. In the context of eternal inflation and the string landscape, we have to face up to the possibility that the Universe that we inhabit is a result of bubble evolution through an inhomogeneous relic Universe. In principle this could give rise to experimental signatures in the form of CMB fluctuations, about which we comment in the final section.  One particular work that stands out in this context is that of Wald \\cite{Wald:1983ky}, who did a comprehensive study of of the late time evolution of anisotropic but homogeneous spaces with a cosmological constant. In this restricted setting, he was able to define the conditions for the cosmic no hair theorem to be satisfied by taking advantage of the Bianchi classification of homogeneous, anisotropic metrics.  Wald studied the time evolution of such metrics when the energy momentum tensor is the sum of two components: a cosmological constant and a term that satisfies the dominant  and strong energy conditions. He found that with the exception of the Bianchi IX, the other models will always asymptote to a de Sitter space, within a time scale $\\sqrt{\\frac{3}{\\Lambda}}$. In the case of the Bianchi IX models, the future behavior depends on the relative sizes of the cosmological constant and the spatial-curvature. Only when the latter exceeds the cosmological constant will the time evolution {\\em not} be asymptotically de Sitter.   It is often claimed that inflation solves  the horizon and flatness problems but,  as already remarked in its early years, these claims are based on the assumption that the pre-inflationary state of the universe can be described by a homogeneous and isotropic Friedmann-Robertson-Walker (FRW) space. Under these strong assumptions, inflation does indeed reduce considerably the amount of fine tuning  necessary to explain the current experimental measurements. In general, however,  one expects a high degree of inhomogeneity in the initial conditions, with some regions of the universe expanding while others contract.  Holographic cosmology  might be the only exception to this rule because it naturally produces a homogeneous and almost flat space-time   \\cite{Banks:2001px, Banks:2003ta, Banks:2004cw}.   The question of whether an inflating region can continue to inflate if the ambient region is inhomogeneous, has been explored by several authors for different pre-inflationary backgrounds. In a set of papers \\cite{Goldwirth:1989pr,Goldwirth:1991rj,Goldwirth:1989vz}, Goldwirth and Piran sought to answer the question of whether large inhomogeneity in the very early universe could prevent it from entering a period of inflation. In their work, inflation was not driven by a cosmological constant but by an inflaton field. Their analysis is numerical and restricted to an inhomogeneous but isotropic and  closed-universe background. They found that a large initial inhomogeneity could indeed suppress the onset of inflation, and that  the inflaton field must have a suitable value over a region of several horizon sizes in order for inflation to start.  This last finding raises the issue of an acausal initial condition. Other numerical simulations disagree with these results claiming that inhomogeneous backgrounds will have enough inflation to explain current observations \\cite{Kurki-Suonio:1993fg,Berera:2000xz}. Analytical computations that use the long wave-length approximation seem to confirm that large inhomogeneities of the spatial curvature prevent the onset of inflation \\cite{Deruelle:1994pa, Iguchi:1996jy, Iguchi:1996rh}.  In this work, we will not be able to conclusively settle this issue once and for all. Instead, our goal here is to study a particular problem that we believe captures some of the same physics by playing local curvature effects against the might of the cosmological constant. The problem we consider is the classical evolution of a vacuum bubble in an inhomogeneous background, containing a cosmological constant and dust. This background, which we will refer to as the Lema\\^{i}tre-Tolman-Bondi (LTB)  space-time \\footnote{ We slightly generalize it from the original form to allow for a positive cosmological constant.}, will be the simplest possible inhomogeneous space, spherically symmetric and with only one center \\cite{Lemaitre:1933gd}-\\cite{Bondi:1947av}. The bubble will be created in a region of space that is expanding but will encounter through its evolution regions of varying curvature. In a nutshell, the problem is that the critical size for bubble nucleation depends locally on the scale factor. As this factor changes as a function of the position of the shell, it is possible for a supercritical bubble to become subcritical \\footnote { A bubble that can become subcritical in the future is a fluctuation, not a phase transition. }. Our analysis will be confined to the evolution of the bubble, it will not address the interesting issue of tunneling probabilities in inhomogeneous backgrounds. This latter issue is very complicated, some attempts to address it have been made in \\cite{Abbott:1987xq}.  Before embarking on describing the details, we emphasize that we are {\\em not} dealing with the most general inhomogeneous bubble evolution imaginable. There are three major assumptions we make to render the problem tractable, and we list them below. The results of this paper should therefore be regarded as exploratory, rather than conclusive. \\begin{itemize} \\item The first assumption we make is that the ambient spacetime into which the bubble evolves is the radially inhomogeneous LTB form. This allows us to work with an exact solution of Einstein's equation, but the restricted form of the metric and matter in this solution implies that we are not dealing with the most general radially inhomogeneous cosmology admissible. But as we explain in more detail in the next section, this still leaves us with a great deal of flexibility in choosing the inhomogeneity profiles. On the inside of course, we are dealing with a vacuum bubble, so we take the matter there to be only a cosmological constant without loss of generality. \\item The second assumption that we make is that the evolution of the new phase can be adequately captured using a thin-wall approximation. In particular this means that we can deal with all the transient phenomena between the two regions using a thin layer. The more general scenario would be to model the matter leakage from one region to the other using a gradual profile, but we believe the thin-wall captures the phenomena involved to a first approximation. \\item The above two assumptions together give us some traction. Once we make them, the problem that we need to solve reduces to that of solving the Israel junction conditions for a thin shell. But the junction conditions do not fully fix the evolution of the shell until we specify an equation of state on the bubble-wall. In the purely general relativistic context that we are working in, these equations of state are not fixed by the dynamics. So what we do here is to take the first step in this direction by considering a few equations of state of the form $p= w \\rho$, for a range of values of $w$. We make some comments elsewhere on how one might try to fix this equation of state using field-theoretic arguments, but the full solution of this issue in the complicated background that we are working in, is likely to be difficult. In any event, as a start, we settle here for an exploration of these equations of state, leaving a more thorough analysis for future.  \\end{itemize}  We investigate a few different scenarios using our numerical formalism. The first is the homogeneous limit of LTB, namely Robertson-Walker cosmologies with or without matter. We reproduce and extend known results in the literature using our methods, and offer some comments and new perspectives. Particular attention is paid to the case of bubble-propagation in an FRW universe with a Big Crunch in the future. We find that there are new ways (details are in a later section) in which the bubble can evolve in this case: it can even recollapse. The final scenario we consider is an LTB space that is truly inhomogeneous. In our examples, we pick our inhomogeneity profiles to have a fluctuation that is as sharp as possible while still respecting the weak energy condition. Surprisingly, we find that the qualitative bubble evolution (as seen in the physically relevant coordinates intrinsic to the shell or in the inside coordinates) is essentially the same with or without the fluctuation. This suggests that the energy conditions work towards the smoothening out of exotic bubble evolution patterns.  In the next section we shall describe the LTB model that represents the background on which the bubbles will propagate. In section 3, we will review the propagation of bubbles in a homogeneous background. We will consider the already known case where the outside background is simply a de Sitter space, either flat or closed, and the novel case where matter is added to the cosmological constant. In this latter case, we shall explore the interesting situation of a bubble moving in a background that evolves from expanding to collapsing. In section 4 we shall study the truly inhomogeneous LTB background. Our results, future directions and cautionary remarks are put together in the Conclusion. We have relegated to the Appendix, an explanation of the known formalism to study the propagation of bubbles in general relativity. ",
        "conclusions": "We have examined the evolution of bubbles of true vacuum in several backgrounds (that always include a cosmological constant) to assess the effect of curvature on the propagation. The bubble evolution changes qualitatively when looked at in terms of the outside coordinates, but when studied in terms of the more physically relevant coordinates on the shell or its interior, the local curvature is largely irrelevant. In the presence of matter, we observe that some bubbles, with an equation of state close to $w=-1$, disappear. Those that survive do behave qualitatively the same.  The explanation for this insensitivity to inhomogeneities is to be found in the weak energy condition: our background spacetime is consistent only when (\\ref{d(t,r)}) is positive, which restricts the sharpness of the inhomogeneity profiles that are allowed. This could of course be a restriction arising from the ans\\\"atz for the background that we are working with. But it should be mentioned that in the work of  Wald \\cite{Wald:1983ky} that was mentioned in the introduction, the energy conditions played a crucial role in the proof of the cosmic no hair theorem for anisotropic (but homogeneous) cosmologies.  We view this work as an attempt to test the requirements to arrive at a universe like ours without making the simplifying assumptions of starting with an FRW metric.  There are many directions in which the work here can be extended. We have examined only a few inhomogeneity profiles here, a more thorough understanding of the evolution curves in more generic situations is certainly of interest. For all the various initial conditions and collapse/inhomogeneity profiles that are allowed, the question of bubble evolution should be tractable with our machinery\\footnote{Though such a scan would admittedly be much more numerically challenging than the few cases we have undertaken here.}. If we think of our setup as a toy model for inflation, it will definitely be of interest to study the {\\em genericity} of inflation in the space of possible inhomogeneity profiles. Another important direction to follow is to evaluate the tunneling amplitudes for inhomogeneous ambient spaces. This is a notoriously difficult problem on which any enlightenment will be welcome.  Another problem that is of interest in this context is whether there are signatures left on the new (inside) Universe due to the propagation of the bubble through the inhomogeneous ambience. The bubble can be thought of as a Casimir cavity with a moving wall and the cavity radiation (the analogue of the CMB in this context) will carry imprints of the external inhomogeneities. Because of the spherical symmetry, this is effectively a two-dimensional QFT problem, with moving mirror boundary conditions \\cite{Birrell}."
    },
    "0711/0711.3129_arXiv.txt": {
        "abstract": "We study structure formation in a phenomenological model of modified gravity which interpolates between $\\Lambda$CDM and phenomenological DGP-gravity. Generalisation of spherical collapse by using the Birkhoff-theorem along with the modified growth equation shows that the overdensity for spherical collapse $\\delta_c$ in these models is significantly lowered compared to $\\Lambda$CDM, leading to enhanced number densities of massive clusters and enhanced cluster merging rates. We find that $\\delta_c(z)$ is well fitted by a function of the form $\\delta_c(z) = a - b\\exp(-cz)$. We examine the sensitivity of PLANCK's and SPT's Sunyaev-Zel'dovich survey to constrain the modified gravity parameterisation and find that these experiments can easily distinguish between models with a cosmological constant and modified gravity, if prior constraints from CMB temperature and polarisation anisotropies are included. ",
        "introduction": "The late-time accelerated expansion of the universe, observed in the magnitude redshift relation of type Ia supernovae, can be explained by three very different theories, either by assuming that the expansion is due to the cosmological constant $\\Lambda$, or due to dark energy as an additional fluid, or due to a modification in the gravitational law. Although there has been rapid progress in developing modified theories of gravity to explain the late-time accelerated expansion of the universe, none of the models is yet convincing. The Dvali-Gabadadze-Porrati (DGP) model \\citep{2000PhLB..485..208D} for example, which is based on a five-dimensional brane world set-up, offers a simple explanation of the accelerated expansion of the universe from 5d physics without dark energy, but it has been shown that the model suffers from a ghost instability \\citet{2007arXiv0709.2399K}. Simple modifications of the 4d Einstein-Hilbert action such as $f(R)$ gravity violate solar system constraints unless $f(R)$ is fine-tuned \\citep{2003PhLB..575....1C}. Structure formation with DGP-gravity was investigated by \\citet{2004PhRvD..69d4005L, 2004PhRvD..69l4015L}, and constraints on DGP gravity have been derived (CMB: \\citet{2007PhRvD..75f4003S}, supernovae: \\citet{2007JCAP...05....3R}, large-scale structure: \\citet{2006PhRvD..74b3004M, 2006JCAP...01..016K}; a review is provided by \\citet{2007arXiv0706.1557K}).  However, we have learnt general features of modified gravity from these attempts: Firstly, any modification of general relativity (GR) on cosmological (horizon) scales results in modifications of GR on sub-horizon scales as well, which influences the growth of perturbations. Whereas in cosmologies with homogeneous dark energy the growth function is uniquely determined by the Hubble-function, modified gravity models generically alter the growth function and it becomes impossible to reconcile measurements of the expansion and growth rates resulting from a modified gravity model with a dark energy model with standard gravity. This leads to the interesting possibility of distinguishing the modified gravity models from dark energy models using probes of the cosmological large-scale structure such as the CMB or the integrated Sachs-Wolfe effect. Secondly, any successful modified gravity models should satisfy the stringent constraints coming from the local tests of gravity. This effectively means that the theory must approach GR on small scales. The DGP-model naturally recovers GR on small scales without introducing any additional phenomena due to the strong coupling. An example of this is $f(R)$-gravity, where a very contrived form of the $f(R)$-function (the so-called Chameleon mechanism) is necessary to recover GR \\citep{2007arXiv0708.1190H}.  An interesting probe of the growth of structure is provided by weak lensing, but the non-linear transition to GR is important for the correct calculation of the weak lensing signal. The common approach in the literature is to parameterise the linear growth rate and to use the non-linear mapping of the linear matter power spectrum developed in GR to calculate nonlinear growth and the nonlinear lensing signal. This is effectively equivalent to assuming that gravity is modified even on small scales which would contradict the local gravity experiments. Recently, attempts have been made to carry out $n$-body simulations in modified gravity \\citep{2006PhRvD..74h4007S, 2007arXiv0709.0307L}, but the mechanism to recover GR due to non-linear dynamics was not properly implemented. An interesting phenomenological attempt was made to address this non-linear transition to GR using the halo-model \\citep{2007arXiv0708.1190H}, but it is still an open question to calculate the non-linear power spectrum.  In this paper we investigate constraints on modified gravity as an alternative to dark energy \\citep{2003astro.ph..1510D, 2006JCAP...03..017K} from structure formation and from the abundance of clusters of galaxies by solving the homogeneous growth and the spherical collapse equations. As the model for modified gravity we employ a phenomenological model which is able to interpolate between $\\Lambda$CDM and DGP-gravity, and predict the number density of clusters of galaxies in a Press-Schechter approach. Due to the spherical symmetry, it is possible to address the non-linear dynamics of the gravity in a tractable way. Similarly to the case of dark energy \\citep{1991MNRAS.251..128L, 2003MNRAS.344..761M}, cluster counts are a promising tool for investigating modified gravity as the collapse threshold for objects forming at late times is lowered, which gives rise to an enhanced spatial density of clusters. The influence of dark energy on cluster formation was extensively studied \\citep[][ to name but a few]{1998ApJ...508..483W,2005A&A...443..819P, 2006A&A...454...27B, 2005JCAP...07..003M}, and we aim to forecast constraints on modified gravity from cluster count experiments.  This paper is structured as follows: The theory is developed in Sect.~\\ref{sect_theory}, where our model for modification of gravity is outlined. The results are compiled in Sect.~\\ref{sect_homogeneous} for homogeneous cosmology, in Sect.~\\ref{sect_collapse} concerning the solution of the spherical collapse equations, and in Sect.~\\ref{sect_number_counts} the impact of modified gravity on cluster number counts and merging rates is discussed. Constraints on cluster formation with modified gravity resulting from future Sunyaev-Zel'dovich surveys are given in Sect.~\\ref{sect_szsurvey}. A summary in Sect.~\\ref{sect_summary} concludes the paper. Parameter choices for the fiducial cosmological model are based on the spatially flat $\\Lambda$CDM cosmology with adiabatic initial perturbations: Specific choices for the cosmological parameters are $\\Omega_m=0.25$, $H_0=100\\: h\\:\\mathrm{km}\\:\\mathrm{s}^{-1}\\:\\mathrm{Mpc}^{-1}$ with $h=0.72$, $\\Omega_b=0.04$, $n_s=1$ and $\\sigma_8=0.8$. ",
        "conclusions": "\\label{sect_summary} In this paper, the influence of modified gravity on Friedmann dynamics, the homogeneous growth of the cosmic density field and the formation of galaxy clusters is considered. The alteration of the gravitational law is motivated by DGP-gravity, and introduced as an additional term in the Friedmann equations, which are consistently applied to the background dynamics as well as for the formation of individual clusters. A parameter $\\alpha$ allows interpolation between the two limiting cases: gravitation with a cosmological constant $\\Lambda$ and DGP-gravity. \\begin{itemize} \\item{The Friedmann dynamics of the universe is remarkably little affected by the modification of gravity, as the Hubble function $H(a)$ for our range of models differs by at most 10\\% at a scale factor of $a=0.5$. This translates to comparable variations on the evolution $\\dd V/\\dd z$ of volume with redshift and the evolution $\\dd t/\\dd z$ of cosmic time with redshift.}  \\item{The growth function shows significant suppression of growth at early times for modified gravity, very similar to dark energy models with evolving equation of state, where the first derivative of the equation of state parameter $w(a)$ is positive at zero redshift. The largest spread between the models considered here occurs around $a=0.3$. The changed behaviour of the growth function is both caused by the decrease in matter density $\\Omega_m(a)$ as well as by the scaling of the Hubble-function with scale factor.}  \\item{The spherical collapse equations can be solved numerically for the modified gravity models motivated by DGP, resulting in the linear overdensity for spherical collapse $\\delta_c$ as a function of collapse redshift $z$. One notices a significant decrease in $\\delta_c$ by almost 20\\% compared to SCDM. The range of values for $\\delta_c$ is quite comparable to those obtained in collapse with early dark energy. An empirical fitting formula of the shape $\\delta_c(z)=a-b\\exp(-cz)$ was found to yield a good approximation.}  \\item{Cluster number counts are significantly enhanced by modified gravity, due to the lower values of the linear overdensity for spherical collapse $\\delta_c$ at low redshifts. At a redshift of $z=0.3$, where the majority of PLANCK SZ-clusters will be situated, as much as 50\\% more clusters with masses of $10^{14}M_\\odot/h$ have formed, whereas at a redshift of $z=0.7$, corresponding to the peak in the cluster redshift distribution forecast for SPT, number counts might be 70\\% larger. At higher masses of $10^{15}M_\\odot/h$, the increase in cluster number density amounts to factors of three (at $z=0.3$) up to six (at z=0.7), but the sparsity of very massive objects at high redshift will make this difficult to observe. At low masses, on the contrary, there is hardly a difference between DGP-like modified gravity models and $\\Lambda$CDM, which is due to the fact that these objects form early well in the matter-dominated phase. The enhancement in cluster number density depends exponentially on the modified gravity parameter $\\alpha$.}  \\item{Merger densities exhibit a similar behaviour at high masses of $10^{15}M_\\odot/h$, where the model predics an increase by factors of roughly up to two at the redshifts considered. At low masses, however, the merger density falls below the prediction of $\\Lambda$CDM, which is carried by the unique weighting of the merger density with the evolution of cosmic time with redshift $\\dd t/dz$, the evolution of comoving volume $\\dd V/\\dd z$ and the Hubble function $H(a)$.}  \\item{The impact of modified gravity on cluster formation is most notable at small redshifts below unity, making it a target for Sunyaev-Zel'dovich cluster surveys. In this paper the cluster yield of the SZ-surveys to be carried out by PLANCK and SPT was modelled with a simple flux threshold as the detection criterion. The Fisher-matrix analysis with   $\\Lambda$CDM as the fiducial cosmological model found that on themselves, the SZ-surveys alone are able to constrain the modified gravity parameter $\\alpha$ with an accuracy of $\\Delta\\alpha\\simeq0.47$, albeit with strong degeneracies. Adding in CMB constraints on $\\Omega_m$, $\\sigma_8$ and $n_S$ significantly lowered the uncertainty to $\\Delta\\alpha=0.04$, making the combined SZ- and CMB-observations a powerful observational probe.} \\end{itemize}  As last points is it worth noting that the above described model for phenomenological DGP-gravity has the same number of parameters as $\\Lambda$CDM and would not be disfavoured on grounds on Bayesian model selection. The feature that cluster formation takes place at higher redshifts compared to $\\Lambda$CDM might have important implications on the arclet problem in gravitational lensing \\citep{2005A&A...442..413M} and on the CBI-anomaly \\citep{2001ApJ...549L...1P, 2005ApJ...626...12B} as it naturally causes a higher level of Sunyaev-Zel'dovich background due to high-redshift clusters with a low contemporary value of $\\sigma_8$."
    },
    "0711/0711.3635_arXiv.txt": {
        "abstract": "{ We discuss the GZK horizon of protons and present a method to constrain the injection spectrum of ultrahigh energy cosmic rays (UHECRs) from supposedly identified extragalactic sources. This method can be applied even when only one or two events per source are observed and is based on the analysis of the probability for a given source to populate different energy bins, depending on the actual CR injection spectral index. In particular, we show that for a typical source density of $4\\times 10^{-5}\\,\\mathrm{Mpc}^{-3}$, a data set of 100 events above $6\\times 10^{19}$~eV  allows one in 97\\% of all cases to distinguish a source spectrum $dN/dE\\propto E^{-1.1}$ from one with $E^{-2.7}$ at 95\\% confidence level. }  \\PACS{98.70.Sa   % }  \\begin{document} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.3897_arXiv.txt": {
        "abstract": "We report V, R, and I band CCD photometry of the radio galaxy 3C 390.3 obtained with the 1.56-m telescope of the Shanghai Astronomical Observatory from March 1995 to August 2004. Combining these data with data from the literature, we have constructed a historical light curve from 1894 to 2004 and searched for periodicities using the {\\it CLEANest} program. We find possible periods of $8.30\\pm1.17$, $5.37\\pm0.49$, $3.51\\pm0.21$, and $2.13\\pm0.08$ years. ",
        "introduction": "Active galactic nuclei (AGNs) can be divided into two classes: radio-loud and radio-quiet. A small subgroup of radio-loud AGNs have displayed flux variability on time scales ranging from hours to decades (see a review by Fan 2005). Examples of long-term variations can be found in the light curves of 3C 120 (Jurkevich et al. 1971), OJ 287 (Sillanp$\\rm\\ddot{a}\\ddot{a}$ et al. 1988), 3C 345 (Kidger et al. 1992), PKS 0735+178 (Pollock, 1975, Fan et al. 1997, Qian \\& Tao 2004), BL Lacertae (Fan et al. 1998), 3C 66A (Lainela et al. 1999), 3C 279 (Fan 1999), S5 0716+714 (Qian et al. 2002), Mrk 335 (Tao et al. 2004), PKS 1510-089 \\& MA 0829+047 (Liller \\& Liller 1975), B2 1101+38 (Miller 1975) and in various other radio-selected BL Lacertae objects (Fan et al. 2002).  3C 390.3 (z=0.0561; Osterbrock et al. 1975) is a well-known broad-line radio galaxy with broad double peaked emission line profile (Sandage 1966; Lynds 1968). It is also an extended double-lobed FR II radio source. Leahy \\& Perley (1995) discussed the jets and hotspots of it and they also provided some basic data of this source. The VLBI observations at 5 GHz show evidence of superluminal motions, with $v/c\\approx4.0$ (Alef et al. 1996). 3C 390.3 has a well known history, both in the continuum light and in emission lines(e.g., Dietrich et al. 1998; Shapovalova et al. 2001; Sandage 1967, 1973; Zheng 1996). Prieto \\& Kotilainen(1997) detected optically-resolved emission coincident with the brightest region of the northern radio lobe (hot spot B) in B, R, and I bands. Dunn et al.(2006) provided a database of Ultraviolet continuum light curves for AGNs, and 3C 390.3 is one of them. {\\it ROSAT PSPC} of 3C 390.3 show resolved X-ray emission spatially located at the position of the northern radio hotspot(Prieto 1997). It is also a highly variable X-ray source (Inda et al. 1994; Eracleous et al. 1996; Wozniak et al. 1998; Leighly et al. 1997; Leighly \\& O'Brien 1997; Gliozzi et al. 2003). Gaskell(2006) discussed the relationship between the X-ray, UV, and optical variability in 3C 390.3. The narrow-line region in 3C 390.3 is probably on the order of 1 lt yr, and a significant part of the NLR is located near the line of sight(Zheng et al. 1995). The black hole mass of 3C 390.3 has been estimated through reverberation mapping (e.g., McLure \\& Dunlop 2001; Botte et al. 2004; Shapovalova et al. 2001), ranging from $2.2\\times 10^{8}$ to $2.1 \\times 2.1^{9} M_{\\bigodot}$ .  3C 390.3 has been one of our primary targets for long-term monitoring in the optical band at the Shanghai Astronomical Observatory(SHAO). In this paper, we present V, R, and I data taken in the period from 1995 to 2004. We also compile historical data from 1894 to 2004. We obtain a total of 786 observational points, to which we adopt the $CLEANest$ method to search for periodicities. ",
        "conclusions": "AGNs are variable throughout the whole electromagnetic spectrum. In our monitoring program, we have observed the galaxy 3C 390.3 from March 1995 to August 2004. The observations clearly show that the source is variable in the optical band with the variation amplitude of 1$^{m}$.372, 1$^{m}$.240, and 1$^{m}$.284 in the V, R, and I bands respectively.  The V band historical light curve is compiled, which has a time span of 111 years. Possible periods of $8.30\\pm1.17$, $5.37\\pm0.49$, $3.51\\pm0.21$, and $2.13\\pm0.08$ years were found in the light curve by means of $CLEANest$ method.  It has been suggested that the long-term periodic outbursts of OJ 287 may be explained by a binary black hole (Sillanp$\\rm\\ddot{a}\\ddot{a}$ et al. 1988) or such outbursts may be due to thermal and viscous instabilities in a thin accretion disk (Meyer \\& Meyer-Hofmeister 1984; Horiuchi \\& Kato 1990). The historical light curve of 3C 390.3 shows strong variability. Several possible periods were found by $CLEANest$ method. The multiple periods we derived may imply the instabilities in the disk. The analysis of spectra of 3C 390.3 covering a period of over 20 yr may indicate the binary black hole model(Gaskell 1996). Shapovalova et al.(2001) have monitored this object between 1995 and 2000. Their results do not support either the models of outflowing biconical gas streams or those of supermassive binary black holes. They conclude that they favor the accretion disk model."
    },
    "0711/0711.4479_arXiv.txt": {
        "abstract": "The generation of parallel electric fields by the propagation of ion cyclotron waves (with frequency 0.3 $\\omega_{ci}$) in the plasma with a transverse density inhomogeneity was studied. Using two-fluid, cold plasma linearised equations, it was shown for the first time that, in this particular context,  $E_{\\parallel}$ generation can be understood by an analytic equation that couples $E_{\\parallel}$  to the transverse electric field of the driving ion cyclotron wave. It was proven that the minimal model required to reproduce the previous kinetic simulation results of $E_{\\parallel}$ generation [Tsiklauri et al 2005, G\\'enot et al 2004] is the two-fluid, cold plasma approximation in the linear regime. By considering the  numerical solutions it was also shown that the cause of $E_{\\parallel}$ generation is  the electron and ion flow separation induced by the transverse density inhomogeneity. We also investigate how $E_{\\parallel}$ generation is affected by the mass ratio and found that amplitude attained by $E_{\\parallel}$ decreases linearly as inverse of the mass ratio $m_i/m_e$. For realistic mass ratio of $m_i/m_e=1836$, such empirical scaling law, within a time corresponding to 3 periods of the driving ion cyclotron wave, is producing  $E_{\\parallel}=14$ Vm$^{-1}$ for solar coronal parameters. Increase in mass ratio does not have any effect on final parallel (magnetic field aligned) speed attained by electrons. However, parallel ion velocity decreases linearly with inverse of the mass ratio $m_i/m_e$. These results can be interpreted as following: (i) ion dynamics plays no role in the $E_{\\parallel}$ generation; (ii) $E_{\\parallel} \\propto 1/m_i$ scaling is caused by the fact that $\\omega_d = 0.3 \\omega_{ci} \\propto 1/m_i$ is decreasing with the increase of ion mass, and hence the electron fluid can effectively \"short-circuit\" (recombine with) the slowly oscillating ions, hence producing smaller $E_{\\parallel}$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.1892_arXiv.txt": {
        "abstract": "% We present a middleware for visualization and exploration of complex datasets in a VO framework, that performs interaction between data archives and existing VO client applications using PLASTIC. It comprises: (1) PLASTIC-enabled Java control applet, integrated into archive web-pages and interacting with VO applications; (2) cross-browser compatible JavaScript part managing PLASTIC-aware VO Clients (launch, data manipulation) by means of Java LiveConnect. This (or similar) solution is an essential for the new generation VO-enabled data archives providing access to complex observational and theoretical datasets (3D-spectroscopy, N-body simulations, etc.) through web-interface. Thanks to PLASTIC capabilities it is possible to start all necessary client software with a single-click in the archive query result page in a web-browser. This simplifies the scientific usage of the VO resources and makes it easy even for users with no experience in the VO technologies. ",
        "introduction": "After several years of intensive development, the International Virtual Observatory has accumulated numerous useful standards, tools, and access protocols allowing users to deal with the astronomical data of almost any kind. The fact that many software products within the Virtual Observatory follow IVOA conventions and standards opens a very promising perspective to attempts of integrating heterogeneous VO tools within end-user's unique environment. One can glue VO tools developed by third parties, providing various feature sets into a VO system based on intercommunication between them, that will dramatically improve usability and/or feature list available to scientific users.  Being somewhat in a contradiction to the ``classical'' development schema, this approach brings all the power of the VO initiatives to data centres and other service providers involved in development of new services for astronomical needs.  The idea described here is a cornerstone of the approach to visualization of complex observational and theoretical datasets. We have implemented it in several VO-enabled data archives significantly increasing the value that ``usual'' archive can bring to an astronomer. ",
        "conclusions": "The presented solution for simplifying retrieval of and manipulation with the datasets available through VO-enabled data archives is proved to be versatile by two implementations providing access to data of completely different origin. This (or similar) solution is an essential for modern new generation data archives. It greatly simplifies the scientific usage of the VO resources and makes it easy even for astronomers with no experience in the VO technologies."
    },
    "0711/0711.4686_arXiv.txt": {
        "abstract": "We compare the gravitational potential profiles of the elliptical galaxies NGC 4486 (M87) and NGC 1399 (the central galaxy in the Fornax cluster) derived from X-ray and optical data. This comparison suggests that the combined contribution of cosmic rays, magnetic fields and micro-turbulence to the pressure is $\\sim$10\\% of the gas thermal pressure in the cores of NGC 1399 and M87, although the uncertainties in our model assumptions (e.g., spherical symmetry) are sufficiently large that the contribution could be consistent with zero.  In the absence of any other form of non-thermal pressure support, these upper bounds translate into upper limits on the magnetic field of $\\sim$10-20 $\\mu$G at a distance of $1'$-$2'$ from the centers of NGC1399 and M87.  We show that these results are consistent with the current paradigm of cool cluster cores, based on the assumption that AGN activity regulates the thermal state of the gas by injecting energy into the intra-cluster medium.  The limit of $\\sim$10-20\\% on the energy density in the form of relativistic protons applies not only to the current state of the gas, but essentially to the entire history of the intra-cluster medium, provided that cosmic ray protons evolve adiabatically and that their spatial diffusion is suppressed. ",
        "introduction": "Both optical and X-ray data are often used to determine the distribution of the gravitating mass in galaxies.  In analysing optical stellar-kinematic data, stars are treated as collisionless particles and the mass distribution is obtained either from the Jeans equations and some assumption about the anisotropy of stellar orbits (Binney \\& Tremaine 2008; Lokas \\& Mamon 2003), or by finding the potentials in which a distribution of orbits reproduces the observed surface brightness profile and stellar kinematics (e.g., Kronawitter et al. 2000; Thomas et al. 2007). In X-rays, hydrostatic equilibrium is usually assumed for the gaseous atmosphere and the observed density and temperature distributions are then used to evaluate the distribution of gravitating mass (Mathews 1978; Forman, Jones \\& Tucker 1985; Fukazawa et al. 2006; Humphrey et al. 2006).  Ideally both methods should yield identical results, and the mismatch between gravitating masses derived from optical and X-ray data can be used, for example, to measure the contribution of non-thermal particles to the gas pressure. Such non-thermal particles are directly observed as bubbles of relativistic plasma (e.g., Boehringer et al. 1993). They can also be present in the ICM due to mixing of thermal and relativistic plasma, or generated by shocks propagating through the ICM. Below we use the newest Chandra data to compare optical and X-ray data in two giant elliptical galaxies. The $\\sim0''.5$ angular resolution of Chandra has brought the X-ray data on a par with optical data and has made it possible to extend the comparison between the kinematics of stars and hot gas from spatial scales of arcseconds up to $\\sim 10'$ as is done below for two giant elliptical galaxies: M87 and NGC1399.   The fundamental assumption that underlies this paper is that the mass profiles derived from optical observations are correct, so that deviations between the optical and X-ray profiles are due to effects on the gas such as non-thermal pressure. While measuring mass profiles from stellar-kinematic data involves only gravity and dynamics, uncertainties may remain, mainly because of the degeneracy between radial variations in mass-to-light ratio and radial variations in the anisotropy parameter of the velocity ellipsoid (Binney \\& Mamon 1982). Breaking this degeneracy is possible when the kinematic data provide sufficient information on higher-order moments of the line-of-sight velocity distribution (LOSVD) and this is used to constrain the distribution of stellar orbits (Gerhard 1993, Merritt 1993). Uncertainties resulting from modeling good quality absorption line spectra are $\\sim 10\\%$ in circular velocity over the range of radii covered (e.g., Kronawitter et al. 2000; Thomas et al. 2007, and references therein). Additional reasons to believe that the optical mass profiles that we use in this paper are reliable are: (i) the profiles from M87 and NGC1399 are derived from quite different methods--- based on stellar absorption line spectra but including also discrete tracer particles (globular clusters) in the former case---yet yield consistent results for the level of non-thermal pressure; (ii) for both galaxies, the mass distributions have been studied in at least two independent studies that yield consistent results; (iii) most of the concerns about the reliability of optical mass profiles arise from studies with limited observational resolution (e.g., determinations of black-hole masses), limited radial coverage of data, or restrictive modeling techniques (e.g., the assumption of isotropic velocity distributions), neither of which is a concern in the studies used here. ",
        "conclusions": "Using Chandra observations of M87 and NGC 1399 we derive the gravitational potential profiles of these two galaxies and compare them with the potentials derived from optical data. Both the X-ray and optical methods -- one based on the hydrostatic equilibrium of hot gas and the other on stellar dynamics -- give consistent changes of the potential to within $\\sim$10-20\\% over a broad range  in radial distance. In M87, several obvious wiggles in the X-ray derived profile can plausibly be attributed to a spherical shock propagating through the ICM.   These results suggest that the combined contribution of cosmic rays, magnetic fields and micro-turbulence $\\lesssim$ 10-20\\% of the gas thermal energy. These findings are consistent with the current paradigm of cool cluster cores, based on the assumption that AGN activity regulates the thermal state of the gas by injecting energy into the ICM.  We also show that the limit of $\\sim$10-20\\% on the energy density in the form of relativistic protons applies not only to the current state of the gas, but essentially to the entire history of the ICM, provided that cosmic ray protons evolve adiabatically and their spatial diffusion is suppressed.   The comparison of $\\varphi_X$ and $\\varphi_{opt}$ is presented here for two objects only. Given the uncertainties in optical and X-ray modelling and peculiarities of the considered objects (e.g. notable shock in M87) it is desirable to extend the analysis to a larger sample. This project is under way and the results will be reported in subsequent publications."
    },
    "0711/0711.0257.txt": {
        "abstract": "As a first step for studying star formation in the extreme outer Galaxy (EOG), we obtained deep near-infrared ($J, H, K$-bands) images of two embedded clusters at the northern and southern CO peaks of Cloud 2, which is one of the most distant star forming regions in the outer Galaxy (galactic radius $R_g \\sim 19$ kpc).  With high spatial resolution (FWHM $\\sim 0\\farcs3$--$0\\farcs35$) and deep imaging ($K \\sim 21$ mag, 5 $\\sigma$) with the IRCS imager at the Subaru telescope, we detected cluster members with a mass detection limit of $< 0.1 M_\\odot$, which is well into the substellar regime.  These high quality data enables a comparison of EOG to those in the solar neighborhood on the same basis for the first time.  Before interpreting the photometric result, we have first constructed the NIR color-color diagram (dwarf star track, classical T Tauri star (CTTS) locus, reddening law) in the Mauna Kea Observatory filter system and also for the low metallicity environment since the metallicity in EOG is much lower than those in the solar neighborhood. The estimated stellar density suggests that an ``isolated type'' star formation is ongoing in Cloud 2-N, while a ``cluster type'' star formation is ongoing in Cloud 2-S. % Despite the difference of the star formation mode, other characteristics of the two clusters are found to be almost identical: (1) $K$-band luminosity function (KLF) of the two clusters are quite similar, as is the estimated IMF and ages ($\\sim 0.5$--1 Myr) from the KLF fitting, % (2) the estimated star formation efficiencies (SFEs) for both clusters are typical compared to those of embedded clusters in the solar neighborhood ($\\sim 10$ \\%). % The similarity of two independent clusters with a large separation ($\\sim 25$ pc) strongly suggest that their star formation activities were triggered by the same mechanism, probably the supernova remnant (GSH 138-01-94) as suggested in Kobayashi et al. and Yasui et al. ",
        "introduction": "\\label{sec:intro}  %%%%% ENGLISH %%%%% %{\\footnotesize FOG, EOG \u00a4\u00ce\u00c6\u00c3\u00c4\u00a7} \\\\ The extreme outer Galaxy (hereafter ``EOG''), defined here as the region at Galactic radius ($R_g$) greater than 18 kpc, could be regarded as the outermost region of the Galaxy because the distribution limits of Population I and II stars are known to be 18--19 kpc and 14 kpc, respectively \\citep{Digel1994}. The EOG is known to have very low gas density \\citep{{BW95},{Nakanishi03}} and low metallicity \\citep{Smartt1997}. Such an extreme environment could serve as a ``laboratory'' for studying the star formation process in a totally different environment from that in the solar neighborhood. % The star-forming process in such environments has not been studied extensively as for the nearby star-forming regions because of the large distance. The EOG does not seem to have long and complex star formation history because there is little or no perturbation from the spiral arms, which is the major and powerful trigger for continuously converting a large amount of molecular gas into stars. %constantly like an endless star-forming ``factory''. Therefore, in the EOG there is a good chance to observe the detail of the each star formation process, such as the triggering process of cluster formation, without being disturbed by the entangled star formation history in space and time.   %{\\footnotesize \u00a4\u00b3\u00a4\u00ec\u00a4\u00de\u00a4\u00c7\u00a4\u00ceFOG, EOG \u00a4\u00c7\u00a4\u00ce\u00c0\u00b1\u00c0\u00b8\u00c0\u00ae\u00a4\u00ce\u00b8\u00a6\u00b5\u00e6 (\u00ca\u00ac\u00bb\u00d2\u00b1\u00c0)} \\\\ % There are some synthetic studies of molecular clouds in EOG and far outer Galaxy (hereafter ``FOG''), whose $R_g$ is not more than 18 kpc but no less than 15 kpc. The pioneering survey of molecular clouds and star formation beyond the solar circle was undertaken by \\citet{WB89}, who identified the $^{12}$CO molecular emission associated with color-selected IRAS point sources and used kinematic distances to place the sources within the Galaxy. Many of the IRAS sources were found to lie in the far outer Galaxy. Follow-up studies of the gas properties of many of these far outer Galaxy molecular clouds were presented in \\citet{{BW94},{BW95}} separately.  In an independent search, \\citet{Digel1994} found eleven molecular clouds in the EOG, based on CO observations of distant \\ion{H}{1} peaks in the Maryland-Green Bank survey \\citep{WW1982}. Following those pioneering works, \\citet{Heyer98} conducted a comprehensive CO survey of molecular clouds in the FOG with Five College Radio Astronomy Observatory (FCRAO) CO survey.  With H$\\alpha$ and radio continuum study, \\citet{Fich84}, \\citet{BrandB93}, and \\citet{Rudolph96} identified a number of \\ion{H}{2} regions associated with the FOG clouds, indicating the presence of recent massive star formation.       %{\\footnotesize \u00a4\u00b3\u00a4\u00ec\u00a4\u00de\u00a4\u00c7\u00a4\u00ceFOG, EOG \u00a4\u00c7\u00a4\u00ce\u00c0\u00b1\u00c0\u00b8\u00c0\u00ae\u00a4\u00ce\u00b8\u00a6\u00b5\u00e6 (\u00c0\u00b1): Cloud2\u00a4\u00ce\u00ce\u00e3} \\\\ % For the study of star formation, it is ultimately necessary to observe YSOs (young stellar objects) at near infrared (NIR) wavelengths. \\citet{KT2000} found YSOs associated with a molecular cloud in EOG for the first time with a wide-field NIR survey. This cloud is denoted as Cloud 2 in the list of \\citet{Digel1994} and is one of the most distant molecular clouds in the outer Galaxy region.  It is located at a very large Galactic radius ($R_g = 18$--28 kpc) in the second quadrant of the Galaxy ($l \\sim 137.75^\\circ$, $b \\sim -1.0$). % It is by far the largest molecular cloud among those found in EOG with a mass comparable to that giant molecular clouds ($M \\sim 2 \\times 10^4$ $M_\\odot$; Digel et al. 1994). % The Galactic radius of Cloud 2 has been estimated by various methods at $R_g = 28$ kpc (heliocentric distance: $D=21$ kpc) from the kinematic distance of the CO cloud \\citep{Digel1994}; $R_g = 23.6$ kpc ($D=16.6$ kpc) from the latest \\ion{H}{1} observation \\citep{Stil2001}; and $R_g = 15$--19 kpc ($D=8.2$--12 kpc) from the detailed optical study of a B-type star, MR1 \\citep{{Muzzio1974},{Smartt1996}}, which is associated with Cloud 2 \\citep{deGeus93}.  Throughout this paper we adopt $R_g = 19$ kpc ($D=12$ kpc) because it is about the mean value of all the estimated distances and because spectroscopic distance of stars should be more accurate than kinematic distance. % Metallicity at the Galactic radius of Cloud 2 is estimated to be $\\sim -1$ dex from the standard metallicity curve by \\citet{Smartt1997}. In fact, the metallicity of the B-type star, MR1, is measured at $-0.5$ to $-0.8$ dex by optical high-resolution spectroscopy \\citep{{Smartt1996},{Rolleston2000}} and that of the Cloud 2 itself is measured at $-0.7$ dex from the radio molecular emission lines \\citep{{Lubowich2004}, {Ruffle07}}. These metallicity values are comparable to that of LMC ($\\sim -0.5$ dex) and SMC \\citep[$\\sim -0.9$ dex;][]{Arnault}. % Using deeper near-infrared imaging with QUIRC at the University of Hawaii (UH) 2.2m telescope, Kobayashi et al. (2007) also found two embedded clusters at the northern and southern CO peaks of Cloud 2. They suggested that the overall star formation activity in Cloud 2 was triggered by the supernova remnant (SNR) \\ion{H}{1} shell, GSH 138-01-94, and that supernova triggered star formation could be one of the major star formation modes in EOG.     Recently Brand \\& Wouterloot (2007) reported comprehensive study of a star forming region, WB89-789, whose distance is estimated at $D = 11.9$ kpc, $R_g \\sim 20.2$ kpc. This region was found in \\citet{WB89} described above.  They obtained NIR imaging with limiting magnitude of $m_K = 16.5$ and 17.5 mag for 10 $\\sigma$ and 5 $\\sigma$, respectively, which correspond to a 2--3 $M_\\odot$ dwarf main sequence star with 10 mag of visual extinction.  They also used molecular line and dust continuum observations for properties of the molecular cloud and showing the existence of an outflow.  They also obtained optical spectroscopy of a star for confirmation of the distance to the region. % % \\citet{Santos2000} discovered two distant embedded young clusters in the FOG ($R_g = 16.5$ kpc; $D=10.2$ kpc) with near-infrared imaging of CO clouds associated with IRAS sources. Their detection limit ($m_K = 16.4$ mag) corresponds to a 1 Myr old star of $\\sim$1--2 $M_\\odot$ seen through 10 mag of visual extinction. % \\citet{Snell2002} have conducted a comprehensive NIR survey of embedded clusters in the FOG ($R_g = 13.5$--17.3 kpc) based on the FCRAO CO survey \\citep{Heyer98}. They found 11 embedded clusters with the detection limit of $m_{K'}$=17.5 mag. For their most distant cluster, this magnitude corresponds to a 1 Myr old, 0.6 $M_\\odot$ star with no extinction, or a 1 Myr old, 1.3 $M_\\odot$ star seen through 10 mag of visual extinction. % These independent near-infrared searches so far found cluster members with the mass detection limit of $\\sim 1$ $M_\\odot$.     As a next step to study the details of the star formation activity in the EOG, we obtained deep and high spatial resolution NIR images of the two most distant embedded clusters in Digel Cloud 2 with Subaru 8.2-m telescope.  Because of the achieved sensitivity to detect substellar mass members, % we could compare embedded clusters in EOG to those in the solar neighborhood for the first time. % In \\citet[][hereafter Paper~I]{My2006ApJ}, we presented the data for the Cloud 2-N cluster and (i) estimated the age of the cluster at less than 1 Myr by comparing the observed KLF to model KLFs based on typical IMFs, and (ii) suggested that the formation of the cluster was triggered by the SNR shell based on the age and the geometry of the cluster.  In this paper, we present data of the Cloud 2-S cluster and also re-analyze the data of Cloud 2-N cluster to investigate the fainter stellar population. % For interpreting the photometric result, we have constructed for the first time the NIR color-color diagram (dwarf track, CTTS locus, and reddening vector) in the Mauna Kea Observatory (MKO) near-infrared filter system \\citep{{Simons2002},{Tokunaga2002}}. % which is used for almost all new instrument in the world but is still % relatively new. To accommodate to the low-metallicity in EOG, we also considered the mass-luminosity relation and the color-color diagram for low metallicity stars.  The established methods here would be useful for future study of star formation in EOG and FOG, and ultimately nearby dwarf galaxies. % We then discuss the KLF and star formation efficiency in the EOG based on the results of these two clusters.        %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%   OBSERVATION AND DATA REDUCTION   %%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  %\\clearpage ",
        "conclusions": "We have conducted a pilot study of star formation in EOG using the two embedded clusters in Cloud 2. Our conclusions can be summarized as follows:   \\begin{enumerate} %1. First deep infrared observation in EOG \\\\ % \\item We have obtained deep near-infrared images of two embedded clusters in EOG, the Cloud 2 clusters, with Subaru 8.2 m telescope and the IRCS infrared imager. This is the first deep infrared imaging of star forming regions in EOG, with the mass detection limit close to the substellar regime ($< 0.1 M_\\odot$).    %2. First advanced study of IMF and SFE in EOG \\\\ % \\item Due to the faint detection-limit, we investigated the IMF and age of the clusters in EOG by identifying the peak magnitude of the KLF. The underlying IMF of the cluster down to the detection limit is not significantly different from the typical IMFs in the field and in the nearby star clusters.  For investigating the behavior of IMF (KLF) in the substellar regime, even deeper imaging is required and we might find a different IMF than the solar neighborhood.  We have estimated the age of the clusters to be $\\sim 0.5$--1 Myr. % Additionally, SFEs of the Cloud 2 clusters were comparable ($\\sim 10$ \\%) within the uncertainties to nearby embedded clusters. % Despite these identical characteristics of the Cloud 2 clusters, the estimated stellar density suggests that different types of star formation are ongoing: an ``isolated type'' star formation is ongoing in Cloud 2-N, while a ``cluster type'' star formation is ongoing in Cloud 2-S. %Despite the difference of the star formation mode, other characteristics %of the two clusters are found to be almost identical.    \\item We summarized the existing models of near-infrared characteristics of low metallicity stars to derive the mass-luminosity relation and color-color diagram for low metallicity stars. % We concluded that dwarf track and reddening vector in $JHK$ color-color diagrams do not strongly depend on the metallicity down to that of the stars in Cloud 2 ($\\sim -0.7$ dex). % Therefore, we did not include the effects of low metallicity in the following analysis using color-color diagram, such as the identification of cluster members, % ($\\S$~\\ref{sec:IDENT}), the estimation of extinction and color excess. % However, the effect could be significant for even lower metallicity and more theoretical studies and compilations of the near-infrared data for low-metallicity normal stars are needed.   %3. Establish the analysis methods for the NIR study of EOG \\\\ % \\item For interpreting the photometric result, we have constructed the NIR color-color diagram (dwarf track, CTTS locus, reddening vector) in the MKO filter system, for which colors for main sequence stars have not been well established before. We find that it is important to carefully compare colors in different photometric systems.  \\end{enumerate}    %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "0711/0711.2319_arXiv.txt": {
        "abstract": "We study the statefinder parameters in the Yang-Mills condensate dark energy models, and find that the evolving trajectories of these models are different from those of other dark energy models. We also define two eigenfunctions of the Yang-Mills condensate dark energy models. The values of these eigenfunctions are quite close to zero if the equation-of-state of the Yang-Mills condensate is not far from $-1$, which can be used to simply differentiate between the Yang-Mills condensate models and other dark energy models. ",
        "introduction": "Physicists and astronomers begin to consider the dark energy cosmology seriously and to explore the nature of dark energy actively since the expansion of our universe is proven to be accelerating at present time by the Type Ia supernovae observations\\cite{sn}. The analysis of cosmological observations, in particular of the WMAP (Wilkinson Microwave Anisotropy Probe) experiment\\cite{map}, indicates that dark energy occupies about $70\\%$ of the total energy of our universe, and dark matter about $26\\%$. The accelerated expansion of the present universe is attributed to that dark energy is an exotic component with negative pressure, and the simplest candidate for dark energy is the cosmological constant. However, two difficulties arise from this scenarios, namely the fine-tuning problem and the cosmic coincidence problem\\cite{problem}. So the dynamical models are considered by a number of authors, such as the quintessence, phantom, k-essence, and quintom models\\cite{models}. The effective Yang-Mills  condensate (YMC) as a kind of candidate for dark energy has been detailed discussed in the Refs.\\cite{z,Zhang,zhao}. The effective lagrangian up to 1-loop order is \\cite{pagels, adler} \\be \\L_{eff}=\\frac{b}{2}F\\ln\\left|\\frac{F}{e\\kappa^2}\\right|, \\label{L} \\ee where  $b=11N/24\\pi^2$ for the generic gauge group $SU(N)$ is the Callan-Symanzik coefficient\\cite{Pol}. $F=-(1/2)F^a_{\\mu\\nu}F^{a\\mu\\nu}$ plays the role of the order parameter of the YMC, $\\kappa$ is the renormalization scale, the only model parameter. The attractive features of this effective lagrangian include the gauge invariance, the Lorentz invariance, the correct trace anomaly, and the asymptotic freedom\\cite{pagels}. The effective YMC was firstly put into the expanding Friedmann-Robertson-Walker (FRW) spacetime to study inflationary expansion \\cite{z} and the dark energy \\cite{Zhang}. We work in a spatially flat FRW spacetime with a metric \\be ds^2=a^2(\\tau)(d\\tau^2-\\delta_{ij}dx^idx^j),\\label{me} \\ee where $\\tau=\\int(a_0/a)dt$ is the conformal time. For simplicity we study the $SU(2)$ group. Compared with the scalar field, the YM field is the indispensable cornerstone to particle physics and the gauge bosons have been observed. There is no room for adjusting the form of effective YM lagrangian as it is predicted by quantum corrections according to field theory. In the previous works, we have deeply investigated the 1-order YMC models and found attractive features: \\emph{a.} this dark energy can naturally get the equation-of-state (EOS) of $w>-1$ and $w<-1$\\cite{zhao}, which is different from the scalar quintessence models; \\emph{b.} in the free field models, with the expansion of the universe, the EOS of the YMC naturally turns to the critical state of $w=-1$\\cite{zhao}, consistent to the observations; \\emph{c.} the cosmic coincidence problem is naturally avoided in the models\\cite{zhao}; \\emph{d.} the EOS of the dark energy can cross $-1$ in the double-field models or coupled models\\cite{zhao}; \\emph{e.} the big rip is naturally avoided in the models\\cite{zhao}; \\emph{f.} the magnetic component of the YMC naturally decreases to zero with the expansion of the universe\\cite{zhao}.  In the recent paper\\cite{xia}, the authors have detailed discussed the 2-loop YMC dark energy. In this model, the effective lagrangian is \\be {\\mathcal{L}}_{eff}=\\frac{b}{2}F\\left[\\ln\\left|\\frac{F}{e\\kappa^2}\\right| +\\eta\\ln\\left|\\ln\\left|\\frac{F}{e\\kappa^2}\\right|+\\delta\\right|~\\right] \\ee where $\\eta\\simeq0.84$ and the dimensionless constant $\\delta$ is a parameter representing higher order corrections. In this 2-loop model, the cosmic coincidence problem is also naturally avoided. This feature is same with the 1-loop models. From the Einstein equation\\cite{xia}, we can easily find that, in the free field models, the late time attractor exists, which satisfies the relation \\be \\beta_c=-\\eta\\left(\\log|\\beta_c-1+\\delta|+\\frac{1}{\\beta_c-1+\\delta}\\right), \\ee where $\\beta=\\ln|F/\\kappa^2|$. It is easily found that $w=-1$, when $\\beta=\\beta_c$. So in these models, the EOS of YMC also naturally turns to the critical state $w=-1$. However, the cosmological constant crossing is also naturally realized, if considering the interaction between the YMC and matter \\cite{xia}. The discussion in \\cite{xia} shows that, although the 2-loop models are much more general and complicated than the 1-loop model, the most physics features  are not unchanged.   In this paper, we only consider the 1-loop models. As in previous works, we only consider the electric case with $B^2 \\equiv0 $. The energy density and pressure of YMC are given by \\be \\rho_y=\\frac{E^2}{2}\\left(\\epsilon+b\\right), ~~~~p_y=\\frac{E^2}{2}\\left(\\frac{\\epsilon}{3}-b\\right), \\ee where the dielectric constant is \\be \\epsilon=b\\ln\\left|\\frac{F}{\\kappa^2}\\right|.\\label{epsilon} \\ee and the EOS is obtained \\be w=\\frac{p_y}{\\rho_y}= \\frac{\\beta-3}{3\\beta+3},\\label{13} \\ee where $\\beta\\equiv\\epsilon/b$. At the critical point with the condensate order parameter $F=\\kappa^2$, one has $\\beta=0$ and $w=-1$. Around this critical point, $F< \\kappa^2$ gives $\\beta<0$ and $w<-1$, and $F> \\kappa^2$ gives $\\beta>0$ and $w>-1$. So in the YMC model, EOS of $w >-1$ and $w<-1$ can be naturally realized. When $\\beta\\gg1$, the YM field has a state of $w=1/3$, becoming a radiation component.      In the free field models, the effective YM equations are \\be \\partial_{\\mu}(a^4\\epsilon~ F^{a\\mu\\nu})+f^{abc}A_{\\mu}^{b}(a^4\\epsilon~F^{c\\mu\\nu})=0, \\label{F1} \\ee which reduces to\\cite{zhao} \\be \\partial_{\\tau}(a^2\\epsilon E)=0. \\ee At the critical point ($\\epsilon=0$), this equation is an identity. When $\\epsilon\\neq 0$, this equation has an exact solution \\cite{zhao}: \\be \\beta~ e^{\\beta/2}\\propto a^{-2},\\label{16} \\ee where the coefficient of proportionality depends on the initial condition. For a fixed initial condition, we can obtain the evolution of the EOS of the YMC by using the YM equation (\\ref{16}). In the previous works, we found the free YMC can be separated into two kinds, the quintessence-like or phantom-like, which only depends on the choice of the initial condition. In order to differentiate between the YMC dark energy models and other models, a sensitive and robust diagnostic for dark energy models is needed. For this purpose a diagnostic proposal that makes use of parameter pair $\\{r,s\\}$, the so-called ``statefinder\", was introduced by Sahni et al.\\cite{sahni}. The statefinder probes the expansion dynamics of the universe through higher derivatives of the expansion factor $\\stackrel{...}{a}$ and is a natural companion to the deceleration parameter $q$ which depends upon $\\ddot a$. The statefinder pair $\\{r,s\\}$ is defined: \\begin{equation} r\\equiv \\frac{\\stackrel{...}{a}}{aH^3},~~~~s\\equiv\\frac{r-1}{3(q-1/2)}~.\\label{rs} \\end{equation} The statefinder is a ``geometrical'' diagnostic in the sense that it depends upon the expansion factor and hence upon the metric describing space-time.  Trajectories in the $s-r$ plane corresponding to different cosmological models exhibit qualitatively different behaviors. The spatially flat LCDM (cosmological constant $\\Lambda$ with cold dark matter) scenario corresponds to a fixed point in the diagram $\\{s,r\\} = \\{ 0,1\\}$. Departure of a given dark energy model from this fixed point provides a good way of establishing the ``distance'' of this model from LCDM \\cite{sahni,alam}. As demonstrated in Refs. \\cite{sahni,alam,gorini} the statefinder can successfully differentiate between a wide variety of dark energy models including the cosmological constant, quintessence, the Chaplygin gas, braneworld models and interacting dark energy models. We can clearly identify the ``distance'' from a given dark energy model to the LCDM scenatio by using the $r(s)$ evolution diagram.  The current location of the parameters $s$ and $r$ in these diagrams can be calculated in models, and on the other hand it can also be extracted from data coming from SNAP (SuperNovae Acceleration Probe) type experiments \\cite{snap}. Therefore, the statefinder diagnostic combined with future SNAP observations may possibly be used to discriminate between different dark energy models. In this letter we apply the statefinder diagnostic to YMC dark energy models. ",
        "conclusions": "In summary, we have investigated the statefinder of the YMC dark energy models in this letter. We analyze two cases of the models, the quintessence-like case and the phantom-like case, and perform a statefinder diagnostic to both cases.  It is shown that the evolving trajectory of this scenario in the $s-r$ plane is quite different from those of other models. We also define two eigenfunctions of YMC dark energy model. If the EOS of the YMC is not far from $-1$, the values of the eigenfunctions are very closer to $0$, which can be used to simply differentiate between the YMC and other dark energy models."
    },
    "0711/0711.3824_arXiv.txt": {
        "abstract": "{Among the components of the infrared and submillimetre sky background, the closest layer is the thermal emission of dust particles and minor bodies in the Solar System. This contribution is especially important for current and future infrared and submillimetre space instruments -- like those of Spitzer, Akari and Herschel -- and must be characterised by a reliable statistical model.} {We describe the impact of the thermal emission of main belt asteroids on the 5...1000\\,$\\mu$m photometry and source counts, for the current and future spaceborne and ground-based instruments, in general, as well as for specific dates and sky positions.} {We used the statistical asteroid model (SAM) to calculate the positions of main belt asteroids down to a size of 1\\,km, and calculated their infrared and submillimetre brightness using the standard thermal model. Fluctuation powers, confusion noise values and number counts were derived from the fluxes of individual asteroids. } {We have constructed a large database of infrared and submillimetre fluxes for SAM asteroids with a temporal resolution of 5 days, covering the time span January 1, 2000 -- December 31, 2012. Asteroid fluctuation powers and number counts derived from this database can be obtained for a specific observation setup via our public web-interface.} {Current space instruments working in the mid-infrared regime (Akari and Spitzer Space Telescopes) are affected by asteroid confusion noise in some specific areas of the sky, while the photometry of space infrared and submillimetre instruments in the near future (e.g. Herschel and Planck Space Observatories) will not be affected by asteroids. Faint main belt asteroids might also be responsible for most of the zodiacal emission fluctuations near the ecliptic. } ",
        "introduction": "\\label{sect:intro} Due to their relatively high apparent brightness at infrared wavelength compared e.g.~to Galactic stars, asteroids are among the dominant sources at infrared wavelengths and can seriously affect the infrared and submillimetre photometry and source counts. Currently, 377\\,328 minor planets are known\\footnote{http://cfa-www.harvard.edu/iau/lists/ArchiveStatistics.html} (as of July~30, 2007) in our Solar System, of which about 99\\% are located in the main belt. On the plane of the sky the vast majority of main belt asteroids (MBAs) are found at ecliptic latitudes below $20^{\\circ}$. They have sizes between a few ten meters up to about $1000~\\mathrm{km}$. With temperatures between 200 and 300\\,K, the asteroids emit predominantly at thermal wavelengths between $5~\\mu\\mathrm{m}$ and the millimetre range. Deep infrared observations close to the ecliptic will therefore always include some of these moving targets \\citep[e.g.][]{Tedesco+Desert,Meadows2004}. Such observations also show that only a small fraction of the existing minor body population is currently known and this population might cause a non-negligible confusion noise contribution at certain wavelengths and for specific instruments. Tedesco \\& Des\\'ert (2002) measured the number of main belt asteroids for the first time in a direct way by using the Infrared Space Observatory (ISO). Based on a statistical asteroid model the authors concluded that there are about 1.2$\\pm$0.5$\\times$10$^{6}$ asteroids ($\\ge$ 1\\,km in diameter) in our solar system, twice as many as previously believed.  A set of celestial sources may affect compact source observations in two ways: (1) they contribute to the confusion noise, the uncertainty in point source photometry due to the fluctuations of the sky background; and (2) they also appear as individual sources, that may add 'false' detections to source counts.  Recently, several authors calculated confusion noise and detection limits for current/future infrared space missions (Spitzer, Akari, Herschel and SPICA). These papers considered the two major confusion noise components: the extragalactic background \\citep{Lagache2003,Negrello}, the Galactic cirrus emission \\citep{Kiss2005,Jeong2005}, or the combination of the two \\citep{Jeong2006}.  It is an important question for infrared space missions whether faint asteroids close to or below the detection limit could contribute significantly to the confusion noise of these instruments or could be present as a significant count of contaminating point sources in the field of view. To consider these asteroids, a reliable statistical model is needed. Recently, \\citet{Tedesco2005} presented the \"Statistical Asteroid Model\" (hereafter SAM). This model is based on a population of $\\sim1.9\\times 10^6$ asteroids obtained from the complete known asteroid sample (as of 1999), plus extrapolation of the size-frequency distribution (SFD) of 15 asteroid dynamical families and three background populations, to a diameter limit of $1~\\mathrm{km}$. The validity of the SAM was demonstrated by comparing SAM predictions with ISO measurements at $12~\\mu\\mathrm{m}$ \\citep{Tedesco+Desert} and Spitzer measurements at the $8$ and $24~\\mu\\mathrm{m}$ photometic bands \\citep{Meadows2004}. The asteroid counts from both deep surveys show good agreement with the SAM predictions.  In this paper we give (1) estimates for the impact of MBAs on infrared (IR) to submillimetre (especially spaceborne) measurements, such as asteroid number counts and confusion noise, for a specific date, sky position, wavelength/instrument; and (2) observable predictions for the infrared regime which can be used to check the parameters of the present model, e.g.~size, albedo and positional distribution. This is the first study that investigates the impact of asteroids on infrared and submillimetre photometry and source counts in detail, based on a realistic asteroid sample. ",
        "conclusions": "The SAM is limited to asteroids with a lower limit in size of $1~\\mathrm{km}$ in diameter. There is, certainly, a population of asteroids with sizes below this limit. Since the confusion noise is calculated using Eq.~\\ref{eq:conf_sum}, the impact of very small boides (a few hundred meters in diameter and below) is minor to the confusion noise.  To test the effect of small ($D<1~\\mathrm{km}$) asteroids on the fluctuation power and confusion noise we first created a size-frequency distribution; for asteroids of $D>1~\\mathrm{km}$ we used the number count values estimated by eq.~3 in \\citet{Tedesco2005}, while for $D\\le1~\\mathrm{km}$ the interpolated values of \\citet{Belton1992} were used. This SFD is a combination of the 'SAM' and 'Galileo team' data points in fig.~4 in \\citet{Tedesco2005}. A simple relationship of $S(\\lambda_i)\\propto D^2$ was assumed between the flux density of an asteroid's thermal emission at a specific wavelength $S(\\lambda_i)$ and its diameter $D$. Then the specific fluctuation power $\\delta F(D)$ was calculated for a $D\\pm\\Delta D$ interval, using Eq.~\\ref{eq:dF0}. \\begin{figure}[!ht] \\includegraphics[width=7.8cm]{./below1km.ps} \\caption[]{Relative contribution of asteroids with a specific size $D$ to the asteroid fluctuation power, using the size distribution and the simple flux model as described in Sect.~4.} \\label{fig:below1km} \\end{figure} It is clear from Fig.~\\ref{fig:below1km} that the fluctuation power is dominated by the bright and large asteroids, and that small ($D\\le1~\\mathrm{km}$) minor planets have a negligible contribution to the fluctuation power. { Presently orbits for almost 400\\,000 asteroids are known, i.e., more than 20\\% of the 1.9$\\cdot$10$^6$ SAM asteroid sample (see Sect.~1), but only about 0.1\\% ($\\sim$2000 asteroids) have known sizes and albedos which are the crucial parameters for the confusion noise estimates. The calculated power fluctuations are therefore dominated by asteroids which have been extrapolated in the SAM via size-frequency distributions. However, knowing only 0.1\\% is far from being statistically relevant and in some parts of the sky bright (i.e.~large) asteroids are completely missing. This implies that confusion noise estimates cannot rely on the known asteroid sample alone and the application of the SAM model for confusion noise calculations is necessary.  On the other hand, fluctuation power and confusion noise are dominated by the largest/brightest asteroids in the actual field, even if the known ones are missing. As suggested by Fig.~\\ref{fig:below1km} the contribution of $D\\le1~\\mathrm{km}$ asteroids remains negligible, and a further extension of the SAM to smaller diameters would not improve the accuracy of the confusion noise calculations significantly. }   Asteroids further out in the Solar System (e.g. trans-Neptunian objects) will not contribute considerably to the confusion noise or number counts, as deduced from their currently known size distribution in Appendix~C.  As discussed in detail in Appendix~D, SAM asteroids contribute to the absolute brightness of the zodiacal emission in a negligible level, however, a set of these asteroids may be the dominant sources of zodiacal emission fluctuation power at some specific wavelengths and sky regions."
    },
    "0711/0711.1821_arXiv.txt": {
        "abstract": "We describe here how the late time behavior of the decaying states, which is predicted to deviate from an exponential form, while normally of insignificant consequence, may have important cosmological implications in the case of false vacuum decay.  It may increase the likelihood of eternal inflation, and may help explain the likelihood of observing a small vacuum energy at late times, as well as arguing against decay into a large negative energy (anti-de Sitter space), vacuum state as has been motivated by some string theory considerations.  Several interesting open questions are raised, including whether observing the cosmological configuration of a metastable universe can constrain its inferred lifetime. ",
        "introduction": "It is not commonly known, but nevertheless well established since the work of Khalfin in 1958\\cite{khalfin},  extended by numerous other authors in the intervening years (i.e. \\cite{sudarshan,peres,patra} that for times long compared to the characteristic decay time of a metastable quantum state, the decay of such states is no longer described by an exponential, but rather by a power law.  One of the reasons that this interesting fact is perhaps not more well recognized is that the characteristic time-scale determined for departures from exponential decay in several sample systems is generally so long that for all intents and purposes all unstable states in any finite system will have decayed before this onset would become relevant.  However, if one generalizes the results from quantum mechanics to quantum field theory then the nature of false vacuum decay can altered, with potentially significant implications for cosmology, as we describe here.  We begin by exploring the generic features of late-time decay, examining the transition from quantum mechanics to quantum field theory in order to explore those implications that are independent of the details of the relevant potential or model, in order to consider their impact upon the dynamics of cosmological false vacuum decay.  We find that the time for departures from exponential growth should depend logarithmically on ratio of the energy gap and the decay width with a prefactor of order unity and then explore the cosmological implications of this.    Several open questions are then raised, including whether gravitational effects may alter the long time non-exponential tail in vacuum decay, and also whether cosmological observation within a given metastable universe may alter its lifetime. ",
        "conclusions": "The possibility that macroscopic quantum mechanical considerations may impact upon the very existence of our own universe at late times is an intriguing one.   As we have seen, at the very least they may require a dramatic alteration of probability considerations in chaotic eternal inflationary universe scenarios, and may explain why universes with unnaturally small vacuum energies may be favored at late times, even as these arguments would tend to disfavor ultimate decay into a state with large negative vacuum energy.  However, these considerations also lead us into unexplored realms associated with macroscopic quantum mechanics.     The first, which we plan to investigate, involves the question of whether late-time power law time dependence persists when gravitational effects in vacuum decay \\cite{coleman2} are taken into account.  The heuristic argument regarding spreading of the  wavefunction following barrier penetration, leaving a power law tail at the origin as the source of late time behavior suggests that the exponential expansion of space accompanying metastability may suppress this term and thus produce a faster falloff.     The second consideration is even more interesting.   Have we ensured, by measuring the existence dark energy in our own universe, that the quantum mechanical configuration of our own universe is such that  late time decay is not relevant?  Put another way, what can internal observations of the state of a metastable universe say about its longevity?"
    },
    "0711/0711.4129_arXiv.txt": {
        "abstract": "{The galaxy cluster A3667 was observed using the Two-degree Field (2dF) multifibre spectroscopic system on the Anglo-Australian Telescope in a program designed to examine the velocity structure in the region. Specifically, we sought evidence from the optical data for the putative cluster merger believed to be responsible for the observed radio and X-ray emission. We present 184 new redshifts in the region, of which 143 correspond to member galaxies of A3667. We find the cluster velocity distribution to be well modelled by a single Gaussian in agreement with previous results. In addition, new redshift-selected isodensity plots significantly reduce the prominence of the previously reported subgroup to the north-west of the main cluster. Instead, we find the galaxy distribution to be elongated and well mixed, with a high velocity dispersion and no significant evidence for substructure. These results are consistent with the axis of the proposed merger being close to the plane of the sky. ",
        "introduction": "\\label{intro}  The ACO galaxy cluster, A3667 \\citep{ACO}, is a rich, X-ray-luminous cluster, which is arguably one of the most interesting and, consequently, most studied clusters in the southern sky.  It has the distinction of being one of only a handful of galaxy clusters to have a cold gas front at the centre \\citep{Vikhlinin01} and the only cluster yet seen with two bright, diffuse, radio emission regions straddling the X-ray gas \\citep{Rottgering97}.  Our physical picture of the cluster is becoming clearer as more multiwavelength data are obtained, but many basic questions about the dynamical history of A3667 remain unanswered.  A3667 is relatively close, with a published mean redshift of 0.055 and a large velocity dispersion ($\\sigma_{\\rm v}$) of 1000--1500 km s$^{-1}$ depending on the number of galaxies used and the redshift limits applied \\citep{Melnick81, Proust88, Sodre92, Girardi98}.  The Abell richness class is 2 and the Bautz-Morgan type is I--II \\citep{ACO}, indicating a high galaxy density in this region. \\citet{Proust88} presented evidence for a bimodal spatial distribution of galaxies selected by apparent magnitude, and a similar conclusion was reached by \\citet{Sodre92}.   {\\it ROSAT\\/} observations revealed a high X-ray luminosity $L_X$, consistent with the $L_X-\\sigma$ relation \\citep{Xue00}, of $8.74 \\times 10^{44}\\, h^{2}$ erg s$^{-1}$ in the 0.4--2.4 keV range \\citep{Ebeling96}, making it one of the brightest X-ray sources in the southern sky\\footnote{$H_0=100h$ km\\,s$^{-1}$\\,Mpc$^{-1}$}. Further, the X-ray isophotes are distorted in the direction of the reported bimodal optical distribution and observations from the {\\it Einstein\\/} satellite are reported as showing clear evidence of substructure \\citep{Sodre92}. Observations of the temperature of the X-ray gas with ASCA by \\citet{Markevitch99} showed the central part of the cluster to be cool and the X-ray surface brightness profile along the elongation axis suggests a shocked region is present.  Recent Chandra observations confirmed this interpretation, revealing the central part of the cluster to contain a cold gas front moving through the warm intracluster medium \\citep{Vikhlinin01}.  Perhaps the most dramatic features of A3667 are the two extended, symmetrically located regions of diffuse radio emission \\citep{Rottgering97}, first detected in 843 MHz images obtained with the Molonglo Observatory Synthesis Telescope, and then confirmed and studied in detail at 1.4 and 2.4 GHz with the Australia Telescope Compact Array, ATCA \\citep{mjh03}.  The combination of these features has been interpreted as evidence that the cluster is observed in a post-merger state after close core passage \\citep{Rottgering97, KHB, Roettiger99, Markevitch99, Vikhlinin01}.  Given the strong indicator of merger activity present in the radio, X-ray and previous optical data we undertook a set of spectral observation of the objects in the cluster to further test for signatures of a cluster-wide merger in the optical population of A3667.  In Section 2 we discuss the signatures of mergers that have been observed in galaxy clusters in the optical and present a detailed summary of the previous optical data collected for A3667. Sections 3 and 4 discuss the Two-degree Field (2dF) observations and data reduction. The major results are discussed in Section 5 and the discussion and conclusions are presented in Section 6. ",
        "conclusions": "Cluster mergers are rarely established based solely on evidence from one wavelength regime, generally requiring a multiwavelength approach to reveal the true dynamical nature of the interactions.  The presence of diffuse radio emission in clusters of galaxies points to the existence of cluster-wide magnetic fields of the order of 0.1 --1 $\\mu$G and a population of relativistic electrons with Lorentz factor, $\\gamma \\gg $ 1000 and energy densities of 10$^{-14}-10^{-13}$ erg cm$^{-3}$ \\citep{Feretti04}. Diffuse radio emission is observed most commonly in clusters which are believed to have undergone a recent merger event and this leads naturally to the conclusion that the phenomenon is due to turbulence and shocks in the intra-cluster medium energizing the electron population \\citep{Harris80, Tribble93, Giovannini02}. Simulations of the turbulence and shocks in the cluster environment are able to predict the characteristics of observed diffuse radio emission \\citep{Roettiger99, Miniati01, Pfrommer07} and studies of the number counts suggest that the most likely lifetimes of diffuse radio emission is of the order of 1 Gyr \\citep{Kuo04}. Additionally, theoretical models of merging systems suggest that X-ray substructure will be erased after the sound crossing time and thus, clusters with disturbed X-ray isophotes are likely to be dynamically young and in on-going merger events. Finally, the strongest evidence for clusters being in recent dynamical interactions is the detection of temperature changes across the source \\citep{Govoni04, Pratt07}.  While the number of clusters containing diffuse emission is still comparatively low \\citep{Feretti04}, there are now a number of well studied systems with multiwavelength data including A2256 \\citep{Clarke06}, A754 \\citep{Henry04}, IC1262 \\citep{Hudson04} as well as the numerous examples presented by \\citet{Govoni04}. To date most analysis tends to concentrate on the radio and X-ray properties of the clusters. However, often this is insufficient to determine the merger geometry. For example, the case of A3921 \\citep{Ferrari06} required data from optical, radio and X-ray analysis to determine the dynamical state of the cluster. In this instance optical observations alone were inconclusive but taken together with the other wavelength data the merger axis could be established. Similarly, high resolution X-ray from XMM-Newton analyzed in conjunction with existing multiwavelength data has proved vital to unraveling the cluster dynamics in several cases such as A85 \\citep{Durret05} and A3921 \\citep{Belsole05}.  A3667 has a mean redshift of 0.0555, a high velocity dispersion (1102 $\\pm$ 46 km\\,s$^{-1}$) and an elongated optical axis.  It is typical of rich clusters in that it is dominated by elliptical galaxies. There is marginal evidence for a grouping of spirals in the north-western part of the cluster. The velocity distribution is well modelled by a single Gaussian peaking at $z=0.0555$, with a weak excess around a redshift of 0.0625. There seems to be little or no support for distinct in-falling groups either from the velocity information or the spatial distribution of galaxies with measured redshifts. Moreover, there is no evidence for a velocity gradient across the cluster suggested by the previously reported bimodal optical distribution. However, A3667 is a rich cluster and to date only 231 redshift have been obtained for cluster members.  A larger sample is needed to categorically rule out the presence of substructure.  Several authors have claimed that A3667 is observed in a post-merger state \\citep{Rottgering97, KHB, Roettiger99, Markevitch99, Vikhlinin01}. This may be consistent with the high velocity dispersion and mixing in the cluster. If the observed diffuse radio emission was generated via a cluster-cluster merger it is expected that the collisional velocities would be a few thousand km\\,s$^{-1}$. If such a merger occurred, the current data suggest that the merger axis must be very close to the plane of the sky. This would partly account for the lack of observable subgroup peaks in the velocity data.  Full understanding of the the dynamical situation in A3667 will require multi-wavelength analysis. We have obtained detailed new radio data for the entire cluster at 1.4, \\& 2.4 GHz, we will use this in conjunction with previously published radio (843 MHz) and X-ray data to further investigate this cluster. Findings will be published in a future paper."
    },
    "0711/0711.4773_arXiv.txt": {
        "abstract": "{} {We investigate the origin of X-rays and the  nature of accretion flow in 4 low-ionization nuclear emission-line regions (LINERs) hosted by radio galaxies, namely \\ngc, \\pks, 3C~88, 3C~444, recently observed with \\xmm.} {We combine the results from the time-averaged spectral analysis with model-independent information from X-ray temporal and spectral variability analyses, and with additional broadband information (specifically from the UV band, covered by the Optical Monitor aboard \\xmm, and from archival radio data).} {The values of the Eddington ratios $L_{\\rm bol}/L_{\\rm Edd}$ of our sample span 2 orders of magnitude ranging between $\\sim 1\\times 10^{-5}$ and $1\\times 10^{-3}$. The 4 AGN are adequately fitted by the same continuum model that comprises at least one thermal component ($kT\\sim 0.65-1.45$ keV) and a partially absorbed power law, whose relative contribution and photon index vary substantially from source to source. \\ngc\\ and \\pks\\ have fairly steep power-law components ($\\Gamma\\sim 2.5-2.9$), perhaps indicative of synchrotron emission from the base of a jet. Conversely, the flat photon index derived for 3C~88 ($\\Gamma\\sim 1.1$) may be indicative of a heavily absorbed object. Finally, the time-averaged spectral properties of \\3c4 ($\\Gamma \\sim1.9$ and an apparent line-like excess around 6.7 keV) are more in line with Seyfert-like objects. The temporal analysis reveals that \\pks\\ and 3C~88 are significantly variable in the soft (0.2--1 keV) energy band. \\pks\\ also shows suggestive evidence of spectral variability on timescales of months, with a spectral softening associated with the source brightening. \\ngc\\ is only marginally variable in the soft band, whereas \\3c4\\ does not show significant variability on short timescales. The main findings from the broadband analysis can be summarized as follows: 1) \\3c4, \\pks, and \\ngc\\ have $\\alpha_{OX}$ values consistent with the  $\\alpha_{OX} - l_{UV}$ correlation found by Steffen et al. (2006) for Seyfert-like objects. 2) No positive correlation is found between $L_{\\rm X}$ and the inclination angle, suggesting that the X-ray emission is not beamed. 3) The values of the radio-loudness are inversely proportional to the Eddington ratio and locate our objects in between the ``radio-loud'' and ``radio-quiet'' branches in the $R- l_{UV}$ plane proposed by Maoz (2007). } {} ",
        "introduction": "The recent discovery of \\hst\\ and ground-based observations that most nearby galaxies harbor supermassive black holes (e.g., Magorrian et al. 1998; Gebhardt et al. 2003), provides convincing evidence that active galactic nuclei (AGN) and normal galaxies are fundamentally connected. It is now generally accepted that there is not a sharp distinction between AGN and normal galaxies, but rather a continuous distribution of activity levels between the two extremes. Therefore, of crucial importance is the role played by low-power AGN, because they may represent the link between powerful AGN and normal galaxies. \\begin{table*}[htb] \\caption{ Source properties} \\begin{center} \\begin{tabular}{llcccccc} \\hline \\hline \\noalign{\\smallskip} Source &  Class  &   z  & $  N_{\\rm H,Gal}$ & $ \\log M_{\\rm BH} $  &  $S_{\\rm core,5GHz}$&  $S_{\\rm tot,5GHz}$ &  $\\Re$ \\\\ \\noalign{\\smallskip} &         &      & $[{\\rm 10^{20}~cm^{-2}}]$ &$[M_{\\odot}]$ & [Jy]  & [Jy]     &  $[S_{\\rm core}/S_{\\rm ext}]$ \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\noalign{\\smallskip} NGC~1692       & FRI &  0.035 &3.54 & 9.01 & 0.04 & 1.81 & 0.023 \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} PKS~0625-35   & FRI &  0.055 &7.08 & 9.19 & 0.60 & 2.12 & 0.395 \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 3C~88        & FRI/FRII & 0.030  &8.81 & 8.70 & 0.16 & 2.14 & 0.087 \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 3C~444        & FRII& 0.153  &2.61 & 9.28 & $<$0.002& 1.99 & $<$0.0009\\\\ \\noalign{\\smallskip} \\hline \\hline \\end{tabular} \\end{center} \\footnotesize {\\bf Columns Table 1}: 1= Source name. 2= Fanaroff--Riley classification. 3= Redshift. 4= Galactic column density. 5= Black hole mass estimate from Bettoni et al. (2003); the uncertainty is $\\sim$0.4 dex. 6= Core radio flux from Morganti et al. (1993). 7= Total  radio flux from Morganti et al. (1993). 8= Radio core dominance, $S_{\\rm core}/(S_{\\rm tot}-S_{\\rm core})$, from Morganti et al. (1993). \\label{tab1} \\end{table*}  The relevance of low-power AGN is bolstered by the finding that they make up a very substantial fraction of the local galaxy population. Indeed, according to the Palomar survey (Ho et al. 1997), over 40\\% of nearby galaxies contain low-power  AGN, and are optically classified as low-ionization nuclear emission-line regions (LINERs) or as ``transition objects'', i.e. objects with optical spectra intermediate between those of pure LINERs and H II regions (Heckman 1980).  In the past, several studies investigated the nature of the ionizing mechanism that produces LINER-like optical spectra (i.e., [\\oii]/[\\oiii]$>$1 and [\\oi]/[\\oiii]$>$0.33) and yielded different viable solutions. For example, Heckman (1980) and Dopita \\& Sutherland (1995) proposed shock-heated gas as main ionizing mechanism. Alternatively, the ionization can be provided by very hot stars (such as Wolf-Rayet stars) clustered around the host galaxy nucleus (e.g., Terlevich \\& Melnick 1985). Finally, it has been proposed that LINERs can be powered by low-power AGN (e.g., Filippenko \\& Halpern 1984).  Currently, the dominant ionization mechanism responsible for these unusual line ratios is still a matter of debate. Nevertheless, it is generally accepted that a sizable fraction of LINERs is associated with AGN. This conclusion is supported by the discovery of broad emission lines in many LINER spectra (e.g., Eracleous \\& Halpern 2001), by X-ray imaging and spectral studies (e.g., Terashima \\& Wilson 2003; Dudik et al. 2005), and by the fact that LINERs are often associated with compact radio sources with high brightness temperatures (Nagar et al. 2005).  However, the nature of the AGN producing LINER-like spectra is still controversial. For example, based on the study of non-simultaneous spectral energy distributions (SEDs) of seven LINERs, Ho (1999) outlines the lack of a ``big blue bump'' (usually interpreted as signature for standard efficient accretion disks), arguing that LINERs have intrinsically different SEDs compared to ``normal'' AGN and that they are likely hosts of radiatively inefficient accretion flows (RIAFs). On the other hand, Satyapal et al. (2004) combining mid-IR spectroscopy with X-ray imaging conclude that LINERs are often heavily absorbed and that the intrinsic SEDs in LINERs are not necessarily different from those of standard AGN. This conclusion finds support from the recent findings of Maoz (2007): Using high-quality UV data for a sample of 13 unabsorbed LINERs observed with \\hst, Maoz argues that the SEDs of low-luminosity AGN (LLAGN) are similar to those of Seyfert galaxies.   Previous X-ray studies of low-power AGN with \\chandra\\ demonstrate the importance of X-rays to make progress in this field (e.g., Di Matteo et al. 2001,2003; Pellegrini et al. 2005; Evans et al. 2006; Balmaverde et al. 2006). However, the nature of the central engine in low-power AGN is still a matter of debate, with some authors favoring a jet-dominated scenario and others an accretion-dominated scenario, which in turn might be either radiatively efficient or inefficient. One of the reasons for this controversy is the poor discriminating power of the low signal-to-noise (S/N) spectra, which do not allow one to choose between the competing scenarios.   In order to break this spectral degeneracy, higher S/N spectra combined with additional model-independent constraints are required. This is the approach that we (Gliozzi et al. 2003) adopted to investigate the origin of the X-rays in the nuclear region of the nearby FRI NGC~4261, which is classified as  Weak-Line Radio Galaxy (WLRG) and LINER (Lewis et al. 2003). By definition, WLRGs are radio galaxies with very weak [\\oiii] emission: EW([\\oiii])$<$10\\AA) (Tadhunter et al. 1998). However, no formal constraints are placed on the optical line ratios, therefore, an object classified as WLRG is not necessarily a LINER and vice versa. Nevertheless, a systematic study of the optical properties of a sample WLRGs has shown that the majority of the objects can be robustly classified as LINERs (Lewis et al. 2003). For the sake of simplicity, in the following we will use the term WLRGs to indicate radio galaxies  with weak optical lines, which are optically classified as LINERs.  Exploiting the superior sensitivity of \\xmm, in NGC~4261 we found evidence for significant temporal and spectral variability on short timescales, and the presence of an unresolved Fe K line at $\\sim6.9$ keV of nuclear origin. We thus concluded that the inner jet is unlikely to dominate the X-ray emission. Similar conclusions were reached by Rinn et al. (2005), using \\chandra, \\xmm, and \\sax\\ archival data of 9 WLRGs. Another relevant result from Rinn et al. (2005) is the detection of intrinsic absorption in 6 out 9 objects, with 4 sources having $N_{\\rm H}> 10^{22}$ \\nh. This is in contrast with the systematic lack of absorption found in a larger sample of non--LINER FRIs observed with \\chandra\\ or \\xmm\\ (Donato et al. 2004), and suggests that WLRGs may form a distinct class compared to other low--power radio galaxies.  To summarize, there is still no clear consensus on what is the dominant source of X-rays in low-power AGN/LINERs and what is the nature of accretion flow onto their black holes. Are the different conclusions just a consequence of the lower photon statistics, or are they reflecting intrinsic differences among objects of the same class? And if so, are these differences related to the nuclear power and/or to the beaming of the sources? To answer these questions and make progress in the study of accretion in AGN, it is important to expand the sample of low-power radio galaxies with high-quality X-ray data.  Here, we investigate the nature of  4 WLRGs (optically classified as LINERs) \\ngc, \\pks, 3C~88, and 3C~444, whose properties are described in $\\S$2 and summarized in Table 1. Specifically,  using observations from \\xmm\\ (described in $\\S$3), we try to shed light on the origin of X-rays and the nature of the accretion mechanism. To this end, we combine several pieces of information obtained from the time-averaged X-ray spectroscopy (discussed in $\\S$4), from the X-ray variability (described in $\\S$5), and from broad-band analysis, which combines the simultaneous X-ray/UV coverage with archival radio data (in $\\S$6). Hereafter, we adopt $H_0=71{\\rm~km~s^{-1}~Mpc^{-1}}$, $\\Omega_\\Lambda=0.73$ and $\\Omega_{\\rm M}=0.27$ (Bennet et al. 2003). ",
        "conclusions": "We have investigated the nature of 4 WLRGs, \\ngc, \\pks, 3C~88, \\3c4, using observations from \\xmm, and combining information from the time-averaged spectral analysis, from the temporal and spectral variability, and from broadband studies. The main results can be summarized as follows:  \\begin{itemize}  \\item The 4 AGN have Eddington ratios $L_{\\rm bol}/L_{\\rm Edd}$ that span 2 orders of magnitude ranging between $\\sim 1\\times 10^{-5}$ (\\ngc) and $2-3\\times 10^{-3}$ (\\pks\\ and \\3c4), with 3C~88 having an intermediate value of $\\sim 2\\times 10^{-4}$.  \\item The 4 targets are adequately fitted by the same continuum model that comprises at least one thermal component ($kT\\sim 0.65-1.45$ keV) and a partially absorbed power law. However, the relative contributions of the individual spectral components and the photon indices change significantly from source to source. For example, \\ngc\\ and \\pks\\ have fairly steep power-law components ($\\Gamma\\sim 2.5-2.9$), whereas 3C~88 has a flat  photon index ($\\Gamma\\sim 1.1$), and \\3c4\\ is the only one showing a typical Seyfert-like $\\Gamma \\sim1.9$ accompanied by an apparent line-like excess around 6.7 keV. If this excess is fitted with a Gaussian model, the resulting line is centered around 6.7 keV, fairly strong ($EW\\sim550$ eV) and possibly broad ($\\sigma\\sim0.35$ keV). However, observations with a better coverage of the hard X-ray band and with longer duration are necessary to firmly establish the presence of a broad line and constrain the line parameters  \\item There is evidence only for moderate absorption ($N_{\\rm H}\\sim10^{21}$\\nh) in \\ngc, \\pks, and \\3c4. On the other hand, the flat photon index derived for 3C~88 and the $N_{\\rm H}-\\Gamma$ contour plots suggest that this source might have a larger column density of the order of $N_{\\rm H}\\sim10^{22}$\\nh.   \\item The results of the variability analysis indicate that \\pks\\ is significantly variable in the 0.2--10 keV and 0.2--1 keV energy bands, whereas 3C~88 is variable in the soft band only, and \\ngc\\ is marginally variable in the soft band. In general, there is a suggestive evidence that the sources are more variable in the soft band, although the limited statistics due to the short observations does not allow to draw firm conclusions.  \\item For \\pks, the split of the exposure into two observations taken 4 months a part, makes it possible to investigate the spectral variability using the hardness ratio. The spectrum becomes softer as the source brightens, following the typical Seyfert-like spectral behavior.  \\item The broadband SED of \\pks, comprising (non-simultaneous) radio data combined with (simultaneous) UV and X-ray measurements, can be adequately fitted with a single synchrotron component with little contribution from the putative accretion disk. However, the paucity of data-points does not allow one to rule out alternative scenarios.  \\item  For \\3c4, \\pks, and \\ngc,  the values of $\\alpha_{\\rm OX}$ plotted versus the UV monochromatic luminosity follow the best-fit linear regression found by Steffen et al. (2006) for Seyfert galaxies. This suggests a similarity (or at least a smooth transition) with Seyfert-like objects. No firm conclusions can be drawn for the remaining  object 3C~88, due to the large uncertainty of its UV luminosity.  \\item The lack of a positive correlation between  $L_{\\rm 2-10~keV}$ and the viewing angle, parameterized by the radio core dominance, suggests that the X-ray emission is not beamed and hence more likely to be related to the accretion flow.  \\item The study of the X-ray and UV radio loudness indicates that \\3c4\\ lies well within the radio-quiet region/branch, \\pks\\ is in between the 2 branches, whereas \\ngc\\ and 3C~88 are more consistent with the radio-loud branch.  \\end{itemize}  Before concluding, we can try to combine the different findings from our analysis in an attempt of answering the fundamental question on the origin of the X-ray emission in WLRGs.  Several pieces of evidence seem to favor an accretion-related origin. First, relatively high values of $L_{\\rm X}/L_{\\rm Edd}$, as those found for 3 of our targets, are commonly interpreted as Comptonized emission from a RIAF, based on studies of LLAGN and GBHs in the low/hard state (Yuan \\& Cui 2005; Wu et al. 2007). Second, the lack of a positive correlation in the  $L_{\\rm X}-R$ plane (see  Figure 9), suggests that the X-ray emission associated with the power law component is not beamed and hence unlikely to be produced by a relativistic jet. Third, the results from the temporal and spectral variability analyses (more pronounced variability in the soft band and the spectral softening associated with the source brightening observed in \\pks) are more consistent with the trends observed in Seyfert galaxies and at odds with the typical blazar-like behavior. Finally, the time-averaged spectral results of \\3c4\\ ($\\Gamma\\sim 1.9$ and possibly broad \\feka), coupled with its radio-quiet nature and the relatively high luminosity and Eddington ratio, suggest that not only this source is unlikely jet-dominated but also that the accretion mode might be similar to the Seyfert-like one (i.e., radiatively efficient).  On the other hand, a significant (or even dominant) contribution from the jet in the X-ray regime cannot be a priori excluded at least for some objects. For example, the main properties of \\ngc\\ --$L_{\\rm X}/L_{\\rm Edd}\\sim8\\times 10^{-7}\\,,\\Gamma\\sim2.9$ and strong radio-loudness-- seem to naturally fit the jet-dominated scenario. Similarly, the SED and the steep $\\Gamma$ of \\pks, coupled with the relatively small viewing angle, might be adequately interpreted in the framework of jet-dominated models.   In summary, based on the analysis of time-averaged spectra combined with model-independent information from X-ray temporal and spectral variability, and with inter-band information, we cannot derive a general conclusion valid for the entire sample. In fact, the main conclusion of this work is that LINERs represent a very heterogeneous class. Not only they comprise non-AGN and AGN objects, not only the AGN class can be divided into radio-quiet and radio-loud LINERs, but also the latter subclass seems to encompass quite different objects. These may  range from intrinsically low-luminosity objects (in terms of Eddington ratio), to objects potentially strongly absorbed, to brighter objects with properties more in line with Seyfert galaxies. Our results are based on a very small sample, which needs to be extended to derive more general conclusions. Nevertheless, it demonstrates the need of high-quality X-ray observations to disentangle the relative contribution of the different components in order to shed some light on the nature of low-power radio galaxies."
    },
    "0711/0711.2363_arXiv.txt": {
        "abstract": "Neutral fluorine (F\\,{\\sc i}) lines are identified in the optical spectra of several R\\,Coronae Borealis stars (RCBs) at maximum light. These lines provide the first measurement of the fluorine abundance in these stars. Fluorine is enriched in some RCBs by factors of 800 to 8000 relative to its likely initial abundance. The overabundances of fluorine are evidence for the synthesis of fluorine. These results are discussed in the light of the scenario that RCBs are formed by accretion of an He white dwarf by a C-O white dwarf. Sakurai's object (V4334\\,Sgr), a final He-shell flash product, shows no detectable F\\,{\\sc i} lines. ",
        "introduction": "R Coronae Borealis (RCB) stars comprise a sequence of hydrogen-deficient supergiants with effective temperatures from about $T_{\\rm eff} = 3500$ K, as represented by Z\\,UMi and DY\\,Per, to about 19,500 K, as represented by DY\\,Cen. The characteristic of H-deficiency is shared by the H-deficient cool carbon (HdC) stars at low temperatures and by the extreme helium (EHe) stars at high temperatures. A common assumption is that the sequence HdC - RCB - EHe in the ($T_{\\rm eff},\\log g$) plane reflects a close evolutionary connection. In this sequence, the RCBs are distinguished from HdC and EHe stars by their second principal defining characteristic: their propensity to fade, in visual light, unpredictably as a cloud of carbon dust obscures the star. This propensity is not universally shared: XX\\,Cam has yet to be observed below maximum light. In addition, DY\\,Cen might be considered as an EHe star known to experience RCB-like fadings.  If HdC, RCB, and EHe stars share a common heritage, the expectation is that their atmospheric compositions should show some common features \\citep{pan2004a,rao2005}. It is through the compositions that one hopes to test theoretical ideas about the origins of these extremely rare stars; just five HdC, about 40 RCB \\citep{zanie2005}, and 21 EHe stars are known in the Galaxy. Currently, two scenarios remain in contention to account for these H-poor high luminosity stars. In the first, a final He-shell flash in a post-AGB star on the white dwarf cooling track creates a H-poor luminous star. This is dubbed the `final flash' (FF) scenario. In the second, the H-poor star is formed from a merger of a He white dwarf with a C-O white dwarf. In a close binary system, accretion of the He white dwarf by the C-O white dwarf may lead to a H-poor supergiant with the C-O white dwarf as its core. This is called the `double degenerate' (DD) scenario.  Products of both the FF and DD scenarios may be presumed to exist. A determination of which scenario provided which star rests in large part on the observed chemical composition of a star's atmosphere and theoretical predictions about the FF and DD products. Evidence from elemental abundances, especially the H, C, N, and O abundances, suggests that the RCB and EHe stars evolved via the DD rather than the FF route \\citep{asp00,pan2001,saio02,pan2006}. Convincing, essentially incontrovertible, evidence that the DD scenario led to the HdCs and some cool RCBs was presented by \\citet{clay2007} with their discovery that the $^{18}$O was very abundant in their atmospheres. This usually rare isotope of oxygen was attributed to nucleosynthesis occurring during and following accretion of the He-rich material onto the C-O white dwarf.  Determination of the oxygen isotopic ratios demands a cool star with the CO vibration-rotation bands in its spectrum. The majority of RCBs and all of the EHes are too hot for CO to contribute to their spectra \\citep{tenen2005}.\\footnote{A possibility exists that circumstellar CO may be detectable for the RCBs. Cool gas containing C$_2$ molecules has been seen for V854\\,Cen, R\\,CrB, and V\\,CrA when the stars were in decline \\citep{raolamb2000,raolambsh2006,raolamb2007}.} An alternative tracer of nucleosynthesis during a merger may be provided by the fluorine abundance. Considerable enrichment of EHe stars with F was discovered by \\citet{pana2006} from detection and analysis of about a dozen F\\,{\\sc i} lines in optical spectra of cool EHe stars. Clayton et al.'s (2007) calculations suggest that F synthesis is possible in the DD scenario. In this paper, we report on a search for F\\,{\\sc i} lines in spectra of RCBs and discuss the F abundances in light of the results for EHes and the expectations for the DD and FF scenarios. ",
        "conclusions": "Lithium and fluorine are not the sole or even the principal concerns for the DD scenario. One wonders why the C/He $\\sim 0.01$ ratio, as measured for the EHes, spans a narrow range.  The He is contributed  primarily by the He white dwarf and the C by the He-intershell of the C-O white dwarf. What sets the mass ratio of these two regions rather precisely?  The evolutionary connections between the H-deficient stars remains unproven but with obvious potential links. Both possible sequences -- HdC $\\rightarrow$ RCB $\\rightarrow$ EHe and EHe $\\rightarrow$ RCB $\\rightarrow$ HdC are possible because the DD scenario leads to a very hot star evolving to lower temperatures before the evolutionary track reverses direction. An extension EHe $\\rightarrow$ sdO $\\rightarrow$ O(He) to higher temperatures and higher surface gravities  seems plausible \\citep{rauch2006,werner2006}. The stars in this sequence are all candidates for an origin via the DD scenario, i.e., they have the low C/He ratio and not the high ratio characteristic of the the FF scenario; a similar sequence for the FF scenario exists  linking objects like FG\\,Sge and V4334\\,Sgr (Sakurai's object) through the [WC] to the PG1159 stars.   Continued study of the compositions of stars along these sequences will serve to sharpen understanding of the evolutionary links and the entry points  for stars into the sequence. New data such as that for F and $^{18}$O stimulate continued study. The F abundances of  RCB  and EHe stars are quite similar in value and spread and thus support the presumed evolution from RCB to EHe. On the other hand, there is a hint that the $^{16}$O/$^{18}$O ratios of the HdC and RCB stars may be different with the ratio being systematically higher in the (cool) RCBs than in the three HdCs. Is the difference merely a distraction caused by small number statistics or is it symptomatic of a change introduced as a HdC evolves to a cool RCB or does it reflect that the DD scenario can place stars on the evolutionary sequence at different starting points with different compositions?    We thank the referee Geoff Clayton for the nice report. GP thanks the VBT staff for their assistance. This research has been supported in part by the Robert A. Welch Foundation of Houston, Texas.   \\clearpage"
    },
    "0711/0711.0827_arXiv.txt": {
        "abstract": "The X-ray source \\src\\ was discovered close to the supernova remnant CTB\\,1 in a \\rst\\ observation performed in 1992. The source phenomenology (soft spectrum, apparent lack of counterparts, possible pulsations at 242 ms, hints for surrounding diffuse emission) led to interpret it as an isolated neutron star in a new supernova remnant. We have analysed an archival \\xmm\\ observation performed in 2001. The source coordinates, as computed on the \\xmm\\ images, coincide with those of a bright source listed in optical and infrared catalogues. The X-ray spectrum is well described by an optically thin plasma model. No fast pulsations are seen, nor clear evidence of a supernova remnant associated to the source. Thus, we conclude that \\src\\ is not an isolated neutron star, but the X-ray counterpart of the bright optical/infrared source, most likely a F7 spectral class star located at about $0.2$ kpc. ",
        "introduction": "Most of the observed isolated neutron stars are identified as pulsars, whose emission derives either from rotational energy loss, or as magnetars, powered by magnetic field decay \\citep[e.g.,][]{manchester04}. The central compact objects \\citep[CCOs; see][]{pavlov02} remain perhaps the least understood members of the isolated neutron stars family. These X-ray sources are located within supernova remnants (SNRs) and, despite intensive campaigns, have not been detected as radio or optical sources so far. CCOs are seemingly young (\\mbox{$\\la$10$^4$ years}) isolated neutron stars, with steady X-ray fluxes (with the notably exception of 1E\\,161348$-$5055 at the centre of the SNR RCW\\,103; \\citealt{deluca06}), soft thermal spectra, and lack of surrounding pulsar wind nebulae. Two CCOs out of seven are pulsating sources: 1E\\,1207.4$-$5209 in G\\,296.5+10.0, with period \\mbox{$P=424$ ms} \\citep{zavlin00}, and PSR\\,J1852+0040 in Kes\\,79, with \\mbox{$P=105$ ms} \\citep{ghs05}. Both sources have small spin-down rates with period derivatives $\\dot{P}<2\\times10^{-16}$ s s$^{-1}$ \\citep{gotthelf07,halpern07}.\\\\ \\indent The current sample of CCOs includes seven ``confirmed'' sources and four ``candidates''. The confirmed CCOs are the central sources in RCW~103 \\citep{tuohy80}, G\\,296.5+10.0 \\citep{helfand84}, Pup~A \\citep{petre96}, Vela~Jr. \\citep{aschenbach98}, G\\,347.3$-$0.5 \\citep{slane99}, Cas~A \\citep{tananbaum99}, and Kes~79 \\citep{seward03}. The candidates are those in G\\,349.7+0.2 \\citep{lazendic05}, G\\,15.9+0.2 \\citep{reynolds06cco}, G\\,330.2+1.0 \\citep{park06}, and \\src\\ in G\\,117.9+0.6 \\citep{hailey95}, which is the object of our research.\\\\ \\indent The X-ray point source \\src\\ was discovered with the PSPC instrument on board \\rst\\ near the supernova remnant CTB\\,1, during a \\mbox{$\\sim$9 ks} long observation carried out on 1992 August 16--17. \\citet{hailey95} reported the position \\mbox{$\\rm{R.A.}=00^{\\rm{h}}02^{\\rm{m}}54\\fs1$}, $\\rm{Decl.}=62^{\\circ}46'23''$ (epoch J2000). The observation showed a hint of a faint shell of soft X-ray emission (G\\,117.7+0.6), proposed as a SNR associated with \\src\\ \\citep{hailey95,craig97}. The spectrum of the point source was fitted using a blackbody attenuated by interstellar absorption, with \\mbox{$k_BT\\simeq0.15$ keV}. Assuming a distance of 3 kpc \\citep{hailey95}, this corresponds to a 0.5--2 keV luminosity of \\mbox{$\\sim$$2\\times10^{32}d_3^2$ \\lum} (where we indicate with $d_{\\rm{N}}$ the distance in units of N kpc). \\citet{hailey95} also found some evidence for a possible periodicity in the X-ray emission of \\src\\ with period \\mbox{$P=242$ ms}. Based on these results,  \\citet{hailey95} proposed that \\src\\ is an isolated neutron star in a SNR. Furthermore, the absence of counterparts at other wavelengths \\citep[][]{hailey95,brazier99} suggested that \\src\\ could be a CCO.\\\\ \\indent In 2001 \\src\\ has been observed for 33 ks with \\xmm. \\citet{pavlov04} reported results of the analysis of those data. The spectrum was fitted by a two-component model: a soft blackbody with temperature \\mbox{$k_BT\\simeq0.1$ keV} and a hard component, either a second blackbody with $k_BT\\simeq0.5$ keV or a power-law with photon index $\\Gamma\\simeq2.6$. Any periodicity was excluded, as well as the presence of a SNR around the source. They concluded that \\src\\ is most likely a middle-aged pulsar rather than a CCO.\\\\ \\indent Here we report on a re-analysis of the \\xmm\\ observation of \\src. This paper presents evidence that, contrary to previous claims, the X-ray source \\src\\ is neither a CCO nor a pulsar, but rather a non-degenerate star. We have also identified its likely stellar counterpart using near-infrared data. ",
        "conclusions": "The X-ray source \\src\\ is clearly detected in the \\xmm\\ images and its position is consistent with a rather bright star (\\tmsrc). \\citet{hailey95} ruled out this star as a possible counterpart of \\src\\ mainly for the angular separation of $12\\arcsec$ from their X-ray position. However, they relied upon a positional uncertainty of $10\\arcsec$, a value that in subsequent releases of the WGACAT was conservatively increased to $13\\arcsec$ (1\\,$\\sigma$). Moreover, the source coordinates in the WGACAT are affected by a systematic error\\footnote{See Haberl F., Pietsch W., and Voges W., ``Differences in the two ROSAT catalogs of pointed PSPC observations'' (1994), and comments by White N. E., Angelini L., and Giommi P.. The document is available at ftp://ftp.xray.mpe.mpg.de/rosat/catalogues/sourcecat/wga\\_rosatsrc.html\\,.}. The recent (2001) Second ROSAT Source Catalog of Pointed Observations with the Position Sensitive Proportional Counter\\footnote{See http://www.mpe.mpg.de/xray/wave/rosat/rra/.} (ROSPSPCCAT/2RXP) using the same observation of \\citet{hailey95} provides more reliable coordinates: $\\rm{R.A.} = 00^{\\rm{h}}02^{\\rm{m}}55\\fs4$, $\\rm{Decl.}=62^{\\circ}46'21\\farcs0$ (epoch J2000). Adopting this position, the offset between the \\xmm\\ and \\rst\\ positions decreases to $4\\farcs0$, and that from \\tmsrc\\ to $4\\farcs0$.\\\\ \\indent With the J, H, and K$_{\\rm{s}}$ magnitudes of \\src/ \\tmsrc\\ at hand (see \\mbox{Section~\\ref{ircounterpart}}), we used the relation between the \\nh\\ of the X-ray best-fits (Table~\\ref{fits}) and the interstellar extinction $A_{\\rm{V}}$ of \\citet[][]{predehl95}, as well as the relations between the extinctions at different wavelengths of \\citet{cardelli89} to derive the intrinsic colours of the source \\mbox{(J--H)$_{\\circ}\\simeq0.30$} and (H--K)$_{\\circ}\\simeq0.06$. In a similar way, taking optical photometric data from the Tycho-2 Catalogue\\footnote{See http://www.astro.ku.dk/$\\sim$erik/Tycho-2/.} \\citep{hog00}, we derived also the colour \\mbox{(B--V)$_{\\circ}\\simeq0.50$}. These values are consistent with a F or G type star in the case of a main-sequence star, or with a G type in the case of a supergiant \\citep[e.g.,][]{cox00}. In particular, the intrinsic colours point to a F7-type main-sequence star.\\\\ \\indent In the reasonable frame of a F7-type main-sequence star, the expected absolute optical magnitudes is $M_{\\rm{V}}\\simeq3.4$ \\citep[e.g.,][]{cox00,zombeck07}, implying a distance of \\mbox{$\\sim$230 pc}. Such a relatively small distance is well consistent with the measure of the photoabsorption derived from the best-fitting models of the X-ray spectrum (\\mbox{$N_{\\rm H} \\simeq 1.5 \\times 10^{21}$ cm$^{-2}$}, see Table~\\ref{fits}). This value is in fact significantly smaller than the measurements of the interstellar hydrogen in this direction by \\citet{dickey90} and \\citet{kalberla05}, that give \\nh\\ values of \\mbox{$\\sim$(6--$7)\\times 10^{21}$ cm$^{-2}$}. The X-ray-to-optical flux ratio is \\mbox{$\\log(f_{\\rm{X}}/f_{\\rm{V}})\\simeq -3.3$}, in good agreement with the value of \\mbox{$\\langle\\log(f_{\\rm{X}}/f_{\\rm{V}})\\rangle=-3.7\\pm0.7$} obtained by \\citet{krautter99} averaging \\rst/PSPC and optical data on a sample of 53 F-type stars. This scenario is further confirmed by the X-ray spectrum of \\src\\ measured by \\xmm\\ that is well fit by either the \\textsc{APEC} or \\textsc{MEKAL} codes, with temperatures typical of non-degenerate stellar atmospheres.\\\\ \\indent Based on the accurately identified counterpart (thanks to \\xmm\\ imaging capabilities) as well as on the spectral properties of \\src, together with the lack of an associated SNR and the absence of X-ray pulsations, we conclude that the source is not a neutron star (in any of its manifestations, including a CCO) and its properties are clearly consistent with a non-degenerate star."
    },
    "0711/0711.3564_arXiv.txt": {
        "abstract": "% We report the results of a phase-referencing study aimed at uncovering precession of the VLBI jet of BL Lac. The observations were conducted at 8, 15, 22, and 43 GHz and consist of seven epochs spanning about two years. We investigated the change in the absolute position of BL Lac's radio core by means of phase-referencing with two nearby sources, 2151+431 and 2207+374. The shift in the position of the core perpendicular to the jet is a signature of precession. However, the periodic variations with an amplitude of $\\sim$0.15 mas and a period of 1 year can be attributed to seasonal weather variations. We also detect a trend in position of the core on the scale of $\\sim$0.1 mas over two years. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.0751_arXiv.txt": {
        "abstract": "The temperature of the intergalactic medium (IGM) is set  by the competition between photoheating and adiabatic cooling, which are usually assumed to define a tight equation of state in which the temperature increases monotonically with density. We use a semi-analytic model, accurate at low to moderate IGM densities ($\\la 5$ times the mean), to show that inhomogeneous reionization can substantially modify these expectations.  Because reionization is driven by biased sources, dense pockets of the IGM are likely to be ionized first.  As a result, voids initially remain cool while dense regions are heated substantially.  However, near the end of reionization, dense regions have already cooled from their initially large temperature while voids have only just been heated.  Thus, near the end of helium reionization the equation of state can invert itself, with the hottest gas inside voids.  The degree to which this happens depends on the magnitude of the density bias during reionization:  if rare, bright sources dominate reionization, so that each ionized region contains a typical volume of the IGM, the equation of state will remain monotonic.  We also show that the distribution of temperatures at a fixed density  has significant scatter and evolves rapidly throughout and even after reionization.  Finally, we show that the observed temperature jump at $z \\sim 3.2$ is consistent with the behavior at the end of helium reionization, although it requires a somewhat larger temperature increase than expected. ",
        "introduction": "\\label{intro}  The two most dramatic events in the history of the intergalactic medium (IGM) are the reionization of hydrogen (by the first generations of galaxies) and helium (by quasars).    These have become key landmarks for both observational and theoretical cosmologists in the past several years.  Evidence for hydrogen reionization comes from a number of directions, none of them clear but all consistent with (possibly extended) reionization at $z \\sim 6$--$10$ (see \\citealt{fan06-review, furl06-review} for recent reviews).  The clearest evidence for helium reionization comes from \\ion{He}{2} \\lya forest spectra in the extreme ultraviolet, which seem to show a rapid evolution in the characteristics of the transmission at $z \\sim 3$ \\citep{davidsen96,anderson99, heap00, smette02, zheng04, shull04, reimers04, zheng04-sdss, reimers05}, although there is also some indirect evidence (see below).  These two reionization epochs are largely responsible for determining the thermal history of the IGM.  Before hydrogen reionization, the neutral IGM cooled adiabatically until the first structures formed, probably reaching temperatures $T \\sim 30 \\kel$.  X-rays from the first galaxies most likely slowly heated the neutral IGM, to $T \\la 1000 \\kel$ \\citep{oh01, venkatesan01, kuhlen06-21cm, furl06-glob}.  However, hydrogen reionization was a much more dramatic change:  the $\\sim 1 \\eV$ leftover from each ionizing photon heated the IGM to $\\sim 1$--$2 \\times 10^4 \\kel$ \\citep{miralda94, abel99, tittley06}.  The harder photons responsible for helium reionization would have reheated the IGM to similar, or even larger, temperatures \\citep{hui97}.  Once reionization is complete, this heating channel slows dramatically -- because only the relatively small fraction of ions that recombine will couple to the photoionizing background.  The subsequent temperature evolution is determined by a combination of adiabatic heating and/or cooling, photo-heating, and Compton cooling \\citep{miralda94, hui97}.\\footnote{More precisely, several other factors contribute (see the Appendix of \\citealt{hui97} for a summary) but these three are by far the most important.}  The competition between these processes forces the gas temperature to approach an asymptotic form set by the background ionizing spectrum \\citep{hui97, hui03}.  Because the magnitude (and indeed sign) of the adiabatic term depends on whether the gas is over- or underdense, the IGM assumes an equation of state $T \\approx T_0 (1 + \\delta)^{\\gamma-1}$, where $\\delta$ is the fractional overdensity of the gas element and $\\gamma$ is nearly independent of $\\delta$ in most models.  This equation of state inevitably affects many observables of the \\lya forest, including the power spectrum of the transmitted flux \\citep{croft98, zald01, viel04, mcdonald06}, the optical depth distribution \\citep{lidz06-lya}, and the line widths of individual features \\citep{schaye99, schaye00, ricotti00, mcdonald01}.  In particular, the observed temperature evolution has been used to constrain the epochs of helium and hydrogen reionization:  most dramatically, the velocity widths of \\lya forest absorbers seem to increase sharply at $z \\sim 3.2$ \\citep{schaye00}, while the equation of state simultaneously flattens \\citep{schaye00, ricotti00}.  This may be a result of helium reionization.  If so, the temperature at higher redshifts could provide a clean ``fossil\" record of hydrogen reionization:  it appears to be relatively large and so requires that epoch to have occured at $z \\la 10$ \\citep{theuns02-reion, hui03}.  The cartoon history described above has been developed from simple uniform reionization models.  These existing treatments have ignored one extremely important facet of reionization:  it is highly inhomogeneous, especially because the sources driving it are themselves highly clustered \\citep{barkana04, furl04-bub, furl07-helium}.  As a result, overdense regions in the IGM are ionized before the empty voids (at least on average), and their thermal histories will evolve differently.  This affects the mean equation of state and also adds an element of ``stochasticity\" to it, because patches at identical densities can have different reionization histories and hence different temperatures.  To this point, analytic models have ignored this aspect for simplicity, and numerical simulations have not been able to include both the inhomogeneity of reionization (which occurs on large $\\ga 10 \\Mpc$ scales) while also resolving features of the \\lya forest.  The exception is the Monte Carlo, semi-analytic helium reionization model of \\citet{gleser05}, who showed that this inhomogeneity significantly affects the equation of state.  However, they used a highly approximate prescription for inhomogeneous reionization.  In this paper, we present a model that incorporates the latest models of inhomogeneous reionization for calculating the IGM equation of state. We apply this model to the helium reionization epoch and show that it has an enormous impact on the thermal properties of the IGM during the best-observed epochs, $z \\sim 2$--$4$.  We describe our model for inhomogeneous reionization in \\S \\ref{inhom} and our model for the subsequent temperature evolution in \\S \\ref{temphist}.  We examine the resulting temperature distributions in \\S \\ref{tempdist}, the equation of state in \\S \\ref{eos}, and the evolution of the IGM temperature in \\S \\ref{temp-evol}.  We compare our models to observations in \\S \\ref{comp-obs} and conclude in \\S \\ref{disc}.  In our numerical calculations, we assume a cosmology with $\\Omega_m=0.26$, $\\Omega_\\Lambda=0.74$, $\\Omega_b=0.044$, $H=100 h \\hunits$ (with $h=0.74$), $n=0.95$, and $\\sigma_8=0.8$, consistent with the most recent measurements \\citep{spergel06}.  Unless otherwise specified, we use comoving units for all distances. ",
        "conclusions": "\\label{disc}  We have examined how inhomogeneous reionization affects the equation of state of the IGM.  Because reionization is accompanied by substantial photoheating, it affects the thermal state long after its completion.   We have considered models in which the reionization process proceeds from large-scale overdensities to large-scale voids and in which reionization occurs independently of the underlying density.  In the former case, dense pockets of gas are likely to lie near galaxies (and quasars) and so to be ionized relatively early; voids, on the other hand, are likely to be the last elements to be ionized. Our semi-analytic approach ignores shock-heating in the IGM, as well as galactic winds, so it is only accurate at moderate or low overdensities ($\\delta \\la 4$), where such phenomena are relatively rare, but it does describe the bulk of the gas contributing to the Ly$\\alpha$ forest (which sits in photoionization equilibrium).  Hydrogen reionization occurred early enough that the implications for the thermal state in the observable era ($z \\la 4$) are relatively modest, and we defer a closer look at fluctuations from that process to the future.  This is primarily because the IGM rapidly approaches a (fairly steep) thermal asymptote, allowing only weak constraints on $z_{\\rm H}$ \\citep{theuns02-reion, hui03}.  However, \\ion{He}{2} reionization at $z \\sim 3$--$4$ has a large effect.  If reionization occurs in a density-independent fashion, the temperature always increases with density, but it is still nearly isothermal during and near the end of reionization.   If underdense gas is ionized last, the equation of state can actually become inverted, with the hottest gas elements in the deepest voids.  In practice, the ``median equation of state,\" defined with reference to the median temperature at each density,  is reasonably close to isothermal near the end of helium reionization, but with a factor of a few spread in the actual temperatures.  It also develops a break that moves to lower density as the process unfolds (although the break is a result of taking the median; the ``mean\" equation of state will be smoother).  Gas elements denser than the break are nearly isothermal (because they have already undergone helium reionization), while those that are less dense have temperatures near the thermal asymptote ($T \\la 10^4 \\kel$).  As expected, the temperature increases rapidly near the end of helium reionization, with the magnitude of the jump and its precise timing depending on how it is measured.  The magnitude of the jump depends on $T_i^{\\rm He}$, but it is significantly \\emph{smaller} than the jump experienced by individual gas elements, because some gas elements must have been ionized at significantly higher redshift and so have cooled toward the new thermal asymptote.  For many types of observations (for example, those sensitive to contours in the temperature distribution), the observed temperature jump will also occur somewhat before reionization is complete.  Measurements of \\lya forest lines do show such a temperature jump at $z \\approx 3.2$, accompanied by a trend toward a more isothermal equation of state \\citep{ricotti00, schaye00}.  This is qualitatively consistent with $z_{\\rm He}=3.2$ in our model, but the magnitude of the jump is somewhat higher than one might expect if the \\citet{schaye00} method measures the minimum IGM temperature. If so, we would require an extremely hard ionizing background that can produce $T_i^{\\rm He} \\sim 4 \\times 10^4 \\kel$.  It is possible, however, that a more realistic assessment of the method with numerical simulations that include inhomogeneous reionization would show that it tends to measure temperatures closer to the median values, in which case a large but not unreasonable initial temperature would be required ($T_i^{\\rm He} \\sim 3 \\times 10^4 \\kel$).  On the other hand, the temperature can also be measured from the \\lya forest power spectrum, which shows no evidence for an evolving equation of state and tends to find higher temperatures at $z \\sim 4$ \\citep{zald01, viel04, mcdonald05, mcdonald06}.  The discrepancy probably arises because the power spectrum method is sensitive to the mean temperature, which evolves more slowly with redshift and has a significantly smaller overall jump than statistics like the median.  Thus inferences about the temperature evolution and equation of state will depend on precisely how a measurement is made (if, for example, the minimum, median, or mean temperature at each density is measured, and which density range is selected).  Careful comparison to observations will probably require detailed simulations of the reionization process.  Of course, our model is far from perfect.  We have ignored fluctuations in the ionizing background (in both its amplitude and spectrum), which will create an even larger spread in the gas temperatures.  Such fluctuations may also depend on the local gas density and introduce biases into the equation of state.   For example, voids lie far from the ionizing sources, so photons that reach them may have been strongly filtered and hence hardened by the IGM; on the other hand, after reionization voids appear to show systematically softer spectra than dense filaments \\citep{shull04}.  We have also ignored X-ray heating before a gas element undergoes helium reionization.  In reality, X-rays can travel large distances through the IGM and deposit a fraction of their energy as heat.  This will lift gas off of the thermal asymptote by at least a small amount, especially because such photons will tend to have higher energies than average, and consequently decrease the magnitude of the temperature jump from full reionization.  It may help to reconcile the overall discrepancy of our post-hydrogen reionization temperatures (at $z \\ga 4$) with the observations.  For example, if a fraction $\\bar{x}_{\\rm HeIII}$ of the helium is fully ionized by photons with mean excess energies (above the photoionization edge at $E_{\\rm HeII}$) of $\\VEV{E}$, we would have \\bq \\Delta T \\sim 1500 \\kel \\left( {\\bar{x}_{\\rm HeIII} \\over 0.1} \\right) \\left( {\\VEV{E} \\over E_{\\rm HeII}} \\right). \\label{eq:simmer} \\eq Before helium reionization is complete, this ``simmering\" background is likely to be quite hard (because the soft photons would be absorbed near the sources), so the heat injection could be substantial.  There are, at least in principle, a variety of other ways to test our model.  One could search for density-dependent temperature fluctuations in the \\lya forest lines, for example through correlations between small-scale and large-scale structure in \\lya forest lines; the small-scale structure becomes a proxy for temperature through Jeans smoothing.  Such searches have so far been unsuccessful \\citep{zald02, theuns02-temp}, but larger samples at a range of redshifts may show a signature.  The substantial increase in the temperature variance during helium reionization will probably be the easiest feature to see (see also \\citealt{hui03}), and can provide a diagnostic of when helium reionization occurs and, if the density dependence can be measured, how it proceeds.  Finally, the equation of state also affects the power spectrum of the \\lya forest, both directly (because the optical depth $\\tau \\propto \\delta^{2 - 0.7(\\gamma-1)}$) and possibly indirectly through spatial fluctuations in the temperature (although \\citealt{lai06} have shown that the latter are likely to be small)."
    },
    "0711/0711.3034_arXiv.txt": {
        "abstract": "We consider the spatial offsets of short hard gamma--ray bursts (SHBs) from their host galaxies. We show that all SHBs with extended--duration soft emission components lie very close to their hosts. We suggest that NS-BH binary mergers offer a natural explanation for the properties of this extended--duration/low offset group. SHBs with large offsets have no observed extended emission components and are less likely to have an optically detected afterglow, properties consistent with NS-NS binary mergers occurring in low density environments. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.3805_arXiv.txt": {
        "abstract": "We present analyses of observations of $\\epsilon$\\,Eridani (K2~V) made with the Low Energy Transmission Grating Spectrometer on {\\it Chandra} and the {\\it Extreme Ultraviolet Explorer}, supplemented by observations made with the Space Telescope Imaging Spectrograph, the {\\it Far Ultraviolet Spectroscopic Explorer} and the Reflection Grating Spectrometer on {\\it XMM-Newton}. The observed emission lines are used to find relative element abundances, to place limits on the electron densities and pressures and to determine the mean apparent emission measure distribution. As in the previous paper by \\cite{simjordan03a}, the mean emitting area as a function of the electron temperature is derived by comparisons with a theoretical emission measure distribution found from energy balance arguments. The final model has a coronal temperature of $3.4 \\times 10^6$~K, an electron pressure of $1.3 \\times 10^{16}$~cm$^{-3}$K at $T_{\\rm e} = 2 \\times 10^5$~K and an area filling factor of 0.14 at $3.2 \\times 10^5$~K. We discuss a number of issues concerning the atomic data currently available. Our analyses are based mainly on the latest version of CHIANTI (v5.2). We conclude that the Ne/O relative abundance is 0.30, larger than that recommended from solar studies, and that there is no convincing evidence for enhanced coronal abundances of elements with low first ionization potentials. ",
        "introduction": "\\label{S1} $\\epsilon$~Eri (K2~V) is a nearby dwarf star that has been observed over a wide spectral range, from the infrared to X-ray wavelengths. Its fundamental parameters have been discussed by \\cite{drakesmith93}, where references to earlier work can be found, and more recently by \\cite{difolco04} and \\cite{ap04}. \\cite{difolco04} have derived the diameter of $\\epsilon$~ Eri from interferometric measurements and have also made models to check the self-consistency of their adopted parameters. A more recent measurement by \\cite{difolco07} gives a radius that is smaller by only 0.008~$R_{\\odot}$. The values of the parameters adopted here are given in Table~\\ref{tab1} and are from \\cite{difolco04}, although these are also consistent with those found by \\cite{drakesmith93}, to within the combined uncertainties. Compared with the Sun, $\\epsilon$~Eri has a shorter rotational period (11.68~d) \\citep{donahue96}, a larger spatially averaged magnetic field (165~G) \\citep{rueedi97} and larger stellar surface emission line fluxes. It is, therefore, ideal for studies of a stellar outer atmosphere under conditions of higher mean magnetic activity than the Sun.  Observations with {\\it IRAS} showed that $\\epsilon$~Eri has an infrared excess, indicating the presence of dust; more recently a planet and a debris disk have been detected \\citep{hatzes00,greaves05}. Further studies of the dust have been made by \\cite{difolco07}. Combining the above rotational period with $v \\sin i = 1.7$~km~s$^{-1}$ and $R_* = 0.743 R_{\\odot}$ \\citep{difolco04}, leads to $i = 32^\\circ$, similar to the value of 25$^\\circ$ suggested by \\cite{greaves05} for the inclination of the debris disc to the plane of the sky.  Here we make use of observations at ultraviolet to X-ray wavelengths, as summarized in Table~\\ref{tab2}. In previous work we used observations with the Space Telescope Imaging Spectrograph (STIS) on the {\\it Hubble Space Telescope} to identify forbidden lines of Fe\\,{\\sc xii} \\citep{jordan01a} and to determine the electron pressure ($P_{\\rm e}$) in the transition region \\citep{jordan01b}. Including observations with the {\\it Far Ultraviolet Spectroscopic Explorer} ({\\it FUSE}), we have modelled the chromosphere and transition region, up to an electron temperature ($T_{\\rm e}$) of $\\simeq 3 \\times 10^5$~K \\citep{simjordan05}.  Observations with the {\\it Extreme Ultraviolet Explorer} ({\\it EUVE}) have been analysed by a number of authors. \\cite{schmitt96} used lines of iron to derive the emission measure distribution (EMD) at above $T_{\\rm e} \\simeq 6 \\times 10^5$~K and also placed limits on the electron density ($n_{\\rm e}$) in the upper transition region from lines of Fe\\,{\\sc xiii} and Fe\\,{\\sc xiv}. \\cite{laming96} used the same spectra to derive emission measures from lines of a range of elements and concluded that if any trends in abundances with the first ionization potential (FIP) were present, they were not significantly larger than in the solar corona. \\cite{sf03} included further {\\it EUVE} spectra and also derived the EMD. They found much larger values of $n_{\\rm e}$ from lines of Fe\\,{\\sc xix} and Fe\\,{\\sc xxi} that are now superseded. \\cite{simjordan03a} used the line fluxes measured by \\cite{schmitt96} to derive the EMD from the iron lines alone, adopting a more recent (smaller) value of the interstellar absorption by \\cite{dring97} and using CHIANTI (v4) \\citep{dere97,young03} for line excitation atomic data. They compared the apparent EMD with that calculated from an energy balance in which the net thermal conductive flux was set equal to the local radiative energy loss and hence derived the fractional emitting area as a function of $T_{\\rm e}$ ($\\simeq 0.2$ at $2 \\times 10^5$~K).  High-resolution X-ray observations of $\\epsilon$~Eri have been made with both the Low Energy Transmission Grating Spectrograph (LETGS) on {\\it Chandra} and the Reflection Grating Spectrometer (RGS) on {\\it XMM-Newton}. The former observations have been reduced and analysed by \\cite{sf04} and \\cite{wood06}, who derived the EMD and relative element abundances, making use of the Astrophysical Plasma Emission Database (APED v1.3) \\citep{smith01} and CHIANTI (v4.2), respectively. Because there have been changes in the atomic data available, we have made our own analysis of the above spectra and the {\\it EUVE} spectra, using CHIANTI (v5.2) \\citep{landi06}. We have also made our own measurements of the X-ray line fluxes. $\\epsilon$~Eri was included in a survey of stellar coronal densities by \\cite{ness02}, using LETGS spectra, and in a study of coronal opacities by \\cite{ness03a}, using LETGS and RGS spectra. Comparisons with these earlier results are made in later sections.  The first aim of the present work is to use all the available observations and up-to-date atomic data to derive the mean EMD, the relative element abundances and the electron densities and pressures. The mean EMD is then used in conjunction with new theoretical EMDs to refine the earlier values of the mean emitting areas.  Section~\\ref{S2} describes the X-ray observations used and the data reduction. Limits on $n_{\\rm e}$ and $P_{\\rm e}$ are discussed in Section~\\ref{S3}, since these are required in deciding on the identification of some lines and in the analysis of some line fluxes. The identifications of particular lines are discussed in Section~\\ref{S4}, in conjunction with the emission measures derived and some issues related to the atomic data. Section~\\ref{S5} describes how the mean EMD and the relative element abundances were derived. In Section~\\ref{S6}, the theoretical EMD and the fractional areas occupied by the emitting regions are derived, using the method set out in \\cite{simjordan03a}. Our main conclusions are summarized in Section~\\ref{S7}.  \\begin{table} \\caption{\\label{tab1}Stellar parameters adopted.} \\begin{tabular}{p{.8cm}cccc} \\hline Distance$^a$ & Mass$^b$       & Radius$^b$         & $\\log g_{*}^b$  & $\\log N_{\\rm H}^c$ \\\\ (pc)       & ($M_{\\odot}$)  & ($R_{\\odot}$)      &                 & \\\\ \\hline 3.218        & $0.9\\,\\pm\\,0.1$ &  $0.743\\,\\pm\\,0.01$ & $4.65\\,\\pm\\,0.1$ & $17.88\\,\\pm\\,0.07$\\\\ \\hline \\end{tabular}  $^a$ The {\\it Hipparcos} catalogue \\cite{esa97}. \\\\ $^b$ \\cite{difolco04}; $g_*$ in cm~s$^{-2}$. \\cite{ap04} give a value of $\\log g_{*} = 4.621$.  \\\\ $^c$ \\cite{dring97}; $N_{\\rm H}$ in cm$^{-2}$. \\end{table} ",
        "conclusions": "\\label{S7}  We have analysed line fluxes measured with the LETGS on {\\it Chandra} to obtain an apparent EMD. As part of this work we have examined the self-consistency of the results from lines of a given ion. There remain inconsistencies in the results from different lines of Fe\\,{\\sc xvi} and, to a lesser extent, Fe\\,{\\sc xv}. One source of these differences could be excitations to n-states not yet included in the atomic models, followed by cascades. Although the lines are weak, there is a significant difference between the observed and predicted fluxes for the 3p -- 2s transition in Ne\\,{\\sc viii}. The atomic models and data do not appear to have been updated in CHIANTI since v3. There also needs to be a proper treatment of recombination (including cascades) to the 1s2s~$^3$S level in the He\\,{\\sc i}-like ions and di-electronic recombination needs to be included. The pressure indicated by the f/i ratio in O\\,{\\sc vii} is currently somewhat larger than expected from the final model. Although blends have been taken into account in analysing the Fe\\,{\\sc xiii} lines at around 203.8~\\AA\\, the observed flux is lower than predicted by CHIANTI (v5.2), unless $\\log P_{\\rm e}$ is lower than 15.30, which is not consistent with the results from other lines.  Line optical depths can be estimated using the final model that includes the variable area factors. Several lines, in particular the resonance line of Fe\\,{\\sc ix} at 171~\\AA, are estimated to have line-centre optical depths approaching 1. The effects on the measured fluxes will depend on the geometry; scattering {\\it out of} the line of sight would be expected for lines formed in supergranulation boundaries, but to find the expected flux when integrated over the whole star would require detailed radiative transfer calculations.  Relative element abundances have been determined in the upper transition region/corona. These agree with those found previously by \\cite{sf04} to within the expected uncertainties, and quite well with those of \\cite{wood06}, in spite of differences in the mean EMDs. This reproducibility lends support to the individual line-based methods of deriving abundances. The EMD found here is based on line fluxes measured from the LETGS spectrum, the STIS spectrum (for O\\,{\\sc vi}) and on the {\\it EUVE} counts measured by \\cite{schmitt96} for Fe\\,{\\sc ix} to {\\sc xii} and Fe\\,{\\sc xiv} to {\\sc xvi}. This EMD gives a consistent relative element abundance of iron for all stages of ionization included, apart from Fe\\,{\\sc ix}, for which the line flux is observed to be weaker than predicted. Using the EMD found by \\cite{sf04} leads to a larger difference between the abundances found from Fe\\,{\\sc xi}, {\\sc xii} and {\\sc xiv}. On the basis of the relative abundances of oxygen and iron, we conclude that there is no clear evidence of any difference between photospheric and coronal abundances for low and high FIP elements. Indeed, none of our previous studies of stellar transition regions has shown any clear evidence of FIP effects. The spatially integrated X-ray line fluxes are dominated by the first pressure-squared isothermal scale height; if changes in relative element abundances are occurring at greater heights in the corona, they might not be detectable. (See also discussion by \\citealt{wood06} regarding correlations with stellar mass-loss rates.)  The important Ne/O relative abundance is found to be -0.53 (on a logarithmic scale), slightly smaller than the value of -0.44 found by \\cite{sf04} and \\cite{wood06}, but larger than the values of -0.82 and -0.75 recommended by \\cite{asplund05} and \\cite{grevesse98}, respectively, for the solar atmosphere. It is also larger than that found for the solar transition region by \\cite{young06} (-0.77). \\cite{draketesta05} derived an even larger value (-0.39) from studies of a range of active stars, but used a more approximate method, not a full study of the EMD. Since the relative abundances of O and Fe found in the stellar corona agree with the range of stellar photospheric values to within $\\pm 0.09$ in the logarithm, it appears that it is the abundance of Ne in the stellar corona that differs from the recommended solar values. \\cite{draketesta05} have suggested that a larger Ne/O abundance ratio could resolve the difficulties introduced by the adoption of the lower C, N, and O abundances proposed by \\cite{asplund05}.  The similarity of the variation of the EMD with $T_e$ found by using lines from individual isoelectronic sequences or just one element can be seen from Figs. 3 and 4. We are of the opinion that line-based analyses, using the full emission measure contribution functions is the best method for determining the mean EMD.  Recent work by the other authors mentioned above has concentrated on deriving a mean EMD and element abundances but models of the atmosphere were not produced. We regard the main purpose of deriving the mean EMD is to use it in comparisons with theoretical models based on assumptions about the energy balance. Early work on modelling in terms of loop structures was not entirely successful \\citep[e.g. that by][]{giampapa85}. We intend to study a larger sample of stars in future work, to investigate the systematic behaviour with stellar activity.  The mean (apparent) EMD has been compared with the predictions of models based on the assumption of an energy balance between the divergence of the thermal conductive flux and the radiation losses. This allows the fractional area of the emitting material to be found as a function of $T_{\\rm e}$. When the variation of the fractional area is not included, we find a fractional area that is constant at about 0.24 up to $\\log T_{\\rm e} = 5.8$ and then increases with $T_{\\rm e}$ to reach about 1 by $\\log T_{\\rm e} = 6.4$. Allowing for the variation of the fractional area with $T_{\\rm e}$ increases the calculated values of $P_{\\rm e}$ and reduces the base fractional to 0.14 and the coronal area factor to 0.43. The solutions derived here do not have the problem of fractional areas that are greater than 1 that occurred in the earlier analyses by \\cite{simjordan05}. The derived behaviour is similar to the trend in the area occupied by the supergranulation network boundaries in the solar transition region and inner corona, taking into account that the stellar transition region extends to higher temperatures, because the coronal temperature is higher. At the lower end of the temperature range, the surrounding material would be at near coronal temperatures; the coronal filling factor of less than one is consistent with the additional presence of active regions."
    },
    "0711/0711.1388_arXiv.txt": {
        "abstract": "We solve spherically symmetric radiation flows under full special relativity with the help of a variable Eddington factor $f(\\tau, \\beta)$, where $\\tau$ is the optical depth and $\\beta$ is the flow velocity normalized by the speed of light. Relativistic radiation hydrodynamics under the moment formalism has several complex problems, such as a closure relation. Conventional moment equations closed with the traditional Eddington approximation in the comoving frame have singularity, beyond which the flow cannot be accelerated. In order to avoid such a pathological behavior inherent in the relativistic moment formalism, we propose a variable Eddington factor, which depends on the flow velocity as well as the optical depth, for the case of the sperically symmertic one-dimensional flow. We then calculate the relaticistic spherical flow with such variable Eddington factors to investigate the case that gas is accelerated by radiative force. As a result, it is shown that the gas speed reaches around the speed of light by radiation pressure. ",
        "introduction": "Relativistic outflows from a luminous central object are observed in various active phenomena; e.g., relativistic jets and winds from microquasars (Mirabel, Rodr\\'\\i guez 1999; Fender et al. 2004), such as SS~433, GRS~1915$+$105, GRO~J1655$-$40, jets in active galactic nuclei, such as 3C~273 and gamma-ray bursts (M\\'esz\\'aros 2002). Several mechanisms were proposed to explain these relativistic outflows, including hydrodynamical, radiative, and magnetic drives. When the luminosity highly exceeds the Eddington one, the relativistic outflow seems to be driven by radiation pressure of the central object.  So far relativistic outflows or winds driven by radiation pressure in the spherically symmetric case have been studied by several researchers (Castor 1972; Ruggles, Bath 1979; Mihalas 1980; Quinn, Paczy\\'nski 1985; Turolla et al. 1986; Paczy\\'nski 1990; King, Pounds 2003). As for numerical calculation, the radiation transfer has been sovled in two or three dimensions by a Newtonian treatment, but has not been resolved sufficiently for a highly relativistic case yet.  On the other hand, under the traditional moment formalism, the relativistic outflows driven by radiation pressure have pathological behavior (e.g., Turolla, Nobili 1988; Nobili et al. 1991; Turolla et al. 1995; Dullemond 1999; Fukue 2005). That is to say, moment equations for relativistic radiation transfer can have unphysical critical points. For example, in one-dimensional relativistic radiation flow using the Eddington approximation in the comoving frame, where the moment equations are truncated at the second order, the singularity appears when the flow velocity becomes $c/\\sqrt{3}$. This is understood as follows (Turolla, Nobili 1988; Nobili et al. 1991; Dullemond 1999; Fukue 2006). The radiative diffusion may become {\\it anisotropic} even in the comoving frame of the gas as a result of what the velocity gradient becomes very large in the direction of the flow when the gaseous flow is radiatively accelerated up to the relativistic regime. Hence, in a flow that is accelerated from subrelativistic to relativistic regimes, the Eddington factor should be different from $1/3$ even in the optically thick diffusion limit.   As already stressed in the literature (e.g., Nobili et al. 1991), the location of the critical point in the moment equations depends on the choice of a closure relation, and with a suitable choice of the closure relation, the critical point may disappear. For example, Nobili et al. (1991) adopted a variable Eddington factor, which depends on the optical depth. However, for the present transfer flow, the critical condition, where the denominator of moment equations vanishes, contains the flow velocity (e.g., Nobili et al. 1991; Fukue 2006). Hence, as a natural extension, in the present study we have proposed a variable Eddington factor which depends on the ``flow velocity'' as well as the optical depth. By adopting such a velocity-dependent Eddington factor, we intend to send the critical point away toward the edge of the speed of light.      In this paper we propose a velocity-dependent Eddington factor $f(\\tau, \\beta)$ for a spherically symmetric case, and solve the fully special relativistic spherical outflows driven by radiation pressure using such $f(\\tau, \\beta)$.   In the next section we propose a variable Eddington factor for spherical relativistic radiative flows. In section 3 we describe the basic equations for relativistic outflow driven by radiation pressure under the spherical symmetry. In section 4 we show our numerical results of the radiative flow. The final section is devoted to concluding remarks. ",
        "conclusions": "In the present paper, we examine the relativistic radiation flow in the spherically symmetric case with the velocity- and optical depth- dependent variable Eddington factors within the framework of special relativity. We showed that in the relativisic spherical flow the Eddington factor is no longer constant, but depends on the velocity as well as the optical depth. In particular, when the gas is accelerated up to the relativistic speed, there exists a strong velocity gradient, and the velocity dependence of the Eddington factor becomes important. In addition, such a variable factor can avoid the pathological singularity in the moment equations. We emphasize that we should use such a generalized Eddington factor to treat the relativistic radiation hydrodynamics under the moment formalism.  We can find several solutions for the relativistic spherical flow. The results, however, are slightly different for the Eddington factor adopted. In order to determine the precise form of the Eddington factor, we must solve the relativistic transfer equation rigorously. However, the functional form of the variable Eddington factor is usuful for the study of the relativistic jets, black-hole winds, and the gamma-ray bursts.   It should be commented on the current works on the related topics. Current works are divided mainly into two categories, as refered in the introduction. One type solved the relativistic radiation hydrodynamical equations under the diffusion approximation (e.g., Ruggles, Bath 1979; Quinn, Paczy\\'nski 1985; Paczy\\'nski, Pr\\'oszy\\'nski 1986; Turolla et al. 1986; Paczy\\'nski 1990; Nobili et al. 1994). In these current works the flow is restricted in the subrelativistic region on the order of $\\sim 0.1~c$. However, the diffusion approximatioon may be valid only in the sufficiently optically thick regime, and further, there is no justification that the diffusion approximation can be used in the relativistic regime, since there exists a causality problem. Another type examined the pathological behavior of the traditional moment formalism in the relativistic regime (e.g., Turolla, Nobili 1988; Nobili et al. 1991; Turolla et al. 1995; Dullemond 1999; Fukue 2005), which is one of the motivation of the present study. However, there is no proposal to use a variable Eddington factor, which depends on the flow velocity as well as the optical depth, in order to solve the moment equations in the relativistic regime in the spherically symmetric case. We thus tried to solve the relativistic moment equations with an approximate form of the variable Eddington factor.   In this paper, we considered only the one-dimensinal case without gravity under special relativity. In order to clarify the physics of relativistic jets around a black hole, we must treat the problem within the framework of general relativity. Such a case is a next work.   \\vspace*{1pc}  This work has been supported in part by a Grant-in-Aid for Scientific Research (18540240 JF) of the Ministry of Education, Culture, Sports, Science and Technology."
    },
    "0711/0711.4278_arXiv.txt": {
        "abstract": "Some pulsars have their maximum observable energy output in the gamma-ray band, offering the possibility of using these high-energy photons as probes of the particle acceleration and interaction processes in pulsar magnetospheres. After an extended hiatus between satellite missions, the recently-launched AGILE mission and the upcoming Gamma-ray Large Area Space Telescope (GLAST) Large Area Telescope (LAT) will allow gamma-ray tests of the theoretical models developed based on past discoveries. With its greatly improved sensitivity, better angular resolution, and larger energy reach than older instruments, GLAST LAT should detect dozens to hundreds of new gamma-ray pulsars and measure luminosities, light curves, and phase-resolved spectra with unprecedented resolution.  It will also have the potential to find radio-quiet pulsars like Geminga, using blind search techniques.  Cooperation with radio and X-ray pulsar astronomers is an important aspect of the LAT team's planning for pulsar studies. ",
        "introduction": "Balloon and small satellite observations revealed gamma radiation from the Crab (e.g.\\cite{Browning71}) and Vela \\cite{djt75} pulsars in the 1970's, becoming early examples of multiwavelength pulsar studies. The number of gamma-ray pulsars grew to at least seven during the Compton Gamma Ray Observatory (CGRO) mission in the 1990's (for a summary, see \\cite{djt04}). These high-energy photons are produced by primary interactions of the energetic particles accelerated in the pulsar magnetosphere.  Their potential for studying some basic interaction processes was one stimulus for the development of pulsar models that included gamma rays as a characteristic.  With the recent launch of AGILE and the upcoming launch of GLAST, the gamma-ray window on pulsars offers new promise. ",
        "conclusions": ""
    },
    "0711/0711.3344_arXiv.txt": {
        "abstract": "We analyze the properties of the star S2 orbiting the supermassive black hole at the center of the Galaxy. A high quality SINFONI H and K band spectrum obtained from coadding  23.5 hours of observation between 2004 and 2007 reveals that S2 is an early B dwarf (B0--2.5V). Using model atmospheres, we constrain its stellar and wind properties. We show that S2 is a genuine massive star, and not the core of a stripped giant star as sometimes speculated to resolve the problem of star formation so close to the supermassive black hole. We give an upper limit on its mass loss rate, and show that it is He enriched, possibly because of the presence of a magnetic field. ",
        "introduction": "\\label{intro}  The central parsec of our Galaxy hosts a large population of young massive stars \\citep{krabbe95,pgm06}. Their presence is puzzling since according to standard theories, star formation should not happen so close to the supermassive black hole SgrA*: the tidal forces are so large that any molecular cloud should be disrupted before being able to collapse \\citep{morris93}. For most of the young stars, the solution seems to be star formation in dense accretion disks as witnessed by the two counter-rotating flat stellar structures around SgrA* \\citep{lb03,genzel03,pgm06}. However, this attractive solution does not explain the so-called ``S stars'', the group of objects located within one arcsecond (=0.038 pc) of the black hole. Their orbits are randomly oriented \\citep{frank05} and thus do not fit the accretion disk scenario. Spectroscopically identified as B stars \\citep{ghez03,frank05}, their true nature as massive stars has been challenged: they have been proposed to be the hot, luminous cores of evolved red giant and/or AGB stars the envelope of which has been stripped by tidal interactions with SgrA* \\citep{tal05,kd05}. In that case, the problem of their formation process vanishes since they might very well have formed away from the hostile environment of the black hole before having migrated towards the Galactic Center. Such a process is too slow to explain the presence of short lived massive stars close to SgrA*. Hence, these S stars constitute a ``paradox of youth'' \\citep{ghez03}.  In this letter, we analyze the properties of the brightest S-star -- S2 -- and show that it is a genuine massive star. ",
        "conclusions": ""
    },
    "0711/0711.3491_arXiv.txt": {
        "abstract": "A useful crude approximation for Abelian functions is developed and applied to orbits.  The bound orbits in the power-law potentials $A r^{-\\alpha}$ take the simple form $(\\ell/ r)^k= 1+e \\cos(m \\phi)$, where $k =2-\\alpha > 0$ and $\\ell$ and $e$ are generalisations of the semi-latus-rectum and the eccentricity.  $m$ is given as a function of `eccentricity'.  For nearly circular orbits $m$ is $\\sqrt{k}$, while the above orbit becomes exact at the energy of escape where $e$ is one and $m$ is $k$. Orbits in the logarithmic potential that gives rise to a constant circular velocity are derived via the limit $\\alpha\\rightarrow 0$.  For such orbits, $r^2$ vibrates almost harmonically whatever the `eccentricity'. Unbound orbits in power-law potentials are given in an appendix. The transformation of orbits in one potential to give orbits in a different potential is used to determine orbits in potentials that are positive powers of $r$.  These transformations are extended to form a group which associates orbits in sets of six potentials, e.g. there are corresponding orbits in the potentials proportional to $r, r^{-{2\\over 3}}, r^{-3}, r^{-6}, r^{-{4\\over 3}}$ and $r^{4}$. A degeneracy reduces this to three, which are $r^{-1}, r^2,$ and $ r^{-4}$ for the Keplerian case.  A generalisation of this group includes the isochrone with the Kepler set. ",
        "introduction": "Since schooldays when we encountered the rigid pendulum, most of us have been frustrated by our inability to integrate in elementary terms Abelian expressions of the form $\\int \\left[S(u) \\right]^{-{1\\over2}} du$, where $S(u)$ has simple zeros at $u_a$ and $u_p \\geq u_a$ but is not quadratic.  In practice $S$ usually depends linearly or quadratically on parameters which we shall call $\\varepsilon$ and $h$, and its zeros $u_p$ and $u_a$ depend on $\\varepsilon$ and $h$ often in quite complicated ways.  In Appendix A \\label{A} we relieve this frustration by showing how for each pair of $u_a$ and $u_p$, $S$ may be replaced at lowest order by a quadratic function {\\bf{with the same zeros}} and the integral may be approximately evaluated parametrically via perturbation theory.  We do not have to solve $S(\\varepsilon, h, u)$ for its zeros $u_p$ and $u_a$.  Instead we regard $u_p$ and $u_a$ as parameters and then easily find the $\\varepsilon (u_a, u_p)$ and $h (u_a, u_p)$ to which they correspond.  The process of replacing $S$ by a different quadratic function for each pair of zeros $u_p$ and $u_a$ we call {\\bf{quadrating}} (after the old verb `to quadrate' which means `to make square').  Surely using `quadrate' to mean `to make quadratic' is not too great an extension!  As the simple, though crude, method developed can be applied to a far wider class of problems than those encountered here, we have mentioned it first in the introduction.  Orbits of the general form \\begin{equation} \\label{eq1} \\left( \\ell/ r\\right)^k = 1+ e \\cos \\left( m \\phi \\right) \\end{equation} have a long history.  Newton in Principia (\\citeyear{1687Newton}) showed that orbits of this form with $k=1$ occurred when the central force was an inverse square law supplemented by an inverse cube force. His celebrated theorem on revolving orbits demonstrates that if $r=r(\\phi)$ is an orbit of angular momentum $h$ under any central force $F(r){\\hat{\\bf{r}}}$, then $r=r(m \\phi)$ is an orbit of angular momentum $m h $ under the central force $\\left[ F-\\left(m^2-1 \\right) h^2 r^{-3} \\right] {\\hat{\\bf{r}}}$.  Newton also pointed out that $r(t)$ was the same for both orbits, so, if the new orbit were viewed from axes revolving at the rate $(m-1)\\dot{\\phi}$, then the two orbits would have the same shape.  But notice that $(m-1) \\dot{\\phi}$ is not a constant rotation rate, but speeds up when $r$ is small and slows down when $r$ is large.  When we view the orbit from axes that rotate uniformly at the same mean rate $(m-1) <\\dot{\\phi}>$, the orbits can have very different shapes involving figures of eight for the more eccentric ones \\citep{1995Notes_RS}.  Recently in a fine paper, \\cite{2006AJ....131.1347S} showed that orbits of moderate or low eccentricity in logarithmic or power-law potentials with or without cores were well approximated by analytic orbits of the form (\\ref{eq1}).  His approximate orbits are surprisingly accurate.  Struck was, in part, stimulated to find this result by a paper by \\cite{1997MNRAS.292..905T} that demonstrated the richness of the resonances in the perturbation theory of these systems.  \\cite{2005MNRAS.358.1273V} have emphasised that the apsidal precession found in non-inverse-square orbits is significantly dependent on the eccentricity of the orbit involved.  Surprisingly, we have been led to orbits of the form (\\ref{eq1}) by looking at orbits of extreme eccentricity $e=1$, where Struck's methods did not give accurate results.  Standard works on orbits, \\cite{1996thor.book.....B}, \\cite{2002ocda.book.....C} and \\cite{BT87}, do not point out that the orbits of zero energy in power-law potentials can be exactly solved analytically.  The same variables can be used to solve the nearly circular orbits.  Since both highly eccentric and small eccentricity orbits can be so solved, it would be surprising if there were {\\bf{not}} a good approximation, based on the same variables, that interpolated between $e=1$ and $e=0$.  Struck's methods do this well for small and moderate eccentricities.  Here we show that all orbits are well approximated by analytic orbits of the form (\\ref{eq1}) for $0 \\leqslant e \\leqslant 1$.  Here $\\ell$ and $e$ are generalisations of the semi-latus-rectum and the eccentricity.  The potential is $Ar^{-\\alpha} = Ar^{k-2}$; this parameterisation, using $k$ rather than $\\alpha$, is chosen in this section to simplify equations (\\ref{eq2}) and (\\ref{eq3}) below. If $r_a$ and $r_p$ are the apocentric and pericentric distances the generalised eccentricity is given by  \\begin{equation}\\label{eq2} e= {\\displaystyle{r_p ^{-k} -r_a ^{-k} \\over r_p ^{-k}+r_a ^{-k}}= {r_a ^k - r_p ^k \\over r_a ^k + r_p ^k}}~, \\end{equation} and the generalised semi-latus-rectum is given by $\\ell$ where  \\begin{equation}\\label{eq3} \\ell^{-k} = {1 \\over 2}\\left(r_p ^{-k} + r_a ^{-k} \\right)~. \\end{equation}  We note that for the Kepler case $k=1$ and the above formulae all reduce to the usual ones.  We write $\\ell = Lr_c$, where $r_c(h)$ is the radius of the circular orbit of angular momentum $h$.  We have $r^k _c = (2-k)^{-1} h^2/A$.  The dimensionless parameters $L$ and $m$ are functions of $e$. Orbits of small eccentricity have $m=\\sqrt{k},$ while those with $e=1$ have $m=k$.  We find the orbits in the $-V^2 \\ln r$ potential from the limiting case $k \\rightarrow 2$.  In Appendix~A we show how to improve the accuracy of our orbits via perturbation theory; however for most purposes the simplicity of the initial approximation (\\ref{eq1}) outweighs the extra complication that accompanies greater accuracy.  Our methods can be applied to non-power law potentials (see \\cite{1979AJ.....84.1697K}) but here, for simplicity, we limit ourselves to power laws.  While our methods can be extended to unbound orbits, the results are less pleasing so they are consigned to Appendix~B.  In section \\ref{sec:transf_theory} we use the transformation theory of Newton, \\citeauthor{1911BuAsI..28..113B}, \\citeauthor{1990Arnold} and others to transform our orbits for $0<k<2$ into orbits in potentials with positive powers of $r$.  We show how that theory can be extended naturally to give a set of transformations that form a group.  We develop the subgroup of switch transformations and show that orbits in the potentials $\\psi \\propto r^{k-2}$ are conjugate to orbits in the potentials $r^{2(2-k)/k}, r^{-k}, r^{2k/(2-k)}, r^{-4/k}$ and $r^{-4/(2-k)}$.  These transformations are not restricted to power laws, although special simplifications occur for them.  Applications are made to Plummer's law.  It is shown that the full group has a transformation that connects the Keplerian potential to the isochrone. ",
        "conclusions": "We have found crude, but useful, approximations to Abelian functions by quadrating the expression under the surd while keeping the end points as the constant parameters.  For our problem, these methods give accuracies to better than 1\\%.  We have shown that the $e=1$ `parabolic' orbits at the energy of escape can be solved exactly, and we have given analytic expressions for $m(e)$ which hold for all eccentricities $0 \\leqslant e \\leqslant 1$.  We have thus illuminated why Struck found these orbits to be such good approximations at low and moderate eccentricities.  The transformation theory has allowed us to extend these results to orbits in potentials which are positive powers of $r$ and we have extended the transformations to form a group."
    },
    "0711/0711.3458_arXiv.txt": {
        "abstract": "We report the discovery of a new class of hydrogen-deficient stars: white dwarfs with an atmosphere primarily composed of carbon, with little or no trace of hydrogen or helium. Our analysis shows that the atmospheric parameters found for these stars do not fit satisfactorily in any of the currently known theories of post-asymptotic giant branch (AGB) evolution, although these objects might be the cooler counter-part of the unique and extensively studied PG 1159 star H1504+65. These stars, together with H1504+65, might thus form a new evolutionary post-AGB sequence. ",
        "introduction": "Canonical stellar evolution predicts that the majority of white dwarfs have a core made of carbon and oxygen which itself is surrounded by a helium layer and also, for $\\sim$ 80 $\\%$ of known white dwarfs, by an additional hydrogen layer. Thus, all white dwarfs have been traditionally found to belong to one of these two categories: those with a hydrogen-rich atmosphere (the DAs) and those with a helium-rich atmosphere (the non-DAs). The latter are believed to be the result of a late He-shell flash \\citep{herwig99} at the end of the post-AGB phase (the so-called born-again scenario) which almost completely removes hydrogen and mixes interior material (mostly carbon and oxygen) with the helium envelope. With cooling, because of gravitational settling, helium gradually diffuses upward and the stars then appear as helium-rich white dwarfs (spectral type DO for $T_{\\rm eff} >  \\sim 40,000$ K and DB below that). When these white dwarfs cool to approximately 13,000 K, carbon makes one final appearance as it is convectively dredged-up from the diffusion tale by the deep superficial helium convection zone \\citep[DQ spectral type,][]{pelletier86}. Evolutionary models \\citep{fontaine05} predict that the carbon abundance reaches a maximum at $\\sim 10,000 $ K and decreases monotonically with decreasing effective temperature (note that even at maximum contamination, helium always remains the dominant constituent of the atmosphere with log C/He $\\sim -3$).  Detailed model atmosphere analysis of DQ white dwarfs \\citep{dufour05} have confirmed the theory on the cool side of the maximum contamination but very little has been done to test the theory on the hot ascending side of the curve. The reason for this is simple: only two objects \\citep[G227-5 and G35-26,][]{wegner85,thejll90} were known to exhibit traces of carbon in the optical in that temperature regime. Fortunately, the situation has recently changed drastically thanks to the discovery of several new hot white dwarfs with carbon lines in the Sloan Digital Sky Survey \\citep{liebert03}. However, appropriate atmospheric models were not available at the time, so a detailed analysis of these stars could not be undertaken. We have thus recently developed the appropriate atmospheric models and performed the analysis of these objects. ",
        "conclusions": ""
    },
    "0711/0711.0989_arXiv.txt": {
        "abstract": "{} {We explore the color dependence of the radial profile of satellite galaxies around isolated parent galaxies.} {Samples of potential satellites selected from large galaxy redshift surveys are significantly contaminated by interlopers -- objects not bound to the parent galaxy.  We use the Sloan Digital Sky Survey to estimate the interloper fraction in samples of candidate satellite galaxies.}{We show that samples of red and blue satellites have different interloper populations: a larger fraction of blue galaxies are likely to be interlopers compared to red galaxies.  Both with and without interloper subtraction, the radial profile of blue satellites is significantly shallower than that of red satellites.  In addition, while red and blue primaries have different interloper fractions, the slope of the corrected radial profiles are consistent after interloper correction.  We discuss the implications of these results for galaxy formation models.} {} ",
        "introduction": "In cold dark matter (CDM) models of the universe, large numbers of dark matter (DM) subhalos lie within the virial radius of larger dark matter halos.  In some halos and subhalos, baryonic material has cooled and formed stars, resulting in a central galaxy and satellite galaxies.  The spatial distribution of satellite galaxies in galaxy-sized halos, then, reflects the evolution of satellite galaxies and the mass accretion history of their parent halos.  For example, dark matter simulations suggest that the angular distribution of subhalos follows the shape of the DM halo, which is indicative of infall of subhalos along filaments \\citep[e.g.,][]{zentner_etal05,libeskind_etal05}, while observations suggest that satellites lie along the major axis of the light distribution for early-type galaxies \\citep[e.g.,][]{sales_lambas04,brainerd05,yang_etal06,azzaro_etal07}.  In dark matter simulations, the radial distribution of subhalos is biased with respect to the density profile of DM halos;  at small separations from the center of the halo (within $\\sim 20-50$\\% of the virial radius), the distribution of DM subhalos has a lower concentration, but it follows the dark matter density profile at larger radii \\citep{ghigna_etal98,colin_etal99,ghigna_etal00, springel_etal01,delucia_etal04,diemand_etal04,gao_etal04,nagai_kravtsov05,maccio_etal06}.   Studies that include baryons, star formation and cooling, however, show that the distribution of galaxies associated with subhalos has a steeper inner profile than the subhalo distribution, both at cluster and at galaxy scales \\citep{nagai_kravtsov05,maccio_etal06}.  For samples of subhalos selected by tidally-truncated mass, objects near the halo center lose a greater percentage of their dark matter mass than objects near the virial radius.  Stellar mass selected samples of satellite galaxies are resistant to this effect since baryonic components are located in the centers of dark matter subhalos and are tightly bound.  The observed radial profile of satellites in galaxy-sized halos is generally more centrally concentrated than the subhalo distribution  \\citep{chen_etal06} but consistent with the dark matter profile \\citep{vandenbosch_etal05,chen_etal06}, although \\citet{sales_lambas05} have contrary results.  The color dependence in the spatial distribution of satellite galaxies has been studied extensively in angular distributions.  The angular distribution of satellites is anisotropic and aligned with the major axis for red host galaxies and is consistent with isotropic for blue hosts \\citep{yang_etal06,agustsson_brainerd08,azzaro_etal07}.  These results are consistent with the picture where the orientation of the major axis of elliptical galaxies is determined by the direction along which subhalos are falling into the galaxy -- along filaments.  The orientation of spiral galaxies is determined by the angular momentum vector, which, in simulations, has shown only poor alignment with the minor axis of the dark matter halo.  A secondary result found in these studies is that red satellites of red hosts show a stronger anisotropy than blue satellites.  This result might be explained by a scenario where satellite color is determined by accretion time;  red satellites were accreted earlier, while blue satellites represent recent infall.  If the major axis of the host galaxy is set early and the orientation of infalling subhalos with respect to the host galaxy changes over time, then objects which accreted early would show a stronger alignment with respect to the host galaxy.  These observational results have been analyzed by comparison to mock galaxy catalogs derived from semi-analytic models  -- simulating galaxies using the mass accretion histories of dark matter halos in a simulation and a variety of assumption about the physics of galaxy formation and evolution.  \\citet{agustsson_brainerd08} use mock galaxy catalogs to suggest that the difference in degree of anisotropy between blue and red host galaxies is real and not due to differences in the interloper populations, while \\citet{kang_etal07} come to an opposite conclusion using a very different technique --  the color dependence is due to interlopers in the group catalog.  \\citet{kang_etal07} also conclude that the difference in alignment between red and blue satellites is due to the masses of the satellites;  red satellites are larger and associated with subhalos which were more massive at the epoch of accretion and which in simulations are accreted more preferentially along the major axis of the halo \\citep{libeskind_etal05,wang_etal05}.  The semi-analytic model of \\citet{agustsson_brainerd08} suggests that the amount of interloper contamination is much greater for blue satellites than red satellites ($\\sim 50\\%$ to $\\sim15\\%$, defining interlopers as objects with a physical separation greater than 500 kpc from the parent galaxy), although the color dependence in the anisotropy of satellites is still observed after subtracting interlopers.  The color dependence in the radial distribution of satellites has been studied by \\citet{sales_etal07} using semi-analytic galaxy catalogs constructed from the Millennium Simulation.  They note that the radial distribution of red satellites is significantly more centrally concentrated than the distribution of blue satellites.  This trend is attributable to accretion time in a scenario similar to one explanation for greater anisotropy in the angular distribution of red satellites compared to blue satellites:  early accreting satellites are stripped of hot gas, stop forming stars, and become redder.  Early accreters are also likely to orbit closer to the center.  Observationally, the effects of color selection on the radial distribution of satellite galaxies has been previously examined by \\citet{sales_lambas04} using data from the Two Degree Field Galaxy Redshift Survey (2dFGRS).  They find a steeper outer slope for satellites of blue parent galaxies than for red primaries.  In addition, red satellites have a distribution that requires a larger core radius than blue satellites and they attribute this result to correlations between the primary and satellite properties.  However, in observational searches for satellites, candidate satellites are chosen based upon their projected separation and velocity difference from the parent galaxy.  Samples of objects chosen in this manner are heavily contaminated by objects which are not satellites -- {\\it interlopers}.  In \\citet{chen_etal06}, we discussed the importance of interlopers in making relatively unbiased estimates of the projected radial distribution of satellite galaxies and developed a reliable method of interloper subtraction.  Other methods to account for interlopers have been investigated.  For example, \\citet{vandenbosch_etal04} exclude interlopers using an iterative, adaptive selection criterion for satellites, and \\citet{chen08} model satellites and interlopers together using a halo occupation distribution (HOD) based analytic model for the galaxy correlation function.  The \\citet{chen_etal06} approach has the advantage that it may easily be compared to previous results.  In this paper, we add color selection to the interloper estimation method described in \\citet{chen_etal06} and apply it to data from the Sloan Digital Sky Survey (SDSS) spectroscopic sample.  We show that samples of red and blue candidate satellites have different levels of contamination by interlopers, while the model interloper samples for red and blue primaries are similar in color distribution.  We discuss these results in terms of the environmental dependence of satellite galaxies and wider applications to the angular distribution of satellite galaxies and to galaxy formation models. ",
        "conclusions": "\\label{sec:conclusions}  In \\citet{chen_etal06}, we constrained the projected radial distribution for isolated galaxies and found that their power-law slopes are steeper than the expected slopes for the distribution of dark matter subhalos and may be as steep as the density profile of the host dark matter halos.  We reproduce this result with a survey area that is $\\sim 50\\%$ larger.  However, the distribution of satellite galaxies shows some significant dependence on color.  When samples of candidate satellites are split by color, we see that blue objects are more likely to be interlopers than red objects.  The observed estimated interloper contamination of red objects in the satellite sample is 16\\%, while that of blue objects is 30\\%.  \\citet{agustsson_brainerd08} produce a difference of 15\\% and 50\\%, respectively, in the interloper contamination of red and blue satellites samples, using a semi-analytic galaxy formation model.  It is, then, not appropriate to assume that red and blue satellites have the same level of interloper contamination and this must be taken into account in testing the color dependence in the angular distribution of satellite galaxies.  This result also suggests that the prevalence of faint red galaxies could be developed into a method to find small groups of galaxies in a manner similar to the method by which the red sequence of early-type galaxies is used to find galaxy clusters \\citep[e.g.,][]{gladders_yee00}.  Both with and without interloper subtraction, the radial profile of blue satellites is significantly shallower than that of red satellites.  The best-fit power-law slope of interloper-subtracted blue satellites is $\\alpha =1.34 \\pm 0.10$ and the that of red satellites is $\\alpha = -1.80 \\pm 0.10$.  This result is consistent with the trend found by \\citet{sales_etal07} who test color dependence in semi-analytic galaxy catalogs constructed from the Millennium Simulation, finding a more centrally concentrated radial distribution of red satellites than of blue satellites.  This trend is repeated when they select by a proxy for accretion time -- whether the satellite retains a DM halo or whether it has been tidally destroyed.  Blue satellites are generally fainter than red satellites.  Correspondingly, \\citet{chen_etal06} found that the best-fit power-law slope for bright satellites is steeper than that of faint satellites (cut at $M_r = -18.28$) -- although without statistical significance -- in a volume-limited sample.  In a flux-limited sample, the reverse relation was found;  however, this sample used a luminosity cut as faint as the faintest satellites in our sample, at $M_{r} = -17.76$, and so is not directly comparable.  When we cut the satellites samples by color and luminosity, we find that the dominant effect is from color;  the best-fit power-law slopes of bright and faint samples of red satellites are consistent as are those of bright and faint blue satellites.  The power-law slope for red satellites is as steep as might be expected for the dark matter density distribution of halos which host galaxies like those found in our primary sample.  On the other hand, the power-law slope for blue satellites is as shallow as the expected subhalo distribution.  The shallowness of the subhalo profile is attributed to tidal stripping. This is unlikely to effect satellite galaxies, since they are located at the centers of the dark matter subhalos.  The shallower profile for blue satellites as compared to red satellites, then, might be interpreted as consistent with the scenario where satellite color is determined by accretion time:  red satellites were accreted earlier, while blue satellites represent more recent infall.  Satellites are expected to be redder in the inner regions of parent halos due to environmental processes that shut off star formation (ram pressure stripping, strangulation, etc.).  This morphological segregation has also been observed in more massive clusters and in the fainter satellites found in the Local Group.  When splitting the satellite sample by primary color, red primaries have a significantly larger fraction of red satellites and somewhat smaller interloper fractions than blue primaries.  After interloper subtraction, the best-fit power-law slopes of satellites of red and blue primaries are consistent within errors.  Red primaries are, on average, more luminous than blue galaxies.  Correspondingly, \\citet{chen_etal06} found that the best-fit power-law slope for satellites bright primaries is consistent with that of satellites of faint primaries (cut at $M_r = -21$).  When we cut the primary sample by luminosity and color, the trend with luminosity is different for red and blue primaries.  The slope of satellites of bright, red primaries is shallower than that of faint, red primaries, a relation probably dominated by the mass of the parent halo as brighter primary galaxies reside in bigger parent halos which have mass distributions characterized by smaller concentrations.  The slope of satellites of bright, blue primaries is steeper than that of faint, blue primaries, as the fraction of red satellites drops with primary luminosity.  The color distribution of objects in model interloper samples are similar, regardless of their environment (as measured by average surface density).  While robust conclusions cannot be drawn as to how similar these distributions are, it suggests that there are fundamental differences between satellite galaxies and faint galaxies in the field.  Intriguingly, this blue-tilted color distribution also resembles that of the satellites sample of blue primaries.  Better understanding of the processes that effect the color of faint objects and the radial distribution of satellite galaxies will require further studies."
    },
    "0711/0711.0806_arXiv.txt": {
        "abstract": "{% From high-precision MOST photometry spanning 35 days the existence of two spots rotating with slightly differing periods is confirmed. From the marginal probability distribution of the derived differential rotation parameter $k$ its {\\it expectation} value as well as confidence limits are computed directly from the data. The result depends on the assumed range in inclination $i$, not on the shape of the prior distributions. Two cases have been considered: (a) The priors for angles, inclination $i$ of the star and spot latitudes $\\beta_{1,2}$, are assumed to be constant over $i$, $\\beta_1$, and $\\beta_2$. (b) The priors are assumed to be constant over $\\cos i$, $\\sin\\beta_1$, and $\\sin\\beta_2$. In both cases the full range of inclination is considered: $0^\\circ\\le i\\le 90^\\circ$. Scale-free parameters, i.\\,e.  periods and spot areas (in case of small spots)  are taken logarithmically. Irrespective of the shape of the prior, $k$ is restricted to $0.03\\le k\\le 0.10$ (one-$\\sigma$ limits). The inclination $i$ of the star is photometrically ill-defined. } ",
        "introduction": "In late-type stars with their deep convection zones non-uni\\-form rotation is driven by the action of the Coriolis force on that convective turbulence (cf. Kitchatinov \\& R\\\"udiger 1993). Now quantitative models of differential rotation for the Sun and solar-like stars are available (R\\\"udiger \\& Kitcha\\-tinov  2005; K\\\"uker \\& R\\\"udiger 2005; K\\\"uker [these proceedings]) which should be compared with real stars.  The outcome of these theoretical efforts is a lapping time which is for a given main-sequence star nearly independent of its rotation period (in the case of the Sun, a G2 dwarf, roughly 100 days).  Precision photometry of a spotted star with spots differing in latitude allows a direct measurement of the differential rotation parameter $k$. It parameterizes the surface rotation. With $P_{\\rm eq}$ denoting the equatorial rotation period that at latitude $\\beta$ is $P_\\beta = P_{\\rm eq}/(1-k\\,\\sin^2\\beta)$.  Here the MOST data (Croll et al. 2006; Croll 2006) of the star $\\varepsilon$ Eri have been reanalyzed in a Bayesian framework. The motivation was to get realistic error estimates by computing $k$'s marginal distribution and to find out, how it depends on the chosen prior. In the essence a Bayesian approach explores the whole likelihood mountain in the $N$-dimensional parameter space. It looks not only to the most probable set of parameter values, but provides {\\em expectation\\/} values. Moreover, from a marginal distribution reliable confidence limits can be easily derived.  The nearby K2 dwarf is young ($\\le$ 1 Gyr), rotates twice as fast as the Sun, and shows strong chromospheric activity (cf. Biazzo et al. 2007). ",
        "conclusions": "As expected, with $0.03\\le k\\le 0.10$, the estimated differential rotation parameter $k$ proves smaller than that of the Sun ($k_\\odot\\approx 0.2$).\t% The horizontal shear is $0.017\\le\\delta\\Omega\\le 0.056$ rad/d, the lapping time 130 d. An independent estimate of the star's inclination would exclude either the low-$i$ solution or the high-$i$ one and would help to constrain $k$ even more.   \\bigskip Acknowledgments: The author thanks the MOST team for the $\\varepsilon$ Eri data and W.W. Weiss from Vienna for valuable discussions."
    },
    "0711/0711.0037_arXiv.txt": {
        "abstract": "We present the results of observations obtained using the MASTER robotic telescope in 2005 --- 2006, including the earliest observations of the optical emission of the gamma-ray bursts GRB~050824 and GRB~060926. Together with later observations, these data yield the brightness-variation law $t^{-0.55\\pm0.05}$ for GRB 050824. An optical flare was detected in GRB 060926 --- a brightness enhancement that repeated the behavior observed in the X-ray variations. The spectrum of GRB 060926 is found to be $F_E \\approx E^{-\\beta}$, where $\\beta = 1.0\\pm0.2$. Limits on the optical brightnesses of 26 gamma-ray bursts have been derived, 9 of these for the first time. Data for more than 90\\% of the accessible sky down to $19^m$ were taken and reduced in real time during the survey. A database has been composed based on these data. Limits have been placed on the rate of optical flares that are not associated with detected gamma-ray bursts, and on the opening angle for the beams of gamma-ray bursts. Three new supernovae have been discovered: SN~2005bv (type Ia) --- the first to be discovered on Russian territory, SN~2005ee --- one of the most powerful type II supernovae known, and SN~2006ak (type Ia). We have obtained an image of SN~2006X during the growth stage and a light curve that fully describes the brightness maximum and exponential decay. A new method for searching for optical transients of gamma-ray bursts detected using triangulation from various spacecraft is proposed and tested. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.0215_arXiv.txt": {
        "abstract": "Non--gravitational processes, such as feedback from galaxies and their active nuclei, are believed to have injected excess entropy into the intracluster gas, and therefore to have modified the density profiles in galaxy clusters during their formation.  Here we study a simple model for this so--called preheating scenario, and ask (i) whether it can simultaneously explain both global X--ray scaling relations and number counts of galaxy clusters, and (ii) whether the amount of entropy required evolves with redshift.  We adopt a baseline entropy profile that fits recent hydrodynamic simulations, modify the hydrostatic equilibrium condition for the gas by including $\\approx$20$\\%$ non--thermal pressure support, and add an entropy floor $K_0$ that is allowed to vary with redshift. We find that the observed luminosity--temperature ($L-T$) relations of low-redshift ($\\langle z\\rangle=0.05$) HIFLUGCS clusters and high-redshift ($\\langle z\\rangle =0.80$) WARPS clusters are best simultaneously reproduced with an evolving entropy floor of $K_0(z)=341(1+z)^{-0.83}h^{-1/3} {\\rm keV}$ $\\rm cm^2$.  If we restrict our analysis to the subset of bright ($kT\\gsim 3$ keV) clusters, we find that the evolving entropy floor can mimic a self-similar evolution in the $L-T$ scaling relation. This degeneracy with self-similar evolution is, however, broken when ($0.5 \\lsim kT\\lsim 3$ keV) clusters are also included.  The $\\sim 60\\%$ entropy increase we find from $z=0.8$ to $z=0.05$ is roughly consistent with that expected if the heating is provided by the evolving global quasar population.  Using the cosmological parameters from the {\\it WMAP} 3-year data with $\\sigma_8=0.76$, our best--fit model underpredicts the number counts of the X-ray galaxy clusters compared to those derived from the 158deg$^2$ ROSAT PSPC survey. Treating $\\sigma_8$ as a free parameter, we find a best--fit value of $\\sigma_8=0.80\\pm 0.02$, in good agreement with the results from a recent combined analysis of the Lyman-$\\alpha$ forest, 3D weak lensing and {\\it WMAP} 3-year data.  For the flux--limited cluster catalogs, we include an intrinsic scatter in log--luminosity at both fixed temperature ($\\sigma_{lnL|T}\\approx 0.3$) and at fixed mass ($\\sigma_{lnL|M}\\approx 0.6)$, but we find this does not have a big effect on our results. ",
        "introduction": "Galaxy clusters, the most massive bound objects in the universe, provide several methods to constrain cosmological models, for example through their abundance \\citep[e.g.,][]{Evrard89,HA91,WEF93,ECF96,VL99,Mantz07}, or their spatial distribution \\citep{schuecker01,RVP02,HH03,BG03,SE03,Linder03}, or both \\citep{schuecker03}. In large future surveys, with tens of thousands of clusters, percent--level statistical constraints are expected to be available on dark energy parameters \\citep{H&M&H01}, including constraints on the evolution of its equation of state parameter $w_a\\equiv -dw/da$ \\citep{weller2002,weller2003,wang2004}.  In order to fully realize the cosmological potential of large cluster samples, it is important to understand the cluster mass-observable relations accurately, at least statistically.  It is very unlikely that the structure of clusters will be understood from ab--initio calculations to the level of precision required for the theoretical uncertainties not to dominate over the exquisite statistical errors ~\\citep[e.g.][]{levine02}. However, in principle, when multiple observables depend on the same mass, the mass--observable relation can be accurately determined from the data itself, simultaneously with cosmological parameters. Several works have proposed and quantified the constraints from such `self--calibration'' \\citep{majumdar2004, wang2004, lima2005}, using parameterized phenomenological relations for the mass--observable relations (for example, power--law scalings, or arbitrary evolution in pre-specified redshifts bins).  It has been argued recently \\citep{YHBW06} that even if cluster structure is not precisely predictable, parameterized physical models can further improve on such phenomenological self--calibration, especially when multiple observables (such as X--ray flux and Sunyaev-Zel'dovich [SZ] decrement) can be predicted from the same physical model \\citep{YHBW06}.  In light of this potential, it is important to fit physically motivated cluster models to as many cluster observables as possible; one then hopes that future observations of larger cluster samples will require further fine--tuning of these models, and, at the same time, deliver useful cosmological constraints \\citep{OBB05,YHBW06}.  The gravitational potential of clusters is dominated by dark matter, whose behavior is determined by gravity alone, and is therefore robustly predictable.  The dark matter profiles of galaxy clusters, apart from the innermost regions, are indeed well understood from three--dimensional numerical simulations \\citep{NFW97,MGQSL98}, and are nearly self-similar, as expected.  The physics of gas, on the other hand, involves complicated non-gravitational processes such as radiative cooling and star formation, galaxy evolution, and various forms of feedback.  If these processes were unimportant, the intracluster gas would trace the self--similar dark matter profile, and its global properties should obey simple scaling relations \\citep{Kaiser86}.  Specifically, its X-ray luminosity $L$, if dominated by thermal Bremsstrahlung, as for clusters with temperature $T>2$ keV, should scale as $L\\propto T^2$. This relation is indeed obeyed by clusters in hydrodynamic simulations without non-gravitational processes \\citep{EMN96,BN98}.  However, the observed $L-T$ scaling relation is significantly steeper than the self-similar prediction, closer to $L\\propto T^3$ \\citep{Markevich98,AE99}. This demonstrates that the effect of non-gravitational processes on the intracluster gas is not negligible, even for ``bulk'' observables.  A long--standing proposal for the dominant such non--gravitational effect is that the intracluster gas is heated by some energy input (from star formation, supernovae explosion, galactic winds and/or active galactic nuclei [AGN]), raising the gas to a higher adiabat before the clusters collapse.  Many authors have investigated the effect of such a preheating, and have shown that simply imposing a minimum ``entropy floor'' for the intracluster gas naturally breaks the self--similarity, and steepens the $L-T$ relation as required by the data \\citep{Kaiser91, EH91, CMT97, TN01, BBLP02, VBBB02}.  The pre--heating idea is further supported by the discovery of excess entropy in the inner regions of low--temperature clusters, which suggests the existence of a universal entropy floor \\citep{PCN99, LDPC00}, and by several other independent lines of evidence \\citep[for a brief summary and a list of references, see, e.g.,][]{BEM01}.  A simple model of pre--heating consists of shifting the entropy profile by an overall additive constant, representing the cumulative effect of non--gravitational processes, assumed to be roughly uniform throughout the gas~\\citep[e.g.][]{VBBB02}.  Recent work has tested this simple model, by comparing its predictions with hydrodynamical simulations \\citep[][hereafter YB07]{YB07}.  The model reproduces the simulation results very well, but comparisons with observations show that although it can predict the global X--ray scaling relations, the model can not reproduce the observed entropy profiles \\citep{PSF03,PA05,PAP06} in detail. This requires the model to be further developed, but as far as the global properties are concerned, it appears to be successful, and it is therefore useful to understand the average properties of the intracluster gas.  In this paper, we adopt this simple preheating model, and focus on comparisons with both the observed $L-T$ scaling relations in the redshift range $0\\lsim z\\lsim 1$, and the observed cumulative number counts of the X-ray clusters.  A previous study \\citep{BEM01} calculated the impact of preheating on the X--ray scaling relations, using a sample of 12 simulated clusters, and found a good fit to the data on local clusters (but has not explicitly compared the expected evolution to observations, and has not made simultaneous predictions for the number counts).  Our work is also somewhat similar to a more recent study by \\cite{OBB05}, who present a more detailed physical model for the intracluster gas, and show that it can reproduce local X--ray scaling relations (this paper also did not study evolution).  Our goal here is to clarify (i) whether the model can simultaneously explain both the scaling relations {\\it and} number counts of galaxy clusters, and (ii) whether the amount of entropy required evolves with redshift.  In comparing our predictions to the $L-T$ scaling relations and the number counts, we also study the effects of scatter in the $L-T$ and $L-M$ relations, and the corresponding selection biases that arise in flux limited survey \\citep{NSRE07}.  Our first goal is motivated by our earlier study \\citep{YHBW06}, in which we found that a similar preheating model, with the entropy adjusted to reproduce observed X--ray and SZ scaling relations, tends to overpredict the number counts of bright clusters, even with a relatively low normalization ($\\sigma_8=0.7$) of the power spectrum. A similar discrepancy was found by \\citet[][although they used a higher normalization, $\\sigma_8=0.84$, and suggested that agreement can be recovered by lowering this value]{OBB05} .  The rest of this paper is organized as the follows. In \\S~\\ref{sec:phmodel}, we describe in detail the formalism to implement the preheating model. In \\S~\\ref{sec:LTscaling}, we compare the predicted $L-T$ scaling relations to observations, and find the best--fit entropy level $K_0$ at two different redshifts. In \\S~\\ref{sec:counts}, we further test the model by comparing predictions for the number counts with observations.  In \\S~\\ref{sec:scatter}, we then study the effect of intrinsic scatters and the corresponding selection effects in flux-limited cluster surveys. In \\S~\\ref{sec:discussion}, we discuss our results, and in \\S~\\ref{sec:conclusions}, we offer our conclusions. ",
        "conclusions": "\\label{sec:conclusions}  There is ample evidence that non--gravitational processes, such as feedback from stars and BHs in galaxies, have injected excess entropy into the intracluster gas, and therefore have modified its density profiles.  While in the simplest scenario, the excess entropy is injected at high redshift, well before clusters actually form, and results in a universal entropy floor in galaxy clusters. A more realistic expectation is that the amount of excess entropy evolves with cosmic epoch, tracking on--going star and BH--formation.  Here we studied a simple model of this preheating scenario, and found that it can simultaneously explain both global X--ray scaling relations and number counts of galaxy clusters.  The level of entropy required between $z=0-1$ is $\\sim 200-300\\, {\\rm keV}$ $\\rm cm^2$, corresponding to $\\approx 0.6-0.9 [(1+\\delta)/100]^{2/3}[(1+z)/2]^2$ keV per particle if the energy is deposited in gas at overdensity $\\delta$ at redshift $z$.  This overall level of enrichment is in agreement with previous studies.  Here we find, additionally, evidence that the entropy floor evolves with redshift, increasing by about $\\sim 60\\%$ from $z=0.8$ to $z=0.05$. This fractional increase is in rough agreement with the evolution expected if the heating rate follows the global evolution of the AGN.  The normalization $\\sigma_8=0.8$ preferred when X--ray cluster number counts are fit with our model is somewhat higher than the best--fit value from the three--year {\\it WMAP} data. For a flux--limited cluster catalog, we also find that including an intrinsic scatter in log--luminosity at both fixed temperature and at fixed mass does not have a big effect on the results.  The models presented in this paper should be improved in future work, refined to fit detailed cluster profiles, in addition to the evolution of global observables, and allowing a cluster-to-cluster variation of the level of heating, with possible systematic dependence on cluster mass, in addition to redshift. It will soon be possible to confront this type of more detailed modeling with forthcoming SZ observations. We expect this comparison to securely establish whether the level of entropy is indeed increasing with cosmic epoch, and to place further interesting constraints on both cluster structure models and cosmology."
    },
    "0711/0711.2210_arXiv.txt": {
        "abstract": "{}{With the Solar Optical Telescope on Hinode, we investigate the basic properties of high-degree solar oscillations observed at two levels in the solar atmosphere, in the G-band (formed in the photosphere) and in the \\ion{Ca}{ii}\\,H line (chromospheric emission).} {We analyzed the data by calculating the individual power spectra as well as the cross-spectral properties, i.e., coherence and phase shift. The observational properties are compared with a simple theoretical model, which includes the effects of correlated noise.} {The results reveal significant frequency shifts between the \\ion{Ca}{ii}\\,H and G-band spectra, in particular above the acoustic cut-off frequency for pseudo-modes. The cross-spectrum phase shows peaks associated with the acoustic oscillation (p-mode) lines, and begins to increase with frequency around the acoustic cut-off. However, we find no phase shift for the (surface gravity wave) f-mode. The observed properties for the p-modes are qualitatively reproduced in a simple model with a correlated background if the correlated noise level in the \\ion{Ca}{ii}\\,H data is higher than in the G-band data. These results suggest that multi-wavelength observations of solar oscillations, in combination with the traditional intensity-velocity observations, may help to determine the level of the correlated background noise and to determine the type of wave excitation sources on the Sun.}{} ",
        "introduction": "\\label{intr} The Solar Optical Telescope \\citep[SOT,][]{Tsuneta2007} on the Hinode spacecraft \\citep{Kosugi2007} provides stable series of high-resolution images of the Sun in several different spectral intervals. These data offer a unique opportunity for investigating properties of solar oscillation modes of high angular degree $\\ell$ (high horizontal wave number). \\citet{Sekii2007} and \\citet{Nagashima2007} have used the SOT Broad-Band Filter Imager (BFI) data in the \\ion{Ca}{ii}\\,H line and the molecular G-band for high-resolution, time-distance helioseismology of supergranulation and for studying the distribution of the oscillation power in a sunspot region. \\citet{Kosovichev2007}, using the \\ion{Ca}{ii}\\,H data, made initial observations of oscillations in the sunspot umbra, excited by a solar flare.  A remarkable feature of the SOT Focal Plane Package \\citep[FPP;][]{Tarbell2006} is its capability of providing simultaneous images in several spectral intervals with sufficiently high cadence (1 minute or shorter), thus enabling multi-wavelength, high-resolution helioseismology. Such studies are important for investigating the physical properties of solar oscillations, their excitation mechanism and interaction with the turbulent plasma with radiative losses of the photosphere and chromosphere. This knowledge is important for more accurately measuring oscillation frequencies of solar resonant modes as well as travel times for global and local helioseismology \\citep{JCD2002,Thompson2004}, and also for a better understanding of how the oscillations leak into the solar atmosphere and contribute to the chromospheric and coronal heating \\citep{Erdelyi2007}.   \\begin{figure*}[t] \\centerline{ \\vbox{\\hbox{ \\psfig{figure=powergband.ps,width=7.9cm} \\hspace{-2cm} \\psfig{figure=powercaIIh.ps,width=7.9cm} }\\vspace{.1cm} \\hbox{ \\psfig{figure=phasediff.ps,width=7.9cm} \\hspace{-2cm} \\psfig{figure=coherence.ps,width=7.9cm} }}} \\caption{From top left to bottom right: (a) the power spectrum from G-band observations, (b) the power spectrum from \\ion{Ca}{ii}\\,H line observations, (c) the phase shift between \\ion{Ca}{ii}\\,H and G-band (colour units are in radians), and (d) the coherence spectrum between the two lines. } \\label{fig1} \\end{figure*}   The goal of this paper is to investigate the spectral and cross-spectral properties of solar oscillations of high angular degree of the quiet Sun, using Hinode/SOT observations in the G-band and the \\ion{Ca}{ii}\\,H line, which sample the oscillation signals in the photosphere and lower chromosphere. Previously, such cross-spectral studies were carried out mostly for simultaneous intensity-velocity data at the photospheric level, using the Michelson Doppler Imager (MDI) \\citep{Straus1999} onboard the Solar and Heliospheric Observatory (SOHO) and Global Oscillation Network Group (GONG) data \\citep{Oliviero2000,Barban2004}. The cross-spectra reveal important information about the oscillations and the correlated and background associated with the wave excitation and propagation processes \\citep[e.g.,][]{Severino2001}.  One of the interesting properties is the asymmetry of line profiles in the oscillation power spectra \\citep{Duvall1993}, which is reversed for the velocity and intensity power spectra. This is explained as an effect on the background noise correlated with the oscillation modes \\citep{Nigam1998a}. The asymmetry may have a significant effect on the helioseismic diagnostics \\citep[e.g.,][]{JCD1998}. The physics of the correlated background is not completely understood as the theoretical models are mostly phenomenological \\citep[e.g.,][]{Nigam1998a, Rosenthal1998,Rast1999,Kumar1999,Magri2001}. The observed line asymmetry is reproduced in realistic numerical simulations \\citep{Georgobiani2003}, and was linked to radiative transfer effects. Determining the level of the correlated background and relating it to the excitation mechanism (determining the depth of the excitation sources and their type) is a difficult and important problem of the Sun's oscillation physics \\citep[e.g.,][]{Jefferies2003,Wachter2005}.  In this paper, we present the cross-spectral analysis of solar oscillation observed in intensity signals at two different levels in the solar atmosphere, and show that the results are qualitatively consistent with a simple analytical model that includes the correlated background, the strength of which increases with height. ",
        "conclusions": "\\label{disc}  The initial high-resolution observations of solar oscillations with Hinode, simultaneously in two spectral intervals, reveal new interesting properties. The mode ridges are extended in both the high angular degree (up to $\\ell=4000$) and high-frequency regions (above the Nyquist frequency of 8.33~mHz). For $\\ell > 2500$, the frequencies of all modes become higher than the acoustic cut-off frequency, but the oscillation power is still concentrated along the mode ridges. There is an indication that the p-mode power in \\ion{Ca}{ii}\\,H is higher at high frequencies than in the G-band spectrum. This is consistent with traveling wave properties as expected from solar oscillation theory.  The cross-spectrum phase difference between the G-band and \\ion{Ca}{ii}\\,H oscillations shows rather unexpected properties. It is striking that the phase difference is quite large along the p-mode ridges below the acoustic cut-off frequency, but is completely absent for the f-mode. The naive standing wave picture for the resonant modes trapped below the acoustic cut-off frequency suggests that the phase difference should also be close to zero along the p-mode ridges. We do not have any explanation for the different phase behaviors of f- and p-modes. Above the acoustic cut-off frequency, the phase difference for the p-modes increases with frequency, which is also consistent with the picture of traveling waves.  Some basic properties of these observations can be qualitatively explained in a simple analytical model of solar oscillations, excited by a localized source \\citep{Abrams1996}, if an extra term corresponding to the so-called correlated background is added to the wave signal \\citep{Nigam1998a}. The correlated background affects the asymmetry of the modal lines below the acoustic cut-off frequency, and shifts the frequencies of the pseudo-modes above the cut-off. To explain this frequency shift, the relative level of the correlated background in the \\ion{Ca}{ii}\\,H data should be higher than in the G-band signal. The physics of the correlated background is not yet fully understood, but it plays an important role in defining the line profile of the resonant modes, and the positions of the pseudo-modes. There have been many attempts to measure the level of the correlated background (and disentangle it from the uncorrelated noise) by fitting the intensity-velocity cross-spectra obtained from the same spectral line. Hinode/SOT multi-wavelength observations open a new opportunity for studying this important aspect of the physics of solar oscillations."
    },
    "0711/0711.2026_arXiv.txt": {
        "abstract": "This paper presents galaxy source counts at 24 microns in the six {\\it Spitzer} Wide-field InfraRed Extragalactic (SWIRE)  fields. The source counts are compared to counts in other fields, and to model predictions that have been updated since the launch of {\\it Spitzer}.  This analysis confirms a very steep rise in the Euclidean-normalized differential number counts between 2 mJy and 0.3 mJy.  Variations in the counts between fields show the effects of sample variance in the flux range 0.5-10 mJy, up to 100\\% larger than Poisson errors. Nonetheless, a ``shoulder'' in the normalized counts persists at around 3 mJy. The peak of the normalized counts at 0.3 mJy is higher and narrower than most models predict. In the ELAIS N1 field, the 24 micron data are combined with {\\it Spitzer}-IRAC data and five-band optical imaging, and these bandmerged data are fit with photometric redshift templates.  Above 1 mJy the counts are dominated by galaxies at $z<0.3$.   By 300 microJy, about 25\\% are between $z\\sim$0.3-0.8, and a significant fraction are at $z\\sim$1.3-2.  At low redshifts the counts are dominated by spirals, and starbursts rise in number density to outnumber the spirals' contribution to the counts below 1 mJy. ",
        "introduction": "Galaxy counts from IRAS first hinted at strong evolution of luminous infrared galaxies \\citep{hacking87, lonsdale93, gregorich95}. Mid-infrared surveys made with the ISOCAM instrument at 15 microns revealed an excess population at z$\\sim$0.8 \\citep{Elbaz99,Elbaz02,franceschini01}. This population was detected in part by the redshifting of the strong 7.7 micron PAH feature into the 15 micron band.  The high sensitivity at 24 microns of the MIPS instrument \\citep{rieke04} onboard {\\it Spitzer} allows infrared-luminous galaxies to be traced to even higher redshifts.  Sensitivity to obscured starbursts at z$\\sim$2 is enhanced due to redshifting of the 7.7 micron feature into the 24 micron band.  First source counts have been presented by \\cite{Chary04} for the deep Great Observatory Origins Deep Survey (GOODS) test field, by \\cite{marleau04} for the First-Look Survey (FLS) and the  GOODS test field, and by \\cite{Papovich04} for a range of Guaranteed Time Observer (GTO) fields plus the GOODS test field.  These counts show a steep bump in the normalized source counts as fluxes decrease from several mJy to about 300 $\\mu$Jy, that is attributed to the population discovered by ISOCAM at z$\\sim$0.8 plus a previously unseen population of galaxies at z$\\sim$1-2.  At still fainter fluxes, the normalized counts drop off to the {\\it Spitzer} confusion limit \\citep{Rodighiero06, Chary04, marleau04, Papovich04}.  This paper presents a detailed analysis of 24 micron source counts from the large fields observed for the {\\it Spitzer} Wide-area InfraRed Extragalactic (SWIRE) program.  The SWIRE survey is the largest of the extragalactic {\\it Spitzer} Legacy Science programs, comprising mapping of 49 square degrees at wavelength from 3.6 to 160 microns \\citep{lonsdale03,lonsdale04}.  This survey is designed to dramatically improve our understanding of galaxy evolution, including the history of star formation, the assembly of stellar mass in galaxies, the nature and impact of accretion processes in active nuclei, and the influence of environment on these processes at angular scales up to 3 deg (such as the study of clustering of ultraluminous infrared galaxies by \\citet{Farrah06}).  The SWIRE dataset encompasses six fields and so is an excellent resource for quantifying the effects of sample variance.  The goal of this paper is to present accurate {\\it Spitzer} 24$\\mu$m counts in the flux ranging from 50 mJy to $\\sim$300$\\mu$Jy. A detailed model-based interpretation of the counts is beyond the scope of this paper.  However, the composition of the counts is explored for the ELAIS N1 field using the IRAC data at 3.6-8$\\mu$m \\citep{Surace06} as well as high-quality optical imaging. This merged dataset provides insights into galaxy SEDs and populations \\citep{mrr05, Polletta07}, and is used in this paper to break down the 24 $\\mu$m source counts by population and redshift.  The paper is organized as follows: The SWIRE observations and data processing are described in Section \\ref{sec:data}. The derivation and validation of source counts are presented in Section \\ref{sec:counts}.   Comparison with models, including recent models updated using {\\it Spitzer} counts, is in Section \\ref{sec:models}.  For the ELAIS N1 field, optical identifications, photometric redshifts, and template-fitting lead to source counts subdivided by redshift and galaxy type in Section \\ref{sec:composition}. ",
        "conclusions": "\\label{sec:summary}  The variation between fields of the SWIRE 24$\\mu$m number counts shows the importance of sample variance.  For fluxes near a few mJy, the field-to-field variation in the counts is 50\\% to 100\\% larger than is accountable solely from Poisson uncertainties.   Based on redshift estimates from template fitting and from models, the counts variations near 3 mJy may result from large-scale structures around $z\\sim 0.5$.   At fainter flux levels, the SWIRE 24$\\mu$m catalogs will be useful for tracing structures at $z\\sim 1$ and beyond.  The 24 $\\mu$m counts seem exceptionally interesting and most models published before the launch of {\\it Spitzer} miss badly.  There may be new populations--\\cite{mrr05} claim there is a large population of Type 2 AGN that have not been accounted for in previous models.  The SWIRE counts confirm a steep rise in the counts from about 2 to 0.3 mJy that is quite difficult for models to fit.  Template fitting to the SWIRE EN1 sources is successful for at least 85\\% of the galaxies.  The photometric redshifts indicate that at least a quarter of the sources in the 1-10 mJy interval have photometric redshifts in the range 0.3 to 0.8.   The 25$\\mu$m band onboard IRAS is closest in wavelength space to the MIPS 24$\\mu$m passband.  IRAS counts are complete to about 250 mJy \\citep{Shupe98} and the SWIRE counts go to about 1000x fainter. SWIRE results should be  the best 24 $\\mu$m counts above 250 $\\mu$Jy because it will cover the most area of the {\\it Spitzer} extragalactic surveys.  The full area of {\\it Spitzer} surveys solidify the counts up to about 30 mJy, almost allowing linking up with the all-sky IRAS counts, but leaving a gap of a factor of ten in flux.  It will be left for the {\\it Akari (ASTRO-F)} and {\\it WISE} missions to bridge this gap."
    },
    "0711/0711.1355_arXiv.txt": {
        "abstract": "Stellar population synthesis codes provide the description of the total luminosity resulting from the combination of an ensemble of sources as a function of given physical parameters. Their final aim is to enable users to make inferences about the physical parameters (star formation history, mass transformed into stars, etc.) from the light observed in stellar clusters or galaxies. However, synthesis codes results cannot be interpreted in an one-to-one relation between an observed total luminosity and a theoretical model result since (a) we have not access to the intimate composition of the ensemble of stars and (b) the very theoretical method implicit in populations synthesis makes use of (probabilistic) distributions and there is not an unique model result that describes the integrated luminosity of an ensemble. In general, synthesis models provide the {\\it mean} value of the distribution of possible integrated luminosities, this distribution (and not only its mean value) being the actual description of the integrated luminosity.  Therefore, to obtain the closest model to an observation only provides confidence about the precision of such a fit, but not information about the accuracy of the result. In this contribution we show how to overcome this drawback and we propose the use of the theoretical mean-averaged dispersion that can be produced by synthesis models as a metric of fitting to infer accurate physical parameters of observed systems. ",
        "introduction": "\\label{sec:Intro}  Stellar populations synthesis models have been widely used to make inferences about the evolutionary status of observed systems by means of a comparison between observations and model results. In general it is assumed that the observed system is {\\it well mixed} in terms of stellar populations and that models provide an one-to-one relation between the integrated luminosity and the physical conditions under the model is performed. In other words, it is assumed that the distribution of the possible integrated luminosities of a cluster with given physical conditions (including its number of stars) is a Dirac's delta-like distribution.  However, this assumption is not compatible with the very method implicit in synthesis codes, when they make use of probability distributions, like the Initial Mass Function, to describe the contribution of different stellar types to the integrated light. In the most optimistic case, and following the central limit theorem, the distribution of integrated luminosities must follow a gaussian like probability distribution in the asymptotic case of complete well mixed stellar types in the cluster (i.e. theoretical clusters with infinite number of stars). See \\cite{CLCL06} for more details.  In this contribution we discuss the distributed nature of theoretical integrated luminosities of stellar clusters. We take advantage of such distributions to propose a metric that allows to establish the accuracy of fits of observations with models. We will assume that models are not affected by systematic errors (but see \\cite{CLcancun,CLVO} for detailed discussions on this subject), and we will consider here only the intrinsic dispersion of the theoretical integrated luminosities. ",
        "conclusions": "Four conclusions can be drawn from this work: (a) the comparison of models and observations must be done taking into account the intrinsic metric of each model and not an absolute metric, (b) not all wavelengths must have the same weight in the fitting, (c) fits of line profiles have, in general, a larger dispersion than fits of the continuum, and (d) the more sensitive is the index (line profile, equivalent width, color...) about {\\it a particular} stellar population, the less useful is the index to describe the {\\it overall} evolutionary status of the population."
    },
    "0711/0711.0329_arXiv.txt": {
        "abstract": "We report on the results of an $I$-band time-series photometric survey of NGC 2547 using the MPG/ESO 2.2m telescope with WFI, achieving better than $1\\%$ photometric precision per data point over $14 \\la I \\la 18$.  Candidate cluster members were selected from a $V$ vs $V-I$ colour magnitude diagram over $12.5 < V < 24$ (covering masses from $0.9\\ \\msun$ down to below the brown dwarf limit), finding $800$ candidates, of which we expect $\\sim 330$ to be real cluster members, taking into account contamination from the field (which is most severe at the extremes of our mass range). Searching for periodic variations in these gave $176$ detections over the mass range $0.1 \\la M/\\msun \\la 0.9$.  The rotation period distributions were found to show a clear mass-dependent morphology, qualitatively intermediate between the distributions obtained from similar surveys in NGC 2362 and NGC 2516, as would be expected from the age of this cluster.  Models of the rotational evolution were investigated, finding that the evolution from NGC 2362 to NGC 2547 was qualitatively reproduced (given the uncertainty in the age of NGC 2547) by solid body and core-envelope decoupled models from our earlier NGC 2516 study without need for significant modification. ",
        "introduction": "\\label{intro_section}  NGC 2547 is a relatively well-studied nearby young open cluster; we adopt the parameters of \\citet{nj2006}, who give an age of $38.5^{+3.5}_{-6.5}\\ {\\rm Myr}$, distance modulus $(M - m)_0 = 7.79^{+0.11}_{-0.05}$ and reddening $A_V = 0.186$.  The age derived by these authors puts the cluster at the approximate point where G-dwarfs reach the zero age main sequence (ZAMS; see for example Figure 1 of \\citealt{i2007b}), and slightly younger than the typically-used benchmark of $\\alpha$ Persei ($\\sim 50\\ {\\rm Myr}$, e.g. \\citealt{bm99}).  The small distance modulus makes it possible to probe down to very low masses, with relatively modest-sized telescopes.  The cluster has been surveyed in X-rays using ROSAT \\citep{jt98} and XMM-Newton \\citep{j2006}.  The latter study found that the solar type stars in NGC 2547 exhibited a similar relation between X-ray activity and Rossby number to field stars and older clusters, but with saturated or super-saturated X-ray activity levels, and median levels of X-ray luminosity and X-ray to bolometric luminosity ratio similar to T-Tauri stars in the ONC, but much higher than in the Pleiades. These measures are consistent with the relative youth of NGC 2547.  Surveys of lithium depletion (\\citealt{o2003}; \\citealt{j2003}; \\citealt{j2005}) have placed constraints on the cluster age, and conversely, on stellar evolution models by comparison with isochrone fitting results.  \\citet{j2005} find ages based on lithium depletion of $34 - 36\\ {\\rm Myr}$, reasonably consistent with the isochrone-based estimate we assumed earlier.  \\citet{jtj2000} performed a $v \\sin i$ survey of the solar-type stars in the cluster, finding a distribution indistinguishable from that in the similar-age clusters IC 2391 and IC 2602, with a similar rotation-activity relationship to older clusters such as the Pleiades. \\citet{l2003} found weak evidence for mass segregation in the cluster, confirmed by \\citet{j2003} using a new photometric survey, where they also derived the cluster mass function down to $\\sim 0.05\\ \\msun$, finding a result very similar to the Pleiades mass function over the range $0.075 < M/\\msun < 0.7$.  Finally, NGC 2547 has been observed with \\spitzer using the IRAC and MIPS instruments by \\citet{y2004} to investigate the cluster disc frequency.  \\citet{gorlova07} subsequently extended the survey coverage, finding only $2-4$ out of $\\sim 600$ stars of B---mid-M spectral type to show excesses in the IRAC $8\\ \\micron$ band, a fraction of $\\la 1$ per cent.  In the MIPS $24\\ \\micron$ band, they found $\\sim 40$ per cent of the stars of B---F spectral types to show excesses.  This striking result is most likely due to the discs clearing from the inside outwards, with the lack of short-wavelength excesses caused by the discs being substantially cleared within the central $\\sim 1\\ {\\rm AU}$ parts by the age of NGC 2547.   \\subsection{Evolution of stellar angular momentum} \\label{amevol_section}  For a discussion of the context of this work, the reader is referred to our M34 and NGC 2516 publications in \\citet{i2006} and \\citet{i2007b}.  The present survey probes a different open cluster of younger age ($\\sim 40\\ {\\rm Myr}$ for NGC 2547, compared to $\\sim 150\\ {\\rm Myr}$ for NGC 2516 and $200\\ {\\rm Myr}$ for M34).  Due to the small distance modulus and young age, the NGC 2547 survey probes a relatively wide range of masses, covering $0.1 \\la M/\\msun \\la 0.9$, compared to $0.1 \\la M/\\msun \\la 0.7$ in NGC 2516 and $0.25 \\la M/\\msun \\la 1.0$ in M34.  This will allow a more detailed comparison with models of angular momentum evolution, and in particular to test our results in \\citet{i2007b}, where we found that the evolution of the fast rotators appeared to be better-reproduced by solid body models, and the slow rotators by differentially rotating models.  At the age of NGC 2547, stars with masses $M \\la 0.4\\ \\msun$ should be fully-convective, and at $0.4\\ \\msun$, just beginning to develop a radiative core \\citep{cb97}.  \\subsection{The survey}  We have undertaken a photometric survey in NGC 2547 using the MPG/ESO 2.2m telescope with the Wide Field Imager (WFI; \\citealt{b99}).  Our goals are two-fold: first, to study rotation periods for a sample of low-mass members, covering K and M spectral types, down to $\\sim 0.1\\ \\msun$, and second, to look for eclipsing binary systems containing low-mass stars, to obtain dynamical mass measurements in conjunction with radial velocities from follow-up spectroscopy.  Such systems provide the most accurate determinations of fundamental stellar parameters (in particular, masses) for input to models of stellar evolution, which are poorly constrained in this age range.  We defer discussion of our eclipsing binary candidates to a later paper once we have obtained suitable follow-up spectroscopy.  These observations are part of a larger photometric monitoring survey of young open clusters over a range of ages and metalicities (the Monitor project; \\citealt{hodg06} and \\citealt{a2007}).  The remainder of the paper is structured as follows: the observations and data reduction are described in \\S \\ref{odr_section}, and the colour magnitude diagram (CMD) of the cluster and candidate membership selection are presented in \\S \\ref{memb_section}.  The method we use for obtaining photometric periods is summarised in \\S \\ref{period_section} (see \\citealt{i2006} for a more detailed discussion), and our results are summarised in \\S \\ref{results_section}.  We discuss the implications of these results in \\S \\ref{disc_section}, and our conclusions are summarised in \\S \\ref{conclusions_section}. ",
        "conclusions": "\\label{conclusions_section}  We have reported on results of an $I$-band photometric survey of NGC 2547, covering $\\sim 0.6\\ {\\rm sq. deg}$ of the cluster.  Selection of candidate members in a $V$ versus $V-I$ colour-magnitude diagram using an empirical fit to the cluster sequence found $800$ candidate members, over a $V$ magnitude range of $12.5 < V < 24$ (covering masses from $0.9\\ \\msun$ down to below the brown dwarf limit).  The likely field contamination level was estimated using a simulated catalogue of field objects from the Besan\\c{c}on Galactic models \\citep{r2003}, finding that $\\sim 470$ objects were likely field contaminants, an overall contamination level of $\\sim 59 \\%$, implying that there are $\\sim 330$ real cluster members over this mass range in our field-of-view.  From $\\sim 100\\ {\\rm hours}$ of time-series photometry we derived light curves for $\\sim 130\\,000$ objects in the NGC 2547 field, achieving a precision of $< 1 \\%$ per data point over $14 \\la I \\la 18$. The light curves of our candidate cluster members were searched for periodic variability corresponding to stellar rotation, giving $176$ detections over the mass range $0.1 < M/\\msun < 0.9$.  The rotation period distribution a a function of mass was found to show a clear mass-dependent morphology, intermediate between those found in NGC 2362 (Irwin et al., in preparation) and NGC 2516 \\citep{i2007b}, with a median relation of $J \\propto M$, or $j = {\\rm constant}$, as seen in the other clusters we have surveyed.  In \\S \\ref{model_section}, simple models of the rotational evolution were considered, both for the solid body case, and including differential rotation between a decoupled radiative core and convective envelope, as in our earlier work on NGC 2516 \\citep{i2007b}.  Allowing for the uncertainty in the age of NGC 2547, the results indicate that it is possible to qualitatively reproduce the shape of the observed rotation rate distribution by evolving the NGC 2362 distribution forward in time, using the parameters fit to the NGC 2516 data of \\citet{i2007b}.  It is difficult to examine the effects of differential rotation at this age since they are still relatively small. Some small discrepancies were noted, caused by what appear to be `outliers' in the distributions for NGC 2362 and NGC 2547, but it is difficult to resolve these issues at the present time. Follow-up spectroscopy will be required to make further progress."
    },
    "0711/0711.2430_arXiv.txt": {
        "abstract": "In this paper, genus statistics have been applied to an HI column density map of the Small Magellanic Cloud in order to study its topology. To learn how topology changes with the scale of the system, we provide the study of topology for column density maps at varying resolution. To evaluate the statistical error of the genus we randomly reassign the phases of the Fourier modes while keeping the amplitudes. We find, that at the smallest scales studied ($40 \\mbox{ pc}\\leq\\lambda\\leq 80 \\mbox{ pc}$) the genus shift is in all regions negative, implying a clump topology. At the larger scales ($110 \\mbox{ pc}\\leq\\lambda\\leq 250 \\mbox{ pc}$) the topology shift is detected to be negative in 4 cases and positive (``swiss cheese'' topology) in 2 cases. In 4 regions there is no statistically significant topology shift at large scales. ",
        "introduction": "The Small Magellanic Cloud (SMC) is an irregular dwarf galaxy\\footnote[3]{We assume a distance of 60kpc to the SMC for all calculations throughout this paper.}. Although the SMC is believed to be gravitationally bound to the Milky Way, recent studies suggest that this may not be true \\citep{B07}. The SMC is relatively metal poor and very gas-rich, and represents an ideal galaxy to study the formation of stars in low-metalicity environments. Its proximity to the Large Magellanic Cloud (LMC) and Milky Way has resulted in a turbulent history. Recent evidence points to a close encounter with the LMC some 1.5 GYr ago which deformed the SMC, creating a long, thin filament of HI called the Magellanic Stream \\citep{M74,MF,S02}. The interaction between the SMC and LMC is still dynamic, with observations showing that metal poor star clusters with ages ~$<$ 200 MYr and a metalicity ratio [Fe/H] $<$ -0.6 within the LMC originated from the infalling SMC gas \\citep{BK07}.  Statistical techniques are indispensable while studying turbulence in the interstellar medium (ISM). Images of neutral hydrogen (HI) distribution are frequently used as in most cases it is possible to ignore self-absorption \\citep{L95}. HI also occupies a large portion of the Galactic disc (roughly a 20\\% filling factor) and its movements should reflect large-scale turbulence \\citep{L99}. Furthermore, the prevalence of HI means that it can be studied not only in our galaxy but in nearby galaxies as well.  Several studies have performed statistical analyses of the HI distribution in the SMC \\citep{S99,S02}. These studies determined the spatial power spectrum of the HI intensity and employed the velocity channel analysis (VCA) technique (Lazarian \\& Pogosyan 2000) to reveal shallower-than-Kolmogorov velocity and density spectra. However, power spectra, while being informative about the distribution of the energy with scale, are not sensitive to gas topology. For instance, HI surveys of the SMC have shown numerous filamentary structures and shells of expanding gas. These surveys have detected 501 shells dispersed throughout the SMC (see Fig. \\ref{Shells}), six of which have radii $>$ 350pc and are large enough to be classified as supergiant shells (SGSs) \\citep{SS97}. Many of these shells were created by massive associations of OB stars, but some show no spatial correlation to any young stellar population. These `orphan' shells may be a result of gamma ray bursts or collisions between high velocity clouds and the SMC \\citep{H05}. Alternatively, they could be produced by the ISM turbulence.  This calls for employing other techniques for the statistical studies of SMC. Recent years have been marked by an increased interest to statistical studies of astrophysical turbulence (see Lazarian 2004, for a review). In particular, two techniques Velocity Channel Analysis (VCA) and Velocity Coordinate Spectrum (VCS) which are capable of studying  power spectra of velocity and density have been developed (Lazarian \\& Pogosyan 2000, 2004, 2006). However, the topology of the gas distribution cannot be described by the power spectrum. On the contrary, genus analysis, with its ability to accurately describe and quantify the topology, is a promising synergetic tool in the quest to better understand turbulence.  Genus statistics was developed to study the topology of the universe and distribution of galaxies in three \\citep{GMD86,GWM87} and two (Coles, 1986, 1991,  \\citep{M89},  Plionis et al. 1992, Davis \\& Coles, 1993, Coles et al. 1993) dimensions.  Subsequent projects have used genus statistics to study the topology of the temperature variations within the  Cosmic Microwave Background \\citep{S94,CG03} and have also  been applied to a systematic study of the variations of  the two-dimensional genus with MHD simulations \\citep{KLB07}.  The use of genus statistics for the study of HI was first discussed in \\cite{L99}, and subsequent studies presented the first  genus curves for the SMC \\citep{L02,L04}.  A recent paper by Kim \\& Park  (2007) provided a more thorough study of the topology of  the HI peak brightness temperature distribution in the LMC.  The goal of the present paper is to extend the  technique in Lazarian (2004) and Kim \\& Park  (2007) by providing a quantitative measure of the uncertainties involved in the genus studies. We then apply the technique on the SMC HI data set to quantitatively describe the topology of the HI distribution. Kowal  et al. (2007) have shown that genus statistics of a 3D density distribution from synthetic observations obtained with MHD simulations agrees with the genus results for the 2D integrated column density of the same data set. Therefore, genus statistics complements the power spectrum analyses in providing insights into the physical processes that shape the ISM.  The structure of this paper is organized as follows. \\S 2 discusses our approach to the analysis of data. \\S 3 summarizes the observational procedure and subsequent data analysis of the HI column density map. Section 4 is an analysis of the cropped regions within the SMC, focusing on the genus shift and its topological implications. Section 5 discusses the results and posits astrophysical connections between the genus analysis and the SMC.  In cosmological studies the distributions in question, e.g. the distribution of CMB intensity are nearly Gaussian and genus is used to study small deviations from the Gaussian. Dealing with the ISM, in particularly, as in this paper, with the distribution of column densities of SMC, one cannot expect deviations from symmetry to be small, a priori. ",
        "conclusions": "The genus shift estimated in the previous section can give us insight into the underlying physical processes of the SMC. Referring to the genus shift of the entire SMC (see Figure \\ref{SMC_Entire}), it is readily apparent that the shift varies according to the smoothing scale $\\lambda$. At the smallest scales studied, the genus shift has a negative value, implying a strong clump topology. We can infer that the clumps are caused by clouds of HI gas, as well as the  numerous knots and filaments that compose the SMC. At medium scales (120-200 pc), the genus shift takes on a neutral or slight positive value. This increase in the genus shift can be connected to the abundance of shells that compromise the SMC, as the $\\sim500$ shells which are interspersed throughout the SMC have a mean radius of $\\sim100$ pc \\citep{SS07}. This is consistant with our results, as the largest positive values of the genus shift occur between 120 and 150 pc.  At the largest scales studied ( $>$ 170 pc), the genus shift takes on a slight negative value;  this is due to the prominent `wing' and `bar' features \\citep{S99}, indicating that at even the largest scales, (large) clumps dominate the topology.  From our results, we reached the conclusion that the SMC tends to exhibit a clump topology. The dominance of the ``meat ball'' topology is expected in the case of supersonic turbulence (Kowal et al. 2007\\footnote{The adiabatic approximation, used in the latter work, may not be valid for all the studied here scales, but we still can use its results as a guidance.}).  Several of the surveyed regions display characteristics similar to those of supersonic (low $\\beta$ case) -- an asymmetrical genus curve with a tail that extends into the high-density portion of the plot (see \\cite{KLB07}, figure 18 for more information). However, other possible reasons for a clump topology could be cooling instability and self-gravity of the HI gas.   However, our results differ from the results in Kim \\& Park (2007) who found that the HI distribution in the LMC shows mainly hole topology at intermediate scales, despite the fact that fewer shells are present in the LMC (124 as compared to $\\sim500$ in the SMC). There are several important factors that could contribute to this difference. Firstly, Kim \\& Park (2007) performed genus analysis on the HI peak brightness temperature image, while our work used the HI column density distribution. Peak brightness images emphasize more small-scale and shell structure, while the column density distribution emphasizes density distribution washing out small-scale fluctuations that are caused both by density and velocity fluctuations. Secondly, while the LMC has an almost face-on HI disk, the SMC has a larger inclination and a non-disk-like morphology. In fact, several authors have claimed a large line-of-sight depth of the SMC (see Stanimirovic, Staveley-Smith, Jones 2004 for details). Therefore, any line-of-sight through the SMC integrates over a longer physical depth. It would be interesting to apply our procedure on the LMC HI column density image for a more direct comparison.  Interestingly, for several regions where a hole topology is detected, the estimated shell size is comparable to the one from Kim \\& Park (2007). A possible reason for a hole topology in regions 7 and 9 could be related to the shells created by supernovae explosions and stellar winds. It is interesting to note that those two regions are both off the main `bar' of the SMC and most likely correspond to areas of low line-of-sight depth. Nigra et al. (2008, in preparation) study the Eastern Wing region (our region 9) and find a small line-of-sight thickness. Our region number 7 has been identified in Stanimirovic et al. (1999) and \\cite{H05} as containing several `orphan' shells with high luminosity. \\cite{H05} hypothesize that these shells may be associated with an ancient chimney.  We stress, that it is important to gauge the statistical techniques against numerical simulations. \\cite{KLB07} recently undertook a extensive investigation of density statistics in MHD turbulence, which included different statistical tools, including genus. They concluded that the genus statistic \\textit{is} is sensitive to the sonic Mach number \\textit{$M_{s}$}. In the case where the magnetic pressure dominates, i.e. the high-$\\beta$ case, and subsonic Mach number (e.g. $M_{s}$ $\\approx$ 0.3) the genus curve  is highly symmetrical. For the low-$\\beta$ cases and supersonic (e.g. $M_{s}$ $\\approx$ 2.1, 6.5), the curve stretches into the positive (high-density) side and becomes increasingly non-symmetrical as $M_{s}$ increases. The end result is that it is possible to obtain $M_{s}$ from the plot of genus curve: the Mach number is directly related to the length of the high-density tail. Furthermore, the genus statistic also gives topological information. For values of low $M_{s}$, the genus curve is symmetrical, implying that there are equal numbers of high-density clumps and low-density holes. As $M_{s}$ increases, turbulence creates more high-density structures, which causes the tail of the genus curve to extend towards the high-density side. This corresponds well to the results of studying fractal dimension of density while varying the density threshold in Kowal \\& Lazarian (2007).  In our analysis we assumed that the HI gas is optically thin. This is an important assumption, as the analysis of the effects of absorptions in turbulent gas in Lazarian \\& Pogosyan (2004) suggests that absorption introduces a critical spatial scale for plane-of-sky statistics, with the fluctuations larger than this critical scale being strongly affected by absorption effects. This means that the genus the way we use it above is not applicable to $^{12}$CO data, but should be applicable to C$^{18}$O data, with $13$CO data is in limbo, as this isotope is frequently thick. Potentially, the topology of gas at scales less than the critical scale is also of interest. However, Lazarian \\& Pogosyan (2004) study suggests that fluctuations of the integrated self-absorbing spectral line may be dominated by velocity caustics, provided that the spectrum of density is sufficiently steep, e.g. $E(k)\\sim k^{-\\aleph}$, $\\aleph>1$. The Kolmogorov spectrum corresponds to $E(k)\\sim k^{-5/3}$ and is steep according to the aforementioned definition. The topology studies are, as we discussed above, are different from the studies of spectrum.  However, we believe that the criteria for density fluctuations being observable is the same. Therefore, we expect that for $\\aleph<1$ one can analyze genus for the scales less than the critical one. Note, that in terms of the present study, the constancy of the spectral indexes observed in Stanimirovic \\& Lazarian (1999) provides an additional evidence of HI not being substantially affected by absorption."
    },
    "0711/0711.0435_arXiv.txt": {
        "abstract": "{% Nowadays, g-mode detection is based upon {\\it a priori} theoretical knowledge.  By doing so, detection becomes more restricted to what we can imagine.  {\\it De facto}, the universe of possibilities is made narrower.  Such an approach is pertinent for Bayesian statisticians.  Examples of how Bayesian inferences can be applied to spectral analysis and helioseismic power spectra are given.  Our intention is not to give the full statistical framework (much too ambitious) but to provide an appetizer for going further in the direction of a proper Bayesian inference, especially for detecting gravity modes.} ",
        "introduction": "Since the beginning of helioseismology, the detection of g modes has been the most challenging quest in our field.  There were claims of g-mode detection \\citep[][]{PDPS83, DT95}, none of which were confirmed.  Since the conception of the SOHO mission, one of the goals of this mission was to detect g modes. In 1997, following the lack of g-mode detection by SOHO experimenters, the Phoebus group was formed, with the aim of detecting g modes.  The group set an upper limit to the g-mode amplitude of 10 mm/s at 200 $\\mu$Hz \\citep{TA2000}.  Since then, this lower limit has been even decreased down to about 4.5 mm.s$^{-1}$ for a 10-year observation with a singlet, and down to 1.5 mm.s$^{-1}$ for a multiplet \\citep{YE2006}.  Over the years, the Phoebus group has developed several techniques for g-mode detection.  Some of which are based upon having no knowledge of the structure and dynamics of the Sun.  There are also other techniques using {\\it a priori} knowledge of the Sun such as using rotational splitting patterns \\citep[][and references therein]{YE2006}.  Recent g-mode detection claims rely upon a new detection technique derived from the asymptotic properties of g-mode periods derived from theoretical models \\citep{Rafa2007}.  Although these approaches look promising, they are all ba\\-sed on what we believe we know about the Sun.  From that point of view, it is time that we turn to an approach that has been talked about a lot: a Bayesian approach to statistical inference.  First, I will present a short introduction to what I believe I understand about a Bayesian approach versus a frequentist approach.  In the second section, I will show simple examples of this can be applied to Fourier analysis.  In the third section, the Bayesian approach is applied to well known example in helioseismology.  I then conclude and give a tentative roadmap for the future. ",
        "conclusions": "The Bayesian approach to the analysis of solar power spectra seems extremely promising.  One must not forget that many aspects have been neglected such as marginalization.  On this latter aspect, either one is fond of integration and derives the integrals of Eq. (4), or one uses the so-called Markov Chain Monte Carlo (MCMC) algorithm as used by \\citet{Gregory}.  So far, I have not delved into the subject but I sense that the MCMC algorithms will need to be understood before we can proceed to the proper application of the Bayesian approach to g-mode detection.  Last but not least, this paper could be very far from a {\\it proper} and {\\it accurate} description of Bayesian inference.  Here I wanted to attract the {\\it helioseismic} reader to a field that has recently been fast developing in astrophysics.  I hope this paper will trigger the interest of the reader for a genuine treatment of the Bayesian inference."
    },
    "0711/0711.0745_arXiv.txt": {
        "abstract": "A measurement of the absorption of neutrinos with energies in excess of 10~TeV when traversing the Earth is capable of revealing its density distribution. Unfortunately, the existence of beams with sufficient luminosity for the task has been ruled out by the AMANDA South Pole neutrino telescope. In this letter we point out that, with the advent of second-generation kilometer-scale neutrino detectors, the idea of studying the internal structure of the Earth may be revived using atmospheric neutrinos instead. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.0090_arXiv.txt": {
        "abstract": "We show that the constraint algebra of Ashtekar's Hamiltonian formulation of general relativity can be non-trivially deformed by allowing the cosmological constant to become an arbitrary function of the (Weyl) curvature. Our result implies that there is not one but infinitely many (parameterized by an arbitrary function) four-dimensional gravity theories propagating two degrees of freedom. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0711/0711.2095_arXiv.txt": {
        "abstract": "We present observations of the occultation of Umbriel by Oberon on 4 May, 2007. We believe this is the first observed mutual event between satellites of Uranus. Fitting a simple geometric model to the lightcurve, we measure the mid-event time with a precision of 4 seconds. We assume previously measured values for the albedos of the two satellites \\citep{kar01}, and measure the impact parameter to be $500 \\pm 80$~km. These measurements are more precise than estimates based on current ephemerides for these satellites. Therefore observations of additional mutual events during the 2007--2008 Uranian equinox will provide improved estimates of their orbital and physical parameters. ",
        "introduction": "The planetary satellite systems of the giant planets undergo seasons of mutual eclipses and occultations twice during a planet's orbital revolution, when the Sun and the Earth respectively pass through the planet's equatorial plane. Jovian and Saturnian mutual events have been observed since 1973 \\citep{aks84,arl92,arl97,thu01} resulting in very precise measurements of the satellites' positions from so-called ``photometric astrometry'' \\citep{vas03,nvd03}.  The Uranian system, although it resembles in many respects the Jovian and Saturnian systems, has not yet benefitted from such circumstances. The last Uranian equatorial plane crossing occurred in February 1966, well before the advent of CCD technology.  The 2007--2008 Uranian equinox presents the only opportunity to observe the mutual events of the Uranian satellites until the late 2040s \\citep{chr05,atl06}.  Apart from their value in improving the satellite ephemerides and system constants, mutual event lightcurves can provide information on large-scale albedo variations across the northern hemispheres of the satellites that were in darkness during the Voyager 2 flyby on Uranus in 1986 \\citep{chr05}. Combined with Voyager 2 imagery, they may enable compilation of the first global, albeit crude, albedo maps of these bodies.  Here we report our observations and analysis of an occultation of Umbriel (Uranus II) by Oberon (Uranus IV).  To our knowledge, this constitutes the first ever observation of a mutual event between two satellites of Uranus.  In Section~2, we describe our observing strategy, the equipment used and the data reduction process.  In Section~3 we present the results of our lightcurve analysis.  We discuss the implications of our results for Uranian satellite science in Section~4. ",
        "conclusions": "We have carried out the first observation of a mutual event between two satellites of Uranus, an occultation of Umbriel by Oberon. The parameters of the occultation as estimated from the data have been compared to two different sets of predictions (Table~\\ref{compare}).  One employs GUST86, a Voyager-era ephemeris while the other makes use of the more recent LA06 ephemeris which incorporates post-1986 astrometry of the satellites.  The errors in these predictions reflect the observational uncertainties in the satellite positions used to derive said ephemerides. Typical satellite-to-satellite relative positional errors of 0.03~arcseconds \\citep{chr05} translate to $\\sim 400$~km at the distance of Uranus. For the mutual event observed here, the relative velocity of the satellites is 7~km/s, so the mid-event time predictions are uncertain by $\\sim 60$~seconds. Also, the unusual orientation of the Uranian satellite system renders precise determination of the inclination of the orbit planes difficult when the system is pole-on to the Earth . This was the case until the early 1990s, leading to increased uncertainties in the predicted impact parameters.  We find that our observations are in closer agreement with the LA06 predictions. In this case, considering the above errors, the predicted and observed mid-times are in agreement. Using \\cite{kar01} to fix the relative albedo between the two satellites, we estimate the impact parameter to be $500\\pm80$~km or $0.036\\pm0.006$~arcsec compared to a value of 0.047~arcsec predicted by LA06. The formal errors of our results are smaller than those achieved by conventional astrometry \\citep[e.g.][]{jtw98,vvm99,she02}. We thus expect a considerable improvement in the ephemerides of the satellites to result from observing a large number of such events predicted to occur throughout the rest of 2007 and into 2008. This should also improve our knowledge of some poorly-known physical parameters of the system such as the masses of the inner three satellites Miranda, Ariel and Umbriel \\citep{jac92} and result in a better understanding of the Uranian system as a whole.  Finally, we note that our observations of this event do not strongly constrain the relative albedo of the two satellites. This is due to the degeneracy between the albedo and the impact parameter for a single event. This degeneracy can be lifted either by (a) simultaneous fitting of lightcurves for multiple events sampling different satellite aspects assuming good a priori knowledge of the orbits or (b) a global fit of the albedo and orbit model together.  Although both problems are sensitive to noise, they do not contain fundamental degeneracies and have been successfully used in the past to derive large-scale maps of Pluto \\citep{you99}.  Such a fit can only be attempted when observations of as many events as possible have been successfully acquired.  If successful, it will yield regional to hemispherical albedo information on the unimaged hemispheres of the major uranian satellites."
    },
    "0711/0711.4399_arXiv.txt": {
        "abstract": "%  We have recently completed an observing program with the Australia Telescope Compact Array towards massive star formation regions traced by 6.7\\,GHz methanol maser emission. We found the molecular cores could be separated into groups based on their association with/without methanol maser and 24\\,GHz continuum emission. Analysis of the molecular and ionised gas properties suggested the cores within the groups may be at different evolutionary stages. In this contribution we derive the column densities and temperatures of the cores from the $\\nhthree$ emission and investigate if this can be used as an indicator of the relative evolutionary stages of cores in the sample.  The majority of cores are well fit using single-temperature large velocity gradient models, and exhibit a range of temperatures from $\\sim$10\\,K to $>$200\\,K. Under the simple but reasonable assumption that molecular gas in the cores will heat up and become less quiescent with age due to feedback from the powering source(s), the molecular gas kinetic temperature combined with information of the core kinematics seems a promising probe of relative core age in the earliest evolutionary stages of massive star formation. ",
        "introduction": "Although short in astronomical terms, the formation time-scale for massive stars is long enough that observations of an individual massive star formation (MSF) region can only provide a snap-shot of the formation process. Given the paucity of young Galactic MSF regions, developing an end-to-end model of the formation process will be greatly facilitated by the ability to accurately determine the relative evolutionary stage of different regions. Observationally, MSF regions tend to be placed into broadly defined groups such as ultra-compact and hyper-compact \\hii~regions and hot/cold molecular cores. However, it is not clear if the discrete nature of the formation process implied by this grouping is a physical reality (suggesting each of currently defined groups represent relatively long-lived stages which quickly transition to the next stage), or if the formation process is in fact more continuous.  In an attempt to develop our understanding of the early evolutionary stages of the MSF process, \\citet{L07A} (L07A) have completed an observing program with the Australia Telescope Compact Array (ATCA) towards MSF regions traced by 6.7\\,GHz methanol maser emission (a powerful sign-post of very early massive star formation). The aim of these observations was to derive the physical properties of the gas in these regions and see how accurately it is possible to determine the relative evolutionary stages of the cores. The initial results suggested cores could be separated into different evolutionary stages based on their association with/without methanol maser and 24\\,GHz continuum emission. In this contribution we focus on deriving the core temperatures and investigate the use of molecular gas kinetic temperature as an indicator of the cores' evolutionary stage. ",
        "conclusions": "\\label{sec:conclusions} Based on the $\\nhthree$ observations of \\citet{L07A}, we have performed a preliminary analysis of core temperature towards a sample of very young MSF regions, to try and determine their relative evolutionary stages. We find, \\begin{itemize} \\item a large number of cold ($\\sim$ 10\\,K) cores with very quiescent gas. These are likely to be MSF regions at the earliest evolutionary stages. \\item above $\\sim$20\\,K the cores exhibit a smooth range in temperature suggesting the transition from `cold' to `hot' cores is continuous rather than discrete. \\item two lines of evidence suggesting some of the cores have more than one temperature component. Although more detailed modelling is required to accurately derive the core structure, these are potentially interesting for future study as regions in which the central powering sources has only recently begun to heat up its surrounding environment. \\end{itemize}  \\noindent In summary, molecular gas kinetic temperature combined with information of the core kinematics appears to be a promising way to compare the evolutionary state of very young MSF regions."
    },
    "0711/0711.3436_arXiv.txt": {
        "abstract": " ",
        "introduction": "The WMAP 3--years data \\cite{Peiris:2003ff,Spergel:2006hy} on the cosmic microwave background radiation (CMB) represent a milestone in the history of observational cosmology due to their great precision. The challenge for theoretical cosmology and particle physics is to understand these data within some theoretical framework and to make predictions to be tested by future observations. In this respect, inflation stands as the most successful and promising theoretical scenario. Proposed in 1981 as a solution to the flatness and horizon problems~\\cite{Guth:1980zm} it was later shown to predict an almost scale invariant spectrum of matter perturbations~\\cite{Mukhanov:1981xt, Starobinsky:1982ee, Hawking:1982cz, Guth:1982ec, Bardeen:1983qw}, making it a very successful idea. However, we still lack a concrete description of the mechanism producing inflation in the very early universe. The simplest implementation postulates a real scalar field, $\\phi$, whose homogeneous part drives inflation during its way towards the minimum of its potential, $V(\\phi)$. Therefore, to provide a description of the inflationary epoch of the universe we need to elucidate the functional form of the inflaton potential \\cite{Lidsey:1995np}. In this respect it is fair to say that, despite interesting developments, there is still no convincing inflationary scheme based on particle physics.  The main goal of this work is to compare a broad class of models, which are physically well motivated, with CMB and large scale structure (LSS) data. As shown below, these models predict a scalar spectral index with non-constant slope, in contrast with the ans\\\"atze used in standard analyses. Interestingly, these models are very predictive: the number of independent parameters is comparable to the free parameters used in more phenomenological approaches.  This paper is organized as follows. In section 2 we review the standard way of comparing inflationary potentials with cosmological data through parametrizations of the primordial scalar and tensor power spectra that involve the spectral indices and their derivatives (up to a given order). In section 3 we introduce the class of models we are interested in, namely, models with flat tree-level potentials which are generically lifted in a controlled way by radiative corrections. This type of models is then extended by taking into account possible effects from non-renormalizable operators that appear after integrating out some heavy new physics. In both cases we discuss the physical motivation and significance of such models. Section 4 is devoted to the derivation of  the power spectra for the two classes of models as a function of suitable model parameters. The data analysis procedure is discussed in section 5; we present parameter constraints within each class of models from our analysis in section \\ref{results} while section 7 is devoted to the issue of model comparison, assessing which model is in better agreement with the data. Finally, our conclusions are presented in section 8. In addition, the appendix contains some technical details relevant to the analysis of the model with non-renormalizable operators. ",
        "conclusions": "In this paper we have compared a broad and physically well-motivated class of inflationary models with CMB and LSS observational data. Namely, we have considered models with flat tree-level potentials, which typically appear in supersymmetric theories, where $V_{\\rm tree}^{\\rm SUSY}$ ordinarily has plenty of accidental flat directions. These models, beside being very well motivated from the physical point of view (on a similar footing to monomial potentials), lead to very model-independent cosmological predictions. The reason is that the potential derivatives $V'$, $V''$,... arise from the radiative corrections to $V$, which has a characteristic logarithmic dependence on the inflaton field. This scenario has been labelled ``LOG\" throughout the paper. In addition, we have considered the possible presence of new physics beyond a certain high-energy cut-off. This physics does not need to respect the flat directions of the ``low-energy\" theory, and thus it will show up as non-renormalizable operators (NRO) in the inflaton field, which will be dominated by the lowest--order one. This modified scenario (labelled ``LOG+NRO\") is also very well motivated and still quite model-independent.  We have studied the performance of these scenarios when compared with CMB and LSS. We have made first a detailed study of the features of these models, working out both numerically accurate results and approximate analytical expressions for $P_s(k)$, $P_t(k)$ and other relevant quantities, such as the spectral index, $n(k)$, as a function of suitably defined model parameters. We also discussed the number of independent parameters and the theoretical and phenomenological constraints on them (to be imposed a priori). As a matter of fact, one (combination) of the parameters is almost irrelevant, which makes these models even more predictive. Another parameter is essentially the number of e-folds, $\\Ne$, since the time when the largest observable scales crossed out the horizon until the end of inflation. This allows to perform the fits in a twofold way: either leaving $\\Ne$ free, and let the data determine its value, or imposing a prior on its value according to the usual theoretical prejudice ($\\Ne=50-60$). Both approaches (labelled ${\\cal F}$ and ${\\cal G}$ respectively) are interesting and complementary.  In the analysis we also study the performance of standard parametrizations of the power spectrum, based on a Taylor expansion of $\\ln P_s(k)$ and $\\ln P_t(k)$ around an (arbitrary) pivotal scale, $\\ln k_0$. At first (second) order these parametrizations correspond to a constant (constantly running) spectral index, $n(k)=$~constant ($d n/d\\ln k=$~constant). They have been used in reference analyses, in particular by the WMAP collaboration. It is important to keep in mind that, although useful, the standard parametrizations are not inspired by any particularly well-motivated inflationary physics. E.g. the results of the fit with the $d n/d\\ln k=$~constant assumption are not consistent with a number of e-folds in the required range. Still we have also studied them (going one order beyond the WMAP analysis) to facilitate the discussion of the performance of the LOG and LOG+NRO scenarios. As a general comment, care must be taken to test inflationary models which predict a non-negligible  scale dependence. In many cases the standard Taylor series parametrization  of equations (\\ref{stp}) and (\\ref{stp2}) cannot be accurately used in  such a situation unless a high number of terms is taken in the expansion.   Our main results are the following:  \\begin{itemize}  \\item Both the LOG and LOG+NRO scenarios predict small tensor perturbations: $r_0\\le {\\cal O}(10^{-3})$.  \\item The LOG scenario has essentially two parameters, $P_s(k_0)$ and $\\Ne$, and implies a nearly constant $n(k)$.  \\item Leaving $N_e$ as a free parameter (LOG${\\cal F}$ fit), one gets $24 < \\Ne < 49$ ($16 < \\Ne < 84$) at 68\\% (95\\%) c.l. while the corresponding spectral index is close to $n_0=0.96$. This result is consistent with the theoretical prejudice $\\Ne\\sim 50-60$, required to solve the horizon problem, which is remarkable. In the LOG${\\cal G}$ fit (i.e. imposing $\\Ne=50-60$) one gets $n_0\\simeq 0.98$. Note that this fit has only one parameter and still works very well.  \\item The LOG+NRO scenario has two more parameters, which arise from the order and the suppression scale of the NRO. We have fixed the order of the NRO (denoted $4+2N$ throughout the paper) to two representative values ($N=2, 10$) that reasonably encompass the sensible physical range. Thus in practice we are playing with just one additional parameter. This scenario can produce a sizeable running of the spectral index, and still be consistent with the data and a reasonable number of e-folds, especially if $N$ is not very small. The impact of the NRO (driving a running of the spectral index) is relevant at small $k$, corresponding to the first stages of the inflationary period, and then it quickly converges (especially for not too small $N$) to the LOG scenario.   \\item The model comparison is delicate for several technical and fundamental reasons. In the paper we give a fully Bayesian discussion of the relative quality of the various scenarios considered. Qualitatively, it can be said that the goodness of the LOG (NRO+LOG${\\cal F}$) fits is similar to the standard $n=$~constant ($dn/d\\ln k=$~constant) parametrization. The improvement in the goodness of the fit obtained by the inclusion of an extra parameter (as the LOG+NRO scenario implies) is not enough (with the present data) to justify such modification, but still it remains an interesting theoretical possibility. On the other hand, a rigorous comparison between the evidences for LOG, LOG+NRO scenarios and for the standard parametrizations (which are phenomenological by construction) is not feasible. However, the prior--independent method of the Bayesian calibrated p--values still indicates that no higher order term than a tilt is required in the standard Taylor expansion.  \\end{itemize}  As a final conclusion, the LOG and LOG+NRO scenarios analyzed in this paper (based on flat tree-level potentials without or with the presence of extra physics) are not only very well motivated from the physical point of view, but they also fit remarkably well the CMB and LSS data, with very few parameters (the predictions are quite model independent). In addition they are naturally consistent with a reasonable number of e-folds. Therefore, they can be considered as a standard physical class of inflationary models, on a similar footing as monomial potentials.  \\newpage   \\newpage \\appendix"
    },
    "0711/0711.4811_arXiv.txt": {
        "abstract": "The transparent Sun is modeled as a spherically symmetric and centrally condensed gravitational lens using recent Standard Solar Model (SSM) data.  The Sun's minimum focal length is computed to a refined accuracy of 23.5 $\\pm 0.1$ AU, just beyond the orbit of Uranus.   The Sun creates a single image of a distant point source visible to observers inside this minimum focal length and to observers sufficiently removed from the line connecting the source through the Sun's center.  Regions of space are mapped where three images of a distant point source are created, along with their associated magnifications.  Solar caustics, critical curves, and Einstein rings are computed and discussed.  Extremely high gravitational lens magnifications exist for observers situated so that an angularly small, unlensed source appears near a three-image caustic.  Types of radiations that might undergo significant solar lens magnifications as they can traverse the core of the Sun, including neutrinos and gravitational radiation, are discussed. ",
        "introduction": "Our Sun is known to act as a gravitational lens.  The angular shift of the apparent position of a star located behind the Sun was first observed during the Solar eclipse of 1919 \\citep{eddington}, in conformity with the predictions of general relativity \\citep{einstein16}.  Many years later, Einstein (1936) discussed the lens like action of the Sun's gravitational field more generally.  The influence of that discussion has caused the ring seen by the observer during perfect alignment of the source, lens and the observer to be called the Einstein ring around a point lens.  For an opaque lens, the result of exterior focusing of the parallel light rays incident with an impact parameter $b$ on the lens plane is a deflection $\\alpha$. The deflection $\\alpha$ is inversely proportional to the impact parameter $b$ according to the formula first derived by Einstein (1911) \\begin{equation} \\alpha=\\frac{4GM_b}{b c^2}  , \\label{def} \\end{equation} where $b$ is the impact parameter of the passing light ray, $M_b$ is the deflecting mass, $G$ is the gravitational constant and $c$ is the speed of light in vacuum.  Assuming the deflection angle is small, the minimum focal length for the opaque sun is approximated by \\begin{equation} F_{min}=\\frac{R_\\odot^2 c^2}{4 G M_\\odot}=548.30(\\pm .01)\\;\\rm{AU} , \\label{focal} \\end{equation} where $M_\\odot$ and $R_\\odot$ are the total mass and radius of the Sun, respectively.  The uncertainty in $F_{min}$ is based on statistical errors on $R_{\\odot}$ and $G M_{\\odot}$ \\citep{cox}.  Slight offsets of the lens or source from the optic axis will break the Einstein ring into two bright images as seen by the observer.  Therefore the opaque Sun produces multiple images only beyond $F_{min}$.  The possibility of using the {\\it transparent} Sun as a gravitational lens was first discussed by Cyranski and Lubkin (1974) and more recently by Nemiroff and Ftaclas (1997). Demkov and Puchkov (2000) computed a minimum focal length of the transparent Sun to be about 24 AU.  Starting with the Standard Solar Model as it was known in 1989 \\citep{bahcall89}, they computed the gravitational lens deflection angle as a function of impact parameter and then approximated this with a Taylor's series near the Sun's center.  Throughout the present analysis, the more recent 2005 Standard Solar Model (SSM) data is used \\citep{bahcall05}.   A more complete model of the gravitational lens characteristics of the transparent Sun is computed, including a more accurate minimum focal length (23.5 $\\pm$ 0.1 AU), regions of multiple images, and the locations of caustics and critical curves.  The plan of the paper is as follows: $\\S $ 2 outlines two alternative approaches to obtain the minimal focal length of the transparent Sun. The strength of a composite lens capable of producing multiple images is given a through treatment in $\\S$ 3. The critical curves, caustics, magnification and multiple image zones are computed in $\\S$ 4, and conclusions summarized in $\\S$ 5. ",
        "conclusions": "The effects of plasma and core rotation in the interior of the Sun, thought to be small, were ignored during this analysis.   The minimum focal length, the critical curve and caustics were computed numerically. The magnification of point sources for typical source separations was analyzed. The code can be modified to simulate galaxies and dark matter halos with an assumed density profile as input parameters.  To date, four spacecraft have traveled past the minimum transparent focal length of our Sun, just beyond the orbit of Uranus.  Launching a spacecraft to this distance with sufficiently sensitive gravitational wave or neutrino detectors remains a dream, however.  In a different paper, the possibility of detecting strong gravitational radiation and a wide range of neutrino energy spectrum will be examined in some detail.  With the advancement of modern technology and the improvement in the resolution power of detectors, such an endeavor could shed new light on sources emitting neutrinos, high energy gravitational waves or hitherto hypothetical particles that are predictions of at least some theories.   \\begin{appendix}"
    },
    "0711/0711.1119_arXiv.txt": {
        "abstract": "We study the equation state of strongly interacting quark matter within a NJL-like model in which the chiral condensates and the color superconducting gaps are computed self-consistently as a function of the baryon density. A vector interaction term is added to the Lagrangian in order to render the quark matter equation of state stiffer.  For the low density hadronic phase we use a relativistic mean field model. The phase transition to quark matter is computed by a Maxwell construction.  We show that stable CFL cores in hybrid stars are possible if the superconducting gap is sufficiently large. Moreover we find stable stellar configurations in which two phase transitions occur, a first transition from hadronic matter to 2SC quark matter and a second transition from 2SC quark matter to CFL quark matter. ",
        "introduction": "The possibility that quark matter, and eventually color superconducting quark matter, is present in the center of neutron stars has stimulated many theoretical investigations in the last years both on the modelling of the equation of state (EoS) of quark matter and on the phenomenological signatures of the presence of quark matter in neutron stars \\cite{Alford:2007xm}. Presently the ``state of the art'' for the EoS of quark matter is represented by the three-flavor Nambu-Jona-Lasinio (NJL) model in which both the chiral condensates and the diquark condensates are self-consistently computed as a function of the chemical potential and temperature. Chemical equilibrium and charge neutrality (both electric and color charge) conditions necessary to describe neutron star matter are also imposed in this model \\cite{Steiner:2002gx,Abuki:2004zk,Ruester:2005jc,Blaschke:2005uj,Ippolito:2007uz}. Recently, also the effect of a finite neutrino chemical potential has been included \\cite{Ruester:2005ib,Sandin:2007zr}. The structure of the QCD phase diagram within this model turns out to be very rich, with many different possible quark phases. One of the most striking feature, on which we will focus here, is the first order phase transition between the two flavor phase, 2SC or normal quark matter depending on the diquark coupling constant, and the three-flavor superconducting phase, the Color-Flavor-Locking phase (CFL), at vanishing temperature.  Concerning the phenomenological signatures, the differences between the mass-radius relation for neutron stars and quark or hybrid stars is currently studied. The astrophysical data so far are still affected by large uncertainties but will improve considerably with the advent of new satellite missions, as XEUS, Constellation-X, SKA, JWST and LISA. From recent theoretical studies it turns out that the values of the maximum mass of neutron stars and compact stars containing quark matter are very similar \\cite{Alford:2004pf,Alford:2006vz}. Therefore, it seems difficult with the present knowledge to rule out quark matter from just a mass measurement \\cite{Nice:2005fi,Ozel:2006bv}. Other interesting quantities have been calculated (see Ref.~\\cite{SchaffnerBielich:2007mr} for a recent review) for the different possible phases of quark matter, as the neutrino emissivity and the heat capacity which are important for the cooling of compact stars \\cite{Shovkovy:2002kv,Reddy:2002xc,Aguilera:2005tg,Popov:2005xa,Anglani:2006br} or the bulk viscosity which determines stability with respect to gravitational waves emission via r-modes \\cite{Madsen:1999ci,Sa'd:2006qv,Alford:2006gy,Alford:2007rw,Drago:2003wg,Drago:2007iy}. Also in explosive phenomena, as supernovae and gamma-ray-bursts (GRBs), quark matter can play an important role. For instance, the possibility of a double phase transition, first from hadronic matter to 2-flavor quark matter and then from 2-flavor quark matter to the CFL phase has been proposed to explain the complicated time structure of GRBs exhibiting a long quiescent time in their light curve \\cite{Drago:2005qb,Drago:2005rc,parenti}.    When the above mentioned NJL-EoS is used for the applications on compact stars, hybrid stars become unstable at the onset of the CFL phase and therefore CFL phase can not be present in the core of neutron stars. This conclusion was obtained first in Ref.~\\cite{Baldo:2002ju} where quark matter does not occur at all in compact stars because there is a direct transition from hadronic matter to CFL matter in the model used. In Refs.~\\cite{Buballa:2003et,Klahn:2006iw,Grunfeld:2007jt} for different EoSs for the hadronic matter and different parameters for the NJL model for quark matter, again the CFL phase was ruled out because it renders the stars unstable. The 2SC phase could appear; the conclusion is therefore that only 2-flavor superconducting quark matter can be realized in compact stars. A similar result about the allowed quark phases in compact stars was also found in Ref.~\\cite{Lawley:2006ps} where a modified 2-flavor NJL model, which simulates confinement at low density, is proposed. The dependence of the stability of a quark core from the momentum cut-off of the NJL model has been analysed using density dependent cut-offs, but the instability still persists \\cite{Baldo:2006bt} \\footnote{Very recently, stable hybrid stars configurations have been also obtained by adopting the three flavor LOFF phase for the quark matter EoS \\cite{nicola}.} .  On the other hand, completely different results are obtained using MIT-bag-like models as shown in Refs.~\\cite{Alford:2002rj,Drago:2004vu,Bombaci:2006cs} where the appearance of CFL cores does not compromise the stability of the star. Moreover, the absence of the 2SC phase in compact stars has been demonstrated in Ref.~\\cite{Alford:2002kj}.    In this paper we want to consider again the NJL-EoS as computed in Ref.~\\cite{Ruester:2005jc} to study the structure and composition of compact stars. We will investigate larger windows of the model parameters with respect to previous work, with particular attention to the diquark coupling. Furthermore, we discuss the importance of the procedure used to fix the effective bag constant within the NJL model for the stability of a star when the phase transition to quark matter is considered.  We will investigate a new procedure for fixing the effective bag constant by requiring that the chiral symmetry restoration coincides with the transition from the hadronic to the quark matter description \\footnote{In the following, we adopt the term deconfinement transition for the switch from the hadronic to the quark model, as usually done in the literature. However, the Polyakov loop is not a good order parameter for the deconfinement transition at finite chemical potential anymore but only the quark condensate for the chiral phase transition \\cite{Fraga:2001id,McLerran:2007qj}.}. We use then different EoSs to compute the mass-radius relations of hybrid stars showing that in some cases a stable CFL core is possible and, even more intriguing, that two phase transitions, from hadronic matter to the 2SC and then to the CFL phase, can take place in compact stars. The double phase transition is particular intriguing in connection with the deconfinement quark model of GRBs in which the interpretation of bursts presenting two emission periods is due to a double phase transition in compact stars \\cite{Drago:2005qb,Drago:2005rc,parenti}.  The paper is organised as follows. In Sec.~II we discuss a simple toy model for a first order phase transition between hadronic matter and quark matter in compact stars and we study the stability of the star by varying the parameters of the quark matter EoS. In Sec.~III we compute the quark matter EoS within the NJL model for different sets of parameters and finally in Sec.~IV we discuss the stability of the stars obtained using our EoSs with particular emphasis on the CFL core stability.  In Sec.~V we draw our conclusions. ",
        "conclusions": "We have studied the phase transition from hadronic matter to quark matter and the corresponding impact on the mass-radius relation of compact stars. By using first a toy-model for the quark matter equation of state we have analysed how the stability of hybrid stars depends on the properties of the phase transition. In particular a low value of the pressure at the onset of the phase transition seems to be the most crucial quantity for the stability of the hybrid star. We have then computed the equation of state of quark matter within the NJL model by including effects from the chiral condensates, the diquark coupling pattern and a repulsion vector term. For large enough values of the diquark coupling strength, hybrid stars containing a CFL core are found to be stable. Moreover, exciting stable hybrid stars containing both a layer of a 2SC phase and a core of a CFL phase appears to be possible. This result opens the possibility to find new signatures of the presence of quark matter in compact stars. The formation of different quark phases in compact stars can release a huge amount of energy powering both energetic bursts of neutrino and gamma-ray signals. A detailed study of the time evolution of a proto-neutron star including the possible appearance of the 2SC phase followed by the formation of the CFL phase could reveal some new features and opportunities for detecting the chiral phase transition of QCD in the sky.       \\vspace{0.15cm} {\\bf Acknowledgements}:\\\\ We thank S. R\\\"{u}ster and D. Rischke for providing us the code to compute the equation of state of quark matter. G.P. thanks M. Buballa and D. Nickel  for useful discussions and the INFN for financial support. \\vspace{0.5cm}"
    },
    "0711/0711.4216_arXiv.txt": {
        "abstract": "{} {We have observed the deuterated gas in the high-mass star formation region IRAS 05345+3157 at high-angular resolution, in order to determine the morphology and the nature of such gas.} {We have mapped the \\H\\ (1--0) line with the Plateau de Bure Interferometer, and the \\D\\ (3--2) and \\H\\ (3--2) lines with the Submillimeter Array.} {We have detected two condensations in \\D , with masses of $\\sim 2-3$ and $\\sim 9$ \\solm\\ and diameters of 0.05 and 0.09 pc, respectively. The high deuterium fractionation ($\\sim 0.1$) and the line parameters of the \\D\\ condensations indicate that they are likely low- to intermediate-mass pre--stellar cores, even though other scenarios are possible.} {} ",
        "introduction": "\\label{Introduction}  The initial conditions of the star formation process are still matter of debate. Studies have begun to unveil the chemical and physical properties of starless low-mass cores on the verge of forming low-mass stars (Kuiper et al.~\\citeyear{kuiper}; Caselli et al.~\\citeyear{casellia};~\\citeyear{casellib}; Tafalla et al.~\\citeyear{tafalla02};~\\citeyear{tafalla06}), demonstrating that in the dense and cold nuclei of these cores C-bearing molecular species such as CO and CS are strongly depleted (e.g.~Caselli et al.~\\citeyear{casellib}; Tafalla et al.~\\citeyear{tafalla02}), while N-bearing molecular ions such as \\H\\ and \\D\\ (Caselli et al.~2002a; 2002b; Crapsi et al.~2005) maintain large abundances in the gas phase, and their column density ratio reaches values of $\\sim 0.1$ or more, much higher than the cosmic [D/H] elemental abundance ($\\sim 10^{-5}$, Oliveira et al.~2003).  The characterisation of the earliest stages of the formation process of high-mass stars is more difficult than for low-mass objects, given their shorter evolutionary timescales, larger distances, and strong interaction with their environments. In order to check if the chemical properties peculiar of the earliest stages of low-mass stars are valid also for high-mass stars, Fontani et al.~(2006) observed the deuterated gas in 10 sources selected from two samples of high-mass protostar candidates (Molinari et al.~1996; Sridharan et al.~2002), that are believed to be the closest to the earliest stages of the high-mass star formation process. One of these, IRAS 05345+3157 (hereafter \\i ), stands out because of its very interesting characteristics: it is a luminous ($1.38\\times 10^{3}$\\soll\\ at a distance of 1.8 kpc, Zhang et al.~2005) young stellar object, embedded in a massive ($\\sim 180$ \\solm) dusty clump (Fontani et al.~2006) in which Molinari et al.~(2002) revealed a complex structure in the molecular gas observed in the HCO$^{+}$(1--0) line. From IRAM-30m spectra of the \\H\\ (1--0) and \\D\\ (2--1) lines, Fontani et al.~(2006) have derived an average deuterium fractionation of $\\sim 0.01$, quite close to the values found by Crapsi et al.~(2005) in low-mass starless cores. These results indicate the presence in the source of molecular gas with physical conditions similar to those of low-mass starless cores (i.e. $T\\sim 10$ K and $n_{\\rm H_{2}}\\sim 10^{6}$\\cmc\\ ). However, the precise location and the distribution of this gas, required to understand its nature, can be determined only through higher angular resolution observations.  In this letter, we present observations of \\H\\ and \\D\\ towards \\i , obtained with the Submillimeter Array (SMA) and the Plateau de Bure Interferometer (PdBI), and we report on the detection of two compact condensations of \\D\\ that have chemical features typical of low-mass pre--stellar core candidates. A full report and a more detailed analysis of the data obtained will be presented in a forthcoming paper. ",
        "conclusions": ""
    }
}